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1403 | 1403.6407_arXiv.txt | {In this paper we present a comprehensive study of the mass of the intermediate mass black hole candidate HLX-1 in the galaxy ESO 243-49. We analyse the continuum X-ray spectra collected by {\it Swift}, {\it XMM-Newton}, and {\it Chandra} with the slim disc model, {\tt slimbh}, and estimate the black hole mass for the full range of inclination (inc = $0^{\circ} - 85^{\circ}$) and spin ($a_* = 0 - 0.998$). The relativistic {\tt slimbh} model is particularly suited to study high luminosity disc spectra as it incorporates the effects of advection, such as the shift of the inner disc edge towards smaller radii and the increasing height of the disc photosphere (including relativistic ray-tracing from its proper location rather than the mid-plane of the disc). We find for increasing values of inclination that a zero spin black hole has a mass range of 6,300 - 50,900 $\Msun$ and a maximally spinning black hole has a mass between 16,900 - 191,700 $\Msun$. This is consistent with previous estimates and reinforces the idea that HLX-1 contains an intermediate mass black hole. } \authorrunning{Straub et al.} \titlerunning{Investigating the mass of HLX-1 with {\tt slimbh}} | \label{sec:intro} There is a limit to how luminous an object of a given mass can be. When a star or an accretion disc is in hydrostatic equilibrium it supports itself against gravity by its own internal radiation pressure. The critical luminosity (assuming isotropic emission) is thus given by the Eddington limit, $\LEdd = 4 \pi c \, G M / \kappa_{\rm es} = 1.26 \, \times \, 10^{38} \, (M/\Msun) \,\, {\rm erg/s}$, where c is the speed of light, G is the gravitational constant, M is the mass of the gravitating body, $\Msun$ is the solar mass, and $\kappa_{\rm es} = 0.2 \,(1+X) \, {\rm cm}^2/{\rm g}$ (with $X = 1$) is the electron scattering opacity of a pure hydrogen plasma. There are, however, objects whose luminosities exceed this natural limit. Ultraluminous X-ray sources (ULXs) are sources with X-ray luminosities $\gtrsim 10^{39} \, {\rm erg/s}$. Most ULXs are thought to be powered by super-Eddington accretion onto a stellar mass black hole which can be accomplished (i) by powering strong disc winds \citep{sha+73, lip99}, (ii) by advecting the radiation along with the flow as in radiation pressure dominated disc models like Polish doughnuts \citep{abr+78,jar+80} and slim discs \citep{abr+88}, or (iii) both, advection and outflows \citep{pou+07, dot+11}. Luminosities up to $10^{41} \, {\rm erg/s}$ can therefore still be explained by super-Eddington mass accretion rates onto stellar mass black holes which can have maximum masses up to $\sim 80 \, \Msun$. These higher mass black holes can be explained by direct collapse of metal poor stars \citep{bel+10}. The brightest known ULX in the sky is \object{2XMM J011028.1-460421} in the lenticular galaxy ESO 243-49 \citep[z = 0.0223, ][]{wie+10}. With peak luminosities $\sim 10^{42} \, {\rm erg/s}$ this object, dubbed HLX-1 \citep{far+09}, belongs to a subclass of ULXs called the hyper-luminous X-ray sources \citep[HLX,][]{gao+03}. Like many X-ray binaries, HLX-1 shows transitions from low/hard to high/soft states \citep{god+09,ser+11}, transient radio emission that can be associated with hard-to-soft transitions \citep{web+12} and a weak optical counterpart \citep{sor+10}. The extremely high luminosity of HLX-1 suggests, however, the presence of an intermediate mass black hole (IMBH) with a mass of about $100 \Msun$ to $\sim 10^5 \Msun$. \begin{table*}[htp] \centering \caption{Observational data} \begin{tabular}{ccccc} \hline \hline Obs. Name & Instrument & Obs. ID & Start Date & End Date \\ \hline {\it Swift} & XRT & 00031287(001-252) & 2008-Oct-24 & 2012-Sep-16 \\ [5pt] & & 00032577(001-011) & 2012-Oct-02 & 2012-Nov-11 \\ [5pt] {\it XMM - Newton} & {\em pn}, \em MOS 1 $\&$ 2 & 0560180901 & 2008-Nov-28 & 2008-Nov-28 \\ [5pt] {\it Chandra} & ACIS & 13122 & 2010-Sep-06 & 2010-Sep-07 \\ \hline \label{tab:data} \end{tabular} \end{table*} Spectra of HLX-1 have already been studied with various disc models. These were either limited to one particular inclination and/or black hole spin, non- or semi-relativistic, or based on the standard \citet{sha+73} disc and therefore only valid in the lowest luminosity regime, $L \lesssim 0.1 \LEdd$. All previously used models agree that HLX-1 contains an IMBH. \citet{ser+11} predict a black hole mass $M > 9000 \Msun$ from fitting the non-relativistic {\tt diskbb} model to multi-epoch data collected by {\it Swift}, {\it XMM-Newton} and {\it Chandra}. \citet{dav+11} improve the mass constraints using a relativistic thin disc model with full radiative transfer, {\tt bhspec}, and find $3000 \Msun < M < 3 \times 10^5 \Msun$. They assume a mass range of $1778 - 316228 \Msun$, consider spins $a_* = -1$ to 0.99 and luminosities between $0.03 - 1 \LEdd$ and fit simultaneously for the degenerate mass and spin parameters. As a consequence, a large fraction of their fits peg at the boundary value of at least one of these free parameters, and the results from different spectra are inconsistent with each other. In addition, most of their fits require luminosities far higher than the standard disc models allows. \citet{god+12} address the last point by employing a simplified slim disc model that includes Comptonisation and some relativistic corrections \citep[][]{kaw03} and estimate $M \sim 2 \times 10^4 \Msun$. They fit for a large mass range between $1 - 10^5 \Msun$ and have a disc structure that allows for high luminosities, but their study is limited to a face-on disc around a Schwarzschild black hole. In this work we resolve the previous shortcomings regarding parameter space, luminosity regime and consistency among results from different spectra. We use a \emph{fully} relativistic slim disc model \citep{sad+11, str+11} that accounts on the one hand for effects related to high mass accretion rates such as advection of radiation, relocation of the inner disc edge towards radii smaller than the innermost stable circular orbit (ISCO), and extended disc height. On the other hand, the model incorporates full vertical radiative transfer by integrating the emission of local annuli spectra and ray-tracing from the proper photosphere location. This slim disc model, {\tt slimbh}, is valid for luminosities up to 1.25 $\LEdd$ and covers all inclinations and prograde spins (see the discussion of the model limits in Section~\ref{sec:modelling}). It differs from the well-known disc model {\tt bhspec}, which is based on a standard thin disc, only in the structure of the underlying disc \citep[the spectral differences between these two models have been studied in][e.g. their Section 3.2 and the top panel in their Fig.4]{str+13}. In comparison to the \citet{kaw03} slim disc, {\tt slimbh} is fully relativistic and uses full radiative transfer. We present results from fitting three thermal X-ray spectra of HLX-1 taken by {\it Swift}, {\it XMM-Newton}, and {\it Chandra}. With the employed slim disc model we are not only able to study the whole parameter plane spanned by black hole spin and inclination, we also get a consistent mass estimation for all spectra. This improves and solidifies the previous mass estimates. The paper contains a brief summary of the data selection in Section~\ref{sec:data}, a discussion about spectral fitting with the slim disc model in Section~\ref{sec:modelling}, and a summary of the results in Section~\ref{sec:results}. We finish with conclusions in Section~\ref{sec:discussion} | \label{sec:discussion} We analysed three spectra of the IMBH candidate HLX-1 that were collected by {\it Swift}, {\it XMM-Newton}, and {\it Chandra} during different missions between 2008 and 2012. We estimate the black hole mass using the fully relativistic slim disc model, {\tt slimbh} \citep{sad+11, str+11}, which allows us to self-consistently probe the trans-Eddington luminosity regime in the whole parameter plane spanned by black hole spin and inclination. This addresses and remedies the deficits of previously used models which were either not relativistic \citep{ser+11}, only valid at lowest luminosities \citep{dav+11}, or only valid for one particular inclination and spin \citep{god+12}. Assuming a low disc viscosity ($\alpha = 0.01$) we find that a Schwarzschild black hole has a mass of about 6,300 - 50,900 $\Msun$ (increasing with disc inclination), whereas a maximally spinning black hole has a mass between 16,900 - 191,700 $\Msun$. A high viscosity disc ($\alpha =0.1$) has black hole masses that are roughly 10\% higher. This result is consistent among all three observations with {\it Swift}, {\it XMM-Newton}, and {\it Chandra}. Moreover, it is also in good agreement with earlier measurements based on {\tt bhspec} \citep{dav+11} and other slim disc models \citep{god+12}. The continuum fitting method that we have applied here determines the inner edge of the accretion disc given its effective temperature and flux, i.e. it is designed to measure the black hole spin and relies strongly on the knowledge of the binary parameters, M , D, and i. Given that the inclination is only constrained to $i < 75^{\circ}$ the constraints on the black hole mass are necessarily fairly weak. Nonetheless, our results clearly place HLX-1 in the regime of intermediate mass black holes. Future dynamical measurements of the binary parameters of HLX-1 will allow us to apply the continuum fitting method as it has been intended, namely to assess the spin of the IMBH. | 14 | 3 | 1403.6407 |
1403 | 1403.1820_arXiv.txt | Cosmic rays (CRs) are protons and atomic nuclei that flow into our Solar system and reach the Earth with energies of up to $\sim10^{21}\,\rm{eV}$. The sources of ultra-high energy cosmic rays (UHECRs) with $E \gtrsim 10^{19}\,\rm{eV}$ remain unknown, although there are theoretical reasons to think that at least some come from active galactic nuclei (AGNs). One way to assess the different hypotheses is by analysing the arrival directions of UHECRs, in particular their self-clustering. We have developed a fully Bayesian approach to analyzing the self-clustering of points on the sphere, which we apply to the UHECR arrival directions. The analysis is based on a multi-step approach that enables the application of Bayesian model comparison to cases with weak prior information. We have applied this approach to the 69 highest energy events recorded by the Pierre Auger Observatory (PAO), which is the largest current UHECR data set. We do not detect self-clustering, but simulations show that this is consistent with the AGN-sourced model for a data set of this size. Data sets of several hundred UHECRs would be sufficient to detect clustering in the AGN model. Samples of this magnitude are expected to be produced by future experiments, such as the Japanese Experiment Module Extreme Universe Space Observatory (JEM-EUSO). | \label{section:intro} Cosmic rays (CRs) are high-energy particles that flow into our Solar system and reach the Earth. They consist mainly of protons and atomic nuclei, and have energies in the range $10^{9}\,{\rm eV}$ to $10^{21}\,{\rm eV}$, which makes them the most energetic particles observed in nature (see e.g.\ \citealt{Stanev2011} for a review). A number of open issues remain in this field, especially with respect to ultra-high energy cosmic rays (UHECRs) with arrival energies $E_{\rm{arr}} \gtrsim 10^{19}\,{\rm eV}$. In particular, no consensus has been reached on the sources of UHECRs. A number of candidates, such as active galactic nuclei (AGNs) and pulsars have been proposed, but lack empirical verification. The strongest demonstration of the origin of the UHECRs would be if they could be associated with their progenitors, something which is made plausible by the fact that the most energetic CRs can only travel for cosmologically short distances before losing energy. UHECRs with energies of $E \ga 5\times 10^{19}\, {\rm eV}$ scatter off the cosmic microwave background (CMB) radiation via the Greisen-Zatsepin-Kuzmin (GZK) effect (\citealt{Greisen66}, \citealt{ZK66}). The resultant energy loss is very significant: the mean free path of the GZK effect at high energies is a few ${\rm Mpc}$ and the energy loss in each collision is 20-50 \% depending on energy (\citealt{Stanev2009}). The GZK effect is expected to cause an abrupt cutoff in the flux of UHECRs at $\sim$ $4 \times 10^{19}\, {\rm eV}$, for which there has been observational support (\citealt{CutoffPAO,CutoffHiRes}). UHECRs that arrive at Earth with energies above the GZK limit can only have come from within a limited radius (the GZK horizon) of $\sim 100 \,{\rm Mpc}$. Due to the low flux, the number of detected UHECRs is small: the largest currently available sample is the 69 events with $E \geq 5.5\times 10^{19}\, {\rm eV}$ recorded by the Pierre Auger Observatory (PAO) between 2004 January 1 and 2009 December 31 (\citealt{PAO2010}). The low number of events is the main reason why any hypothesis about the sources is difficult to investigate. Another difficulty is that CRs are charged particles, and so are deflected by magnetic fields. The deflection due to the extra-Galactic magnetic fields is expected to be $\sim 2$ to $\sim 10$ deg for the highest energy CRs (e.g. \citealt{Medina1998,Sigl2004,Dolag2005}). This complicates the study of UHECR origins because it becomes becomes difficult to directly link arrival directions with possible sources. Nevertheless, a number of attempts have been made to find a correlation between the arrival directions of UHECRs and catalogues of potential sources, although no clear consensus has yet been reached. The Pierre Auger Collaboration reported a strong correlation between the arrival directions of UHECRs with energies $E \ge 5.7 \times 10^{19}\,{\rm eV}$ and the positions of nearby AGNs (\citealt{PAO2007}). The result is supported by Yakutsk data (\citealt{Ivanov2009}), but not by HiRes (\citealt{Abbasi2008}) or Telescope Array (\citealt{AbuZayyad2012}). A more recent analysis of a larger PAO UHECR sample has shown a much weaker correlation than before (\citealt{PAO2010}). All attempts to associate UHECRs with specific sources are hampered to some degree by large magnetic deflections, possibly transient sources and incomplete catalogues. An alternative approach is based on the idea that if the UHECR sources are distributed inhomogeneously inside the GZK horizon, it should be possible to detect a self-clustering in the UHECR arrival directions, independent of any source catalogue. Examples of such work include \cite{Domenico2011} and \cite{PAO2012}. In \cite{PAO2012}, the Pierre Auger Collaboration studied the self-clustering using three statistical methods based on correlation functions (two methods based on the 2-point correlation function, one method based on a 3-point correlation function, developed by \citealt{Ave2009}). No strong evidence of non-uniformity was found based on the $p$-values obtained under the null hypothesis of no clustering. The interpretation of $p$-values is, however, known to be problematic as they have no quantitative link to the (posterior) probability that the null hypothesis is correct (see e.g.\ \citealt{Berger_Delampady1987}). Whereas $p$-values are probabilities conditional on the null hypothesis, what is needed is a method of calculating the probability that the null hypothesis is correct. \cite{Cox1946} proved that Bayesian inference is the only self-consistent method to make probabilistic statements about models based on observations, and Bayesian methods have previously been used to assess whether UHECRs originate from AGNs \citep{Watson_etal2011,Soiaporn2012}. In this paper we present a Bayesian analysis of the self-clustering of the PAO UHECRs. The Bayesian method for assessing non-uniformity is explained in Section~\ref{sec:StatMeth}. In Section~\ref{sec:ApplicationSimul}, the effectiveness of the method is discussed, based on tests of the method on simulated mock UHECR catalogues. The application of the method to data from PAO is discussed in Section~\ref{sec:ApplicationPAO}. Our conclusions are summarized in Section~\ref{sec:Conc}. \vspace{1mm} | \label{sec:Conc} We have developed a Bayesian method for the analysis of the self-clustering of points on a sphere and applied it to the 69 highest energy UHECRs detected by PAO up until 31 December 2009. The method is a three-step Bayesian approach, in which the data are divided into three subsets: the first two subsets of the data are used to generate a model of self-clustered UHECRs; the third subset is used to perform Bayesian model comparison between this self-clustered model and a uniform model of UHECRs. This approach is an extension of the Bayesian model comparison methods that were developed by \cite{Spiegelhalter1982}, \cite{Aitkin1991}, \cite{OHagan1991} and \cite{OHagan1995}. Like the multi-step method that is presented here, those approaches are aimed to evaluate the marginal likelihood in cases when there is weak prior information on the model parameters. The method we have presented here is not specific to the UHECR problem in question and could be applied to anisotropy searches in other areas of astronomy, such as the search for angular anisotropies in the distribution of gamma-ray bursts described by e.g. \cite{Balazs_etal1998} and \cite{Magliocchetti_etal2003}. There is some ambiguity in the partitioning of the full data set. In the present implementation, the total data set is divided into three subsets of equal size. However, it is possible that a different partitioning, or perhaps an average over partitions could make this method more effective. These issues will be explored in future work. We tested our model comparison method on mock catalogues of UHECRs. The results for uniform UHECR arrival directions were compared to the results for UHECRs originating in AGNs from a realistic mock catalogue. UHECR clustering in a realistic AGN centred model is too weak to be detected in a sample of 69 events, but would be detectable in samples of 690 events. This is consistent with the results of \cite{PAO2012}. We assumed a pure proton composition of the cosmic rays, but there are some indications that heavier nuclei are also part of the composition (\citealt{Unger2007}). The effect of including heavier nuclei will be investigated through additional simulations. For the PAO data, Bayes factors were calculated for different random partitions of the data. The geometric and arithmetic means of the Bayes factors were 0.57 and 1.26 respectively, corresponding to posterior probabilities of 0.37 and 0.56 for the clustered model. Thus, we did not find strong evidence for clustering in the PAO data, although the data are also consistent with the AGN-centred simulations. It is expected that future experiments will produce data sets that will be sufficiently large for our Bayesian method (and other statistical approaches; see e.g.\ \citealt{FutureAnis2014}) to detect even the weak clustering expected if the UHECRS have come from nearby AGNs. PAO is continuing to take data and is expected to produce a sample of $\sim 250$ UHECRs over its first decade of operations. Looking further ahead, the planned Japanese Experiment Module Extreme Universe Space Observatory (JEM-EUSO, \citealt{JEMEUSO2013}) on the International Space Station (ISS) is scheduled for launch in 2017 and is expected to detect $\sim 200$ UHECRs annually over its five year lifetime. These data sets should be sufficiently large to detect the self-clustering of UHECRs independent of the source population. | 14 | 3 | 1403.1820 |
1403 | 1403.0221_arXiv.txt | {Using hydrodynamical simulations coupled to a radiative transfer code, we study the additional heating effects in the intergalactic medium (IGM) produced by $z\sim6$ quasars in their near-zones. If helium is predominantly in \HeII\ to begin with, both normalization ($T_0$) and slope ($\gamma$) of the IGM effective equation-of-state get modified by the excess ionization from the quasars. Using the available constraints on $T_0$ at $z\sim6$, we discuss implications for the nature and epoch of \HI\ and \HeII\ reionization. We study the extent of the He~{\sc iii} region as a function of quasar age and show, for a typical inferred age of $z\sim6$ quasars (i.e. $\sim 10^8$ yrs), it extends up to 80\% of the H~{\sc i} proximity region. For these long lifetimes, the heating effects can be detected even when all the \HI\ lines from the proximity region are used. Using the flux and curvature probability distribution functions (PDFs), we study the statistical detectability of heating effects as a function of initial physical conditions in the IGM. For the present sample size, cosmic variance dominates the flux PDF. The curvature statistics is more suited to capturing the heating effects beyond the cosmic variance, even if the sample size is half of what is presently available.} | Unravelling the process of reionization, which signals the end of the `dark ages' of our universe, is one of the current challenges of observational and theoretical cosmology. Two major milestones in the reionization history of the universe are those of hydrogen (\HI) and singly ionized helium (\HeII). Study of the evolution of hydrogen reionization combines observational evidences from various sources; optical probes include the Gunn-Peterson absorption troughs \citep{gunnpeterson} in the spectra of high-redshift bright sources such as (a) quasars \citep{fan, willott2007,mortlock2011}, (b) Lyman-$\alpha$ emitters \citep{kashikawa2006, stark2007, ouchi2010, nakamura2011} and (c) $\gamma$-ray bursts \citep[GRBs;][]{totani2006, kistler2009, ishida2011,robertsonellis2012}. The Thomson scattering optical depth measurements from the Cosmic Microwave Background (CMB) temperature and polarization power spectra are consistent with an instantaneous reionization at redshift $z \sim 11$ \citep{larson,planck, wmap7}, which may be interpreted as an estimate of the mean reionization redshift. At radio frequencies, the redshifted 21-cm hyperfine line of neutral hydrogen promises a unique three-dimensional mapping of the epoch of reionization (EoR) of hydrogen \citep[for a review, see][] {furlanettorev}. All the available observations at present are consistent with an extended \HI\ reionization history that probably began at $z \sim 15$ and ended around $z \sim 6$ \citep{wyithe2003,trc2005, trc2006, pritchard2010, pritchard2010proc,mitra2011,mitra2012}. The current observational probes of \HeII\ reionization include measuring the Gunn-Peterson absorption troughs in the \HeII\ Lyman-$\alpha$ forest \citep{jakobsen1994,zheng2004,reimers2005,shull2010, worseck2011,syphers2014}. These observations suggest that the EoR of \HeII\ is close to $z \sim 2.7$. The reionization of \HeII\ also leaves a thermal imprint on the hydrogen Lyman-$\alpha$ forest due to the additional heating effect on the velocity widths of the Lyman-$\alpha$ lines \citep{huignedin}. The thermal evolution of the intergalactic medium (IGM) from $2 \leq z \leq 4.8$ has been probed using the observations of the Lyman-$\alpha$ forest \citep{ricotti2000a, schaye2000,mcdonald2001}. The velocity widths of the hydrogen Lyman-$\alpha$ forest lines seem to exhibit a sudden increase between redshifts $z \sim 3.5$ and 3, which may represent evidence for the reionization of \HeII. The inferred temperature measurements, taken in conjunction with the adiabatic cooling expected to occur after the reionization of hydrogen, also constrain the EoR of hydrogen to below $z \sim 9$ \citep{theuns2002}. Recently, \citet{becker11} reported measurements of the IGM temperature from $2 \leq z \leq 4.8$ using the curvature statistic to quantify the temperature; their observations indicated gradual heating of the IGM from $z \sim 4.4$ towards lower redshifts, in contrast to the adiabatic cooling expected in single-step models of reionization. These measurements are consistent with an extended epoch of \HeII\ reionization starting probably at $z \gtrsim 4.4$ and terminating around $z \sim 3$. Helium is expected to be singly ionized around the same time as the hydrogen gets ionized, and first-generation galaxies are believed to be the likely sources for completion of hydrogen and \HeI\ reionization. In the single-step model of reionization, it is believed that massive, metal-free Population III stars \citep{ohetal2001,venkatesan2003} may have provided the hard photons required for \HeII\ reionization. In this model, a population of metal-free (Pop III) stars are required at redshifts $z > 6$ to reionize both \HI\ and \HeII. In the absence of a strong ionizing background for \HeII, it may recombine and hence to be reionized again at a lower redshift. Therefore, probes of intergalactic \HeII\ are important for understanding the role of Population III stars in the early reionization of \HeII\ and setting up a \HeII\ ionizing background prior to the quasar era (i.e. $z \sim 6$). Recently, there are indications of the presence of Population III stars even as late as $z \sim 3$ possibly due to inefficient transport of heavy elements and/or poor mixing that leave pockets of pristine gas even in chemically evolved galaxies \citep{jimenez2006, tornatore2007, inoue2011,cassata2013}. If, on the other hand, reionization took place as a two-step process (hydrogen first and \HeII\ later), quasars\footnote{{Strictly speaking, the term `quasar' is reserved for describing radio-loud quasi-stellar objects. However, as frequently done in the literature, we will use the term `quasars' in this paper to indicate quasi-stellar objects, irrespective of their radio properties.}} are believed to be the most likely candidates for reionization of \HeII\ since their spectra are sufficiently hard. However, the number density of bright quasars peaks at $z \sim 2 - 3$ and decreases rapidly above $z \sim 4$ \citep{assef2011, masters2012}. Hence, in the two-step model of reionization, the final stages of \HeII\ reionization are expected to coincide with the peak of the quasar activity at $z \sim 2 - 3$. Quasar proximity zones\footnote{{Here, and in what follows, the term ``proximity zone'' or ``\HI\ proximity zone'' describes the region in the vicinity of the quasar where the ionizing flux from the quasar dominates the background flux.}}, where the excess ionization by the quasar allows the measurement of the velocity width of the Lyman-$\alpha$ line, have been used to probe the thermal state of the IGM at $z \sim 6$ \citep{bolton10}. This, in turn, can be used to probe the role of quasars in \HeII\ reionization. The IGM temperature in the near-zone\footnote{{Here, and in what follows, the term ``near zone'' refers to the region in the vicinity of the quasar within the \HeIII\ ionization front, where the heating effects are significant.}} is influenced by both the existing background radiation as well as the additional radiation from the quasar itself. A first measurement of the near-zone temperature around a quasar at redshift 6 has been reported in \citet{bolton10} using Keck/HIRES data in combination with hydrodynamical simulations. Recently, an additional source of heating has been observed in the ionized near-zones of high-redshift quasars at $z \sim 6$, which is attributed \citep{bolton12} to the initial stages of helium reionization around that redshift, since the excess heating can be easily accounted for if the \HeII\ is ionized by the quasar. The inferred excess temperature in the quasar near-zone can be used to place constraints on the epoch of \HI\ reionization \citep[see, for example,][]{ciardi,raskutti2012}. In this paper, we explore several aspects of the additional heating effect in the near-zones of quasars at $z \sim 6$ using the results of high-resolution hydrodynamical (SPH) simulations with {\sc gadget-2} \citep{gadget2}, and the ionization correction done using a 1D radiative transfer code which we have developed. The gas temperature in the general IGM is given by the assumed equation of state \citep{huignedin} and computed self-consistently for the near-zone of the quasar. We first validate our simulations by computing the additional temperature in the near zone for different initial equations of state of the general IGM, and different assumed values of the \HeII\ fraction prior to the active quasar phase. We obtain the expected relationship between the excess temperature and the initial \HeII\ fraction in the quasar near-zone, and also find a connection between the magnitude of the steepening of the equation of state and the initial \HeII\ fraction. We then use our simulation results to measure the size of the region in the near-zone heated by the quasar in comparison to the \HI\ proximity zone, as a function of the age of the quasar. We also validate the usage of the flux and curvature statistics to measure the increased temperature in the near-zone of the quasar, and, in particular, address the effect of cosmic variance. For the flux statistics tests, we employ a number of pixels typical of the sample sizes in available observations of quasar near-zones at redshifts $\sim 6$. Using the Kolmogorov-Smirnov (KS) statistic to quantify the effect of the additional heating, and examining its variation with the parameters of the equation of state, $T_0$ and $\gamma$, we establish the connection between the thermal evolution of the IGM following the reionization of hydrogen, and the detectability of the additional heating in the quasar near-zone. We also consider the possible dependence of the detectability of the additional heating effect on the assumed values of the background (metagalactic) photoionization rate of \HeII, which translates into varying the \HeII\ fraction in the near-zone of the quasar. This allows a connection to the effect of Population III stars on reionizing \HeII\ at redshifts $z > 6$ (which constrains the initial \HeII\ fraction in the quasar near-zone) in single-step reionization scenarios. The paper is organized as follows: In Sec. \ref{sec:nummodel}, we describe our hydrodynamical simulations and the numerical formalism for obtaining the simulated spectra in the quasar near-zone. In Sec. \ref{sec:validepen}, we provide a validation of our simulations by computing the excess temperature in the quasar near-zone for different values of the equation of state normalization, and the initial \HeII\ fraction, with comparison to the measured average temperature \citep{bolton12} in seven quasar near-zones at redshift $\sim 6$. We also describe the modification to the initial equation of state of the IGM due to the additional heating, and its dependence on the initial \HeII\ fraction in the quasar near-zone. In Sec. \ref{sec:results}, we describe the results obtained from our calculations as regards (a) the extent of the region around the quasar within which the additional heating is expected to contribute significantly, (b) the dependence of the additional heating effect in the near-zone on the initial equation of state of the IGM, quantified by the flux and curvature statistics, and (c) the dependence of the heating effect on the initial \HeII\ fraction in the near-zone, which is related to the single-step reionization by Population III stars. We then summarize our findings in a brief concluding section. Throughout this article, we assume the cosmological parameters $\Omega_m = 0.26$, $\Omega_{\Lambda} = 0.74$, $\Omega_b h^2 = 0.024$, $h = 0.72$, $\sigma_8 = 0.85$, and $n_s = 0.95$, which are consistent with the third-year WMAP and Lyman-$\alpha$ forest data \citep{seljak2006,viel2006}. The helium fraction by mass is taken to be 0.24 \citep{oliveskillman}. | In this paper, we used detailed hydrodynamical simulations to provide an analysis of the features associated with the heating due to the ionization of \HeII\ in the near-zones of high-redshift quasars, and their implications for constraining the epochs of \HI\ and \HeII\ reionization. Our main findings may be summarized as follows: \begin{enumerate} \item We have seen that the measured temperature \citep{bolton12} in the quasar near-zones arises from a combination of two effects : the initial \HeII\ fraction in the quasar vicinity, and the normalization of the initial equation of state of the IGM. If the initial temperature at mean density is $\lesssim 8000$ K, the measured temperature in the quasar near-zones is higher than that expected for the allowed range of initial \HeII\ fractions ($x_{\rm He II} = 0.04 - 1$) in the quasar vicinity. This shows that the temperature measurement can be used to place constraints on (a) the epoch and temperature of hydrogen reionization, and (b) single-step models of reionization that predict the initial \HeII\ fraction. \item We recover the expected linear relationship of $\Delta T_0$ increasing with the initial helium fraction $x_{\rm HeII}$. Akin to the $\Delta T- x_{\rm HeII}$ relation discussed in the literature \citep{furlanetto2008}, we also demonstrate a $\Delta \gamma - x_{\rm He II}$ relation, which shows a decrease in $\Delta \gamma$ with increasing $x_{\rm HeII}$ and a flattening out at the lowest $x_{\rm HeII}$ values, thus illustrating the steepening of the equation of state with decrease in the \HeII\ fraction in the quasar vicinity. Observationally, this steepening effect, which persists even for high initial temperatures where $\Delta T_0$ is low, may also be used to constrain the near-zone \HeII\ fraction. However, the maximum expected increase in the slope may be more difficult to detect observationally than the expected shift in the overall normalization. \item Optical depth effects are coupled to the propagation of the ionization front in the radiative transfer, so that we obtain a handle on the extent of the near-zone of \HeIII, where the additional heating is expected to contribute significantly. If the quasar age is $\sim 100$ Myr, more than 80\% of the \HI\ proximity zone is heated in 78\% of the sightlines. The heated fraction of the \HI\ proximity zone is only about 30\% - 35\% for quasar lifetimes of $\sim 10$ Myr. This indicates that including the entire extent of the \HI\ proximity zone for the temperature enhancement may result in some dilution of the statistics when the quasar lifetimes are short. However, considering the entire proximity zone of \HI\ is a valid approximation if the quasar lifetimes are longer, $\gtrsim 100$ Myr. This is also the timescale for the saturation of the heating effect, making it fairly independent of distance. \item We have quantified the effect of additional heating by using the flux PDF and curvature statistics to compare the real spectra to the simulated spectra without heating. We have noted that the sensitivity of the curvature statistic to the noise in the spectra may be effectively removed by smoothing with a Gaussian filter with a velocity width of 7 km/s. Both these statistics indicate that a higher value of $T_0$ and/or $\gamma$ leads to less detectability of the effect of additional heating. This connects the additional heating due to \HeII\ reionization, to the epoch of hydrogen reionization. \item We find that the curvature statistic provides far more effective distinguishability of the heating effect, which is over and above the cosmic variance of individual samples of 10 lines-of-sight each having 512 pixels (chosen to match the typical sample sizes available in observations of seven quasars at redshift $\sim 6$). We also find that the detectability of the heating effect is dependent on the initial \HeII\ fraction in the quasar vicinity, with a greater \HeII\ fraction leading to greater detectability. \end{enumerate} } | 14 | 3 | 1403.0221 |
1403 | 1403.5760_arXiv.txt | The recent detection of the cosmic microwave background polarimeter experiment BICEP2 of tensor fluctuations in the B-mode power spectrum basically excludes all plausible axion models where its decay constant is above $10^{13}$\,GeV. Moreover, there are strong theoretical, astrophysical, and cosmological motivations for models involving, in addition to the axion, also axion-like particles (ALPs), with decay constants in the intermediate scale range, between $10^9\,{\rm GeV}$ and $10^{13}\,{\rm GeV}$. Here, we present a general analysis of models with an axion and further ALPs and derive bounds on the relative size of the axion and ALP photon (and electron) coupling. We discuss what we can learn from measurements of the axion and ALP photon couplings about the fundamental parameters of the underlying ultraviolet completion of the theory. For the latter we consider extensions of the Standard Model in which the axion and the ALP(s) appear as pseudo Nambu-Goldstone bosons from the breaking of global chiral $U(1)$ (Peccei-Quinn (PQ)) symmetries, occurring accidentally as low energy remnants from exact discrete symmetries. In such models, the axion and the further ALP are protected from disastrous explicit symmetry breaking effects due to Planck-scale suppressed operators. The scenarios considered exploit heavy right handed neutrinos getting their mass via PQ symmetry breaking and thus explain the small mass of the active neutrinos via a seesaw relation between the electroweak and an intermediate PQ symmetry breaking scale. For a number of explicit models, we determine the parameters of the low-energy effective field theory describing the axion, the ALPs, and their interactions with photons and electrons, in terms of the input parameters, in particular the PQ symmetry breaking scales. We show that these models can accommodate simultaneously an axion dark matter candidate, an ALP explaining the anomalous transparency of the universe for $\gamma$-rays, and an ALP explaining the recently reported 3.55 keV gamma line from galaxies and clusters of galaxies, if the respective decay constants are of intermediate scale. Moreover, they do not suffer severely from the domain wall problem. | \label{sec:intro} The axion $A$ is a strongly motivated very weakly interacting ultralight particle beyond the Standard Model (SM). Its existence is predicted in the course of an elegant solution to the strong CP problem \cite{Peccei:1977hh,Weinberg:1977ma,Wilczek:1977pj}, that is the non-observation of flavor conserving CP violation originating from the theta-angle term in the Lagrangian of QCD, \begin{equation} \lag_{\rm QCD}\supset -\frac{\alpha_s}{8\pi} \,\bar\theta\, G_{\mu\nu}^a {\tilde G}^{a,\mu\nu} , \label{qcd} \end{equation} involving the gluon field strength $G_{\mu\nu}^a$ and its dual, ${\tilde G}^{a,\mu\nu} \equiv \epsilon^{\mu\nu\lambda\rho} G_{\lambda\rho}^a/2$, with $\varepsilon^{0123}=1$. This solution consists in adding to the Standard Model a scalar field theory describing a (pseudo) Nambu-Goldstone boson arising from the breaking of a global chiral (Peccei-Quinn (PQ)) $U(1)_{\rm PQ}$ symmetry: that is, the corresponding scalar field $A(x)$ satisfies a shift symmetry $A(x)\to A(x) + {\rm const.}$ which is only violated by the anomalous coupling to gluons, \begin{equation} \lag \supset - \frac{\alpha_s}{8\pi} \,\frac{A}{f_A}\, G_{\mu\nu}^a {\tilde G}^{a,\mu\nu} . \label{qcd} \end{equation} In fact, the $\bar\theta$-term can then be eliminated by absorbing it into the axion field, $A + \bar{\theta} f_A\to A$. Moreover, non-perturbative QCD effects\footnote{Semiclassical instanton methods are not reliable to calculate the potential accurately. One has to use matching to the low-energy chiral Lagrangian instead~\cite{Weinberg:1977ma,Georgi:1986df}.} provide for a non-trivial potential for the shifted axion field $A$ -- minimized at zero expectation value, $\langle A\rangle =0$, thus wiping out strong CP violation -- and predict a mass for the particle excitation around this minimum, the axion $A$, \begin{equation} m_A = \frac{m_\pi f_\pi}{f_A} \frac{\sqrt{z}}{1+z} \simeq {6\, {\rm meV}} \times \left( \frac{10^{9}\, {\rm GeV}}{f_A}\right). \label{axion_mass} \end{equation} in terms of the pion mass, $m_\pi = 135$ MeV, its decay constant, $f_\pi \approx 92$ MeV, the ratio $z=m_u/m_d \approx 0.56$ of up and down quark masses, and the decay constant $f_A$. For large $f_A$, the axion is an ultralight particle with very weak interactions with the Standard Model \cite{Kim:1979if,Shifman:1979if,Dine:1981rt,Zhitnitsky:1980tq}. Its universal and phenomenologically most important interaction with photons \cite{Georgi:1986df,Bardeen:1977bd,Kaplan:1985dv,Srednicki:1985xd}, \begin{equation} \label{photon_coupling_universal} \lag \supset - \frac{g_{A\gamma}}{4}\,A\, F_{\mu\nu} \tilde{F}^{\mu\nu} ; \hspace{6ex} | g_{A\gamma} | \sim \frac{\alpha}{2\pi f_A} \sim \frac{\alpha}{2\pi} \frac{m_A}{m_\pi f_\pi} \sim 10^{-12}\ {\rm GeV}^{-1} \left( \frac{10^{9}\, {\rm GeV}}{f_A}\right) \sim 10^{-12}\ {\rm GeV}^{-1} \left( \frac{m_A}{6\, {\rm meV}}\right) , \end{equation} is tiny (see the yellow band in Fig. \ref{ALP_coupling_limits}), since observations in astrophysics -- in particular the observed duration of the neutrino signal from supernova SN 1987A -- require a large decay constant $f_A$ (small mass $m_A$) \cite{Raffelt:2006cw}, \begin{equation} f_A\gtrsim 4\times 10^8\,{\rm GeV} \Rightarrow m_A\lesssim 16\,{\rm meV} . \end{equation} A further strong motivation for the axion comes from the fact that, for sufficiently large decay constant $f_A$, \begin{equation} 10^9\,{\rm GeV} \lesssim f_A\lesssim 10^{13}\,{\rm GeV} \Rightarrow 10^{-7}\,{\rm eV}\lesssim m_A\lesssim 1\,{\rm meV} , \label{cosmic_axion_window} \end{equation} it is a cold dark matter candidate, being non-thermally produced in the early universe by the vacuum-realignment mechanism and, in some models and under certain circumstances, also via the decay of topological defects such as axion strings and domain walls~\cite{Preskill:1982cy,Abbott:1982af,Dine:1982ah,Sikivie:2006ni,Bae:2008ue,Visinelli:2009zm,Wantz:2009it,Hiramatsu:2010yu,Hiramatsu:2012gg}. The upper bound on the decay constant in Eq. \eqref{cosmic_axion_window} follows from the recent discovery of the cosmic microwave background polarimeter experiment BICEP2 of tensor fluctuations in the $B$-mode power spectrum \cite{Ade:2014xna}. This implies a high value for the Hubble scale during inflation, \begin{equation} H_I\simeq 1.1\times 10^{14}\,{\rm GeV}, \end{equation} which, together with constraints on isocurvature fluctuations \cite{Fox:2004kb,Beltran:2006sq,Hertzberg:2008wr,Hamann:2009yf,Ade:2013uln}, rule out plausible scenarios where inflation occurs after PQ symmetry breaking \cite{Higaki:2014ooa,Marsh:2014qoa,Visinelli:2014twa}, that is where \begin{equation} f_A>{\rm max}\{T_{\rm GH},T_{\rm max}\}, \end{equation} with \begin{equation} T_{\rm GH}=\frac{H_I}{2\pi } \end{equation} the Gibbons-Hawking temperature during inflation and \begin{equation} T_{\rm max}=\epsilon_{\rm eff}\sqrt{\sqrt{\frac{3}{8\pi}}M_{\rm Pl}\, H_I} \end{equation} the maximum thermalization temperature after reheating. Here $M_{\rm Pl}=1.22\times 10^{19}$\,GeV is the Planck mass and $0 < \epsilon_{\rm eff} < 1$ is an efficiency parameter. Ways out of this conclusion have been put forward in Refs. \cite{Higaki:2014ooa,Marsh:2014qoa}. The pre-inflationary PQ symmetry breaking scenario would have allowed, in principle, much higher values of the decay constant without overshooting the dark matter abundance, by invoking small values of the initial misalignment angle. After BICEP2, the plausible axion possibilities have narrowed down to scenarios with post-inflationary PQ symmetry breaking. A conservative upper bound on the axion decay constant is then \begin{equation} f_A < \frac{H_I}{2\pi }\simeq 1.8\times 10^{13}\,{\rm GeV}, \hspace{6ex} m_A > 0.3\ \mu{\rm eV} . \end{equation} A more stringent upper bound can be obtained by requiring that the axion dark matter abundance generated via the vacuum-realignment mechanism should not exceed the observed dark matter abundance, leading to \cite{Bae:2008ue,Hertzberg:2008wr,Visinelli:2009zm,Wantz:2009it,Visinelli:2014twa} \begin{equation} f_A\lesssim ( 1 \div 10 ) \times 10^{11}\,{\rm GeV} , \hspace{6ex} m_A\gtrsim (6 \div 60 )\ \mu{\rm eV} . \end{equation} To be on the very conservative side, the red region in Fig. \ref{ALP_coupling_limits} labelled ``Axion DM" comprises still both cases, the disfavored pre- and the favored post-inflationary PQ symmetry breaking. Two haloscope experiments (ADMX \cite{Asztalos:2011bm} and ADMX-HF) are currently aiming at the direct detection of axion dark matter, based on microwave cavities in the mass region $2\times 10^{-6}\,{\rm eV}\lesssim m_A\lesssim 2\times 10^{-5}\,{\rm eV}$ (see the green regions labelled ``Haloscopes" in Fig. \ref{ALP_coupling_limits}). Further experiments have been proposed to extend this region on both ends of the spectrum \cite{Graham:2011qk,Baker:2011na,Irastorza:2012jq,Horns:2012jf,Graham:2013gfa,Budker:2013hfa,Jaeckel:2013sqa,Jaeckel:2013eha,Horns:2013ira,Sikivie:2013laa,Redondo:2013hca,Rybka:2014cya}. There is also a theoretical motivation for the existence of additional axion-like particles (ALPs) (for reviews, see \cite{Jaeckel:2010ni,Ringwald:2012hr,Ringwald:2012cu}), emerging as pseudo Nambu-Goldstone bosons from the breaking of other global symmetries, such as lepton number \cite{Chikashige:1980ui,Gelmini:1980re} or family symmetries \cite{Wilczek:1982rv,Berezhiani:1990wn,Jaeckel:2013uva}. Most notably, the low-energy effective field theories emerging from string theory generically contain candidates for the axion plus possibly a multitude of additional ALPs. Indeed, when compactifying the six extra spatial dimensions of string theory, Kaluza-Klein zero modes of antisymmetric form fields -- the latter belonging to the massless spectrum of the bosonic string propagating in ten dimensions -- appear as axion and further ALP candidates in the low-energy effective action \cite{Witten:1984dg,Conlon:2006tq,Svrcek:2006yi,Arvanitaki:2009fg,Acharya:2010zx,Cicoli:2012sz}, their number being determined by the topology of the internal compactified manifold. Moreover, the axion and a multitude of additional ALP candidates may also arise as pseudo Nambu-Goldstone bosons from the breaking of accidental global $U(1)$ symmetries that appear as low energy remnants of exact discrete symmetries occurring in string compactifications~\cite{Lazarides:1985bj,Choi:2006qj,Choi:2009jt}. Intriguingly, very light and weakly coupled generic ALPs $a_i$, with decay constants in the same range as Eq. \eqref{cosmic_axion_window}, share with the axion the property of being cold dark matter candidates since they are also produced via the vacuum-realignment mechanism \cite{Arias:2012az,Marsh:2013ywa,Kadota:2013iya}. In fact, their relic abundance, in the now strongly favored post-inflationary symmetry breaking scenario, is roughly given by\footnote{The red line labelled ``ALP CDM" in Fig. \ref{ALP_coupling_limits} corresponds to the ALP photon coupling, where an ALP, according to \eqref{eq:omegaalp}, can account for all of the dark matter in the universe, assuming that it couples with strength $g_{a_i \gamma} = \alpha/(2\pi f_{a_i})$. The region below this line is strongly disfavored after BICEP2 by overdensity constraints \cite{Higaki:2014ooa,Marsh:2014qoa}.} \begin{equation} \Omega_{a_i} h^2 \approx 0.053 \times \left( \frac{m_{a_i}}{\rm eV} \right)^{1/2} \left( \frac{f_{a_i}}{10^{11}\ \rm GeV} \right)^{2} . \label{eq:omegaalp} \end{equation} \begin{figure} \begin{center} \includegraphics[width=0.75\textwidth]{ALPplot3kev} \caption{Prediction of the axion-photon coupling versus its mass (yellow band). Also shown are excluded regions arising from the non-observation of an anomalous energy loss of massive stars due to axion or ALP emission \cite{Friedland:2012hj}, of a $\gamma$-ray burst (in coincidence with the observed neutrino burst) from SN 1987A due to conversion of an initial ALP burst in the galactic magnetic field, of changes in quasar polarizations due to photon-ALP oscillations, and of dark matter axions or ALPs converted into photons in microwave cavities placed in magnetic fields~\cite{DePanfilis:1987dk,Wuensch:1989sa,Hagmann:1990tj,Asztalos:2001tf,Asztalos:2009yp}. Axions and ALPs with parameters in the regions surrounded by the red lines may constitute all or a part of cold dark matter (CDM), explain the cosmic $\gamma$-ray transparency, and the soft X-ray excess from Coma. The green regions are the projected sensitivities of the light-shining-through-wall experiment ALPS-II, of the helioscope IAXO, of the haloscopes ADMX and ADMX-HF, and of the PIXIE or PRISM cosmic microwave background observatories.} \label{ALP_coupling_limits} \end{center} \end{figure} Therefore, there is a strong motivation to search not only for the axion $A$, but also for additional light ALPs $a_i$, for which the low-energy effective Lagrangian describing their interactions with (the lightest Standard Model particles) photons can generically be written as \begin{equation} \lag \supset\frac{1}{2}\, \partial_\mu A\, \partial^\mu A -\frac{1}{2} m_A^2 A^2 - \frac{g_{A\gamma}}{4}\,A\, F_{\mu\nu} \tilde{F}^{\mu\nu} +\sum_{i=2}^{n_{\rm ax}} \left( \frac{1}{2}\, \partial_\mu a_i\, \partial^\mu a_i -\frac{1}{2} m_{a_i}^2 a_i^2 - \frac{g_{a_i\gamma}}{4}\,a_i\, F_{\mu\nu} \tilde{F}^{\mu\nu} \right) . \label{AALPgen_leff} \end{equation} However, unlike the axion, which inherits many of its properties ($m_A$, $g_{A\gamma}$), from non-perturbative QCD effects associated with chiral symmetry breaking, the masses $m_{a_i}$ and the photon couplings $g_{a_i\gamma}$ of the additional ALPs are model dependent. Thus, there exists the possibility that ALPs are hierarchically more strongly coupled to photons than axions with the same mass and therefore easier to detect. Interestingly, there are indications from gamma-ray astronomy, which may point to the existence of at least one ALP beyond the axion. Gamma-ray spectra from distant active galactic nuclei (AGN) should show an energy and redshift-dependent exponential attenuation, $\exp (-\tau (E,z))$, due to $e^+ e^-$ pair production off the extragalactic background light (EBL) -- the stellar and dust-reprocessed light accumulated during the cosmological evolution following the era of re-ionization. The recent detection of this effect by Fermi \cite{Ackermann:2012sza} and H.E.S.S. \cite{Abramowski:2012ry} has allowed to constrain the EBL density. At large optical depth, $\tau\gtrsim 2$, however, there are hints that the Universe is anomalously transparent to gamma-rays \cite{Aharonian:2007wc,Aliu:2008ay,Essey:2011wv,Horns:2012fx,Horns:2013pha}. This may be explained by photon $\leftrightarrow$ ALP oscillations: the conversion of gamma rays into ALPs in the magnetic fields around AGNs or in the intergalactic medium, followed by their unimpeded travel towards our galaxy and the consequent reconversion into photons in the galactic/intergalactic magnetic fields~\cite{DeAngelis:2007dy,Simet:2007sa,SanchezConde:2009wu,Horns:2012kw,Mena:2013baa}. This explanation requires a very light ALP, which couples to two photons with strength \cite{Meyer:2013pny}, \begin{equation} |g_{a\gamma}|\gtrsim 10^{-12} \ {\rm GeV}^{-1} ; \hspace{6ex} m_{a}\lesssim 10^{-7}\ {\rm eV}. \label{dec_const_transp} \end{equation} Note that the entire parameter region \eqref{dec_const_transp} has no overlap with the universal prediction of the axion, see Fig. \ref{ALP_coupling_limits}. In fact, an axion with $m_A\lesssim 10^{-7}$\,eV would have a photon coupling $|g_{A\gamma}|\lesssim 10^{-16}$\,GeV$^{-1}$. Therefore, if this hint is taken seriously, it points to the existence of an ALP beyond the axion. Intriguingly, an observed soft X-ray excess in the Coma cluster may also be explained by the conversion of a cosmic ALP background radiation (CABR) -- produced by heavy moduli decay and corresponding to an effective number $\triangle N_{\rm eff}\sim 0.5$ of extra neutrinos species \cite{Cicoli:2012aq,Higaki:2012ar} -- into photons in the cluster magnetic field \cite{Conlon:2013txa,Angus:2013sua}. This explanation requires that the spectrum of the CABR is peaked in the soft-keV region and that the ALP coupling and mass satisfy \begin{equation} |g_{a\gamma}|\gtrsim 10^{-13} \ {\rm GeV}^{-1}\, \sqrt{0.5/\triangle N_{\rm eff}}\,; \hspace{6ex} m_a\lesssim 10^{-12}\ {\rm eV}, \end{equation} respectively, overlapping with the parameter range \eqref{dec_const_transp} preferred by the ALP solution of the gamma-ray transparency puzzle, as is apparent in Fig. \ref{ALP_coupling_limits}. Astrophysical bounds on ALPs arising from magnetised white dwarfs \cite{Gill:2011yp} and from the non-observation of a $\gamma$-ray burst in coincidence with neutrinos from the supernova SN 1987A~\cite{Brockway:1996yr,Grifols:1996id} provide limits close to $|g_{a\gamma}|\lesssim 10^{-11}~{\rm GeV}^{-1}$, for masses $m_a\lesssim 10^{-7}~{\rm eV}$ and $m_a\lesssim 10^{-9}~{\rm eV}$, respectively, and thus cut into the parameter range \eqref{dec_const_transp} preferred by the cosmic $\gamma$-ray transparency anomaly. Even stronger limits, $|g_{A\gamma}|\lesssim 6.3 \times 10^{-12}\, {\rm GeV}^{-1} (n_e / 10^{-5} {\rm cm}^{-3})^{1.3} (2 \mu G / B_{\rm cell})$ for $m_A\lesssim 10^{-14}~{\rm eV}$, have been obtained by exploiting high-precision measurements of quasar polarizations \cite{Payez:2012vf,Payez:2013yxa}, where $n_e$ is the electron density and $B_{\rm cell}$ is the magnetic field in the neighborhood of the quasars. But there remains still a sizeable region in ALP parameter space motivated by the above anomalies and at the same time consistent with all astrophysical constraints, as can be seen in Fig. \ref{ALP_coupling_limits}. At small masses below $10^{-14}$ eV, a part of the region of interest in ALP parameter space will be probed indirectly by the next generation of cosmic microwave background (CMB) observatories such as PIXIE \cite{Kogut:2011xw} and PRISM \cite{Andre:2013nfa} (see the region labelled ``PIXIE/PRISM" in Fig. \ref{ALP_coupling_limits}, based on the assumption of an extragalactic magnetic field $B$ of nG size; the projected sensitivity scales with the magnetic field as $B^{-1}$), because resonant photon-ALP oscillations in primordial magnetic fields may lead to observable spectral distortions of the CMB \cite{Mirizzi:2009nq,Tashiro:2013yea,Ejlli:2013uda}. A complementary part of the region of interest in parameter space will soon be probed by a pure laboratory experiment: the light-shining-through-a-wall (LSW) \cite{Redondo:2010dp} experiment Any-Light-Particle-Search II (ALPS-II) is designed to detect photon--ALP--photon oscillations in the range \cite{Bahre:2013ywa} (see the green region labelled ``ALPS-II" in Fig. \ref{ALP_coupling_limits}) \begin{equation} |g_{a\gamma}|\gtrsim 2\times 10^{-11} \ {\rm GeV}^{-1} ; \hspace{6ex} m_{a}\lesssim 10^{-4}\ {\rm eV}. \label{ALPSII} \end{equation} Further experimental opportunities covering this region in ALP parameter space will open if the International Axion Observatory (IAXO), a helioscope searching for solar axions and ALPs, is realized \cite{Armengaud:2014gea}. Its projected sensitivity is \begin{equation} |g_{a\gamma}|\gtrsim 5\times 10^{-12} \ {\rm GeV}^{-1} ; \hspace{6ex} m_{a}\lesssim 10\ {\rm meV}. \label{IAXO} \end{equation} The latter instrument has also the possibility to probe the possible coupling of the axion or further ALPs to electrons, \begin{equation} \mathcal{L} \supset \frac{\left( g_{Ae}\partial_\mu A + g_{ae}\partial_\mu a \,\right)} {2 m_e} \,\bar{e} \gamma^\mu\gamma_5 e , \end{equation} via their solar production by atomic axio-recombination, axio-deexcitation, axio-Bremsstrahlung in electron-ion or electron-electron collisions, and Compton scattering \cite{Redondo:2013wwa}. This is of considerable interest because of hints of an extra stellar cooling mechanism not accounted by the Standard Model. In fact, the white dwarf (WD) luminosity function seems to require a new energy-loss channel that can be interpreted in terms of losses due to sub-keV mass axions or ALPs, with Yukawa couplings~\cite{Isern:2008nt}, \begin{equation} |g_{Ae}|\equiv |C_{Ae}| \,m_e/f_A \sim 10^{-13} \ {\rm \ and/or\ }\ | g_{ae}|\equiv |\, C_{ae}| m_e/f_a \sim 10^{-13}, \label{wd_energy_loss} \end{equation} which are well in the range expected for an intermediate scale axion or further ALPs (if the model-dependent couplings $C_{Ae}$ and $C_{ae}$ are of order unity). The same parameter range is preferred to explain the anomalous size of the observed period decrease of the pulsating WDs G117-B15A and R548 by additional axion/ALP losses \cite{Isern:2010wz,Corsico:2012sh}. A third independent hint of anomalous stellar losses has recently been found in the red-giant branch of the globular cluster M5, which seems to be extended to larger brightness than expected within the Standard Model. A possible explanation of this observation is that the helium cores of red giants lose energy in axions or further ALPs with electron couplings of the same order as in Eq. \eqref{wd_energy_loss} \cite{Viaux:2013lha}. Very recently, two groups have found an unidentified X-ray line signal at 3.55 keV in the stacked spectrum of a number of galaxy clusters \cite{Bulbul:2014sua} and the Andromeda galaxy \cite{Boyarsky:2014jta}. It is tempting to identify this line with the expected signal from two photon decay of 7.1 keV mass ALP dark matter \cite{Higaki:2014zua,Jaeckel:2014qea,Lee:2014xua,Cicoli:2014bfa}. To match the observed X-ray flux, but allowing for the likely possibility, that the ALP dark matter makes only a fraction $x_a \equiv \rho_a/\rho_{\rm DM}$ of the total density of dark matter, the required lifetime and thus coupling of the ALP is \cite{Jaeckel:2014qea} \begin{equation} \tau_a = \frac{64 \pi}{g^2_{a\gamma} m_a^3} = x_a\times \left( 4\times 10^{27} - 4\times 10^{28}\right)\,{\rm s} \hspace{2ex} \Rightarrow\hspace{2ex} 3\times 10^{-18}\,{\rm GeV}^{-1} \left( \frac{1}{x_a}\right)^{1/2} \lesssim | g_{a\gamma} | \lesssim 10^{-12}\,{\rm GeV}^{-1} \left( \frac{10^{-10}}{x_a}\right)^{1/2}\,. \end{equation} Therefore, it is timely to have a close look onto ultraviolet extensions of the Standard Model featuring, apart from an intermediate scale axion, also further ALPs and to investigate possible correlations between the low-energy axion and ALP couplings. In fact, as we will show in Sec. \ref{sec:correlations}, such correlations inevitably occur if there are originally two (or more) Nambu-Goldstone fields coupled to $G\tilde G$. The crucial determination of the decay constants and couplings of axion-like fields from original high-scale theories to gluons, photons, and electrons will be done in Secs. \ref{axions_ad_hoc} and \ref{accions_bottom_up}. In the latter section, we construct particularly well-motivated ultraviolet completions of the Standard model featuring accidental Peccei-Quinn symmetries arising from exact discrete symmetries and deduce their low-energy parameters, as summarized in Table \ref{summary_models}. As a first example (model A.1), we present a multi-Higgs model with discrete $\ZZ_{13}\otimes \ZZ_5\otimes\ZZ_{5}^\prime$ symmetry (Sec. \ref{sec:unificaccion}), which predicts gauge coupling unification near the scale of Peccei-Quinn symmetry breakdown, as well as neutrino mass generation through the seesaw mechanism. Another example (model B1.1) features a $\ZZ_{11}\otimes\ZZ_9$ symmetry (Sec. \ref{accion_laccion_model}), has two Higgs doublets and a photophilic ALP which can fit the astrophysical hints such as the anomalous cosmic transparency to gamma-rays, and has also a seesaw mechanism. A third example (model B3) with $\ZZ_{11}\otimes\ZZ_7$ symmetry (Sec. \ref{another_singlets_model}) has similar properties as model B1.1, but with the ALP being a dark matter candidate whose decay into photons can explain the 3.55 keV line, since it has a mass of 7.1 keV generated through effective interactions suppressed by the Planck scale. Last, but not least, we present in \ref{yet_another_singlets_model} a model with $\ZZ_{11}{\otimes} \ZZ_{9}{\otimes} \ZZ_{7}$ symmetry, which has two photophilic ALPs, besides the axion. The field content and the discrete symmetry of this model is such that one of the ALPs is very light and able to match the astrophysical hints previously mentioned, while the other one has a coupling to photons and a mass as required to explain the 3.55 keV line. Finally, we summarize, conclude and give an outlook for further investigations in Sec. \ref{sec:conclusions}. \begin{table} \begin{center} \begin{tabular}{|c||c|c|c|c|c|c|c|}\hline & \multicolumn{7}{c|}{Resulting low-energy parameters} \\ \cline{2-8} \raisebox{1.5ex}{Model} & $f_A$ [GeV] & $m_A$ [eV] & $m_a$ [eV]& $|g_{A\gamma}|$ [GeV]$^{-1}$ & $|g_{a\gamma}|$ [GeV]$^{-1}$ & $|g_{A e}|$ & $|g_{a e}|$ \\ \hline A.1 & $8.3\times 10^{11}$ & $7.2\times 10^{-6}$ & $1.3\times 10^{-22}$ & $8.2\times 10^{-16}$ & $5.4\times 10^{-16}$ & $2\times 10^{-16}$ & $1.7\times 10^{-17}$ \\ \hline B1.1 & $3\times 10^{11}$ & $2\times 10^{-5}$ &$1\times 10^{-7}$ & $1.6\times 10^{-14}$ & $1.4\times 10^{-11}$ & 0 & $1\times 10^{-12}$ \\ \hline B3 & $5.5\times 10^{11}$ & $1.1\times 10^{-5}$ &$7.1\times 10^{3}$ & $3.4\times 10^{-15}$ & $1.3\times 10^{-12}$ & 0 & $0$ \\ \hline \end{tabular} \caption{\label{summary_models} Low-energy parameters for: model A.1 with $\ZZ_{13}\otimes \ZZ_5\otimes\ZZ_{5}^\prime$ symmetry (Sec. \ref{sec:unificaccion}); model B1.1 with $\ZZ_{11}\otimes\ZZ_9$ symmetry (Sec. \ref{accion_laccion_model}); and model B3 with $\ZZ_{11}\otimes\ZZ_7$ symmetry (Sec. \ref{another_singlets_model}).} \end{center} \end{table} | 14 | 3 | 1403.5760 |
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1403 | 1403.7623_arXiv.txt | The BICEP2 collaboration reports a detection of primordial cosmic microwave background (CMB) B-mode with a tensor-scalar ratio $r=0.20^{+0.07}_{-0.05}$ (68\% C.L.). However, this result has a tension with the recent Planck limit, $r<0.11$ (95\% C.L.), on constraining inflation models. In this Letter we consider an inflationary cosmology with a preceding nonsingular bounce which gives rise to observable signatures on primordial perturbations. One interesting phenomenon is that both the primordial scalar and tensor modes can have a step feature on their power spectra, which nicely cancels the tensor excess power on the CMB temperature power spectrum. By performing a global analysis, we obtain the 68\% C.L. constraints on the parameters of the model from the Planck+WP and BICEP2 data together: the jump scale $\log_{10}(k_{\rm B}/{\rm Mpc}^{-1})=-2.4\pm0.2$ and the spectrum amplitude ratio of bounce-to-inflation $r_B\equiv P_{\rm m} / A_{\rm s} = 0.71\pm0.09$. Our result reveals that the bounce inflation scenario can simultaneously explain the Planck and BICEP2 observations better than the standard $\Lambda$CDM model, and can be verified by the future CMB polarization measurements. | 14 | 3 | 1403.7623 |
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1403 | 1403.7309_arXiv.txt | We propose a new method for Point Spread Function (PSF) correction in weak gravitational lensing shear analysis using an artificial image with the same ellipticity as the lensed image. This avoids the systematic error associated with the approximation in PSF correction used in previous approaches. We test the new method with simulated objects which have Gaussian or Cersic profiles smeared by a Gaussian PSF, and confirm that there is no systematic error. | Weak gravitational lensing has been widely recognized as a unique and very powerful method for studying not only the mass distribution of the universe but also cosmological parameters (Mellier 1999, Schneider 2006, Munshi et al. 2008). One of the most interesting aspects of this field is to accurately measure the cosmic shear which is coherent distortion of background galaxies induced by large scale structure of the universe, because this effect depends on the evolution of the structure which is influenced by the nature of dark energy. There have been some detections of cosmic shear (Bacon et al 2000; Maoli et al 2001; Refregier et al 2002; Bacon et al 2003; Hamana et al 2003, Casertano et al 2003; van Waerbeke et al 2005; Massey et al 2005; Hoekstra et al 2006). However, since the distortion is very weak and there are many sources of noise, a very accurate measurement scheme is needed for the shapes of the huge number of background galaxies as well as a sophisticated analysis scheme for the measured shear which avoids systematic errors as much as possible. At the moment there are several plans for wide field surveys with a large enough number of galaxies to reduce the statistical error. This means the systematic errors in shear analysis methods become larger than the statistical errors. In fact, even current survey plans require systematic errors less than 1\% error (Hyper Suprime-Cam http://www.naoj.org/Projects/HSC/HSCProject.html, Dark Energy Survey http://www.darkenergysurvey.org/, Euclid http://sci.esa.int/euclid etc.), and 0.1\% error for future planned surveys (lsst http://www.lsst.org). There is another important point in the analysis for such wide survey data. For example, in the plan for the HSC wide survey, a few hundreds of millions of galaxies will be measured, therefore slow methods are not realistic even if the methods have a high accuracy. Thus it is essential to develop a fast shear analysis scheme which is free from systematic errors. There have been many studies in this direction (Kaiser et al 1995, Bernstein \& Jarvis 2002; Refregier 2003; Kuijken et al. 2006; Miller et al. 2007; Kitching et al. 2008; Melchior 2011), which have been tested using simulated data (Heymans et al 2006, Massey et al 2007, Bridle et al 2010 and Kitching et al 2012). We also have developed new analysis methods, including the E-HOLICs method (E-HOLICs part1 Okura and Futamase 2011; part2 Okura and Futamase 2012; part3 Okura and Futamase 2013). The E-HOLICs method avoids systematic error caused by the approximation in the weight function by adopting an appropriate elliptical weight function for shape measurement. However, the E-HOLICs method uses an approximation for PSF correction similar to other approaches. Therefore, previous approaches including E-HOLICs, cannot correct the PSF effect in some conditions (e.g. large PSF or high elliptical PSF). In this paper we concentrate on PSF correction and propose a new method which is free from the approximations used in previous moments methods and thus free from systematic error associated with PSF correction. The idea is to produce an artificial image with the true ellipticity as the result of re-smearing the PSF smeared image of the lensed image with true ellipticity. We refer to this as the Ellipticity of Re-smeared Artificial images method (ERA). This paper is organized as follows: In section 2, we explain our notation and definitions used in this paper. There we also describe the zero plane used in the ERA method. We then present PSF smearing and PSF correction as used in the ERA method in section 3. In section 4, we test this method using simple test images. Finally we summarize the method and give some comments. | We have developed the ERA method with a possible new PSF correction method for weak gravitational lensing shear analysis. The idea is to construct an artificial image with the same ellipticity as the lensed image by re-smearing the observed image. This approach avoids the approximations in PSF corrections of previous moments method (i.e. Eq.\ref{eq:KSBAP}), therefore there is no systematic error from PSF correction if we choose an appropriate function for re-smearing. Then, we tested the method with simple simulated images. The results of the simulation are as follows. The deconvolution method cannot estimate ellipticity correctly, because deconvolution is not perfect when smearing using a large or high elliptical PSF. The KSB method has systematic errors which cause overestimation in the standard situation (in the simulation, it is about 2-3\%, but it depends on the ellipticity of the image and the size of the PSF, etc.), and is not able to correct the smearing effect with a high elliptical PSF. However, Methods A and B are both able to estimate the ellipticity correctly if the PSF has standard ellipticity. We also confirmed that re-smearing the deconvolved image (method A3) and re-smearing the smeared image (method B2) have no systematic error in PSF correction even if the size and the ellipticity of PSF are large. Although the tests performed here are not entirely realistic, the results are very encouraging. Further studies of systematic errors based on realistic data, (e.g. pixel noise, pixelization etc.) are needed. In particular, pixel noise has the potential to make large systematic errors (see E-HOLICs part3). For example, in measuring cosmic shear, galaxies in high redshift bins are relatively faint and have lower signal to noise ratio, and thus they suffer from more systematic bias due to pixel noise than those at lower redshifts. Thus, the measured shear for high-z galaxies would be underestimated. It is also important to find a more appropriate profile for the re-smearing function and to develop more effective iteration schemes which will result in faster analyses. Future planed surveys such as EUCLID and LSST will treat an enormous number of galaxies and thus not only low systematic errors but also fast analyses are essential. These problems will be approached in future studies. | 14 | 3 | 1403.7309 |
1403 | 1403.0003_arXiv.txt | We study the effects of radio jets on galaxies in their vicinity (satellites) and the role of satellites in triggering radio-loud active galactic nuclei (AGNs). The study compares the aggregate properties of satellites of a sample of 7,220 radio AGNs at $z < 0.3$ (identified by Best \& Heckman 2012 from the SDSS and NVSS+FIRST surveys) to the satellites of a control sample of radio-quiet galaxies, which are matched in redshift, color, luminosity, and axis ratio, as well as by environment type: field galaxies, cluster members and brightest cluster galaxies (BCGs). Remarkably, we find that radio AGNs exhibit on average a 50\% excess (17$\sigma$ significance) in the number of satellites within 100 kpc even though the cluster membership was controlled for (e.g., radio BCGs have more satellites than radio-quiet BCGs, etc.). Satellite excess is not confirmed for high-excitation sources, which are only 2\% of radio AGN. Extra satellites may be responsible for raising the probability for hot gas AGN accretion via tidal effects or may otherwise enhance the intensity or duration of the radio-emitting phase. Furthermore, we find that the incidence of radio AGNs among potential hosts (massive ellipticals) is similar for field galaxies and for non-BCG cluster members, suggesting that AGN fueling depends primarily on conditions in the host halo rather than the parent, cluster halo. Regarding feedback, we find that radio AGNs, either high or low excitation, have no detectable effect on star formation in their satellites, as neither induced star formation nor star formation quenching is present in more than $\sim 1\%$ of radio AGN. | \label{sec:intro} Active galactic nuclei (AGN) are believed to be powered by two fundamentally different modes of accretion. Optically selected and luminous radio-loud AGN (R-AGN) are powered by the radiatively-efficient accretion of gas onto a central supermassive black hole (SMBH) at a rate of one to ten percent of the Eddington limit. In contrast the majority of R-AGN, which have a low radio luminosity, accrete gas in a radiatively inefficient manner at a rate below one percent of the Eddington limit \citep{2012MNRAS.421.1569B}. The mode of accretion may also be related to galaxy morphology. While AGN in general are hosted by galaxies that span a range of morphological types, the hosts of R-AGN are most often massive ellipticals (e.g., Ekers \& Ekers 1973, Best at al. 2005b, Kauffmann et al. 2008). This difference in R-AGN accretion processes could be a result of differences in the fueling mechanism. Optical and luminous radio AGN may be triggered by mergers and one-on-one interactions (Heckman et al. 1986, Barnes \& Hernquist 1996) which could supply the large quantities of cold gas required to maintain a high accretion rate. Low luminosity R-AGN may be triggered by the accretion of hot gas from the surrounding halo (Fabian 1994, Allen et al. 2006), although there could be additional triggering mechanisms \citep{2013MNRAS.430..638S}. The key to understanding these triggering mechanisms therefore lies in the kpc to Mpc-scale environments of R-AGN, which is populated by satellite galaxies. Many studies have examined the relationship between R-AGN incidence and the environment. \cite{2007MNRAS.379..894B} found that R-AGN are more frequently found in central group and cluster galaxies when compared to galaxies of similar stellar mass. \cite{2004MNRAS.351...70B} studied a small sample of 91 R-AGN and found that the fraction of R-AGN shows little dependence on the Mpc-scale local galaxy density, but is dependent on the number of galaxies in the group in which the R-AGN is located, while the slightly larger sample of 212 R-AGN of \cite{2006ApJ...650..717R} showed a slight increase in R-AGN incidence with density at the Mpc scale. \cite{2005MNRAS.362...25B} and \cite{2013MNRAS.430..638S} both found a strong dependence of R-AGN activity on the stellar mass of the host galaxy, but when this is accounted for, R-AGN activity still has a dependence on density at the Mpc scale. \cite{2008MNRAS.384..953K} found that the environments of R-AGN at the $\sim$100 kpc scale are roughly twice as dense as those of radio-quiet AGN or radio-quiet galaxies of similar mass. In contrast to these studies, \cite{2013arXiv1305.2673W} found that at the $\sim$100 kpc scale only powerful R-AGN have higher clustering than radio-quiet galaxies of the same mass, and that for scales greater than about 160 kpc the clustering of R-AGN is similar to that of radio-quiet galaxies. The aforementioned studies find an increase in the R-AGN incidence rate in denser, cluster environments. This should not be surprising considering that R-AGN likely feed on the hot gas that is plentiful in clusters. However, it must be acknowledged that neither all ellipticals or brightest cluster galaxies (BCGs) are radio loud, nor are all massive field ellipticals radio quiet. The approach taken in the current study is therefore different as it separates the general environment (BCG, cluster member, or field galaxy) from the small-scale environment ($<100$ kpc). We compare the small-scale environments (i.e. the satellite populations) of a statistically large sample of R-AGN and a matching control sample of radio-quiet galaxies. Control galaxies are chosen not only to match the mass, type, and star formation history of an R-AGN but also the large scale environment it is found in, such that the controls of field R-AGN are also field galaxies, the controls of cluster member R-AGN are also cluster members, and the controls of BCG R-AGN are also BCGs. This careful matching allows us to compare the satellite populations of R-AGN and otherwise similar radio-quiet galaxies to determine why only a subset of galaxies with similar environments become R-AGN. In our study of satellite populations we will separately focus on radio AGN with high accretion rates. The two modes of R-AGN accretion can be distinguished with emission line ratios \citep{1994ASPC...54..201L}. \cite{2010A&A...509A...6B} defined an `excitation index' composed of four optical emission line ratios which can be used to separate R-AGN with a low accretion rate, which are known as low excitation radio galaxies (LERGs) from those with a high accretion rate, known as high excitation radio galaxies (HERGs). At most radio luminosities HERGs comprise only a few percent of the total R-AGN population, although they become more common at the highest radio luminosities \citep{2012MNRAS.421.1569B}. Many studies employ the Fanaroff-Riley classification, which classifies R-AGN based on their radio morphology \citep{1974MNRAS.167P..31F}. Luminous R-AGN, which tend to be edge brightened, are known as Fanaroff-Riley type 2 (FR2) sources, while core-brightened lower-luminosity sources are referred to as Fanaroff-Riley type 1 (FR1) sources. Although LERGs and HERGs often exhibit FR1 and FR2 morphology, respectively, \cite{1994ASPC...54..201L} demonstrated that a number of FR2 sources have low-excitation spectra (i.e. are LERGs). Since the LERG \& HERG classifications are more closely related to accretion mode and do not require high resolution radio maps, we adopt their use in our study, as determined by \cite{2012MNRAS.421.1569B}. In addition to exploring the role of satellites in R-AGN triggering, the second major goal of this study is to understand if and to what extent satellites can be affected through feedback from the powerful jets that are the hallmark of R-AGN. These highly collimated jets originate in the nucleus of the AGN at relativistic speeds and may extend for kiloparsecs or even megaparsecs beyond the host galaxy (Tremblay et al. 2010, Schoenmakers et al. 2000). Radio AGN deposit most of their energy into the interstellar or intergalactic medium kinetically via their high-velocity jets (see Fabian 2012 for a recent review). Such interactions may be responsible for quenching star formation in the host galaxy, and there is some evidence that R-AGN may lead to quenching of star formation in their satellite galaxies. \cite{2011MNRAS.413.2815S} found that satellite galaxies in the projected jet paths of FR2 sources are redder than satellites outside the jet path, but no such trend was found for the satellites of FR1 sources. This suggests that FR2 jets tend to quench star formation in satellite galaxies while FR1 jets do not. We will revisit these results in our study by using a much larger sample of R-AGN and by contrasting entire satellite populations of R-AGN (regardless of proximity to the jet) to those of the radio-quiet control sample. Somewhat paradoxically, radio jets are also considered as the mechanism behind possible \textit{positive} feedback, both in the host galaxy and in adjacent satellite galaxies. The idea is that interactions with jets may induce star formation by driving shocks into dense clouds which then collapse \citep{2008MNRAS.389.1750A}. Such AGN-driven star formation may have been important in building up the massive spheroid in the early phases of galaxy formation \citep{2009ApJ...700..262S}. The radio jets of both HERGs and LERGs are believed to be capable of producing positive feedback outside of the host as well. The powerful radio jets associated with HERGs pierce external clouds in the IGM, but triggered star formation may proceed in their slowly expanding radio lobes (De Young 1981, Bicknell et al. 2000). Since the jets of LERGs are generally less powerful, they may be capable of inducing star formation even in head-on collisions \citep{2004IAUS..222..485V}. A few candidates of such jet-induced star formation in the vicinity of R-AGN have been observed. The star forming region `09.6' in the eastern lobe of the nearby FR 2 source 3C 285 was first observed by \cite{1993ApJ...414..563V}. \emph{Chandra} observations confirm that this region is indeed experiencing a starburst phase \citep{2007ApJ...662..166H}, while a satellite galaxy of the FR2 source PKS2250--41, which is not currently in the projected jet path, also shows evidence of jet-induced star formation \citep{2008MNRAS.386.1797I}. Induced star formation as a result of interactions with weak R-AGN has also been observed: one example is the LERG Centaurus A. \cite{2012MNRAS.421.1603C} recently showed that the youngest stars in the inner filament of Centaurus A are only a few Myrs old and are probably the result of the shock induced collapse of a molecular cloud. Downstream from the inner star-forming filament, the radio jet interacts with an H I cloud and produces another filament of recently-formed stars \citep{1998ApJ...502..245G}. Minkowski's Object (MO) is another example of LERG-induced star formation. This star forming region is 15 kpc away from its host R-AGN NGC 541. Strong UV and H$\alpha$ emission are indicative of the starburst nature of this object \citep{2004IAUS..222..485V}. Simulations by \cite{2004ApJ...604...74F} have reproduced the observational characteristics of MO, including the star formation rate (SFR) of 0.3 M$_\sun$yr$^{-1}$. While it is possible that MO is an example of star formation being reignited in an already star-forming galaxy, this seems unlikely. Although \cite{2006ApJ...647.1040C} could not rule out an underlying old stellar population in MO, an HI cloud downstream from MO was observed, which suggests that unlike Centaurus A the neutral hydrogen cooled out of a warm and clumpy IGM and then collapsed. While there is observational evidence that in some individual cases intense star formation may have been triggered by R-AGN jets, it is not clear how common this phenomenon is and whether it can be positively stated that the star formation is indeed the result of jet interactions. Furthermore, it is not clear if this induced star formation leads to the emergence of pristine, new satellites or whether it is an enhancement in already existing, star-forming satellites. As in the case of star formation quenching, we will approach this question by studying the colors of satellites of R-AGN and a matched control sample of radio quiet galaxies. In both the study of quenching and the induction of star formation in satellites we will pay special attention to HERGs in which either of these processes may be more pronounced or more common. The layout of our paper is as follows. In Section 2, samples are presented, and in Section 3 we describe our method for selecting the control sample of radio quiet galaxies as well as generating distributions of satellite properties. Our results are presented in Section 4, while in Sections 5 and 6 we discuss the implications of our findings. In this work the cosmological parameters adopted are $\Omega_m\,=\,0.27$, $\Omega_\Lambda\,=\,0.73$, and $H_0\,=\,71\, \textrm{km}\,\textrm{s}^{-1}\,\textrm{Mpc}^{-1}$. | \subsection{R-AGN fraction} \label{fraction} Previous studies (e.g. Best et al. 2005b, Van Velzen et al. 2012) have shown that the fraction of galaxies that are R-AGN (the R-AGN incidence rate) is an increasing function of the stellar mass of the host. We revisit these results by including two new aspects in the analysis. First, we obtain incidence rates separately for R-AGN in each type of environment (BCGs, cluster members, and field). Second, we consider the incidence rate to be the R-AGN fraction in the {\it eligible} galaxy population, not among all galaxies in a given mass (i.e., luminosity) bin. The idea is that one is primarily interested in R-AGN incidence among the galaxies that could (and perhaps did or will) host an R-AGN, so we define the eligible population to occupy the same part of the parameter space in redshift, magnitude, color, and axis ratio as galaxies that are {\it current} R-AGN hosts. Each R-AGN in the sample defines a position in the parameter space around which we determine the R-AGN incidence. We take the vicinity to be $R<0.2$ (Eq. 1) and count the R-AGN used to probe the parameter space as 0.5, which leads to less bias than either not counting it or giving it a full count. These fractions are then binned by $M_r$, and the median of each bin is shown in Fig.~\ref{ragnfrac}. Bins with fewer than five R-AGN were omitted. Figure~\ref{ragnfrac} shows that while the fraction of eligible galaxies that currently host an R-AGN is generally low, with an overall average of 6\%, the incidence increases from close to zero to a high near $\sim$20\% for the most luminous galaxies. The incidence is generally higher in BCGs, as might be expected, while the incidence in cluster members is only slightly greater than that of field galaxies. This suggests that fueling is primarily a \textit{locally} determined process and that being in a more massive cluster halo confers modest additional benefits, unless a galaxy sits in the center of the cluster halo, as is the case for BCGs. \begin{figure}[h!] \includegraphics[width=3.25in]{f7.eps} \caption{Fraction of galaxies that are R-AGN, for different environments. The fraction is calculated as the median of individual incidence estimates. These estimates are the ratio of galaxies that are R-AGN among galaxies that have similar properties (redshift, magnitude, color, axis ratio), i.e. galaxies eligible to be R-AGN. The R-AGN fraction is higher in BCG's, with cluster members having a slightly higher incidence than field galaxies. } \label{ragnfrac} \end{figure} \cite{2005MNRAS.362...25B} show a similar relation for R-AGN incidence but as a function of stellar mass in panel a of their Fig. 2. The mean NVSS radio luminosity of our sample is $2\times10^{24}$ W Hz$^{-1}$ at 1.4 GHz, so our results are best compared to the middle curve of panel a. The R-AGN incidence of \cite{2005MNRAS.362...25B} is expressed as a fraction of all galaxies from SDSS DR2 without regard \begin{verbatim} \end{verbatim} for the galaxies' likelihood of hosting R-AGN, while we have only considered galaxies from the same type of environment that are likely to host R-AGN. Even so, considering that most R-AGN are field galaxies and the majority of massive galaxies are eligible as R-AGN, the results for the field galaxies and cluster members in our sample follow a trend similar to that of \cite{2005MNRAS.362...25B} for $10^{24}$ W Hz$^{-1}$ at 1.4 GHz. \cite{2007MNRAS.379..894B} found that at a stellar mass of $\sim5\times10^{11}M_\sun$ the R-AGN incidence in BCGs is only slightly greater than that of other galaxies, although for masses lower than $\sim\times10^{11}M_\sun$ the R-AGN incidence in BCGs is over an order of magnitude greater than that of other galaxies. This is in agreement with what we find in Fig~\ref{ragnfrac}. \cite{2007MNRAS.379..894B} also found that except for group and cluster galaxies within $0.2r_{200}$ of the center of the system, the R-AGN incidence among non-BCG group and cluster galaxies is similar to that of field galaxies. This was based on a comparison of the R-AGN sample defined by \cite{2005MNRAS.362....9B} to all galaxies in DR4 within the same redshift range as the \cite{2005MNRAS.362....9B} R-AGN sample, regardless of these galaxies' likelihood of hosting an R-AGN. These results are in agreement with our finding that field and cluster member R-AGN have similar incidence rates. \subsection{The role of satellite populations in R-AGN triggering} We have found that in all environments R-AGN have more nearby satellites than radio-quiet galaxies, but our results are not a simple confirmation of what other studies have found regarding the R-AGN clustering. \cite{2008MNRAS.391.1674W}, \cite{2009MNRAS.393..377M}, \cite{2010MNRAS.407.1078D}, and \cite{2008MNRAS.384..953K} find that R-AGN are more clustered than optical AGN or radio-quiet galaxies. The control samples of \cite{2009MNRAS.393..377M}, \cite{2010MNRAS.407.1078D}, and \cite{2008MNRAS.384..953K} were even chosen to match the stellar masses of the R-AGN. However, the results of these studies can be interpreted simply to mean that R-AGN prefer denser, cluster environments or, judging by the results in Section~\ref{fraction} that show that R-AGN incidence among eligible cluster members and field galaxies is the same, that the type of galaxies that typically host R-AGN (massive ellipticals) are more clustered, a well known result \citep{1980ApJ...236..351D}. By including cluster membership in the selection criteria of control galaxies we have added crucial new information: the satellite population of R-AGN is richer, whether the R-AGN is a field galaxy, cluster member, or the very BCG. It is now widely believed that LERGs are triggered by the cooling of small amounts of gas from the hot halos in which they reside \citep{2012MNRAS.421.1569B}. Since $\sim$80\% of the R-AGN in our sample are LERGs, our finding that R-AGN in all environments have an excess of satellites suggests that the availability of hot halo gas is not the only prerequisite for triggering activity in LERGs. For example, the dark matter halo masses of the BCGs in our R-AGN and control samples are similar, which would suggest that they have similar quantities of hot halo gas available. The difference is that BCGs that host R-AGN have more nearby satellites, which therefore probably play some role in triggering. This argument applies to field galaxies and cluster members as well, since the matched pairs of R-AGN and control galaxies have similar halo masses and therefore presumably similar amounts of hot halo gas. Again, the difference between R-AGN and radio-quiet galaxies in these environments is the number of nearby satellites. Results in Section~\ref{results} suggest that the triggering is facilitated by satellites with a wide range of masses and at various projected distances from the host. \subsection{Dark matter halo bias} While it is tempting to interpret the excess of satellites as related to the presence of R-AGN, we must explore the possibility that even if the control sample is matched in stellar mass, it may be systematically different in halo mass. If the abundance of satellites follows the halo mass rather than the stellar mass, the excess of satellites may result from a mismatch in halo masses. We have used the dark matter halo catalog of \cite{2007ApJ...671..153Y} to investigate this possibility. This catalog presents two dark matter halo masses for each galaxy: one calculated based on the ranking of group characteristic luminosity, and the other based on the ranking of group stellar mass: we used those based on the group stellar mass. This catalog uses the fourth SDSS data release (DR4) and extends out to $z=0.2$, so only about 1,400 matched pairs of R-AGN and control galaxies were found. The mean of the log of R-AGN masses is 0.1 dex greater than that of control galaxies. Is this modest systematic offset in halo masses sufficient to explain a factor of 1.5 difference in satellite population? Figure~\ref{halo} shows the number of satellites as a function of halo mass for the field control sample in the redshift range $0.15<z<0.2$. There is a trend that galaxies in more massive halos indeed have more satellites. However, this rise is not so steep. In order to bring the number of satellites from 1.5 (the number observed on average around control sample galaxies) to 2.1 (the number for R-AGN hosts) the mismatch in halo masses would have to be 0.96 dex (a factor of 9) while in reality it is only 0.1 dex (25\%). In other words, the excess of R-AGN satellites cannot simply be a result of R-AGN having dark matter halos that are systematically more massive than those of control galaxies. \begin{figure}[t!] \centering \includegraphics[width=3.25in]{f8.eps} \caption{ Distribution of satellite counts as a function of halo mass for the field sample of control galaxies in the redshift range $0.15<z<0.2$. The solid line is the least-squares fit to the distribution, and the dashed horizontal lines show the average number of satellites for R-AGN (2.1) and control galaxies (1.5). For the excess R-AGN satellites to be a result of mismatched halo masses, a difference of 0.96 dex (dotted vertical lines) between R-AGN and controls is necessary. Since the actual difference is only 0.1 dex, R-AGN must genuinely have more satellites within 100 kpc than radio-quiet galaxies. } \label{halo} \end{figure} \subsection{The effects of R-AGN feedback on the satellite population} \label{feedback} In addition to the triggering of R-AGN, we also wish to examine the satellite population with the goal of understanding the effects of R-AGN interactions (presumably that of the jet) with satellites. In the discussion that follows, we use color distributions of satellites to examine two scenarios. In the first, R-AGN jets frequently quench star formation in their satellites, while in the second interactions with R-AGN jets lead to induced star formation. We study the effects of R-AGN feedback by comparing the difference between $g-r$ color distributions of R-AGN and control satellites. Let us first examine how feedback would modify the difference in color distribution between R-AGN and control galaxies in the absence of an overall excess of satellites around R-AGN. If neither quenching nor induced star formation are present, the difference would be zero, as in panel a of Fig.~\ref{same}. If interactions commonly lead to quenching, there would be a deficit of blue R-AGN satellites accompanied by an increase in the number of red R-AGN satellites relative to the control distribution, as shown in panel b. \begin{figure*}[t!] \centering \includegraphics[width=6.1in]{f9.eps} \caption{ Schematic plots illustrating the effects of R-AGN feedback on the difference color distribution between R-AGN and control galaxies, assuming no net difference in the number of satellites. Panel a shows the case where neither quenching nor induced star formation as a result of jet interactions are common in satellites. Panel b shows the result of star formation being quenched in many R-AGN satellites as a result of jet interactions, leading to a deficit of bluer satellites that is matched by an excess of red ones, while panel c shows the result of radio jets inducing star formation in gas clouds that would not otherwise form stars, producing a net excess of blue satellites. Panel d shows the result of radio jets enhancing star formation in satellites that are already actively forming stars.} \label{same} \end{figure*} \begin{figure*}[t!] \centering \includegraphics[width=6.1in]{f10.eps} \caption{ Schematic plots illustrating the effects of R-AGN feedback on the difference color distribution between R-AGN and control galaxies, for scenario in which R-AGN have more red satellites than control galaxies. Panel a shows the case where neither quenching nor induced star formation as a result of jet interactions are common in satellites. Panel b shows the result of star formation being quenched in many R-AGN satellites as a result of jet interactions, leading to a deficit of bluer satellites that is matched by an excess of red ones, while panel c shows the result of radio jets inducing star formation in gas clouds that would not otherwise form stars, producing a net excess of blue satellites. Panel d shows the result of radio jets enhancing star formation in satellites that are already actively forming stars.} \label{radiomore} \end{figure*} Next, we consider two scenarios of induced star formation. In the first, radio jets induce star formation in gas clouds that would not otherwise be detected as galaxies, i.e. induced star formation forms new galaxies. Induced star formation would cause R-AGN to have an excess of blue satellites, as shown in panel c. Faint satellite galaxies that increase in luminosity as a result of induced star formation would also lead to an excess of blue satellites. In the other scenario, interactions with R-AGN jets may enhance the SFR in satellites that already form stars, making them bluer (panel d). In this scenario there is no net excess, but some moderately blue galaxies become bluer, causing a deficit at moderate value colors and an excess at very blue colors. However, from Section~\ref{results} we know that there is an intrinsic excess of satellites around R-AGN, so the scenarios described above need to be modified to take that into account, as shown in Figure~\ref{radiomore}. As we have seen, this excess is primarily red simply because most satellites are red, so we obtain modified scenarios by adding a red bump to the scenarios described above. Now these modified scenarios can be compared to the observations. \begin{figure*}[t!] \centering \includegraphics[width=4.5in]{f11.eps} \caption{Plots of the difference between the color distributions of R-AGN and control satellites: panels a-c are for the three environment types and panel d is for HERGs. The dotted vertical line at $g-r=0.7$ marks the split between blue and red satellites. Panels a-c resemble panel a of Fig~\ref{radiomore}, indicating that R-AGN neither quench nor induce star formation in their satellites. The point labelled MO in panel b shows what would be 5\% occurrence of satellites with the color of Minkowski's Object. The actual value is much closer to zero, which suggests that R-AGN rarely induce star formation in external gas clouds. Quenching as a result of jet interactions is also uncommon.} \label{minus} \end{figure*} Figure~\ref{minus} shows the difference plots for the three environment types as well as the HERG sample. The vertical dotted line at $g-r=0.7$ indicates the split between blue and red satellites such that these plots can be easily compared to Fig.~\ref{radiomore}. Panels a$-$c of Fig.~\ref{minus} resemble panel a of Fig.~\ref{radiomore}, with no blue deficit accompanying the red bump. This indicates that in any environment, LERGs play no role in quenching star formation in satellites. \cite{2011MNRAS.413.2815S} examined whether R-AGN quench star formation in satellite galaxies that lie in the projected path of the radio jets. These authors manually examined the satellites of a small sample of 21 FR1 and 72 FR2 sources and classified them, based on their spatial relation to the radio jets as lying inside or outside the jet path. It was found that the $u-r$ color distribution of satellites in the path of FR1 radio jets is similar to that of satellites outside the path while satellites in the jet path of FR2s are redder than satellites outside the path. \cite{2011MNRAS.413.2815S} therefore concluded that only the high-power FR2 galaxies quench star formation in satellites. If we assume that most of the HERGs in our sample present FR2 morphology, then we may compare our findings for HERG satellites to the results of \cite{2011MNRAS.413.2815S}. Panel d of Fig.~\ref{minus} shows that, like R-AGN in all environments, HERGs have an excess of red satellites. This excess is not accompanied by a deficit of blue R-AGN satellites, as the sum of the bins bluer than $g-r=0.7$ is 0.04 ($\pm0.07$). Our results indicate that high accretion R-AGN, like low accretion ones, at best only rarely quench star formation in nearby satellites. This non-detection of quenching is in apparent disagreement with the results of \cite{2011MNRAS.413.2815S}, but note that we have considered quenching on entire satellite populations, while \cite{2011MNRAS.413.2815S} only investigated satellites in the path of the jet. If R-AGN frequently induced star formation in their satellites, we would expect an excess of blue satellites between $0\lesssim g-r\lesssim 0.6$ in Fig.~\ref{minus}. Such an excess is not seen for R-AGN in any environment, nor for HERGs. The field sample places a limit on the incidence of star formation induction to no more than 1\%. This provides evidence that star formation is rarely induced by interactions with R-AGN jets. We have included in panel b of Fig.~\ref{minus} a point showing what would be 5\% occurrence of satellites with the color of Minkowski's Object. The actual value is consistent with zero, which indicates that R-AGN rarely induce star formation in external gas clouds, nor do they lead to any significant enhancement in existing gas-rich satellites. | 14 | 3 | 1403.0003 |
1403 | 1403.4620_arXiv.txt | \subsection{A Brief History} It has been nearly a century since Hubble (1925) decisively demonstrated that the Andromeda nebula is a vast island universe of stars similar to our own Milky Way galaxy. His discovery soon thereafter of an expanding cosmos filled with galaxies opened up a vista of an immense and seemingly serene universe. However, this picture of serenity was deceptive. As we write this review, it is the 50 year anniversary of Maarten Schmidt's (1963) realization that the radio source 3C273 was associated with an optically unresolved object with the (then) enormous redshift of 0.158. Subsequent identifications of a whole population of quasi-stellar radio sources at even higher redshifts were soon followed by the discovery of a more numerous population of otherwise-similar but radio-quiet quasi-stellar galaxies (Sandage 1965). [It is fascinating in hindsight that the discovery of the cosmic X-ray background (Giacconi et~al. 1962) and of the Kerr metric (Kerr 1963) occurred at the same time.] The concept of the Violent Universe was born. It is beyond the scope of this article to review the twists and turns that finally led to the consensus that all these quasi-stellar objects (QSOs) were the extremely luminous active nuclei of distant galaxies. Qualitatively similar objects had been found much earlier in the nuclei of some relatively nearby galaxies by Seyfert (1943). It is sobering that the significance of this discovery was long unrecognized: the first investigation into the nature of Seyfert galaxies was published 16 years later by Burbidge et~al.\ (1959) who concluded that NGC 1068 was explosively ejecting ionized gas from its nucleus. Likewise, although other radio sources had been identified with distant galaxies before the discovery of QSOs -- e.g.\ Cygnus A by Baade \& Minkowski (1954) -- there was great confusion as to the origin of this very large-scale radio-emission which had no obvious connection to the galaxy nucleus. For most of the past five decades the communities that studied galaxies and active galactic nuclei (AGNs) remained largely disconnected. AGNs were studied primarily as laboratories in which to probe exotic high-energy processes. There was some effort to understand the role that the environment might play in triggering or fueling the AGN -- for example, see the ancient review in this journal by Balick \& Heckman (1982) -- but there was almost no idea that AGNs played any significant role in the evolution of typical galaxies. Of course, things could hardly be more different today. The notion of the co-evolution of galaxies and AGN has become inextricably engrained in our current cosmogony. Indeed, our review represents the third article in this journal in consecutive years that deals with some aspect of this co-evolution (Fabian 2012, Kormendy \& Ho 2013). The reasons for this change are easy to see. First came the realization that powerful AGNs (as represented by QSOs) were only the tip of the iceberg. Extensive surveys in the radio, optical, and X-ray domains revealed local populations of Seyfert and radio galaxies and established that signs of lower-level activity were commonplace in the nuclei of early-type galaxies. This strongly suggested that the AGN phenomenon -- rather than being simply a rare spectacle -- was a part of the lifecycle of typical galaxies. The demography and physical properties of these low-power AGN has been comprehensively reviewed in this journal by Ho (2008). A second reason followed from the documentation of the overall build-up over most of cosmic time of the populations of galaxies (as traced via star-formation) and of supermassive black holes (SMBHs, whose growth is traced by AGN). The evolution of the two populations is strikingly similar: a steep rise in both the star-formation rate (SFR) and SMBH growth rate by about a factor of 10 from redshift z = 0 to 1, a broad maximum in both rates at $z \sim$ 2 to 3 and then a relatively steep decline at higher redshifts (see Fig.~\ref{shankarhist} -- Shankar et~al.\ 2009 and references therein). For at least the last $\sim$11 Gyr of cosmic history the ratio of these two growth rates has remained roughly constant with a value of-order 10$^3$. Thus, at least in a volume-averaged sense, the growth of galaxies and SMBHs has been synchronized somehow. \begin{figure}[!t] \begin{center} \psfig{file=Shankar_history.ps,width=8.8cm,clip=} \end{center} \caption{\label{shankarhist} The cosmic history of black hole growth and stellar mass growth. The average black hole accretion rate is compared to the SFR as a function of redshift, where the latter is given by Hopkins \& Beacom (2006) and Fardal et~al.\ (2007), scaled by the factor 0.8$\times$10$^{-3}$. The shaded grey area shows the 3$\sigma$ uncertainty region from Hopkins \& Beacom (2006). Figure from Shankar et al.\ (2009). } \end{figure} Finally, and even more remarkably, observations over the past decade have not only established that SMBH exist in the nuclei of (probably) all galactic bulges, they have shown that the properties of these present-day SMBHs and the galaxies in which they live are linked on a galaxy-by-galaxy-basis. This impressive fossil record of co-evolution has been exhaustively reviewed by Kormendy \& Ho (2013). In addition to the accumulating observational evidence for the co-evolution of galaxies and SMBHs, a considerable theoretical motivation to invoke this linkage has developed as well. Without some form of feedback from AGN, neither current semi-analytic models nor numerical simulations can successfully reproduce the properties of massive galaxies. This fascinating subject has been reviewed in this journal by Fabian (2012). In conclusion, while the evidence to date remains indirect, it is hard not to infer that the cosmic evolution of galaxies and of SMBHs have seemingly been driven by a suite of inter-linked physical processes. \subsection{Our Perspective: Large Surveys of the Contemporary Universe} Our goal here is to review the evidence for the co-evolution of galaxies and SMBHs as derived from large surveys of the contemporary (low redshift) universe. The majority of the growth of SMBH and of the stellar components in galaxies occurred between redshifts of roughly 0.5 and 2, and the present-day growth rates of both populations are over an order-of-magnitude smaller than during the peak epoch. A reasonable question is then whether one can learn anything very useful by studying the contemporary universe. We hope that by the end of the review the reader will agree with us that one can actually learn a great deal. We believe that there are a number of reasons why this is the case. The contemporary universe contains a rich fossil record that (once successfully decoded) reveals the processes that produced this record. Moreover, the basic physical processes that operated during the peak epoch of SMBH/galaxy growth are still in place and can be studied in greater detail due to their relative proximity. However, the primary reason to study the contemporary universe is that it is the only place where has it been possible to carry out the surveys of galaxies and SMBHs that are of both sufficiently large size and whose data are of sufficiently high quality to have finally allowed us to be able to fully explore the complex inter-relationship between these two populations in a statistically robust way. This has led to a much clearer picture of both the fossil record of galaxy/SMBH co-evolution and of the processes by which the co-evolution continues to play out. The particular importance of high-quality optical spectroscopy cannot be over-stated: these data form the interpretational backbone of the whole structure defined by extensive multi-band imaging and photometric surveys. Of most relevance to the specific topic of our review, studies of the lower-luminosity objects that dominate the AGN population in the local universe benefitted immensely from galaxy spectra obtained as part of the main galaxy sample of the Sloan Digital Sky Survey (SDSS -- Strauss et~al.\ 2002). The high quality of the spectra enabled them to be used to characterize both the AGN and the stellar populations in the host galaxies. Finally, the uniformity and completeness of the SDSS main galaxy sample rendered it ideal for statistical studies of the multi-parameter characteristics of the galaxy population in the contemporary universe, including their AGN properties. The uniformity and wide sky coverage of the 2dF Galaxy Redshift Survey (2dFGRS; Colless et~al.\ 2001) and SDSS galaxy redshift survey make them an ideal starting point for studies of AGN and their host galaxies in regions of the electromagnetic spectrum other than the optical. This encompasses the study of radio-selected AGN using surveys such as the 1.4 GHz National Radio Astronomy Observatory (NRAO) VLA Sky Survey (NVSS; Condon et~al.\ 1998) by Sadler et~al.\ (2002) and Mauch \& Sadler (2007) and the Faint Images of the Radio Sky at Twenty centimeters (FIRST; Becker et~al.\ 1995) survey by Best et~al.\ (2005) and Best \& Heckman (2012). Recently, samples of tens of thousands of nearby mid-infrared-detected AGN covering large areas of the sky have been constructed by cross-correlating sources detected by the Wide-field Infrared Survey Explorer (WISE; Wright et~al.\ 2010) with the SDSS (Donoso et~al.\ 2012, Shao et~al.\ 2013). X-ray data having the depth and wide-field sky coverage to fully exploit the SDSS galaxy sample do not exist. However, to date the SWIFT/BAT survey has detected over 700 (mostly local) AGN at energies above 15 keV (Baumgartner et~al.\ 2013). Near-IR data from WISE and vacuum-ultraviolet data from the Galaxy Evolution Explorer (GALEX; Martin et~al.\ 2005) provide important additional information about AGN host galaxies. In subsequent sections, we will summarize the new scientific insights that resulted from all these large surveys. \subsection{The Landscape of the Galaxy Population and How It Got That Way} To help frame the main issues addressed in this review it is helpful to briefly summarize the basic properties of the population of galaxies in the contemporary universe and the current thinking about how these galaxies were built. The reader is referred to the review by Madau \& Dickinson in this volume for all the details. Results from the SDSS have shown that the galaxy population in the contemporary universe occupies a very small part of the parameter space defined by the structure, stellar content, and chemical composition of a galaxy. In particular, the existence of clear bimodality in the galaxy population was revealed (Kauffmann et~al.\ 2003b, Blanton et~al.\ 2003, Baldry et~al.\ 2004). One population (blue, for short) consists of galaxies with significant on-going star-formation, small stellar masses (M$_*$), low stellar surface mass densities ($\mu_* = 0.5M_*/(\pi R_{50}^2)$, where $R_{50}$ is the radius containing 50\% of the light; see Section~2.4), and small concentrations ($C = R_{90}/R_{50}$) of their light (late Hubble type). The other (red) consists of galaxies with little on-going star formation, large M$_*$, high $\mu_*$, and large $C$ (early Hubble type). The characteristic parameter values that mark the transition between populations are M$_* \sim 10^{10.5}$\,M$_{\odot}$, $\mu_* \sim 10^{8.5}$\,M$_{\odot}$kpc$^{-2}$, and $C \sim 2.6$. Subsequent work showed that the blue population is characterized by a tight almost linear relationship between the star-formation rate (SFR) and M$_*$ (Brinchmann et~al.\ 2004, Schiminovich et~al.\ 2007). This has come to be called the star-forming main sequence. Fig.~\ref{mass_ssfr} shows the relationship between M$_*$ and the specific star formation rate (sSFR $=$ SFR/M$_*$) in the contemporary universe. \begin{figure}[!t] \begin{center} \psfig{file=mass_ssfr_vol.ps,width=9.5cm,clip=} \end{center} \caption{\label{mass_ssfr} The distribution of galaxies in the SDSS main galaxy sample on the plane of stellar mass {\it vs.}\ specific star formation rate (sSFR $=$ SFR/M$_*$). The greyscale indicates the volume-weighted distribution of all galaxies, with each lighter color band indicating a factor of two increase. Galaxies predominantly fall within two regions: a `main sequence' of star-forming galaxies, and a red sequence of `quenched' galaxies. The blue and red contours show the volume-weighted distributions of high ($>$1\%; mostly radiative-mode) and low ($<1$\%; mostly jet-mode) Eddington-fraction AGN, with contours spaced by a factor of two.} \end{figure} Over the past few years, deep surveys have established that the qualitative distribution shown in Fig.~\ref{mass_ssfr} is characteristic of the galaxy population out to at least a redshift of 2 (e.g.\ Whitaker et~al.\ 2013). The most significant difference is that the actual values of sSFR on the star forming main sequence evolve extremely rapidly with time: Elbaz et~al.\ (2011) find as the age of the universe (t$_{cos}$) increased from 2.2 to 13.6 Gyr ($z = 3$ to 0) the characteristic value of the sSFR declined with time as $t_{cos}^{-2.2}$. The simplest picture for the evolution of a typical galaxy (e.g.\ Lilly et~al.\ 2013) is that it evolves along the (strongly evolving) blue star-forming main sequence, increasing in mass through the accretion of cold gas from the cosmic web and (secondarily) through mergers with other galaxies. As it approaches a critical mass, its supply of cold gas is shut off, the star formation is quenched, and the galaxy then evolves into the red population. It can continue to increase in mass through subsequent mergers with other galaxies. It is presently not clear whether the mass-scale at which quenching occurs pertains to the stellar mass or the dark matter halo mass. The physical process(es) that quench the galaxy are unclear. In part, quenching may be due to a change in the nature of accretion: rapid accretion of cold streams of infalling gas at low mass transitioning to slow accretion of hot gas in hydrostatic equilibrium at high mass (e.g.\ Dekel et~al.\ 2009). In principle these processes are included in numerical and semi-analytic models of galaxy evolution. Nevertheless, these models require some additional process to be at play in order to reproduce the observed properties of massive galaxies. Heating and/or the ejection of surrounding gas by an AGN-driven outflow to suppress the cold accretion is a popular idea. In addition the models also require AGN feedback to keep galaxies that arrive in the red/dead population from forming too many stars from the slow accretion of their hot halo gas. Thus, in the current paradigm, AGN play a crucial role in the evolution of massive galaxies. | 14 | 3 | 1403.4620 |
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1403 | 1403.4550_arXiv.txt | {The H.E.S.S. cameras require a precise and regular calibration over time, to reconstruct the gamma-ray characteristics. The different sub-systems used to determine the gain and the uniformity of the PMTs and their evolution with time are presented. Then, we focus on the absolute energy scale calibration, by using a full reconstruction of isolated muons recorded during normal observation. The method and the evolution of the absolute overall light collection efficiency are shown.} | The H.E.S.S. experiment, located in Namibia, is a system of Imaging Atmospheric Cherenkov Telescopes dedicated to the detection of very high energy gamma-rays. It consists of five telescopes including four medium-sized telescopes (radius $\sim$ 6.5 m) and a large telescope (radius $\sim$ 15.5 m). Thereafter, we will focus on the medium-sized telescopes. A detailed description can be found in \cite{bib:calib}. To analyse the observed gamma-ray characteristics, an accurate calibration of the electronic response of the 960 photomultiplier tubes (PMTs) of each H.E.S.S. camera as well as the evaluation of the instrument optical efficiency need to be performed. The different steps are presented and the method to determine the optical efficiency is described in detail. | The calibration parameters of the H.E.S.S. cameras are well reconstructed and monitored over years. Two independent calibration patterns are in operation in the H.E.S.S collaboration and have shown that the intensity in photoelectrons can be evaluated with a precision better than 5 \%. Calibration with muon ring allows to fully understand the losses of Cherenkov photons in the detector. It needs a Monte Carlo study which induces a good knowledge of the detector itself and of the atmosphere. Especially, a precise understanding of the attenuation of Cherenkov light in the atmosphere, of the development of electromagnetic showers and propagation of muons in the atmosphere is required. \vspace*{0.5cm} | 14 | 3 | 1403.4550 |
1403 | 1403.1375_arXiv.txt | We propose the new dark matter particle candidate --- the ``black hole atom'', which is an atom with the charged black hole as an atomic nucleus and electrons in the bound internal quantum states. As a simplified model we consider the the central Reissner-Nordstr\"om black hole with the electric charge neutralized by the internal electrons in bound quantum states. For the external observers these objects would look like the electrically neutral Schwarzschild black holes. We suppose the prolific production of black hole atoms under specific conditions in the early universe. | \label{introsec} The idea of ``black hole atoms'' goes back a long way in several variations. M. A.~Markov et al. proposed and studied in detail the model of maximons (or friedmons) \cite{Mar66,MarFro70,ManMar73,Mar87}. These objects are the particle-like gravitating systems (semiclosed worlds) with mass close to the Planck mass $M_{\rm Pl}=\sqrt{\hbar c /G}\approx10^{-5}$~g. They may have in principle a large gravitational mass defect. The friedmons in connection with self-energy of elementary particles were discussed in \cite{MarFro72}. Maximons are interesting for cosmological applications, in particular, because they have the particle-like properties and may be the enigmatic dark matter particles. Possible role of maximon clusters in cosmology was considered in \cite{MarFro79}. The idea of micro black hole caring the electric charge and having the orbiting electrons or protons at the outer (outside the horizon) orbits was discussed by S.~Hawking in \cite{Haw71}. He firstly proposed that the charged black holes may play the role similar to the atomic nuclei. Later the idea of black hole atoms was investigated in \cite{FlaBer01,Floetal11,FilLap06}. The possible origin of the such Planck mass black hole is the final stable state of the evaporated primordial black holes (PBH), see e.\,g., \cite{ManMar73,Mar87,CarGilLid94}. The remnants of the evaporated black holes can be stable and also can serve as the dark matter candidates \cite{Mac87,DolNasNov00,AdlCheSan01,CheAdl03,Carr03,DymGal07,DymKor10}. In this paper we discuss the black hole atoms, which are the atoms with the charged black hole as atomic nuclei and with electrons in the bound {\bf internal} quantum states. The quantum bound states of electrons may exist in principle not only outside the event horizon but also inside the Cauchy horizon of the charged black hole. So, the main new idea is that there can be configurations in which the orbiting electrons are {\bf inside} the black hole Cauchy horizon. We propose these black hole atoms as the possible origin of dark matter particles. The quantum levels in the gravitational field of black holes outside the event horizon were studied in \cite{DerRuf74,Teretal78,Kof82,SofMulGre77,Teretal80,GalPomChi83,TerGai88,GaiZas92,Lasetal05,GorNez12,Vroetal13,Dzh12}. The resulting black hole atoms can be the dark matter particles in the case of uncompensated charge (electrons at outer levels) as it was proposed in \cite{Vroetal13}. The similar idea but for the zero total charge $q=-Q$ and for the electrons inside a black hole was proposed in \cite{DokEro13}. In the latter case, the total charge of all the electrons at the inner orbits is equal to the charge of the black hole, which appears at Reissner-Nordstr\"om metric. In the case of compensated charge these systems look for the external observer as having the Schwarzschild metric. Neutral systems interact weakly with other particles, it makes them the good candidates for the dark matter particles. The stationary quantum levels of fermions in the gravitational field of the charged black holes have been found in the work \cite{DokEro13} by solving the corresponding Dirac equation. The Dirac equation in the Riemann geometry was first derived in the papers \cite{Fock29}. The using of only the covariant generalization is not enough for derivation of the corresponding Dirac equation. It is needed the determination of the parallel spinor transport. As it was shown in \cite{DokEro13}, a self-consistent steady-state solution with a finite normalization integral can exist only in the case of extreme black hole, whose charge in the appropriate units is equal to its mass $M=|Q|$. | \label{concsec} In this paper we discuss the new kind of ``black hole atom'' system: the Reissner-Nordstr\"om black holes with the electrons at quantum levels under the Cauchy horizon. If the electric charge of the black hole is neutralized by the internal electrons in bound quantum states, these objects would look like the electrically neutral Schwarzschild black holes for the external observers. Due to extremely small interaction cross-section these neutral systems are almost non-interacting with baryons and behave as collisionless and dissipationless gas. This property makes them the good dark matter candidates. The black hole atoms under consideration could form at the final stages of PBHs evaporation at early universe. The PBHs itself may form in different scenarios: from adiabatic perturbations, during cosmological phase transitions or at the early dust-like stages \cite{Khl85}. The extremal black holes doesn't evaporate in the Hawking process, and they are stable in this sense. But the existence of the internal and external quantum levels gives the possibility of the quantum transitions between the levels with radiation of the photons, and this effect makes the ``black hole atoms'' are observable in principal. This study was supported by the grants RFBR 13-02-00257-a and OFN-17. | 14 | 3 | 1403.1375 |
1403 | 1403.6280_arXiv.txt | {The K-type binary star HD$\,$102077 was proposed as a candidate member of the TW Hydrae Association (TWA) which is a young ($5-15\,$Myr) moving group in close proximity ($\sim 50\,$pc) to the solar system. The aim of this work is to verify this hypothesis by different means. We first combine diffraction-limited observations from the ESO NTT $3.5\,$m telescope in SDSS-i' and SDSS-z' passbands (three epochs) and ESO $3.6\,$m telescope in H-band (one epoch) with literature data to obtain a new, amended orbit fit of the visual binary. We then estimate the spectral types of both components from the i'-z' colours and reanalyse the Hipparcos parallax and proper motion taking the orbital motion into account. Moreover, we use two high-resolution spectra of HD$\,$102077 obtained with the fibre-fed optical echelle spectrograph FEROS at the MPG/ESO $2.2\,$m telescope to determine the radial velocity and the lithium equivalent width of the system. Finally, we use all the information to discuss the kinematic properties of HD$\,$102077 and to estimate the age of the system. The orbital elements of the HD$\,$102077 trajectory are well {constrained} and we derive a total system mass of $2.6 \pm 0.8\,$M$_{\sun}$ and a semi-major axis of ${14.9 \pm 1.6}\,$AU. From the i'-z' colours we infer an integrated spectral type of K2V, and individual spectral types of $K0 \pm 1$ and $K5 \pm 1$ for primary and secondary, respectively. The radial velocity corrected {for} the orbital motion of the system is $17.6 \pm 2\,$km/s. Even though the parallax determination from the Hipparcos data is not influenced by the orbital motion, the proper motion changes to $\mu_{\alpha}*\cos{\delta} = -137.84 \pm 1.26$~mas\,yr$^{-1}$ and $\mu_{\delta} = -33.53 \pm 1.45$~mas\,yr$^{-1}$. With {the resultant} space motion, the probability of HD$\,$102077 being a member of TWA is less than ${1\%}$. Furthermore, the lithium equivalent width of $200 \pm 4\,$m\AA $\,$ is consistent with an age between $30\,$Myr and $120\,$Myr and thus older than the predicted age of TWA. The comparison of HD$\,$102077's temperature and luminosity to isochrones supports this result. In conclusion, HD$\,$102077's age, galactic space motion, and position do {not} fit TWA {or} any other young moving group.} | The TW Hya association (TWA) is a loose group of stars {\citep{Gregorio-Hetem92}} that have common {galactic kinematics}, age, and origin \citep{Kastner97, Weinberger13}. The association has $22$ high probability members with individual distances between $28\,$pc and $70\,$pc \citep{Torres08}. The age of its members has been estimated using different indicators of youth such as activity and lithium abundance as well as their position in the Hertzsprung-Russell diagram and is about $10 \pm 5 \,$Myr \citep{Weinberger13}. Young, nearby associations {like} TWA are of great importance to understanding the local star forming history, and provide a sample of young stars, brown dwarfs, and planets for high-resolution studies. Therefore, several groups have been looking for additional members \citep{Schlieder10, Schlieder12, Zuckerman11, Shkolnik12, Rodriguez13, Malo13, Moor13}. The binary HD$\,$102077 {(also called HIP 57269, V 838 Cen, and RST 3558AB)} was proposed as a candidate member of TWA based on its kinematics by \cite{Makarov01}. Kinematic candidate members, however, need to be verified by spectroscopic measurements. The first spectroscopic follow up was done by \cite{Song02}, who measured the Li$\,6708\,$\AA $\,$ strength of HD$\,$102077. They find a fairly large equivalent width of $196\,$m\AA, which is a strong indication of a young system. The authors argue, however, that an age of more than $30\,$Myr is more likely and thus do not consider HD$\,$102077 as a new member of TWA. Apart from the lithium absorption line, HD$\,$102077 shows strong X-ray emission which also hints at a young age \citep{Makarov03}. The WISE data \citep{Cutri12} shows, however, no infrared excess and thus no sign of a disk. Even though HD$\,$102077 is a young star, several research groups have assigned it low metallicity values. \cite{Holmberg07} used the Stroemgren UVBY photometry by \cite{Olsen94} to derive a [Fe/H] of $-0.9$. \cite{Casagrande11} redid their analysis and find [Fe/H]$=-0.76$. An independent metallicity measurement was done by \cite{Randich93}. They fitted spectral templates to high-resolution spectra with $R=50000$ over the region of $\approx 6680-6730\,$\AA $\,$ and obtained [Fe/H]$=-0.4$. The photometric determination of the metallicity content is not very accurate because of HD102077’s binarity and variability and this may be the cause of the discrepancy. For instance, \cite{Cutispoto90} measured the $UBVRI$ colours over several photometric phases and observed that $U-B$ varies by $0.04\,$mag and that the star gets bluer at light maximum. Even stronger variations have been observed by \cite{Udalski85}. They find brightness changes of $0.08\,$mag with a probable period of $1.84\,$ or $2.2\,$days. {Since the V-band lightcurves at different epochs differ greatly in shape and mean magnitude \citep{Cutispoto90, Cutispoto93, Cutispoto98} the starspot size and distribution seem to vary considerably. Duplicity of one of the two components could also add to the observed variations.} \cite{Konig03} measured the radial velocity of the system and find that the space motion is quite different from the average space velocity of TWA. Nevertheless, they point out that the distance of HD$\,$102077, its location in the sky, its spectral type, its position in the H-R diagram, its $v \sin i$, and X-ray emission are very similar to other confirmed TWA members. {This group also presents evidence of a third component in the system. The proposed star is at a distance of $8.46''$ from the primary and it shows calcium and lithium absorption. The lithium EW of $0.18 \pm 0.2\,$\AA $\,$ and a spectral type of $K5 \pm 1$ point to a similar age to the A component. We checked the imaging archives to see whether there is an indication of common proper motion. The stars do seem to move together with respect to the background stars in the field when comparing an R-band image from $1983$ to an R-band image from $1992$ and the 2MASS K-band image from $1999$. The psfs are, however, blended in all these images and thus no quantitative proper motion measurement is possible. Nevertheless, the calcium and lithium absorption features as well as the likely common proper motion is strong evidences of a bound system. } {The integrated spectral type of HD$\,$102077 of $K0/1Vp$ was first determined by \cite{Anders91} by high-dispersion observations of the Li $6707$ \AA \, doublet and \cite{Konig03} find $K1/2$ using DFOSC at the $1.54\,$m Danish telescope.} The spectral type of the {primary} components is, however, less clear. \cite{Fabricius00} analysed the Hipparcos Tycho data and find that the brighter component has a redder B-V colour, corresponding to spectral types K4 and K2, but they do not give an explanation for this unexpected result. {\cite{Cutispoto98} uses the UBVRI colours of the integrated system and $\Delta H_p$ from the Hipparcos satellite catalog to classify HD$\,$102077 as consisting of a $K0/1V$ and a $K5V$.} The source HD$\,$102077 is a visual binary star and thus allows the determination of the total system mass. The first orbital parameters were derived by \cite{Heintz86}. His best fit to data from seven epochs is a low eccentricity orbit with a period of $32.35\,$years and a semi-major axis of $0.28''$. At that time, the parallax of HD$\,$102077 was not known. Therefore, Heintz used Kepler's Third Law and the mass-luminosity relation to estimate the dynamical parallax. He obtained $42.3\,$pc and deduced a total system mass of $1.6\,M_{\sun}$. In this paper, we present the results of our astrometric, photometric, and spectroscopic study of HD$\,$102077. We derive a new orbit fit of the binary system, and we determine the i'-z' colours and from these the spectral types of the individual components. We measure the radial velocity and the lithium equivalent width of the system. Finally, we use these results to discuss a possible TWA membership of HD$\,$102077. | \subsection{{Additional companions}} In addition to the wide companion candidate described by \cite{Konig03}, our mass estimate of $2.6 \pm 0.8\,$M$_{\sun}$ may indicate that HD$\,$102077 has another, so far unknown companion; this is considered because the two K-type stars should be less massive than $1.6\,$M$_{\sun}$. This would be in accordance with the classification as RS CVn by \cite{Weiler79}, \cite{Gurzadian92}, and \cite{Pallavicini92}. In addition, the 2D cross correlation shows a large radial velocity difference which is in disagreement with the orbital fit, but could be caused by a third component. In contrast, the V-band variability seems to be linked with stellar activity and HD$\,$102077 has also been classified as BY Draconis by several groups \citep{Udalski85, Alekseev96}. Whether HD$\,$102077 is only active or a spectroscopic binary is {thus} not clear. Further evidence is needed to {answer} this question. \subsection{Spectral type derivation} From the i'-z' colours we determined an integrated spectral type of K2V. The spectral type determined with spectroscopic means was $K0/1Vp$ \citep{Anders91} and is thus one subclass earlier. This deviation is common in a spectral classification based on two colours. We found the individual spectral types to be $K0 \pm 1$ and $K5 \pm 1$ {which are in agreement with $K0/1V$ and $K5V$ derived by \cite{Cutispoto98}}. This classification seems to be more reliable than the K4 and K2 determined by \cite{Fabricius00} since the brighter primary star is expected to have the earlier spectral type unless the companion is a white dwarf. We note that the errors of the two spectral types are correlated. An earlier spectral type of the primary star {implies} a later type for the secondary. \subsection{Revisiting the Hipparcos data for HD$\,$102077} In the original version of the Hipparcos Catalog (ESA 1997) the dataset of HD$\,$102077 = HIP$\,$57269 was solved as a two-component system, with identical proper motions and parallaxes. The separation and position angle are given in Table$\,$1 and have been included in the visual orbital fit; the absolute astrometry given for HIP$\,$57269 in the main Hipparcos Catalog corresponds to component A. The solution quality is grade 'A', indicating the highest possible quality. The two components of the system have been recognized separately in the Hipparcos raw data via a deviation from the distribution of photon count rate as a function of time from the distribution expected for a single star. For these component systems in the original Hipparcos Catalog, the intermediate astrometric data are hard to interpret, since the reference point for the data is not known a priori because it is a complicated function of the photometric and geometric characteristics of the system, and is different for the two data reduction consortia FAST and NDAC. In the version of the Hipparcos Catalog by \cite{vanleeuwen97}, the system was recognized as a double star, but again solved with the standard five astrometric parameter model; the astrometry given in the catalog seems to refer to component A, just as in the original version of the Hipparcos Catalog. Our goal is to re-interpret the Hipparcos Intermediate Astrometric Data, taking the orbital solution derived in the present paper into account, and thereby revising the standard astrometric parameters and in particular the parallax. We used the Hipparcos Intermediate Astrometric Data (abscissae) from \cite{vanleeuwen07} and fitted corrections to the standard five astrometric parameters (positions and proper motions in right ascension and declination, respectively, and the trigonometric parallax). Before the fitting, we corrected the abscissae for orbital motion, with the orbital parameters given in Table~3. Since the Hipparcos data are absolute astrometric measurements, the reference point for the orbital motion is the centre of mass of the system, against which both components are seen to be moving. In order to derive the centre of mass of the system, we used mass estimates of 0.9$\,$M$_\odot$ for the primary and 0.7$\,$M$_\odot$ for the secondary. As expected, the largest adjustment occurs in the proper motions when the orbit is taken into account in the analysis of the Hipparcos data. This is because the Hipparcos measurements are spread out over approximately three years, whereas the orbital period is more than a factor of 10 larger, so that the orbital phase coverage of the Hipparcos measurements is relatively small. Thus, some orbital motion has been presumably subsumed into the proper motions before, whereas the corrected proper motions should be free of orbital motion. The new proper motions are $\mu_{\alpha}*\cos{\delta} = -137.84 \pm 1.26$~mas\,yr$^{-1}$ and $\mu_{\delta} = -33.53 \pm 1.45$~mas\,yr$^{-1}$; we reassess the kinematic membership of HD$\,$102077 to a moving group on the basis of these new proper motions in Section~4.4. The parallax of HD$\,$102077 {of $20.59 \pm 2.14\,$mas}, however, does not change at all when the orbit is taken into account in the fitting of the Hipparcos abscissae. We caution that we might have misinterpreted the reference point of the Hipparcos Intermediate Astrometric Data, which would render our newly derived proper motions meaningless. However, we have verified that the parallax result is very robust and does not change at all even if we assume other reference points. So we conclude that the published Hipparcos parallax is correct even in the case of orbital motion. \subsection{Moving group membership} The source HD$\,$102077 was proposed as candidate member of TW Hydrae \citep{Makarov01}. However, objections have been raised based on the space motion and weaker-than-expected lithium absorption \citep{Song02, Konig03}. We now revisit both points. Our new, high-resolution measurements of the Li EW confirm the lithium absorption value and the comparison of our data to stellar isochrones also suggests a higher age of HD$\,$102077 than is predicted for TW Hydrae. \cite{Anders91} measured the radial velocity of HD$\,$102077 in July 1987 and \cite{Konig03} in January 2002; we measured the radial velocity in July 2013. The results are $15.9\,$km/s, $19 \pm 3\,$km/s, and $19.8 \pm {0.1} \,$km/s respectively. Another RV-measurement was done by {\cite{Nordstroem04}} who find $16.8 \pm {0.3}\,$km/s, but they do not give the epoch. These radial velocity values do not take the orbital motion into account. According to our orbit fit the relative motion of both components was $3.8\,$km/s in July 1987, $0.6\,$km/s in January 2002, and $5.0\,$km/s in July 2013. As in section 3.2.2 we adopt a maximum uncertainty of half the relative motion (e.g. $1.9\,$km/s for 1987). For the measurement by {\cite{Nordstroem04}} we assume an uncertainty of $2\,$km/s. The weighted mean of this values gives ${17.6 \pm 2} \,$km/s. This RV value together with the newly derived proper motion can be used to calculate the space velocity as well as the kinematic membership likelihood. We derive UVWXYZ Galactic velocities and positions of the binary to be ${U=-18.9 \pm 1.1}\,$km/s, ${V=-29.7 \pm 2.5}\,$km/s, and ${W=-11.7 \pm 2.9}\,$km/s, as well as $X=17.7 \pm 1.4\,$pc, $Y=-44.0 \pm 2.4\,$pc, and $Z=10.1 \pm 2.3$pc and compare these values to those of TWA ($U=-9.9 \pm 4.2\,$km/s, $V=-18.1 \pm 1.4\,$km/s, ${W=-4.5 \pm 2.8}\,$km/s, and $X=12.5 \pm 7.1\,$pc, $Y=-42.3 \pm 7.3\,$pc, $Z=21.6 \pm 4.2$pc; \cite{Malo13}). We find that even though the Galactic position of HD$\,$102077 is close to most of the TW Hydrae members its space motion is different. The space motion of HD$\,$102077 {also does} not fit other young associations (see Fig.~\ref{movinggroup}). To calculate the moving group membership probability quantitatively we use BANYAN {II} (Bayesian Analysis for Nearby Young AssociatioNs {II}) \citep{Malo13, Gagne13} and find a probability of ${0.4\%}$ of HD$\,$102077 being a member of TWA. Even if the true parallax value happens to be outside of its $3 \, \sigma$ confidence region, the kinematic membership probability does not get higher than ${10 \%}$. Together with the estimated age this strongly suggests that HD$\,$102077 is not a member of TWA. The probability of HD$\,$102077 being a member of another young moving group is only non-zero for AB Doradus. With the newly derived system parameters it is about ${0.03 \%}$ and it is thus unlikely even though the age of the binary system matches the age of this moving group better. {HD$\,$102077 is most likely a member of the young Galactic field.} \begin{figure*} \centering \includegraphics[width=1\textwidth]{HD102077_6D_PLOTS_COLOR.pdf} \caption{2D projections of the UVW space velocities compared to well known moving groups, including TWA. The dashed grey box is taken from \cite{Zuckerman04} and shows the region of UVW space typically occupied by nearby young stars. The UVW velocities of the young moving groups are taken from \cite{Malo13}, the $2 \sigma$ uncertainties are plotted. HD$\,$102077 is clearly not a member of any of these groups.} \label{movinggroup} \end{figure*} | 14 | 3 | 1403.6280 |
1403 | 1403.3693_arXiv.txt | % {} {The aim of this work is to constrain the evolution of the fraction of strong \lae~emitters among UV selected star--forming galaxies at $2<z<6$, and to measure the stellar escape fraction of \lae~ photons over the same redshift range.} {We exploit the ultradeep spectroscopic observations with VIMOS on the VLT collected by the VIMOS Ultra--Deep Survey (VUDS) to build an unique, complete, and unbiased sample of $\sim4000$ spectroscopically confirmed star--forming galaxies at $2<z<6$. Our galaxy sample includes UV luminosities brighter than $M^*_{FUV}$ at $2<z<6$, and luminosities down to one magnitude fainter than $M^*_{FUV}$ at $2<z<3.5$. } {We find that 80\% of the star--forming galaxies in our sample have $EW_0(Ly\alpha)<10$\AA, and correspondingly \esc$<1$\%. By comparing these results with the literature, we conclude that the bulk of the \lae~luminosity at $2<z<6$ comes from galaxies that are fainter in the UV than those we sample in this work. The strong \lae~emitters constitute, at each redshift, the tail of the distribution of the galaxies with extreme \ew~and \esc. This tail of large \ew~and \esc~becomes more important as the redshift increases, and causes the fraction of strong \lae~with \ew$>25$\AA~to increase from $\sim$5\% at $z\sim2$ to $\sim$30\% at $z\sim6$, with the increase being stronger beyond z$\sim4$. We observe no difference, for the narrow range of UV luminosities explored in this work, between the fraction of strong \lae~emitters among galaxies fainter or brighter than $M^*_{FUV}$, although the fraction for the faint galaxies evolves faster, at $2<z<3.5$, than for the bright ones. We do observe an anticorrelation between E(B-V) and \esc: generally galaxies with high \esc~also have small amounts of dust (and vice versa). However, when the dust content is low (E(B-V)$<$0.05) we observe a very broad range of \esc, ranging from 10$^{-3}$ to 1. This implies that the dust alone is not the only regulator of the amount of escaping \lae~photons.} {} | Narrowband surveys targeting the strong \lae~emission from star--forming galaxies (Lyman-$\alpha$ emitters, LAEs; Partridge\&Peebles~1967; Djorgovski~et~al.~1985; Cowie~\&Hu~1998; Hu~et~al.~2004; Kashikawa~et~al.~2006; Gronwall~et~al.~2007; Murayama~et~al.~2007; Ouchi~et~al.~2008; Nilsson~et~al.~2009) and broadband surveys targeting the deep Lyman break (LBG; Steidel~et~al.~1999; Bouwens~\&~Illingworth~2006; Bouwens~et~al.~2010; McLure~et~al.~2011) have been very successful at exploring the high--redshift Universe. However, the overlap between the populations selected by the two techniques is still debated: LAEs are claimed to be forming stars at rates of $1\div10 M_{\odot}yr^{-1}$ (Cowie~\&~Hu~1998; Gawiser~et~al.~2006; Pirzkal~et~al.~2007), to have stellar masses on the order of $10^8\div10^9M_{\odot}$ and to have ages smaller than 50 Myr (Pirzkal~et~al.~2007; Gawiser~et~al.~2007; Nilsson~et~al.~2009), while LBGs have in general a broader range of properties (Reddy~et~al.~2006; Hathi~et~al.~2012; Schaerer,~de~Barros~\&~Stark~2011; but see also Kornei~et~al.~2010). Steidel~et~al.~(2000) and Shapley~et~al.~(2003) showed that only $\sim20$\% of z$\sim3$ LBGs have a \lae~emission strong enough to be detected with the narrowband technique. Recently, many authors have investigated the evolution with the redshift of the fraction of strong \lae~emitters among LBG galaxies. Stark~et~al.~(2010;~2011) showed that this fraction evolves with redshift, and that the overall fraction is smaller (and that the rate of evolution is slower) for UV bright galaxies ($-21.75<M_{UV}<-20.25$) than for UV faint ($-20.25<M_{UV}<-18.75$) galaxies; they find that the fraction of UV faint galaxies with strong (\ew$>25$~\AA) \lae~emission is around 20\% at $z\sim2\div3$ and reaches $\sim50\div60$\% at $z\sim6$. At higher redshift ($z>6\div8$), many authors claim a sudden drop in the fraction of spectroscopically confirmed LBGs with strong \lae~emission (Fontana~et~al.~2010; Pentericci~et~al.~2011; Ono~et~al.~2012; Schenker~et~al.~2012; Caruana~et~al.~2014), interpreting this as the observational signature of the increasing fraction of netural hydrogen between $z\sim6$ and $z\sim7$ due to the tail end of the reionization, although Dijkstra~et~al.~(2014) has argued that the effect can be due to a variation of the average escape fraction over the same redshift range. However, the bulk of studies of the \lae~fraction at $3<z<8$ (Stark~et~al.~2010;~2011; Pentericci~et~al.~2011) are based on a hybrid photometric-spectroscopic technique: the denominator of the fraction (i.e., the {\it total} number of star--forming galaxies at those redshifts) is only constrained by photometry, and thus its determination relies on the strong assumption that the contamination by low--$z$ interlopers and incompleteness are fully understood and well controlled. The numerator of the fraction is the number of the LBGs that are observed with spectroscopy, and for which a strong \lae~ rest--frame Equivalent Width (EW$_0>$25~\AA) is measured. In fact, the LBGs for which this experiment is done have a UV continuum that is generally too faint to be detected, even with the most powerful spetrographs on 10-meter class telescopes. Recently, Mallery~et~al.~(2012) combined a sample of LAEs and LBGs to constrain the evolution of this fraction, confirming earlier results by Stark~et~al.~(2010;~2011). Given the nature of the selection of these samples, it is important to make a robust estimate of the evolution of the \lae~fraction covering as wide a range in redshift as possible, and based on larger samples. The \lae~is interesting not only because it allows for the exploration of the high--redshift universe. In fact, its observed properties can give a lot of information about the physical condition of star--forming galaxies. \lae~is thought to be mainly produced by star formation, as the contribution of AGN activity to the \lae~population at $z<4$ is found to be less than 5\% (Gawiser~et~al.~2006; Ouchi~et~al.~2008; Nilsson~et~al.~2009; Hayes~et~al.~2010). Because of its resonant nature, \lae~photons are easily scattered, shifted in frequency, and absorbed by the neutral hydrogen and/or by the dust. As a result, in general, \lae~emission is more attenuated than other UV photons, with the \lae~escape fraction (i.e., the fraction of the \lae~ photons that escape the galaxies) that depends strongly on the relative kinematics of the HII and HI regions, dust content, and geometry (Giavalisco~et~al.~1996; Kunth~et~al.~1998; Mas-Hesse~et~al.~2003; Deharveng~et~al.~2008; Hayes~et~al.~2014). As a result of their nature, \lae~photons are found to be scattered at much larger scales than UV photons (Steidel~et~al.~2011; Momose~et~al.~2014). Predicting the escape fraction of the \lae~photons as a function of the galaxy properties involves including all the complex effects of radiative transfer of such photons. Developing the first models by Charlot~\&~Fall~(1993), Verhamme~et~al.~(2006; 2008; 2012) and Dijkstra~et~al.~(2006; 2012) made huge progress in predicting the shape of the \lae~emission as a function of the properties of the ISM, the presence of inflows/outflows, and dust. Verhamme~et~al.~(2006; 2008) predicted a correlation between \esc~and E(B-V), with the escape fraction being higher in galaxies with low dust content. Verhamme~et~al.~(2012) and Dijkstra~et~al.~(2012) studied the escape fraction of \lae~photons through a 3D clumpy medium, constraining the dependence on the column density of neutral hydrogen and on the viewing angle. A lot of effort has been recently put to constrain the correlation between the \lae~properties and the general properties of star--forming galaxies (e.g., dust attenuation, SFR, stellar mass) in the local Universe. Hayes~et~al.~(2014) and Atek~et~al.~(2014) have found that \lae~photons escape more easily from galaxies with low dust content. At high redshift, although on samples that are much smaller than the one we use in this paper, a similar trend has been found by Kornei~et~al.~(2010) and Mallery~et~al.~(2012), respectively at $z\sim3$ and at $4<z<6$. In this paper, we look for this correlation using a sample that is respectively five and ten times larger than the ones used by Mallery and Kornei. The aim of this paper is to estimate the evolution of the fraction of strong \lae~emitters as a function of the redshift, exploiting data from the new VIMOS Ultra--Deep Survey (VUDS). The goal is twofold: first, to put on firmer grounds the trends that have been found with photometric LBG samples (Stark~et~al.~2010; 2011) and improve on the knowledge of the evolution of the \lae~fraction; second, to offer the theoreticians a reference sample of galaxies with robust spectroscopic redshifts, with a well measured \ew~distribution. In fact, in this paper, we select a sample of galaxies, sliced in volume limited samples according to different recipes, for which we have a spectroscopic redshift in $\sim$90\% of the cases. The continuum is detected for almost all objects in the sample, thus allowing a robust measurement of the redshift based on the UV absorption features even in absence of \lae. \begin{figure*}[!ht] \centering \includegraphics[width=.85\textwidth]{ex511249489.ps} \centering \includegraphics[width=.85\textwidth]{ex520266763.ps} \centering \includegraphics[width=.85\textwidth]{ex520396679.ps} \centering \includegraphics[width=.85\textwidth]{ex511682756.ps} \centering \includegraphics[width=.85\textwidth]{ex520458995.ps} \centering \includegraphics[width=.85\textwidth]{ex5101052041.ps} \caption{Six examples of spectra for the galaxies in the sample. The left panels show the region around \lae, while the right ones show the full spectrum, with the most common UV rest-frame lines highlighted in red. These examples are chosen to be representative of the $i-$band magnitudes, redshifts and \lae~equivalent widths covered by the sample presented in this work. The red dashed curves show polynomial fits to the continuum: for each spectrum, the region between 912~\AA~and \lae~and the region between \lae~ and 2000~\AA~are fitted separately. We note that the fits are not used at all in the analysis presented in this paper; they only provide a guidance to assess the continuum around \lae. The blue triangles show the points on the continuum bracketing the \lae~ line, shown in green.} \label{ex}% \end{figure*} \begin{table} \centering \begin{tabular} {clr@{$\pm$}rr@{$\pm$}lr@{$\pm$}lr@{$\pm$}lr@{$\pm$}lc} \hline\hline \noalign{\smallskip} \multicolumn{1}{c}{} & \multicolumn{2}{c}{$2<z<2.7$} & \multicolumn{2}{c}{$2.7<z<3.5$} & \multicolumn{2}{c}{$3.5<z<4.5$} & \multicolumn{2}{c}{$4.5<z<6$}\\ \noalign{\smallskip} \hline \noalign{\smallskip} \multicolumn{1}{c} {f=0}& \multicolumn{2}{c}{83(0)} & \multicolumn{2}{c}{90(0)} & \multicolumn{2}{c}{40(0)} & \multicolumn{2}{c}{18(0)}\\ \noalign{\smallskip} \multicolumn{1}{c} {f=1}& \multicolumn{2}{c}{299(2)} & \multicolumn{2}{c}{200(3)} & \multicolumn{2}{c}{88(2)} & \multicolumn{2}{c}{14(0)}\\ \noalign{\smallskip} \multicolumn{1}{c} {f=2}& \multicolumn{2}{c}{614(24)} & \multicolumn{2}{c}{614(17)} & \multicolumn{2}{c}{163(5)} & \multicolumn{2}{c}{47(3)}\\ \noalign{\smallskip} \multicolumn{1}{c} {f=3,4}& \multicolumn{2}{c}{646(106)} & \multicolumn{2}{c}{701(153)} & \multicolumn{2}{c}{205(57)} & \multicolumn{2}{c}{41(22)}\\ \noalign{\smallskip} \multicolumn{1}{c} {f=9}& \multicolumn{2}{c}{28(15)} & \multicolumn{2}{c}{31(10)} & \multicolumn{2}{c}{19(6)} & \multicolumn{2}{c}{20(5)}\\ \noalign{\smallskip} \hline \noalign{\smallskip} \end{tabular} \caption{The final sample of galaxies used in this work, divided in 4 redshift bins, as a function of the spectroscopic quality flag. The number in parentheses indicates the number of objects at that redshift and of that spectroscopic quality flag that have EW$_0>25$~\AA. }\label{tab:sample} \end{table} Our selection is not based on LBG or narroband techniques, that are prone to incompleteness and contamination, but it is rather based on the magnitude in the $i'-$band and on the photometric redshifts measured on the full Spectral Energy Distribution (SED) of galaxies. The most important point to emphasize is that our flux selection is completely independent of the presence of \lae, at least up to $z\sim5$, because it enters the photometric $i'-$band only at $z>5$: since the $i'-$band does not contain the \lae~line, objects with strong \lae~emission have not a boosted $i'-$band magnitude. Moreover, when the photo-$z$ are computed, some variable \lae~flux (as for other lines like OII, OIII and H$\alpha$) is added to the SED: this ensures that even objects with large \lae~flux are reproduced by the template set that is used to compute the photo-z. This also implies that if our selection is incomplete at some redshift, the incompleteness is also independent of the presence (or absence) of \lae. For these reasons, this sample is ideal to study the \lae~properties of a well controlled sample of star--forming galaxies. The fraction of strong \lae~emitters among star--forming galaxies is completely constrained by spectroscopy, as is also the case for non-\lae~ emitters. Throughout the paper, we use a standard Cosmology with $\Omega_M=0.3$, $\Omega_{\Lambda}=0.7$ and $h=0.7$. Magnitudes are in the AB system. | In this paper we used the unique VUDS dataset to build an unbiased and controlled sample of star--forming galaxies at $2<z<6$, selected according to the photometric redshifts determined using the overall SED of the galaxies. This selection is complementary to the classical LBG technique, resulting in more complete and less contaminated samples of galaxies at high--$z$. For the purpose of this paper, even more imporant is that the combination of the selections we use are independent of the presence of \lae~in emission, at least up to $z\sim5$: whatever incompleteness could affect our sample, it would affect galaxies with and without \lae~in the same way. The sample is limited at $m_i<25$, ensuring that the continuum is detected with S/N$\sim10$ per resolution element: this allows an accurate determination of the spectroscopic redshift through the identification of UV absorption features even for galaxies without \lae~in emission. We split this sample in two volume limited samples, using a far-UV luminosity cut that is evolving with redshift, following the observed evolution of $M^*_{FUV}$ (Hathi~et~al.~2010): the bright sample include objects that at each redshift are brighter than $M^*_{FUV}$; the faint one include objects with $M^*<M_{FUV}<M^*+1$. We use these two samples to constrain the distribution of the EW of \lae~of star--forming galaxies, that spans from objects with \lae~in absorption to objects with \lae~in emission. We find that $\sim80$\% of the star--forming galaxies in our sample have a \lae~equivalent width $EW_0(Ly\alpha)<15$\AA. We use our sample to constrain the evolution of the fraction of strong \lae~emitters among star--forming galaxies at $2<z<6$. We showed in Section~4 that the fraction of strong \lae~emitters with $EW_0(Ly\alpha)>25$\AA~ and $EW_0(Ly\alpha)>55$\AA~ monothonically increases with redshift, approximately at the same rate for the two EW thresholds. The evolution is characterized by a slower phase between $z\sim2$ and $z\sim4$, and by a faster evolution between $z\sim4$ and $z\sim5.5$. We see no difference, at $2<z<3.5$ where both samples are well represented, between the fraction of strong emitters in the bright and faint volume limited samples. This is partly in contraddiction with results by Stark~et~al.~(2010; 2011), who found that the fraction is higher, and the rate of evolution with redshift faster, for UV faint galaxies at $4<z<6$. However, this might be due to the narrower range of UV luminosity probed by our work compared to the one probed by Stark~et~al.~(2010; 2011). Moreover, slicing our sample with the same UV luminosity limits used by Stark ($-21.75<M_{FUV}<-20.25$) we see that the evolution of the fraction of strong \lae~emitters (for both $EW_0(Ly\alpha)>25$\AA~ and $EW_0(Ly\alpha)>55$\AA) is in very good agreement with the values by Stark~et~al.~(2010; 2011), despite the different sample selection methods and available spectroscopy. This is a very important result, placing on firmer grounds the measures of the fraction of star--forming galaxies with \lae~in emission. In fact, their sample is LBG based and only the objects with strong \lae~emission are spectroscopically confirmed. In our case, on the other hand, we stress that all the galaxies, with and without \lae, have a spectroscopic redshift. Finally, in Section~5, we have explored the possibility that the evolution of the fraction of strong \lae~emitters is primarly due to a change in the escape fraction of \lae~photons. We have found that, as expected, the strong \lae~emitters are the objects for which \esc~is the largest. We find as well that the median \esc~for the \lae~ emitters (with not much difference between objects with \ew$>25$\AA~and with \ew$>55$\AA) evolves from $\sim$5\% at $z\sim2.5$ to $\sim$20\% at $z\sim5$. If we try to estimate the median escape fraction for the whole population, we find that it is formally zero at all redshifts, since the majority of the galaxies in our sample have \lae~in absorption, and 80\% of our galaxies have \esc$<1$\%. If we estimate at each redshift the \esc~value below which 80\% of the galaxies lie, we find that this value evolves from 1 to 2\% between $z\sim2$ and $z\sim5$. It is interesting to compare these findings with Hayes~et~al.~(2011), who integrated the \lae~and UV luminosity functions from $z\sim0$ to $z\sim8$ and then compared the two to estimate the average \esc~of the Universe at those redshifts. According to Hayes~et~al.~(2011) the average escape fraction is around 5\% at $z\sim2$ and 20\% at $z\sim5$, values that are much higher than those we obtain for our sample. This implies that for the galaxies with UV luminosities that we sample in this paper ($M_{FUV}<M^*$ at $2<z<6$ and $M^*<M_{FUV}<M^*+1$ at $2<z<3.5$) the average escape fraction of \lae~photons is much smaller than the average escape fraction of the Universe. In other words, the bulk of the \lae~luminosity, at least in the redshift range $2<z<6$ that is probed in this paper, is not coming from galaxies with the UV luminosities that are probed in this work, but from galaxies that are much fainter in the UV. In fact, Stark~et~al.~(2011) showed that the fraction of strong (\ew$>$25\AA) emitters is higher in galaxies with $-20.25<M_{FUV}<-18.75$ than in those with $-21.75<M_{FUV}<-20.25$, implying a larger escape fraction for faint UV galaxies. This is also in line with the results by Ando~et~al.~(2006), who found a deficiency of strong \lae~emitters among UV bright galaxies and by Schaerer,~de~Barros~\&~Stark~(2011), who also found that the fraction of \lae~emitters rapidly increases among galaxies with fainter UV luminosities, indicating that the bulk of the \lae~luminosity in the universe comes from galaxies with $M_{FUV}>-20$. Similarly to Kornei~et~al.~(2010) and Mallery~et~al.~(2012), we also find that there is an anti-correlation between \esc~and the dust content E(B-V): galaxies with low \esc~have preferentially a higher E(B-V), and vice versa. This implies that the dust is a crucial ingredient in setting the escape fraction of galaxies. However, we note that galaxies with low extinction ($E(B-V)<0.05$) have a very wide range of \lae~escape fractions, ranging from $10^{-3}$ to 1: this means that the dust content, although important, is not the only ingredient to regulate the fraction of \lae~photons that escape the galaxy. In a forthcoming paper, we will further investigate the dependence of \esc~on other quantities as stellar mass, star formation rate and dust content, and on the evolution with redshift of these correlations. | 14 | 3 | 1403.3693 |
1403 | 1403.5189_arXiv.txt | Grain-surface reactions play an essential role in interstellar chemistry, since dust grain catalyses reactions at its surface allowing for the formation of molecules. We used a chemical model in which both gas-phase and grain-surface reactions occur and studied particularly the diffusion mechanisms on the surface of the grains. Surface reactions can occur via thermal hopping when species cross over a potential barrier or via quantum tunneling when species cross through this barrier. We show that the thermal diffusion (hopping) can be much more efficient after a cosmic ray particle collides with a dust grain, heating it to a peak temperature of 70$\,$K. We present here the results of numerical simulations after including the quantum tunneling mechanism for species H, H$_{2}$ and O and considering the effect of cosmic ray particle collision on the surface reactions. As a consequence, the gas-phase and grain-surface abundances are affected and we show that more complex molecules can be formed in molecular clouds. | Complex Organic Molecules (COMs) have been detected in warm star-forming regions such as OMC-1 \citep{b27} and Sgr B2 \citep{b28,b29,b30} and are usually associated with warm gas-phase and surface chemistries. COMs are expected to be formed at the surface of the grains during the warm-up phase of star formation when the temperature is about 40 - 50$\,$K \citep{b16}. Under such conditions, heavy radicals can move at the surface of the grains and react without being evaporated. When the temperature gets above 80$\,$K, the COMs are then evaporated in the gas-phase \citep{b34}. Recently, acetaldehyde (CH$_{3}$CHO), dimethyl ether (CH$_{3}$OCH$_{3}$), methyl formate (HCOOCH$_{3}$), methanol (CH$_{3}$OH) and formaldehyde (H$_{2}$CO) have also been observed in the cold pre-stellar cores L1689B \citep{b24} and B1-b \citep{b25}. These detections have revived the question of the formation of these species and their presence in the gas-phase with abundances (relative to the total hydrogen) between $10^{-12}$ to more than $10^{-9}$. \citet{b35} proposed that these complex species are formed at low temperature via gas-phase reactions between precursors, such as CH$_{3}$O, formed on the grain surface and released through reactive desorption. At temperature typical of dark clouds or pre-stellar cores (10$\,$K and below), the direct thermal evaporation is only efficient for H and H$_{2}$. Similarly, the diffusion of the species at the surface of the grains, through thermal hopping, is only efficient for atomic hydrogen, so that hydrogen-rich saturated species, such as H$_{2}$O, CH$_{4}$ and NH$_{3}$, are the main constituents of the ices. Current chemical models take into account a number of processes to desorb back into the gas-phase species from the surface. The first one is the desorption induced by stochastic cosmic ray heating \citep{b3}. This process is however only efficient for simple molecules. Another process was introduced by \citet{b11} in which the chemical energy released by exothermic surface reactions contributes to desorb the products. Considering the lack of experimental data on it, the efficiency of the process is not well constrained. \citet{b36} proposed that mantle explosions could be induced by exothermic radical recombination reactions. They concluded that the energy stored by those free radicals being larger than the energy deposits by cosmic rays, the chemical desorption is the dominant non-thermal desorption mechanism in dark clouds. Another desorption mechanism was considered by \citet{b37} in which the energy released by H$_{2}$ formation at the surface of the grains leads to local heating and desorb species back into the gas-phase. However, only weakly bound species (with binding energies less than or equal to that of CO) can be evaporated during this process \citep{b38}. Recent experimental studies by \citet{b39, b43} showed the efficiency of the photodesorption processes. The importance of this mechanism has been underlined by \citet{b45} for the desorption of water at A$_{\text{V}}$ lower than 10 while many observational studies of pre-stellar cores \citep{b46, b44} invoke photodesorption by secondary cosmic ray photons to explain H$_{2}$O gas-phase abundances. Whatever the desorption mechanism included in the models, the gas-phase abundances reflect the surface abundances. The formation of species at the surface of the grains depends on the diffusion rate of the precursors. This diffusion can be thermal when species migrate from one site to another one by thermal hopping or non-thermal when species cross through a potential barrier by quantum tunneling. Diffusion at the surface of the grains could be faster if tunneling effects are included or if the temperature of the grains is higher. \\ In this paper, we revisit the efficiency of tunneling diffusion based on recent experimental studies of oxygen diffusion \citep{b1}. We also study the effect of cosmic ray impacts which cause a stochastic heating of the dust particles \citep{b2} allowing for surface radicals to diffuse quickly and react to form more complex species. We report here the effect of these two mechanisms on molecular abundances and more specially those of complex organic molecules. Note that we do not distinguish between the different layers of ices in our model. As a consequence, we do not take into account diffusion through the bulk of the ice or any differentiation between the surface and the bulk, which can be an important aspect of surface chemistry as shown by \citet{b42}.\\ The paper is organized as follows. In Section 2 we describe the chemical gas-grain model Nautilus and more particularly the grain-surface reactions. Models predictions are presented in Section 3, while comparisons with observations in two dark clouds (TMC-1 (CP) and L134N) are shown in Section 4. We present our conclusions about this work in the last section. | In this paper, we have studied the diffusion of species on the grain surfaces in molecular clouds using our gas-grain model Nautilus. In particular, we revisited the efficiency of diffusion by tunneling effect of H, H$_{2}$ and O based on recent experiments by \citet{b1}. We also introduced a new mechanism to take into account the "boost" of surface species diffusion induced by the stochastic heating by cosmic ray particles. We call this new mechanism Cosmic Rays Induced Diffusion (CRID). Species most affected by those two mechanisms are listed in Table A1. The diffusion by tunnel effect allows surface radicals to move faster at lower temperature which increased the abundances of simple and complex species (in the gas-phase and at the surface of the grains), such as HOOH, CH$_{3}$OCH$_{3}$, for typical dense cloud physical parameters and ages. This effect becomes even more important for temperatures below 10$\,$K. The newly introduced CRID mechanism increases the mobility of radicals at the surface of the grains, which can recombine with each other and form more complex molecules. However, we notice that this mechanism is only efficient when the visual extinction is smaller or equal to 3 for 0.1$\,$$\mu$m grains. For such values of visual extinction, we show that complex species abundances can be increased significantly while those of simpler molecules are decreased. This is because the photodissociation of molecules at the surface of the grains produces more radicals that can then react and their formation rate is then larger than their destruction one. The general agreement of our chemical model with molecular abundances observed in the two clouds TMC-1CP and L134N is not significantly changed by considering the diffusion by tunneling of H, H$_{2}$ and O because most of the observed species are not affected by this process. Only three are significantly sensitive: HNCO, CH$_{2}$CHCN and CH$_{3}$OH. For those species, the predicted gas-phase abundances are larger and come closer to the observations. Furthermore, on the surface, their abundances are much larger than the gas-phase observed ones for CH$_{3}$OH and HNCO. For the complex species observed in pre-stellar cores by \citet{b24} and \citet{b25}, the diffusion of atomic oxygen by tunneling is not important enough to reproduce the large gas-phase abundances for HCOOCH$_{3}$, CH$_{3}$OCH$_{3}$, CH$_{2}$CO and CH$_{3}$O. Even on the surfaces, their abundances stay at a level lower than the gas-phase observed ones. | 14 | 3 | 1403.5189 |
1403 | 1403.7974_arXiv.txt | We have carried out a deep (150 micro Jy rms) P-band, continuum imaging survey of about 40 square degrees of sky in the XMM-LSS, Lockman Hole and ELAIS-N1 fields with the GMRT. Our deep radio data, combined with deep archival observations in the X-ray (XMM/Chandra), optical (SDSS, CFHTLS), near-infrared (UKIDSS, VISTA/VIDEO), mid-infrared (Spitzer/SWIRE, Spitzer/SERVS) and far-infrared (Spitzer/SWIRE, Herschel/HerMES) will enable us to obtain an accurate census of star-forming and active galaxies out to $z\sim 2$. This panchromatic coverage enables accurate determination of photometric redshifts and accurate modeling of the spectral energy distribution. We are using our large, merged photometric catalog of over 10000 galaxies to pursue a number of science goals. | \label{s:intro} Understanding the mass assembly history of galaxies remains a major challenge in astrophysics. This is because a major fraction of galaxy assembly happens at high redshifts, in intense bursts of star formation as well as black hole accretion. Therefore, disentangling starburst and AGN activity, and understanding why the peak in the comoving luminosity density of star formation coincides with that from AGN activity at $z\sim2$, is vital for measuring the stellar/BH mass buildup. Since most of the energy from these activities is absorbed by dust and then re-radiated in the rest-frame infrared, comprehensive multi-wavelength studies are required to achieve this, with emphasis on far-IR/submillimeter observations which probe the peak in the IR emission from moderately warm dust, supplemented by radio surveys which can probe the inner regions of star formation and AGN activity, without being adversely affected by dust obscuration. The last couple of years have seen the arrival of revolutionary FIR/submm surveys, on a large new FIR telescope -- {\em Herschel}. Among those, the Herschel Multi-tiered Extragalactic Survey (HerMES\footnote{http://www.hermes.sussex.ac.uk}) is the largest (900\,hr) guaranteed time program. It aims to chart the evolution of galaxies through cosmic history via a set of nested (wide and shallow as well as deep pencil-beam) surveys over twelve well studied areas of sky, in particular, those extensively studied by {\it Spitzer} at 24 and 70 $\mu$m, such as the ELAIS-N1, Spitzer FLS, CDFS, GOODS-N, Lockman and XMM-LSS fields (total area coverage of $\sim$110 deg$^2$ plus 270 deg$^2$ of shallow survey). As demonstrated by the results already published in about 40 refereed papers (see the HerMES website for a list of these papers), HerMES has provided new insight on distant dusty galaxies and AGNs. It seems that previous phenomenological galaxy populations need revision and it is now anticipated that HerMES will be able to catalogue over 100,000 galaxies with $>5\sigma$ detections at 250 $\mu$m. HerMES will constitute a lasting legacy to the community, providing an essential complement to multiwavelength surveys in the same fields and providing targets for follow-up using many facilities for many years to come. | 14 | 3 | 1403.7974 |
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1403 | 1403.2626_arXiv.txt | The dynamical and radiative properties of the quiescent state (X-ray luminosity $\lesssim10^{34}\ \ergs$) of black hole X-ray transients (BHXTs) remains unclear, mainly because of low-luminosity and poor data quantity. We demonstrate that, the simultaneous multi-wavelength (including radio, optical, ultraviolet and X-ray bands) spectrum of V404 Cyg in its bright quiescent state can be well described by the radiation from the companion star and more importantly, the compact jet. Neither the outer thin disc nor the inner hot accretion flow is important in the total spectrum. Together with recent findings, i.e. the power-law X-ray spectrum and the non-variable X-ray spectral shape (or constant photon index) in contrast to the dramatic change in the X-ray luminosity, we argue the quiescent state spectrum of BHXTs is actually jet-dominated. Additional observational properties consistent with this jet model are also discussed as supporting evidences. | Soft X-ray transients are binary systems, where the primaries (black holes or neutron stars) accrete material from their companion low-mass stars through Roche lobe overflow. The black hole X-ray transients generally undergo occasional outbursts, during which they exhibit distinctive states (soft, hard and intermediate) according to their spectral and timing properties \citep{zg04,hb05,rm06,d07,b10,zh13}, likely the consequence of the changes in geometry and radiation mechanism of the accretion flow (e.g. \citealt{e97}). The soft state can be well described by cold disc model (\citealt{ss73}, hereafter SSD). The hard state is now generally understood under the hot accretion-jet scenario (see \citealt{yn14} [YN14 hereafter] for an up-to-date review on hot accretion flow and its applications on various objects including BHXTs), developed from the truncated disc model originally proposed by \citet{e97}. Three components are involved in this model, 1) an outer SSD, which is truncated at radius $R_{\rm tr}$, 2) an inner hot accretion flow within such radius, and 3) a relativistic jet. The hot accretion flow generally stands for the radiativelly inefficient, advection-dominated accretion flow (ADAF; \citealt{ny94}), and it is updated by the recent progresses in accretion theory (see \citealt{xy12}; YN14 for summaries), i.e. the existence of outflow and the direct viscous heating to electrons. This model provides comprehensive explanations on the spectral and timing properties of BHXTs in their hard states (e.g. \citealt{y05a} [YCN05 hereafter]). For majority of the time BHXTs are extraordinarily faint, characterised by the so-called ``quiescent state'', with the X-ray ($2-10\ {\rm keV}$) luminosity $L_{\rm X} \lesssim 10^{34} \ergs$. Mainly because of the difficulties in detection, the nature of the quiescent states remains unclear \citep{n02,nm08}. For example, although the origin of the X-ray in quiescent state should come from the high-energy electrons near black holes, it is still unclear whether it is the synchrotron radiation from the non-thermal electrons in the jet (e.g. \citealt{f03,y05b}), or the synchrotron self-Comptonization from the thermal electrons in ADAF (e.g. \citealt{n97}), or Comptonization (in jet) of seed photons from cold disc \citep{m05,k08}. From the accretion-jet model, it is shown that \citep{y05b}, because the ADAF emission in X-rays decreases faster than the jet emission with decreasing accretion rate $\mdot$, the X-ray radiation $L_{\rm X}$ may be dominated by the jet emission when $\mdot$ is below a certain threshold. In other words, the X-ray emission of the quiescent state of BHXTs is likely to from the jet. This prediction has passed several observational tests, mainly through spectral fitting (\citealt{p08}; YCN05). However, these tests suffer the shortage of poor observational data (the data are non-simultaneous and/or limited to narrow wavebands), thus lead to some debates. We in this {\it Letter} aim to argue that the quiescent state of BHXTs can be well-characterised by the jet model. We first in \S \ref{jetmodel} give a brief description of the jet model we used, and then in \S \ref{specfit} provide a comprehensive broad band (from radio to X-ray) spectral fitting of the quiescent spectrum of V404 Cyg, where simultaneous X-ray, ultraviolet (UV), optical and radio observations are available. We then in \S \ref{evidence} provide additional observational supports. Finally we provide some discussions and a brief summary in \S \ref{summary}. | \label{summary} One consequence of the jet-dominated quiescent state is that, the radio/X-ray luminosity relationship steepens to $b\approx 1.23$ (where $L_{\rm R} \propto L_{\rm X}^b$; \citealt{y05b}), compared to $b\approx 0.62$ typically observed in the hard states of BHXTs \citep{c03,c13}. However, such steepen has not been observed in BHXTs yet. For V404 Cyg, our spectral modelling indicates it enters the steep correlation regime and the radio flux will be $(3-6)$ times lower to that extrapolated from the $b\sim 0.6$ correlation, a difference not be observed \citep{c08}. Further efforts are still needed to understand this discrepancy. On the other hand, we note that observationally the low-luminosity AGNs (which are analogy to BHXTs in their hard and quiescent states) do exhibit a steeper correlation, with $b\approx 1.22$, in excellent agreement with the theoretical prediction \citep{y09}. Besides, if the X-ray radiation of the quiescent state has a jet-origin, then, we would expect to have relatively high degree of polarization in X-ray band, compared to the typical hard state, where the X-rays are of inverse-Compton origin. Currently the X-ray polarization is nearly an unexplored field in astronomy, and future sensitive spaceborne X-ray polarimeters (e.g. the proposed X-ray Timing and Polarization [XTP] mission and the Gravity and Extreme Magnetism SMEX X-ray polarimetry mission) may detect variable X-ray polarization from synchrotron emission of the jets. We now give a brief summary of this work. With recent advances in X-ray observations, especially the deep sensitive spectral observation and the long-term monitoring, we confirm that the emission of the quiescent state of BHXTs is dominated by the radiation from the compact relativistic jet. The outer thin disc and the inner hot accretion flow generally play negligible roles in radiation, while the companion dominates the emission between mid-IR and optical bands. We also illustrate (cf. \S \ref{evidence}) that the jet-dominated quiescent state model can explain most of the observational features. \vspace{-0.5cm} | 14 | 3 | 1403.2626 |
1403 | 1403.6834_arXiv.txt | To date, more than 750 planets have been discovered orbiting stars other than the Sun. Two sub-classes of these exoplanets, ``hot Jupiters" and their less massive counterparts ``hot Neptunes," provide a unique opportunity to study the extended atmospheres of planets outside of our solar system. We describe here the first far-ultraviolet transit study of a hot Neptune, specifically GJ436b, for which we use \emph{HST}/STIS Lyman-$\alpha$ spectra to measure stellar flux as a function of time, observing variations due to absorption from the planetary atmosphere during transit. This analysis permits us to derive information about atmospheric extent, mass-loss rate from the planet, and interactions between the star and planet. We observe an evolution of the Lyman-$\alpha$ lightcurve with a transit depth of GJ436b from $8.8\pm4.5\%$ near mid-transit, to $22.9\pm3.9\%$ $\sim2$ hours after the nominal geometric egress of the planet. Using data from the time-tag mode and considering astrophysical noise from stellar variability, we calculate a post-egress occultation of $23.7\pm4.5$\%, demonstrating that the signature is statistically significant and of greater amplitude than can be attributed to stellar fluctuations alone. The extended egress absorption indicates the probable existence of a comet-like tail trailing the exoplanet. We calculate a mass-loss rate for GJ436b in the range of $3.7\times10^6 -1.1\times10^{9}$ g s$^{-1}$, corresponding to an atmospheric lifetime of $4\times10^{11}-2\times10^{14}$ years. | Since the mid-1990's, astronomers have been regularly discovering planets orbiting other stars using the radial velocity method \citep{Mayor_1995}. Many of these planets called ``hot Jupiters" are massive ($\gtrsim50~ \mathrm{M}_\oplus \gtrsim 0.15 ~\mathrm{M}_\mathrm{Jup}$), orbit very close to their host stars ($\le$ 0.05 AU), and have orbital periods of only a few days \citep{Seager_2010}. Less massive planets ($10-50~ \mathrm{M}_\oplus \approx 0.03-0.15~ \mathrm{M}_\mathrm{Jup} \approx 0.6-3~ \mathrm{M}_\mathrm{Nep}$), that also orbit very close to their host stars, are similarly named after their solar system analogs, ``hot Neptunes.'' The majority of the first exoplanet detections were ``hot Jupiters'' and ``hot Neptunes'' because there is a detection bias in favor of this type of planet. Since they are large (in mass and radius) and orbit close to their host stars, they are more easily detected in both radial velocity and transit searches. In 1999, \citet{Henry_2000} discovered the first transiting exoplanet, HD209458b. Transiting exoplanets provide an opportunity to study the composition and structure of their atmospheres, because as these planets pass in front of their host star, their atmosphere blocks a portion of the starlight. Spectroscopy at IR and optical wavelengths has led to the discovery of \ion{Na}{1}, H$_2$O, CH$_4$, and CO (\citealt{Charbonneau_2002}; \citealt{Tinetti_2007}; \citealt{Swain_2008}; \citealt{Swain_2009}) in the atmospheres of exoplanets, shifting the emphasis of exoplanet studies toward detecting and characterizing their atmospheres \citep{Koskinen_2010}. While many studies utilize IR ($700\mathrm{~nm}-3~\mu$m) (e.g. \citealt{Deming_2007}) and optical ($400-700$ nm) (e.g. \citealt{Butler_2004}) wavelengths, UV ($1000-4000$ \AA) observations provide a unique opportunity to characterize exoplanet exospheres. Stellar FUV and EUV radiation heats, accelerates, and, when the photon energies exceed 13.6 eV, ionizes the hydrogen in the upper atmosphere \citep{Murray_2009}. When the thermal energy exceeds the gravitational potential energy ($3kT/2 > GMm/r$), the heated gas expands \citep{Lammer_2003} and likely escapes from the planet. These inflated atmospheres contain many atomic species, such as H, C$^+$, O, and Si$^{2+}$, that absorb the stellar radiation in UV resonance lines \citep{Vidal-Madjar_2004,Linsky_2010}. Because the envelopes of these exoplanets are inflated, the geometric area in their UV resonance lines is larger than that of molecules detected in the IR and optical, and the transit depth will be larger. Such data permit us to understand the composition of extended exoplanet atmospheres. \citet{Vidal-Madjar_2003} used the \ion{H}{1} Lyman-$\alpha$ 1215 \AA~emission line of HD209458 to obtain the first evidence of an inflated exoplanetary atmosphere, revealing a transit depth in Lyman-$\alpha$ ($15\pm4\%$) several times that of the geometric depth determined from the optical transit ($1.58\pm0.18\%$). The Lyman-$\alpha$ transit depth indicates an atmosphere that extends beyond its Roche lobe, likely indicating atmospheric escape. Subsequent UV observations have demonstrated the existence of enlarged envelopes in the spectral lines of O, C$^+$ \citep{Vidal-Madjar_2004}, and Si$^{2+}$ \citep{Linsky_2010} around HD209458b, and later observations identified a similar absorption spectrum from the atmosphere of HD189733b. Observations of HD189733 have been used to detect extended atmospheres in lines of \ion{H}{1} \citep{Lecavelier_2010}, \ion{O}{1}, and possibly \ion{C}{2} \citep{Ben-Jaffel_2013}. The non-detection of an extended Lyman-$\alpha$ envelope in April 2010, indicated significant variations in the evaporation rate of \ion{H}{1} \citep{Lecavelier_2012}. UV transit observations are often used to infer mass-loss rates of hot Jupiters, from which we may learn about the evolution of exoplanets and their atmospheres. \citet{Vidal-Madjar_2003} used a particle simulation to limit the mass loss rate of HD209458b to $\ge 10^{10}$ g s$^{-1}$. \citet{Linsky_2010} performed a theoretical calculation to find mass-loss rates in the range of $(8-40) \times 10^{10}$ g s$^{-1}$. \citet{Lecavelier_2010} used their data of HD189733b and a numerical simulation to find a best fit mass-loss rate of $10^{10}$~g~s$^{-1}$. \citet{EhrenreichDesert_2011} compare the planet gravitational potential energy to the stellar X/EUV energy deposited in the atmosphere and estimate the mass-loss rate for WASP-12b to be $2.5\times10^{11}$ g s$^{-1}$. Many theoretical models have been developed to explain the evaporation process and to predict the mass-loss rates for various hot Jupiters. \citet{Lammer_2003} used the heating rate from stellar X-ray and EUV radiation to estimate a mass-loss rate for HD209458b of $\sim10^{12}$ g s$^{-1}$. \citet{Yelle_2004} (later revised in \citealt{Yelle_2006}) included chemical calculations to estimate a mass-loss rate. Cooling from H$^+_3$ and ionization of H to H$^+$, reducing the amount of stellar energy available for heating, led to a lower mass-loss rate of $4.7\times10^{10}$ g s$^{-1}$. Similar calculations by \citet{Garcia_2007} determined mass-loss rates in the range $6-15\times10^{10}$ g s$^{-1}$ depending on the level of stellar activity. \citet{Holmstrom_2008} attributed a portion of the Lyman-$\alpha$ absorption to protons from the stellar wind that have been neutralized by charge exchange with hydrogen atoms in the planetary atmosphere and calculated a mass-loss rate of $7\times10^8$ g s$^{-1}$. \citet{Guo_2013} calculated a mass-loss rate of $4.3\times10^9$ g s$^{-1}$ using a two-dimensional model that assumed HD209458b is tidally locked with one side of the planet always facing the star. \citet{Murray_2009} modeled the atmospheric escape of a theoretical hot Jupiter (similar to HD209458b) that includes realistic heating and cooling rates, ionization balance, tidal gravity, and pressure confinement by the stellar wind. They found a mass-loss rate of $2\times10^{10}$ g s$^{-1}$. In the above examples, the calculated mass-loss rates vary by several orders of magnitude depending on what physics is considered and on the values of system parameters (for example the amount of stellar EUV flux). The manner in which the atmospheric envelope interacts with the stellar wind determines the structure of the gas around the planet. High-velocity neutral hydrogen escaping from the atmosphere trails the planet absorbing Lyman-$\alpha$ photons until the hydrogen is ionized by the stellar EUV flux or charge exchanges with stellar wind protons. The trailing material may cause the hydrogen Lyman-$\alpha$ transit to last longer than the optical occultation. The formation of a comet-like tail trailing the planet was first suggested by \citet{Schneider_1998}. Models of neutral hydrogen escaping from the hot Jupiter HD209458b and HD189733b by \citet{Bourrier_2013} support this structure. The evaporating super-Mercury exoplanet KIC 12557548b likely has a dusty comet-like tail \citep{Rappaport_2012, Budaj_2013} and the extreme hot Jupiter WASP-12b is losing sufficient mass to completely obscure its host star's emission in the cases of the \ion{Mg}{2} h and k lines \citep{Haswell_2012}. \citet{Vidal-Madjar_2003} observed absorption out to Doppler velocities $\pm100$ km s$^{-1}$ from the Lyman-$\alpha$ line center of HD209458b, and absorption has been seen at velocities as large as -230 km s$^{-1}$ from the center of Lyman-$\alpha$ for HD189733b \citep{Lecavelier_2012}. Thermal velocities can only account for absorption out to $\sim$10 km s$^{-1}$ \citep{Murray_2009}. Models by \citet{Lecavelier_2012} suggest that stellar radiation pressure can accelerate particles up to 120 km s$^{-1}$, but an additional mechanism is necessary to explain the large observed radial velocities. Charge exchange with hot, slow ($<50-100$ km s$^{-1}$) stellar wind protons can produce the observed velocities (\citealt{Holmstrom_2008, Ekenback_2010, Tremblin_2013}). \citet{Holmstrom_2008} proposed that the absorption in neutral hydrogen at high velocities is due to charge exchange between protons from the stellar wind and planetary neutral hydrogen. The planetary hydrogen is ionized and the stellar wind proton becomes neutral hydrogen maintaining its large velocity. \subsection{Previous Studies of GJ436b} A promising exoplanet for UV transit observations is GJ436b, a hot Neptune, ($\mathrm{M}=0.07~\mathrm{M}_\mathrm{Jup}$), orbiting an M2 dwarf at 0.029 AU with a period of 2.6 days. GJ436b is particularly promising because it is very close to Earth, at a distance of only 10.2 pc. The properties of the system are summarized in Table~\ref{star_dat}. GJ436b was discovered by \citet{Butler_2004} using the radial velocity method, but was later found to be transiting its host star \citep{Gillon_2007}. In 2010 \emph{Spitzer} observed GJ436 during several secondary transits in 6 bands from $3.6-24~\mu$m. Using a Metropolis-Hastings Markov-chain Monte Carlo (MCMC) model, \citet{Stevenson_2010} explored a wide range of parameter space to determine the best-fit compositional models. They found a high CO abundance and a deficiency of CH$_4$ relative to thermochemical equilibrium. \citet{Madhusudhan_2011} also used MCMC to confirm the overabundance of CO and CO$_2$, and a slight underabundance of H$_2$O, as compared to equilibrium chemistry with solar metallicity. They explained the observed abundances by a combination of high metallicity ($\sim10\times$ solar) and vertical mixing. Observations by \citet{Knutson_2014} indicate an effectively featureless transmission spectrum, ruling out cloud-free, hydrogen-dominated atmosphere models. The measured spectrum is consistent with either a high cloud or haze layer or with a relatively hydrogen-poor atmospheric composition. \citet{Hu_2014} find that hot Neptunes, like GJ436b, are likely to have thick atmospheres that are not hydrogen dominated, but are water-rich or hydrocarbon-rich depending on their C/O ratio. Using limited \emph{HST}/STIS data of the stellar Lyman-$\alpha$ flux, \citet{Ehrenreich_2011} developed numerical simulations to determine the transit signature of GJ436b for various assumed mass-loss rates. They predicted an 11\% transit depth in Lyman-$\alpha$ for a mass-loss rate of $10^{10}$ g s$^{-1}$. The analysis we have conducted is the first to place observational constraints on the mass-loss rate of GJ436b or any other Neptune-mass exoplanet. In this paper, we present and analyze transit observations of GJ436b in Lyman-$\alpha$ observed with Space Telescope Imaging Spectrograph (STIS) on the \emph{Hubble Space Telescope} (\emph{HST}). In Section 2 we describe the data sets used and the reduction process, and how we created the lightcurves and velocity profiles. We present the lightcurve and velocity profile for Lyman-$\alpha$ in Section 3 along with the measured absorption depths. In Section 4 we discuss the structure of the system and calculate a range of mass-loss rates for GJ436b. We summarize our results in Section 5. \begin{deluxetable*}{lccc} \tablecaption{Properties of GJ436 and GJ436b\label{star_dat}} \tablewidth{0pt} \tablehead{ \colhead{Property} & \colhead{Value} & \colhead{Reference} } \startdata Host Star Spectral Type & M2 V &\citet{Butler_2004}\\ Distance (pc) & $10.14\substack{+0.25 \\ -0.23}$ &\citet{vanLeeuwen_2007}\\ $M_*/M_\odot$ & $0.452\substack{+0.014\\-0.012}$ &\citet{Torres_2008}\\ $R_*/R_\odot$ & $0.464\substack{+0.009\\-0.011}$ &\citet{Torres_2008}\\ $P_{orbit}$(days) & $2.643850\pm9\times10^{-5}$ &\citet{Pont_2009}\\ Transit Center (JD) & $2454279.436714\pm1.5\times10^{-5}$ &\citet{Pont_2009}\\ RA (h:m:s) & +11:42:11.18 &\citet{Zacharias_2012}\\ Dec (d:m:s) & +26:42:22.64 &\citet{Zacharias_2012}\\ $R_{planet}/R_{Jup}$ & $0.3767\substack{+0.0082\\-0.0092}$ &\citet{Torres_2008}\\ $M_{planet}/M_{Jup}$ & $0.0727\pm0.0032$ &\citet{Butler_2006}\\ Semimajor axis (AU) & $0.02872\pm0.0048$ &\citet{Butler_2006} \\ Transit duration (hours) & $0.7608\pm0.012$ &\citet{Pont_2009}\\ Transit depth for $R_{planet}$ & $0.00696\pm0.000117$ &\citet{Torres_2008}\\ Escape speed from GJ436b (km s$^{-1}$) & 26.4 &\\ Orbital velocity amplitude (km s$^{-1}$) & 118 &\\ Radial Velocity (km s$^{-1}$) & $9.6\pm0.1$ &\citet{Nidever_2002} \enddata \end{deluxetable*} | \subsection{GJ436b Transit} We show the Lyman-$\alpha$ lightcurve for GJ436 in Figure \ref{LyA_int} and the Lyman-$\alpha$ velocity profile in Figure \ref{diff}. Table~\ref{GJ436_tr_dep} shows the transit depths extracted from these figures. The transit depths are identical for the two procedures (because the reference spectrum and the integration method are the same), however, the spectral analysis results in somewhat larger errors, due to the extra step of subtracting prior to normalizing and integrating. We will use the higher time resolution lightcurve data to analyze the transit depth. For GJ436b we see a mid-transit depth of $16.6\pm7.2\%$ in the blue wing and $4.5\pm5.7\%$ in the red wing. This corresponds to an occulting disk of 5.0 R$_\mathrm{p}$ and 2.6 R$_\mathrm{p}$ respectively, both smaller than the Roche lobe radius of 6.1 R$_\mathrm{p}$. When both wings are combined, the mid-transit depth is $8.8\pm4.5\%$, corresponding to an equivalent opaque occulting disk of 3.6 R$_\mathrm{p}$. The asymmetry in the absorption can be explained by charge exchange of the stellar wind with the atmosphere. As viewed from Earth, only the stellar wind traveling towards us can be observed, causing excess absorption blue-ward of line center. This type of asymmetry, with more absorption in the blue wing, is also seen in the Lyman-$\alpha$ absorption during transit from HD209458b \citep{Vidal-Madjar_2003} and HD189733b \citep{Lecavelier_2012}. Interestingly, the Lyman-$\alpha$ transit extends much later in phase than the optical transit. We examine these data for the possibility of extended egress, finding post-egress depths of $29.9\pm6.4\%$ in the blue wing, $19.1\pm5.0\%$ in the red wing, and $22.9\pm3.9\%$ for both wings combined $\sim2$ hours after mid-transit and about 1.5 hours after fourth contact. The time-tag data points are presented in Table~\ref{GJ436_tr_dep_tt}. These data corroborate the detection of both the transit and extended egress, as three of the four transit data points and all of the post-egress data points show a transit detection, although large variations in the blue-wing time-tag data are observed near mid-transit. \begin{figure} \epsscale{1} \plotone{fig2.eps} \caption[Normalized Lyman-$\alpha$ count rates for GJ436]{ Normalized Lyman-$\alpha$ count rates for GJ436. Blue points show the flux from the blue wing of Lyman-$\alpha$ and red points show the flux from the red wing. Filled points are calculated from the time-tag data, while the open circles are from the entire exposure. Error bars indicate $\pm1\sigma$. The black dotted line indicates the normalized flux level, while the black dashed line shows the transit curve as calculated from the optical transit parameters. The green and pink lines shows the predicted transit signature of GJ436b as calculated by \citet{Ehrenreich_2011} for mass-loss rates of $10^9$ and $10^{10}$ g s$^{-1}$ respectively. Vertical dashed lines indicate first and fourth contacts. } \label{LyA_int} \end{figure} \begin{figure} \plotone{fig3.eps} \caption[Lyman-$\alpha$ velocity profile for GJ436]{ Lyman-$\alpha$ velocity profile for GJ436. The top panel compares the pre-ingress spectrum, in dark red, to the post-egress spectrum, in dark blue, and the bottom panel shows the difference between these spectra, pre-ingress minus post-egress. Error bars indicate $\pm1\sigma$. The regions of integrated flux are also shown; the blue wing is between the blue dashed lines and the red wing is between the red dashed lines. During the post-egress time interval, we find an occultation depth of $29.9\pm8.3\%$ in the Lyman-$\alpha$ blue wing, $19.1\pm5.8\%$ in the Lyman-$\alpha$ red wing, and $22.9\pm4.8\%$ in the combined wings of the Lyman-$\alpha$ line. } \label{diff} \end{figure} \begin{deluxetable}{lcc} \tablecaption{Transit Depths for GJ436 Full Exposures\label{GJ436_tr_dep}} \tablewidth{0pt} \tablehead{ & \colhead{Transit Depth from} & \colhead{Transit Depth} \\ \colhead{Species}& \colhead{Difference Spectrum}& \colhead{from Lightcurve} } \startdata Lyman-$\alpha$ Blue Wing mid-transit & $16.6\pm8.2$\% & $16.6\pm7.2$\% \\ Lyman-$\alpha$ Red Wing mid-transit & $4.5\pm5.9$\% & $4.5\pm5.7$\% \\ Lyman-$\alpha$ coadded mid-transit & $8.8\pm4.8$\% & $8.8\pm4.5$\% \\ \\ Lyman-$\alpha$ Blue Wing post-egress & $29.9\pm8.3$\% & $29.9\pm6.4$\% \\ Lyman-$\alpha$ Red Wing post-egress & $19.1\pm5.8$\% & $19.1\pm5.0$\% \\ Lyman-$\alpha$ coadded post-egress & $22.9\pm4.8$\% & $22.9\pm3.9$\% \enddata \end{deluxetable} \begin{deluxetable}{lcc} \tablecaption{Transit Depths for GJ436 Time-Tag Points\label{GJ436_tr_dep_tt}} \tablewidth{0pt} \tablehead{ & \colhead{Transit Depth in} & \colhead{Transit Depth in} \\ \colhead{Phase (hr)}&\colhead{Blue Wing}& \colhead{Red Wing} } \startdata -1:08:54.7 & $ 5.6\pm12.3$\% & $-5.7\pm10.7$\% \\ -0:56:54.7 & $-5.6\pm13.4$\% & $ 5.7\pm10.9$\% \\ 0:03:12.3 & $-1.2\pm12.4$\% & $ 4.2\pm9.2$\% \\ 0:15:12.3 & $ 18.0\pm10.6$\% & $-5.9\pm10.0$\% \\ 0:27:12.3 & $ 15.5\pm11.8$\% & $15.7\pm10.1$\% \\ 0:39:12.3 & $ 40.2\pm11.5$\% & $13.9\pm10.4$\% \\ 1:38:55.3 & $ 36.5\pm10.0$\% & $19.7\pm8.5$\% \\ 1:50:55.3 & $ 31.3\pm10.7$\% & $11.9\pm8.8$\% \\ 2:02:55.3 & $ 30.0\pm10.7$\% & $20.8\pm9.1$\% \\ 2:14:38.5 & $ 31.9\pm12.3$\% & $24.0\pm9.3$\% \\ 3:14:38.4 & $ 27.6\pm10.7$\% & $-3.2\pm9.7$\% \\ 3:26:38.4 & $ 16.1\pm11.9$\% & $19.5\pm8.7$\% \\ 3:38:38.4 & $ 16.4\pm11.7$\% & $19.7\pm9.6$\% \\ 3:50:38.4 & $-15.8\pm15.0$\% & $ 9.7\pm10.7$\% \enddata \end{deluxetable} \subsection{Extended Egress in GJ436\label{egress}} To determine whether or not the extended egress is real, we consider the time-tag data. We form the null hypothesis that the data are Gaussian distributed and the apparent occultation is random noise. We use the average of the error bars on the time-tag data as the standard deviation for the Gaussian distributions ($\sigma_{blue}=0.1092, \sigma_{red}=0.08930$), both with a mean of unity. Assuming these distributions, we determined the probability of finding four consecutive points lower than the highest point in the set of four at the deepest transit depth, which is located at phase $\sim$2 hours ($x_{blue}=0.3003, x_{red}=0.1187$ below the mean). We did this by randomly picking 14 points from the specified Gaussian distribution and counting how many trials out of $10^7$ had four consecutive points outside the requisite range. For the parameters determined for the blue wing, none of the $10^7$ trials had four consecutive points. For the red wing distribution, we found a probability of 9.31$\pm0.76\times10^{-5}$ to randomly produce the result. We therefore conclude that the deep, extended egress signal seen in Figure \ref{LyA_int} is real and not due to random statistical variations. \subsection{Stellar Variability of GJ436\label{variability}} While we have determined that the post-egress detection is not due to statistical noise, we have not yet considered whether the drop in flux could be due to stellar variability, as opposed to atmospheric absorption from GJ436b. To address this issue, we look at a resonance line from \ion{N}{5}, an ion that we do not expect to find in the atmosphere of the exoplanet. For the \ion{N}{5} time-tag data points we find $\mathrm{RMS}=0.2765$, while the average value of the $1\sigma$ error bar for those points is 0.4010. We conclude that the signal to noise is too low for the \ion{N}{5} doublet to be a suitable tracer of stellar variability. Similarly, there was not enough flux in the \ion{Si}{3} line to be measurable above the noise. Instead we look to the literature to assess the potential magnitude of stellar variability. \citet{Loyd_2014} studied time variability in the \ion{C}{2}, \ion{Si}{3}, and \ion{Si}{4} resonance lines of 38 cool stars, including GJ436. They did not attempt to characterize Lyman-$\alpha$ line variability because geocoronal airglow cannot be removed reliably from their COS data. Instead we use their \ion{C}{2} emission line variability, because \ion{C}{2} has a similar formation temperature to Lyman-$\alpha$ (T$_{form}\approx(1-3)\times10^4$ K, \citealt{Dere_2009}). \citet{Loyd_2014} found that the mean-normalized chromospheric \ion{C}{2} line variability, excluding flare periods, in GJ436 is $0.20^{0.09}_{0.12}$ on 60 second timescales. Our post-egress data cover a 48 minute time span. Assuming that the stellar variability is uncorrelated over time, the noise associated with stellar fluctuations is estimated to be $0.20/\sqrt{48}=0.0289$. Combining the 8 post-egress time-tag data points and the photon noise with the upper limit on the noise expected from chromospheric variability, we find a 23.7\% occultation with an uncertainty of 4.5\%. From this we conclude that the post-egress detection is very likely real. Future observations over several transit cycles would be very valuable. We have analyzed new observations of GJ436. We used \emph{HST}/STIS data to detect and characterize the extended atmosphere of GJ436b for the first time. We detected $8.8\pm4.5\%$ absorption in the Lyman-$\alpha$ line at mid-transit, and used this transit depth to calculate a mass-loss rate in the range $3.7\times10^6-1.1\times10^{9}$ g s$^{-1}$, corresponding to an atmospheric lifetime of $4\times10^{11}-2\times10^{14}$ years. We also detected strong absorption after the optical transit with a depth of $23.7\pm4.5\%$. We confirmed that this extended egress is not a statistical fluctuation, and showed that it is unlikely to be due to stellar variability; the most likely explanation is that GJ436b is trailed by a comet-like tail of neutral hydrogen. This material is based upon work supported by the National Science Foundation Graduate Research Fellowship Program under Grant No. DGE 1144083. The data presented here were obtained as part of \emph{HST} Observing program \#12034. KF acknowledges support through a NASA Nancy Grace Roman Fellowship during this work. | 14 | 3 | 1403.6834 |
1403 | 1403.1175.txt | We have identified ionized outflows in the narrow line region of all but one SDSS type 2 quasars (QSO2) at $z\la$0.1 (20/21, detection rate 95\%), implying that this is a ubiquitous phenomenon in this object class also at the lowest $z$. The outflowing gas has high densities ($n_e\ga$1000 cm$^{-3}$) and covers a region the size of a few kpc. This implies ionized outflow masses $M_{outf}\sim$(0.3-2.4)$\times$10$^6$ \msun~ and mass outflow rates $\dot M<$few \msun~ yr$^{-1}$. The triggering mechanism of the outflows is related to the nuclear activity. The QSO2 can be classified in two groups according to the behavior and properties of the outflowing gas. QSO2 in Group 1 (5/20 objects) show the most extreme turbulence, they have on average higher radio luminosities and higher excess of radio emission. QSO2 in Group 2 (15/20 objects) show less extreme turbulence, they have lower radio luminosities and, on average, lower or no radio excess. We propose that two competing outflow mechanisms are at work: radio jets and accretion disk winds. Radio jet induced outflows are dominant in Group 1, while disk winds dominate in Group 2. We find that the radio jet mode is capable of producing more extreme outflows. To test this interpretation we predict that: 1) high resolution VLBA imaging will reveal the presence of jets in Group 1 QSO2; 2) the morphology of their extended ionized nebulae must be more highly collimated and kinematically perturbed. | Evidence for an intimate connection between supermassive black hole (SMBH) growth and the evolution of galaxies is nowadays compelling. Not only have SMBHs been found in many galaxies with a bulge component, but correlations exist between the black hole mass and some bulge properties, such as the stellar mass and velocity dispersion (e.g. Ferrarese \& Merritt \citeyear{fer00}). The origin of this relation is still an open question, but quasar induced outflows might play a critical role. Hydrodynamical simulations show that the energy output from quasars can regulate the growth and activity of black holes and their host galaxies (e.g. di Matteo, Springel, Hernquist \citeyear{dim05}). Such models also show that the energy released by the strong outflows associated with major phases of accretion expels enough gas to quench both star formation and further black hole growth. Growing observational evidence for the dramatic impact that quasar induced outflows may have on their host galaxies is accumulating. Such impact is likely to depend on the luminosity of the active galactic nucleus (AGN) and to be more efficient at the highest luminosities (Page et al. \citeyear{pag12}). Studies of powerful radio galaxies show that the radio structures can induce powerful outflows (Humphrey et al. \citeyear{hum06}) which may have enormous energies sufficient to eject a large fraction of the gaseous content of the galaxy (Nesvadba et al. \citeyear{nes06}, Morganti, Tadhunter \& Oosterloo \citeyear{mor05}). However, only $\sim$10\% of active galaxies are radio loud. How AGN feedback works in radio quiet objects is still an open question. Several mechanisms could be at work, including stellar winds, radio jet induced outflows and accretion disc winds. Even in radio quiet quasars the presence of radio jet induced winds cannot be discarded. Indeed, the existence of such jets in many (all?) radio quiet AGN has been proposed by different authors (e.g. Ghisellini, Haardt \& Matt \citeyear{ghi04}, Ho \& Peng \citeyear{ho01}, Falcke \& Biermann \citeyear{fal95}, Malzac et al. \citeyear{mal98}) even if in many of these systems the jet is likely to be aborted. This is supported by the VLBI imaging of radio quiet Seyfert galaxies (Ulvestad \citeyear{ulv03} and references therein) which revealed the presence of mini-jets (sub parsec scale) in many of them. Mass loss via an accretion disc wind (driven out by radiation pressure, magneto-hydrodynamic or magneto-centrifugal forces) has also been proposed (see Hamman et al. \citeyear{ham13} and references therein). Type 2 quasars (QSO2) are unique objects for investigating the way feedback works in the most powerful radio quiet AGN. The active nucleus in QSO2 is occulted by obscuring material, which acts like a convenient ``natural coronograph'', allowing a detailed study of many properties of the surrounding medium. This is very complex in type 1 quasars (QSO1) due to the dominant contribution of the quasar point spread function. QSO2 have been discovered in large quantities only in recent years. In particular, \cite{zak03} and \cite{rey08} have identified $\sim$900 objects at redshift $z \lesssim$0.8 in the Sloan Digital Sky Survey (SDSS, York et al. \citeyear{york00}) with the high ionization narrow emission line spectra characteristic of type 2 AGN and narrow line luminosities typical of QSO1 (log(L\oiii/L$_{\odot}$)$>$8.3). Based on a spectroscopic study of 13 SDSS QSO2 at 0.3$\lesssim z \lesssim$0.6, Villar-Mart\'\i n et al. (2001b, hereafter \cite{vm11b}) found clear evidence for ionized outflows in the majority of objects and argue that this is a ubiquitous phenomenon in the nuclear region ($r<$several kpc) of this object class. \cite{liu13b} propose that the ionized outflows extend much farther, and can reach distances $\ga$15 kpc (see also Humphrey et al. 2010). The ionization, kinematic and morphological properties of the outflowing gas (\cite{vm11b}, Liu et al. \citeyear{liu13b}) suggest that the {\it ionized} outflows are preferentially induced by AGN related processes. \cite{lal10} studied the radio properties of 59 SDSS QSO2 at $z\ga$0.3. The detection rate of their survey is 59\% (35/59). They find that 15\%$\pm$5\% can be considered radio loud, according to the radio-to-[OIII]$\lambda$5007 luminosity ratios, while the vast majority of their detected sources fall in a region intermediate between those traditionally occupied by radio loud and radio quiet quasars. They detect a high fraction (75\%) of compact cores, which confine the radio emission to typical physical diameters of 5 kpc or less. Thus, both the radio jet and disc wind modes are possibly present. We present in this paper the results of our spectral analysis of the nearest SDSS QSO2 (21 objects at $z\la$0.1 from Reyes et al. \citeyear{rey08}). Based on the spectral decomposition of the most important optical emission lines using the SDSS spectra, we search for signatures of ionized outflows. By isolating the emission from the quiescent and turbulent (outflow) components, we compare their kinematic properties and line ratios and characterize how the outflows alter the properties of the ambient gas. We also investigate whether the properties of the turbulent gas depend on radio loudness. Our goal is to understand whether the radio structures play a role in triggering the outflows. By studying the outflows in the most nearby, luminous radio quiet quasars, we hope to shed light on feedback in the most distant quasars. We assume $\Omega_{\Lambda}$=0.7, $\Omega_{\rm M}$=0.3, H$_0$=71 km s$^{-1}$ Mpc$^{-1}$. At $z=$0.1, 1$\arcsec$ corresponds to 1.8 kpc. | We have investigated the existence and properties of nuclear ionized outflows in the 21 most nearby QSO2 at $z\la$0.1, by means of the spectral decomposition of the main optical emission lines in their SDSS spectra. Most objects are radio quiet according to the radio to [OIII] luminosity ratio. A few are radio-intermediate, but in all cases they have radio luminosities well below typical radio loud AGN. The turbulence parameter $R=\frac{\rm FWHM_{[OIII]}}{\rm FWHM_{\star}}$ provides the basis for an efficient discrimination method between the turbulent ($R>$1.4) and the quiescent ($R\le$1.4) gaseous components. Based on their kinematic properties we identify the turbulent components with emission from outflowing gas and the quiescent components with ambient gas which has not been reached by the outflows. We have detected ionized outflows in all but one object. Thus, the detection rate is 95\%. This result extends to the lowest $z$ the conclusion by \cite{vm11b} for QSO2 at $z\sim$0.3-0.6, that ionized outflows are ubiquitous in optically selected obscured quasars. The ionized outflows are located in the narrow line region. Both the ambient and the outflowing gas are ionized by AGN related processes. The outflowing gas is more highly reddened than the ambient gas (\ha/\hb= 4.91$\pm$0.03 and 3.68$\pm$0.02 respectively). It is also denser $n_e\sim$few$\times$hundred -10$^3$ cm$^{-3}$ for the ambient gas, while $n_e\ga$1000 cm$^{-3}$ is frequent in the outflowing gas. To explain these results, we propose that the bulk of the outflow line emission is originated in a more compact region ($r\sim$1-2 kpc) and closer to the AGN than the ambient gas. Such high densities imply low $M_{outf}\sim$(0.3-2.4)$\times$10$^6$ \msun~ and mass outflow rates $\dot M<$few \msun~ yr$^{-1}$. Larger values could exist if most of the outflow mass is contained in a reservoir of low density gas ($n_e<$few cm$^{-3}$). This might dominate the ionized outflow mass, but not the line luminosity which would be too faint to detect in spatially integrated spectra. The triggering mechanism of the outflows is related to the nuclear activity, rather than star formation. The QSO2 can be classified in two clearly differentiated groups according to the behavior and properties of the outflowing gas. QSO2 in Group 1 (5 objects) show the most extreme turbulence ($R_{max}\ge$4.7), they have on average higher radio luminosities and lower $q$ parameter, which parametrizes the excess of radio emission or radio-loudness. In this group, it is found that the lowest $q$ values (i.e. larger radio excess) are associated with the most extreme gas turbulence (larger $R_{max}$). QSO2 in Group 2 (15 objects) show less extreme turbulence ($R_{max}\sim$2-3), they have lower radio luminosities and, on average, higher $q$ values. Based on these results, we propose that two competing outflow mechanisms are at work: radio jets and another AGN related process (possibly accretion disk winds). Although both mechanisms might be present in at least some objects, the radio jet induced outflows are dominant in Group 1 (5/20 QSO2), while the disk wind dominates in Group 2 (15/20 QSO2). In this scenario, we find that the radio jet mode is capable of producing more extreme (turbulent) outflows. Independently of the outflow origin, our results suggest that the larger the turbulence, the smaller the mass fraction involved in the outflow, in coherence with \cite{vm11b}. To test this interpretation we predict that high resolution VLBA imaging will reveal the presence of jets in Group 1 QSO2. If the outflows expand to distances $>$several kpc from the AGN, the morphologies and kinematic properties could be characterized across scales of $\ga$several arcsec at $z\sim$0.1. We predict that the morphology of the extended ionized nebulae is more highly collimated and kinematically perturbed for objects in Group 1. Our basic analysis reveals no evidence for neutral outflows in QSO2. The ISM might be too highly ionized across large spatial scales (indeed no significant component of neutral ISM absorption is detected). Alternatively, the detection method, consistent of tracing NaID absorption against regions with a bright stellar continuum, may be inadequate for tracing AGN induced outflows. | 14 | 3 | 1403.1175 |
1403 | 1403.6469_arXiv.txt | We compare the kinematics of M31's satellite galaxies to the mass profiles of the subhaloes they are expected to inhabit in \LCDM{}. We consider the most massive subhaloes of an approximately M31-sized halo, following the assumption of a monotonic galaxy luminosity-to-subhalo mass mapping. While this abundance matching relation is consistent with the kinematic data for galaxies down to the luminosity of the bright satellites of the Milky Way and M31, it is \emph{not} consistent with kinematic data for fainter dwarf galaxies (those with $L \lesssim 10^8 \Lsun$). Comparing the kinematics of M31's dwarf Spheroidal (dSph) satellites to the subhaloes reveals that M31's dSph satellites are too low density to be consistent with the subhaloes' mass profiles. A similar discrepancy has been reported between Milky Way dSphs and their predicted subhaloes, the ``too big to fail'' problem (TBTF). By contrast, total mass profiles of the dwarf Elliptical (and similarly bright) satellites are consistent with the subhaloes. However, they suffer from large systematic uncertainties in their dark matter content because of substantial (and potentially dominant) contributions from baryons within their half-light radii. | \label{sec:intro} Dwarf satellite galaxies of the Local Group (LG) are unique cosmological probes. They are close enough that very faint objects can be detected, and their individual stars can be spectroscopically observed. This enables measurements that are impossible with integrated-light observations \citep[e.g.,][]{sandg, stri08commonmass, walker09fnx}. A problem revealed by these observations was recently pointed out by \citet{bkbk11,bkbk12}, dubbed the ``too big to fail'' problem. They demonstrated that the bright satellites of the Milky Way (MW) have internal kinematics inconsistent with the expectations of a \LCDM{} simulation of a MW-sized halo. Specifically, they showed that the most massive subhaloes in \LCDM{} simulations have central masses systematically larger than those measured in the ten brightest dSph satellites of the Milky Way. The population of satellites compatible with observed kinematics of the dSphs contains intermediate-mass subhaloes, leaving the question of why the most massive subhaloes appear to be dark. This implies either that the massive subhaloes of the MW are inexplicably without luminous galaxies, or the central densities of their dark matter haloes are different from the expectations of a dark matter-only simulation in \LCDM{}. A number of explanations for this problem have been suggested, from forming cores in the galaxies' mass profiles to the existence of dark matter particles with unusual properties \citep{dicintio11, lovell12, zolotov12, vogelsberger12, vinas12, maccio13, rocha13, peter13, libeskind13, brooks13, dicintio13}. The simplest possibile explanation, however, is that the MW is an outlier relative to other similar galaxies. That is, if the MW satellite population were different from that of a typical galaxy with the same halo mass, it would cast doubt on the common practice of comparing the MW to typical halos in simulations. The other bright spiral of the LG, M31, provides a second testing ground for this effect. It has a large population of known satellites, thanks in large parts to the efforts of the Pan-Andromeda Archaeological Survey \citep[PAndAS,][]{pandas09nat}. While a few of its brightest satellites are classified as dwarf Ellipticals (dEs), most are dSphs, implying that the satellite system of M31 may be comparable to that of the MW. This (among other goals) has motivated spectroscopic surveys of the M31 satellites, including the dwarf component of the Spectroscopic and Photometric Landscape of the Andromeda Stellar Halo (SPLASH) survey \citep{kalirai10, howley12m32, paper1}, as well as a parallel survey by \citealt{collins13} (and a corresponding kinematics analysis in \citealt{collins13_2}, discussed more in \S \ref{sec:comp}). These and other kinematic investigations of M31 satellites provide a wealth of data for examining the internal dynamics of M31 satellites. In this paper, we make use of these new data sets to compare M31's satellites to subhaloes from a \LCDM{} N-body simulation intended to approximate the expected halo of M31. In \S \ref{sec:sats}, we provide an overview of the M31 satellites and the observational data we use here. In \S \ref{sec:subs}, we describe the comparison simulation datasets and how we map them on to the observations. In \S \ref{sec:comp}, we compare the M31 satellites to the simulated subhaloes. Finally, in \S \ref{sec:conc}, we provide concluding thoughts. Where relevant, we assume a distance to M31 of 783 kpc \citep[e.g.,][]{paturel02, Mcconnachie05, perina09}. | \label{sec:conc} In this paper, we have compared the internal dynamics of M31's satellites (with $L_V > 10^5 \Lsun$) to the subhaloes they are expected to inhabit in a \LCDM{} universe, given a simple set of assumptions for mapping galaxies to haloes. This yields the following main results: \begin{enumerate} \item The bright satellites of both M31 and the MW are consistent with a monotonic luminosity-to-$\vmax$ relation, the primary assumption necessary for abundance matching. However, they are \emph{not} consistent with extrapolating the abundance matching of $\gtrsim 0.1 L_* $ galaxies to lower luminosities. Specifically, low-luminosity dwarfs are significantly less dense than an abundance matching relation matching the LG luminosity function would predict. These results are robust to assumptions about the masses of the M31 and MW halo. \item The dEs and similarly bright satellites of M31 (which have no analogue in the MW) are nominally consistent with the subhaloes they would be expected to inhabit. However, the interpretation of this in the context of their dark matter haloes is hampered by the significant contribution to their mass budget by their baryonic mass and the associated possible systematic uncertainties. \item The dSph satellites of M31 have lower densities (within their half-light radii) than the densities of the most massive subhaloes that are expected to host them in \LCDM{} collisionless dark matter simulations. Their dynamics are thus consistent with the satellites of the MW, and exhibit the ``too big to fail'' problem. Thus, the simplest explanation for the problem -- that the MW is a statistical fluke -- does not seem to be valid, as a statistical anomaly is much less likely to be found in two different haloes. \end{enumerate} These results provide crucial context for interpreting the small-scale puzzles presented by \LCDM{} or alternative cosmological models by moving beyond the MW and its satellites. They demonstrate that these puzzles persist in a very similar form for M31 and its satellite galaxies. However, understanding if they also hold for a larger, statistical sample of galaxies is necessary for interpreting such results in a cosmological context. Unfortunately, resolved star spectroscopic data like those used here are nearly impossible to obtain beyond the Local Group with current spectroscopic capabilities. Fortunately, in the coming era of deep, wide surveys like the Large Synoptic Survey Telescope and extremely large telescopes for spectroscopic follow-up (e.g., the Thirty Meter Telescope and Giant Magellan Telescope), the potential exists to push the boundaries of near-field cosmology well beyond the Local Group. \vskip 1.5\baselineskip \noindent {\bf{Acknowledgements}} The authors thank Marla Geha, Ana Bonaca, Jorge Pe\~narrubia, Alis Deason, Andrey Kravtsov, Carlos Frenk, and the anonymous referee for valuable discussions. We also thank the Aquarius collaboration for providing access to their simulation data. This research made use of Astropy, a community-developed core Python package for Astronomy \citep{astropy}. Support for this work was provided by NASA through Hubble Fellowship grant \#51316.01 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. This work was also supported by NSF Grant AST-1009973. | 14 | 3 | 1403.6469 |
1403 | 1403.4693_arXiv.txt | We present a millimetre-wave site characterisation for the Australia Telescope Compact Array (ATCA) based on nearly 9 years of data from a seeing monitor operating at this facility. The seeing monitor, which measures the phase fluctuations in the signal from a geosynchronous satellite over a 230\,m baseline caused by water vapour fluctuations along their sight lines, provides an almost gapless record since 2005, with high time resolution. We determine the root mean square (rms) of the path length variations as a function of time of day and season. Under the assumption of the ``frozen screen'' hypothesis, we also determine the Kolmogorov exponent, $\alpha$, for the turbulence and the phase screen speed. From these, we determine the millimetre-wave seeing at $\lambda = 3.3$\,mm. Based on the magnitude of the rms path length variations, we estimate the expected fraction of the available observing time when interferometry could be successfully conducted using the ATCA, as a function of observing frequency and antenna baseline, for the time of day and the season. We also estimate the corresponding observing time fractions when using the water vapour radiometers (WVRs) installed on the ATCA in order to correct for the phase fluctuations occurring during the measurement of an astronomical source. | The Australia Telescope Compact Array (ATCA) is a radio interferometer located near the town of Narrabri in northwestern New South Wales, Australia at S30$^{\circ}$ 18'46" E149$^{\circ}$ 33'00". For over 10 years it was the only millimetre interferometer located in the southern hemisphere \citep{ATCAUserGuide:2011fk}. It was originally planned as a centimetre wavelength synthesis telescope with upgrades to the 12, 7 and 3 millimetre observing bands to be implemented at a later stage. This, along with some accounts of how the site selection came to agree on Narrabri, is described in detail in \cite{Frater:1992uq}. Following a proposal to fit out the array with millimetre wave receivers, a site testing programme was initiated to establish its feasibility. Longer term observations however were limited to opacity measurements of the atmosphere at 30 GHz while single snapshot observations were proposed at a wavelength of 3\,cm with the array antennae. This is described in an ATNF internal memo \citep{Hall:1992fk}. From these measurements the fraction of nights when the transmission is suitable for obervation at a given wavelength can be estimated, but they do not allow the phase stability for operating an interferometer as a function of baseline to be determined. In this paper we present a comprehensive analysis of the millimetre site characteristics at the ATCA based on 8.5 years of data and extending the analysis of \citet{Middelberg:2006vn}, which used the initial 1 year of data gathered with a seeing monitor that was installed in 2004. | The magnitude of the path length fluctuations caused by variations in the water vapour columns between two antennae of an interferometer determines when observations may be successfully attempted at a given frequency. We have analysed 8.5 years of data from a seeing monitor at the ATCA site near Narrabri to provide statistics on the site conditions for this telescope, extending the initial 1 year study conducted by \citet{Middelberg:2006vn}. Two monitors located 230\,m apart measure the signal from a geosyncronous satellite. We use the fluctuations in the path differences across this baseline to determine the rms value for the path difference, the Kolmogorov exponent, $\alpha$, for the turbulence and the phase screen speed, under the ``frozen-screen'' assumption \citep{Taylor:1938fk} for the passage of the turbulent cells across the site. We also determine the seeing at 3.3\,mm from the half power points in the visibility. Considerable variations in the rms path differences between the two monitors are found between summer and winter, as well as between day and night time, with 25\% (first) quartile values ranging from a low of 111$\mu$m in May between 03--06 hrs, to a high of 653$\mu$m in February between 12--15 hrs. However, typical fluctuations during summer nights, $\sim 300 \mu$m, are similar to those occuring during winter days; in other words mm-wave interferometry can typically be undertaken during summer nights. Variations in the Kolmogorov exponent are much less, with a median value of $\sim 0.4$ and variations of $\sim \pm 0.04$. Similarly, variations in the phase screen speed are relatively small, with mode values found of $\sim 2$\,m/s. The Kolmorogorov exponent is usually closer to the value expected for 2D turbulence, $1/3$, than that for 3D turbulence, $5/6$, consistent with the turbulent layer containing the fluctuating water vapour being relatively thin, as per the frozen-screen hypothesis. The lowest values for $\alpha$ are found in winter, as would be expected for a less \textbf{agitated} troposphere than in summer. The seeing at 3.3\,mm shows significant seasonal and time of day variations, depending on the path length rms values. It ranges from a low of $\sim 0.3''$ during winter nights to $\sim 3''$ in the middle of summer days. We have calculated the maximum path length fluctuations that can be tolerated for interferometric observations to be conducted with the ATCA on a given antenna baseline and frequency (in particular, the three receivers at 22, 45 and 90\,GHz). Using the measured path length fluctuations, and assuming Kolmogorov turbulence with the median value measured for $\alpha$ of 0.4, we then estimate the useable observing fractions for time of day and month of the year. We also do this assuming that phase variations can be partially rectified using the ATCA water vapour radiometers (WVRs). Gains are possible on all baselines, with useable observing periods typically increased by about 4 months per year for any given combination of frequency/baseline using the WVRs. In particular, observations at 90\,GHz could be conducted out to the currently maximum 3\,km baseline of the ATCA during winter months, or out to 6\,km were the six kilometre antenna to be equipped with 3mm capability. With the use of WVRs, observations will often be undertaken with baselines longer than now being used. This would then facilitate flexible observing as when the conditions are not suitable for mm observations, the longer baselines will be conducive for many centimetre band programmes, as was proposed by \citet{Hall:1992fk} in their original proposal for operation of the ATCA at millimetre wavelengths. | 14 | 3 | 1403.4693 |
1403 | 1403.1280_arXiv.txt | If a $\gamma$-ray line is observed in the near future, it will be important to determine what kind of dark matter (DM) particle could be at its origin. We investigate the possibility that the $\gamma$-ray line would be induced by a slow DM particle decay associated to the fact that the DM particle would not be absolutely neutral. A ``millicharge'' for the DM particle can be induced in various ways, in particular from a kinetic mixing interaction or through the Stueckelberg mechanism. We show that such a scenario could lead in specific cases to an observable $\gamma$-ray line. This possibility can be considered in a systematic model-independent way, by writing down the corresponding effective theory. This allows for a multi-channel analysis, giving in particular upper bounds on the intensity of the associated $\gamma$-ray line from cosmic rays emission. Our analysis includes the possibility that in the two-body decay the photon is accompanied with a neutrino. We show that, given the stringent constraints which hold on the millicharge of the neutrinos, this is not an option, except if the DM particle mass lies in the very light KeV-MeV range, allowing for a possibility of explanation of the recently claimed, yet to be confirmed, $\sim 3.5$~KeV X-ray line. | One of the most promising ``smoking-gun'' signals for establishing the existence of the dark matter particle is the possible observation of a sharp cosmic $\gamma$-ray line from dark matter annihilation or decay \cite{Bergstrom:1988fp}. The forthcoming Cherenkov telescopes \cite{CTA::2013}, the current Fermi large area telescope \cite{Fermi::2013} and the HESS instrument \cite{HESSII::2012} will allow to probe this possibility with further sensitivity. If such a signal is observed in the near future, the question of the identification of the DM particle that could have caused it will become crucial. Such a signal could be induced through annihilation, coannihilation or decay. For all these scenarios, it is generally assumed that the photon is emitted through the loop of a charged particle. Beside this general class of models, there exist other ways along which DM could emit monochromatic photons. One possibility consists in assuming that the $\gamma$-ray line is due to a $Z-Z'-\gamma$ Chern-Simons interaction \cite{Dudas:2012pb}. Another possibility, much less studied, would be to consider a photon directly emitted by the DM particle. This is a priori perfectly possible if DM is not exactly neutral, but is millicharged. For an annihilation such a possibility is not much of an option because the associated $\gamma$-ray line would be in general suppressed with respect to the total cross section, by the square of the millicharge. Given the constraints there are on the total cross section (in particular from the relic density in the thermal freezeout scenario), this would lead to a signal sizeably smaller than present or near future sensitivities. Instead for a decay, there is a priori more freedom because the decay lifetime is not so directly constrained by the relic density. In this work we consider such a decay possibility. In the following we will first consider the two main frameworks that can in a simple way justify a millicharge for the DM particle, kinetic mixing and Stueckelberg scenarios. In such scenarios, in order to justify that the DM particle would have a slow decay, we assume that its stability is due to an accidental symmetry that, being accidental, would be naturally broken by any UV physics. Along these lines, the decay is naturally slow because suppressed by powers of the UV scale, just as expected for the proton. The appropriate language to consider in a model-independent way the possibility of a slow decay is therefore the one of the higher-dimensional operator effective theory. Unlike for an annihilation, the use of an effective theory for a decay is fully justified since one expects a clear scale separation. Consequently, such an effective theory allows for a systematic study of possibilities. We will determine all dimension-five and dimension-six operators that can lead to a two-body radiative decay from a millicharged fermion, scalar or vector DM particle. These operators come in addition to the effective operators which can lead to a $\gamma$-ray line in the case where DM would be exactly neutral, given and studied in Ref.~\cite{Gustafsson::2013}. The former operators involve a covariant derivative of the millicharged field, whereas the latter ones can involve a photon only from the presence of a hypercharge or $SU(2)_L$ field strength $F_{Y,L}^{\mu\nu}$ in the operator. In the following, we will perform a detailed analysis of the constraints that hold on the various ``millicharged operators'' for the fermionic DM case. The scalar and vector cases will be discussed more briefly before concluding. A simple constraint that turns out to be relevant in some cases is that the DM particle lifetime should be larger than the age of the Universe. Another one concerns the emission of cosmic rays (CR) that could be associated to the one of the photon, either from the particle that accompanies the photon in the decay final state, or from other decays that the effective operator unavoidably predicts on top of the radiative one. Gauge invariance in particular predicts decays where the photon is replaced by a $Z$. If the electromagnetic coupling to the $Z$ is not millicharge suppressed, the flux of cosmic rays produced is much larger than the flux of monochromatic photons. In particular, if the particle accompanying the photon in the final state is a neutrino, which is the only Standard Model (SM) particle possibility (a decay of special interest being ``poly-monochromatic'', i.e.~monochromatic for both types of cosmic rays that are the less affected while propagating), we will see that an observable $\gamma$-ray line is not an option, unless the DM mass is quite low. Therefore, except for this case, the possibilities we will find point towards multi-component DM scenarios. Other constraints are related to the fact that along the Stueckelberg scenario the DM particle is charged under a new $U(1)'$ gauge group, which may be at the origin of the unsuppressed emission of the associated $Z'$. | In summary, there are very few ways of probing the DM hypothesis that can really be considered in a systematic and model-independent way. However, for the decay of an absolutely neutral DM particle into a $\gamma$-ray line, this turns out to be feasible \cite{Gustafsson::2013}. This stems from the facts that, on the one hand, the use of an effective theory is fully justified, slow enough decay can naturally be explained from a much higher scale physics, and, on the other hand, it turns out that there are very few operator structures of this kind. Ref. \cite{Gustafsson::2013} considered the usual scenario where the DM particle is absolutely neutral so that the photon appears in the operator through a field strength (i.e. typically from a charged particle in a loop). Here we show that, for the same reasons, such a study can also be systematically carried out in the less considered scenario where the DM is millicharged, having therefore a tree-level coupling to the photon through a covariant derivative, either from an ad-hoc millicharge, or through mixing of the $U(1)_Y$ gauge boson with another $U(1)'$ gauge boson. To the emission of a $\gamma$-ray line from such operators is associated the emission of a continuum of cosmic rays. The monochromatic photon to cosmic ray flux ratio is determined by the SM quantum numbers of the field on which the covariant derivative applies (and in one case also crucially on the DM mass), and if this particle is not a SM singlet on the value of its millicharge. This leads to upper bounds on the intensity of the $\gamma$-ray line produced, given in Fig. \ref{fig:Kin_mix}. This figure shows that if the DM is only charged under the dark sector, it can lead to a line matching the present experimental sensitivities without overshooting the bound on the flux of antiprotons and diffuse photons. On the contrary, when the particle emitting the photon from its millicharge is also charged under the SM, the cosmic rays constraints are much stronger than direct searches for spectral lines. Therefore, in this case, if a line were to be detected with energy above the $Z$ mass and with about the present experimental sensitivity, it could not be explained in such a way. Such a conclusion can also hold for $m_{DM}$ far below the $Z$ mass. For the massless hidden gauge boson case (and the massive case where $m_{Z'}$ is both below the GeV scale and smaller than $m_{DM}$) relevant additional constraints show up imposing that the two-body decay width to a $\gamma'$ (Z') leads to a lifetime longer than the age of the Universe. Combining this constraint with the direct detection bounds on a millicharge, an observable $\gamma$-ray line requires small values of the dark charge $g'Q'$. As for a decay into a neutrino and a photon, given the stringent constraints that exist on the millicharge of a neutrino, the $Z$ mediated decay into three neutrinos, or into a neutrino and a electron-positron pair, this possibility is forbidden unless $m_{DM}$ is below the MeV scale. For lower masses, and down to the KeV scale, an observable line induced in this way is not excluded by these considerations. Such a neutrino millicharge scenario could even be at the origin of the recently reported, yet to be confirmed, 3.5 KeV X-ray line. \vspace{5mm} | 14 | 3 | 1403.1280 |
1403 | 1403.1763_arXiv.txt | { Multi-molecule observations towards an increasing variety of galaxies have been showing that the relative molecular abundances are affected by the type of activity. However, these studies are biased towards bright active galaxies, which are typically in interaction. } { We study the molecular composition of one of the most isolated galaxies in the local Universe where the physical and chemical properties of their molecular clouds have been determined by intrinsic mechanisms. } { We present 3~mm broad band observations of the galaxy CIG~638, extracted from the AMIGA sample of isolated galaxies. The emission of the $J=1-0$ transitions of CCH, HCN, HCO$^+$, and HNC are detected. Integrated intensity ratios between these line are compared with similar observations from the literature towards active galaxies including starburst galaxies (SB), active galactic nuclei (AGN), luminous infrared galaxies (LIRG), and GMCs in M33. } { A significantly high ratio of CCH with respect to HCN, HCO$^+$, and HNC is found towards CIG~638 when compared with all other galaxies where these species have been detected. This points to either an overabundance of CCH or to a relative lack of dense molecular gas as supported by the low HCN/CO ratio, or both. } { The data suggest that the CIG~638 is naturally a less perturbed galaxy where a lower fraction of dense molecular gas, as well as a more even distribution could explain the measured ratios. In this scenario the dense gas tracers would be naturally dimmer, while the UV enhanced CCH, would be overproduced in a less shielded medium. } | Deep systematic multi-molecular studies in galaxies have been carried out for almost a decade now and have been significantly boosted by the recent increase in the instantaneous bandwidth of mm and submm facilities \citep[see][for a review]{Mart'in2011a}. Both targeted and unbiased surveys have been carried out towards the nuclei of the brightest and nearest prototypical active galaxies \citep{Wang2004,Mart'in2006,Mart'in2011,Aladro2011a,Aladro2013} with the aim of studying the effects of the nuclear activity on the overall molecular abundances. Studies using small samples of galaxies have resulted in some potential activity diagnostics based on molecular abundance ratios \citep{Kohno2001,Mart'in2009,Costagliola2011}. In particular, the ratio between HCN and HCO$^+$ has been proposed as a potential discriminator of the activity type whether AGN or SB dominated \citep{Kohno2001,Krips2008}. In this scenario, the HCN abundance would be enhanced by the pervading X-ray radiation from the central black hole, while the HCO$^+$ would be favoured in UV irradiated and/or cosmic ray enhanced star-forming regions. However, all these studies have been naturally limited to the brightest, more active nearby galaxies. The AMIGA isolated galaxy sample \citep{Verdes-Montenegro2005,Verdes-Montenegro2010} has clearly established that parameters expected to be enhanced by interactions, such as $L_{FIR}$, radio continuum emission, AGN rate, or HI asymmetry, are lower in isolated galaxies than in any other sample, even compared with field galaxies. In particular, the comparison between the AMIGA sample and a sample of Hickson compact groups found indications that molecular gas in isolated galaxies could be more extended than in environments with a higher density of galaxies \citep{Lisenfeld2011,Martinez-Badenes2012}. In this paper, we aim to explore the chemistry of the most isolated galaxies in the local Universe in order to add a new type of galaxy to these comparative molecular studies. It is expected that interactions will have a direct effect on the physical conditions of giant molecular clouds (GMCs), the star formation history, and activity within the nuclear regions of galaxies, which will very likely lead to a different chemical evolution of the molecular material in the galaxy. Interestingly, only eight LIRGs and no ULIRG are found in the AMIGA sample. A comparison based on molecular abundances can provide hints to how the secular evolution affects the ISM associated to a nuclear activity not triggered and/or affected by interactions. Therefore, galaxies in isolation might be key as chemical baselines to understand the ways in which the ISM has evolved in isolation as opposed to other more active interacting galaxies. \begin{figure} \centering \includegraphics[width=\linewidth]{figIsolation.eps} \caption{Location of the number density and tidal parameters of CIG~638, as obtained from the isolation parameters in \citet{Argudo-Fernandez2013}, compared to the whole AMIGA sample. \label{fig.isolation}} \end{figure} The star-forming galaxy CIG~638 (NGC~5690) matches the isolation criteria from \citet{Verley2007b} based on revised local number density and tidal parameters (Fig.~\ref{fig.isolation}) from \citet{Argudo-Fernandez2013}. A revision of isolation using spectroscopic data shows CIG~638 is isolated from similar luminosity neighbours, unlike the comparison galaxies used in this paper, for which data exist for evaluating isolation. Also, in the $L_{\rm FIR}-L_{\rm B}$ and $M_{\rm H_2}$ vs $L_{\rm B}$ plots, CIG~638 lies in the same parameter space as the bulk of the AMIGA sample. Its morphological type (Sbc) is the most abundant in AMIGA, and the symmetric HI profile also reinforces its isolation \citep{Espada2011}. While the degree of isolation of nearby galaxies ($v<1500$~\kms) cannot be reliably determined, the more distant galaxies, being quiescent, tend to be fainter at all wavelengths. In fact, the vast majority of isolated galaxies in the sample are not bright enough in CO \citep{Lisenfeld2011} to easily detect other molecular species than CO, typically a factor of 20 to 100 fainter. The bright ($\sim50$~mK) and broad ($FWZI\sim280$~\kms) CO emission of the almost edge-on CIG~638 ($i=78^\circ$) makes it the ideal CO-bright isolated galaxy for deep molecular observations. \begin{figure*} \centering \includegraphics[width=0.9\linewidth]{figSPECTRA.eps} \caption{ Detected $J=1-0$ transitions towards CIG~638. Spectral resolution is smoothed down to 26~\kms. \label{fig.spectra}} \end{figure*} | Based on single-dish line ratio observations of the isolated galaxy CIG~638, molecular abundance differences are found when compared to local interacting active galaxies. The lower tidal force in CIG~638 than in our comparison sample has an influece not only on the nuclear activity \citep{Sabater2013} but also on the overall ISM physical conditions. Assuming that CIG~638 is representative of isolated galaxies, we suggest that in isolation a lower fraction of dense gas and a more homogeneously distributed, low-extinction ISM leads to both a low abundance of dense gas tracers, such as HCN and HCO$^+$, and an overabundance of CCH due to poorer shielding from UV radiation. However, spatially resolved sensitive imaging on a larger sample of isolated galaxies is required to confirm this scenario. | 14 | 3 | 1403.1763 |
1403 | 1403.7091_arXiv.txt | We present an analysis of the diversity of $V$-band light-curves of hydrogen-rich type II supernovae. Analyzing a sample of 116 supernovae, several magnitude measurements are defined, together with decline rates at different epochs, and time durations of different phases. It is found that magnitudes measured at maximum light correlate more strongly with decline rates than those measured at other epochs: brighter supernovae at maximum generally have faster declining light-curves at all epochs. We find a relation between the decline rate during the `plateau' phase and peak magnitudes, which has a dispersion of 0.56 magnitudes, offering the prospect of using type II supernovae as purely photometric distance indicators. Our analysis suggests that the type II population spans a continuum from low-luminosity events which have flat light-curves during the `plateau' stage, through to the brightest events which decline much faster. A large range in optically thick phase durations is observed, implying a range in progenitor envelope masses at the epoch of explosion. During the radioactive tails, we find many supernovae with faster declining light-curves than expected from full trapping of radioactive emission, implying low mass ejecta. It is suggested that the main driver of light-curve diversity is the extent of hydrogen envelopes retained before explosion. Finally, a new classification scheme is introduced where hydrogen-rich events are typed as simply `SN~II' with an `$s_2$' value giving the decline rate during the `plateau' phase, indicating its morphological type. | Supernovae (SNe) were initially classified into types I and II by \cite{min41}, dependent on the absence or presence of hydrogen in their spectra. It is now commonly assumed that hydrogen-rich type II SNe (SNe~II henceforth) arise from the core-collapse of massive ($>$8-10\msun) stars that explode with a significant fraction of their hydrogen envelopes retained. A large diversity in the photometric and spectroscopic properties of SNe~II is observed, which leads to many questions regarding the physical characteristics of their progenitor scenarios and explosion properties.\\ \indent The most abundant of the SNe~II class (see e.g.\ \citealt{li11} for rate estimates) are the SNe~IIP which show a long duration plateau in their photometric evolution, understood to be the consequence of the hydrogen recombination wave propagating back through the massive SN ejecta. SNe~IIL are so called due to their `linear' declining light-curves (see \citealt{bar79} for the initial separation of hydrogen-rich events into these two sub-classes). A further two sub-classes exist in the form of SNe~IIn and SNe~IIb. SNe~IIn show narrow emission lines within their spectra \citep{sch90}, but present a large diversity of photometric and spectral properties (see e.g.\ \citealt{kie12,tad13}), which clouds interpretations of their progenitor systems and how they link to the `normal' SNe~II population. (We note that a progenitor detection of the SN~IIn: 2005gl does exist, and points towards a very massive progenitor: \citealt{gal09}, at least in that particular case). SNe~IIb appear to be observationally transitional events as at early times they show hydrogen features, while later such lines disappear and their spectra appear similar to SNe~Ib \citep{fil93}. These events appear to show more similarities with the hydrogen deficient SN~Ibc objects (see \citealt{arc12}, Stritzinger et al. in prep.). As these last two sub-types are distinct from the classical hydrogen-rich SNe~II, they are no longer discussed in the current paper. An even rarer sub-class of type II events, are those classed as similar to SN~1987A. While SN~1987A is generally referred to as a type IIP, its light-curve has a peculiar shape (see e.g. \citealt{ham88}), making it distinct from classical type IIP or IIL. A number of `87A-like' events were identified in the current sample and removed, with those from the CSP being published in \cite{tad12} (see also \citealt{kle11}, and \citealt{pas12}, for detailed investigations of other 87A-like events).\\ \indent The progenitors of SNe~II are generally assumed to be stars of ZAMS mass in excess of 8-10\msun, which have retained a significant fraction of their hydrogen envelopes before explosion. Indeed, initial light-curve modeling of SNe~IIP implied that red-supergiant progenitors with massive hydrogen envelopes were required to reproduce typical light-curve morphologies (\citealt{gra71,che76,fal77}). These assumptions and predictions have been shown to be consistent with detections of progenitor stars on pre-explosion images, where progenitor detections of SNe~IIP have been constrained to be red supergiants in the 8-20\msun\ ZAMS range (see \citealt{sma09} for a review, and \citealt{van12b} for a recent example). It has also been suggested that SN~IIL progenitors may be more massive than their type IIP counterparts (see \citealt{eli10,eli11}).\\ \indent Observationally, hydrogen-rich SNe~II are characterized by showing P-Cygni hydrogen features in their spectra\footnote{While as shown in \cite{sch96} there are a number of SNe~II which show very weak \ha\ absorption (which tend to be of the type IIL class), the vast majority of events do evolve to have significant absorption features. Indeed this will be shown to be the case for the current sample in Guti\'errez et al. (submitted).}, while displaying a range of light-curve morphologies and spectral profiles. Differences that exist between the photometric evolution within these SNe are most likely related to the mass extent and density profile of the hydrogen envelope of the progenitor star at the time of explosion. In theory, SNe with less prominent and shorter `plateaus' (historically classified as SNe~IIL) are believed to have smaller hydrogen envelope masses at the epoch of explosion (\citealt{pop93}, also see \citealt{lit83} for generalized model predictions of relations between different SNe~II properties). Further questions such as how the nickel mass and extent of its mixing affects e.g.\ the plateau luminosity and length have also been posed (e.g. \citealt{kas09,ber11}).\\ \indent While some further classes of SN~II events with similar properties have been identified (e.g. sub-luminous SNe~IIP, \citealt{pas04,spi14}; luminous SNe~II, \citealt{ins13}; `intermediate' events, \citealt{gan13,tak14}), analyses of statistical samples of SN~II light-curves are, to date uncommon in the literature, with researchers often publishing in-depth studies of individual SNe. While this affords detailed knowledge of the transient evolution of certain events, and thus their explosion and progenitor properties, often it is difficult to put each event into the overall context of the SNe~II class, and how events showing peculiarities relate.\\ \indent Some exceptions to the above statement do however exist: \cite{psk67} compiled photographic plate SN photometry for all supernova types, finding in the case of SNe~II (using a sample of 18 events), that the rate of decline appeared to correlate with peak brightness, together with the time required to observe a `hump' in the light-curve (see also \citealt{psk78}). All available SN~II photometry at the time of publication (amounting to 23 SNe) was presented by \cite{bar79}, who were the first to separate events into SNe~IIP and SNe~IIL, on the basis of $B$-band light-curve morphology. \cite{you89} discussed possible differences in the $B$-band absolute magnitudes of different SNe~II, analyzing a sample of 15 events. A large `Atlas' of historical photometric data of 51 SNe~II was first presented and then analyzed by \cite{pat93} and \cite{pat94} respectively. These data (with significant photometry available in the $B$ and $V$ bands), revealed a number of photometric and spectroscopic correlations: more steeply declining SNe~II appeared to be more luminous events, and also of bluer colors than their plateau companions. Most recently, \cite{arc12} published an analysis of $R$-band light-curves (21 events, including 3 SNe~IIb), concluding that SNe~IIP and SNe~IIL are distinct events which do not show a continuum of properties, hence possibly pointing towards distinct progenitor populations. We also note that bolometric light-curves of a significant fraction of the current sample were presented and analyzed by \cite{ber13}, where similar light-curve characterization to that outlined below was presented.\\ \indent The aim of the current paper is to present a statistical analysis of SN~II $V$-band light-curve properties that will significantly add weight to the analysis thus far presented in the literature, while at the same time introduce new nomenclature to help the community define SN~II photometric properties in a standardized way. Through this we hope to increase the underlying physical understanding of SNe~II. To proceed with this aim, we present analysis of the $V$-band light-curves of 116 SNe~II, obtained over the last three decades. We define a number of absolute magnitudes, light-curve decline rates, and time epochs, and use these to search for correlations in order to characterize the diversity of events.\\ \indent The paper is organized as follows: in the following Section we outline the data sample, and briefly summarize the reduction and photometric procedures employed. In \S\ 3 we define the photometric properties for measurement, outline our explosion epoch, extinction, and error estimation methods, and present light-curve fits to SN~II photometry. In \S\ 4 results on various correlations between photometric properties, together with their distributions are presented. In \S\ 5 we discuss the most interesting of these correlations in detail, and try to link these to physical understanding of the SN~II phenomenon. Finally, several concluding remarks are listed.\\ \indent In addition, an appendix is included where detailed light-curves (together with their derived parameters) and further analysis and figures not included in the main body of the manuscript are presented. The keen reader is encouraged to delve into those pages for a full understanding of our analysis and results. | An analysis of $V$-band photometry of 116 SNe~II has been presented with the aim of characterizing the diversity seen within their light-curves. This has been achieved through defining three magnitude measurements at different epochs: $M_\text{max}$, $M_\text{end}$, $M_\text{tail}$, three photometric decline rates: $s_1$, $s_2$ and $s_3$, together with the time durations $Pd$ and $OPTd$. We analyzed these distributions, and searched for possible correlations. Our main findings are that the SN~II family forms a continuum of events in terms of their light-curve morphologies (in the $V$ band), and that while large dispersion is observed, brighter SNe at maximum generally decline more quickly at all epochs. We speculate that the majority of the diversity of SNe~II can be explained through differences in their hydrogen envelope masses at the epoch of explosion, a parameter which is most directly measured through observations of the optically thick phase duration ($OPTd$). Finally, we list our main conclusions originating from this work.\\ \begin{itemize} \item A continuum of SN~II $V$-band properties is observed in all measured parameters (absolute magnitudes, decline rates, optically thick phase durations), and we observe no clear bimodality or separation between the historically defined SNe~IIP and SNe~IIL. \item SNe which are brighter at maximum decline more quickly at all epochs. \item After making a series of data quality cuts, it is found that the dispersion in the relation between $s_2$ and $M_\text{max}$ can be reduced to 0.56 mag, which opens the way to using SNe~II as photometric distance indicators, independent of spectroscopic information. \item While $M_\text{max}$ is more difficult to define and measure than other magnitudes, it shows the highest degree of correlation with decline rates. Hence, it appears that $M_\text{max}$ is the dominant magnitude parameter describing the diversity of SNe~II. \item We find a large range in $V$-band optically thick, and `plateau' durations ($OPTd$, $Pd$) which implies a large range in hydrogen envelope masses at the epoch of explosion. The fact that these parameters show correlation with a number of other light-curve parameters suggests that one of the most dominant physical parameters that explains the diversity of SNe~II light-curves is the envelope mass at the epoch of explosion. \item There are a significant number of SNe~II which decline more quickly during the radioactive tail, $s_3$, than the rate expected through full trapping of gamma-ray emission. This implies a large range in ejecta masses, with many SNe~II having low mass/density ejecta, through which emission can escape. \item Given the qualitative nature of current discussion in the literature of different SNe~II, we suggest the introduction of a new parameter, $s_2$: the decline rate per 100 days of the $V$-band light-curve during the `plateau' phase. This will enable future studies to make quantitative comparisons between SNe and SNe samples in a standardized way. \item The historically defined SN~IIL class does not appear to be significantly represented within this sample, and therefore it is concluded that truly `linearly' declining hydrogen-rich SNe~II are intrinsically extremely rare events. \item SN~II $V$-band magnitudes show a dispersion at the end of the `plateau', $M_\text{end}$ of 0.81 mag, 0.2 mag lower than that at peak, $M_\text{max}$. \end{itemize} | 14 | 3 | 1403.7091 |
1403 | 1403.2455.txt | Type IIb Supernova (SN) 2011dh, with conclusive detection of an unprecedented Yellow Supergiant (YSG) progenitor, provides an excellent opportunity to deepen our understanding on the massive star evolution in the final centuries toward the SN explosion. In this paper, we report on detection and analyses of thermal X-ray emission from SN IIb 2011dh at $\sim 500$ days after the explosion on {\em Chandra} archival data, providing a solidly derived mass loss rate of an YSG progenitor for the first time. We find that the circumstellar media (CSM) should be dense, more than that expected from a Wolf-Rayet (WR) star by one order of magnitude. The emission is powered by a reverse shock penetrating into an outer envelope, fully consistent with the YSG progenitor but not with a WR progenitor. The density distribution at the outermost ejecta is much steeper than that expected from a compact WR star, and this finding must be taken into account in modeling the early UV/optical emission from SNe IIb. The derived mass loss rate is $\sim 3 \times 10^{-6} M_{\odot}$ yr$^{-1}$ for the mass loss velocity of $\sim 20$ km s$^{-1}$ in the final $\sim 1,300$ years before the explosion. The derived mass loss properties are largely consistent with the standard wind mass loss expected for a giant star. This is not sufficient to be a main driver to expel nearly all the hydrogen envelope. Therefore, the binary interaction, with a huge mass transfer having taken place at $\gsim 1,300$ years before the explosion, is a likely scenario to produce the YSG progenitor. | Evolution of a massive star in the final stage toward the supernova (SN) explosion is one of the main issues in modern stellar astrophysics. Evolutionary paths to SNe IIb/Ib/Ic, sometimes called stripped-envelope (SE-) SNe \citep{filippenko1997}, have been actively debated, since the problem here is directly related to a still-unresolved mass loss mechanism and/or binary evolution. With a progenitor being (nearly) a bare He or C+O star, as analogous to a Wolf-Rayet (WR) star, a question is how the progenitor's envelope has been stripped off for these SE-SNe \citep{nomoto1993,woosley1994,georgy2012,benvenuto2013}. Discovery and confirmation of an unprecedented Yellow Supergiant (YSG) progenitor of SN IIb 2011dh \citep{maund2011,vandyk2011,vandyk2013} brought many questions to the field, but at the same time, it provides a great opportunity to deepen our understanding on stellar evolution toward SE-SNe. With the main-sequence mass of the progenitor estimated as $\sim 12 - 15 M_{\odot}$ \citep{bersten2012}, a single-star evolution model requires a boost in the mass loss rate as compared to the value in the standard prescription \citep{georgy2012}, otherwise it will require a close-binary interaction at some point during the evolution \citep{benvenuto2013}. Observationally deriving properties of circumstellar materials (CSM), namely a mass loss rate, should shed light on this issue. For this purpose (among others), radio and X-ray observations have been actively conducted for nearby SE-SNe. The radio signal is however coupled with still-unresolved relativistic electron acceleration mechanism(s). For SN 2011dh, while a standard mass loss rate of Galactic Wolf-Rayet (WR) stars was {\em assumed} \citep{krauss2012,soderberg2012,horesh2013}, it was also clarified that the radio data (even complemented by optical data) provide only a loose constraint, $4 \lsim A_{*} \lsim 30$ \citep{maeda2012}. Hereafter, $A_{*} \sim (\dot M/10^{-5} M_{\odot} \ {\rm yr}^{-1}) (v_{w}/10^3 \ {\rm km}^{-1})^{-1}$ which is a measure of the CSM density normalized by a typical WR mass loss rate and wind velocity \citep{chevalier2006,chevalier2010}. This is defined by $\rho_{\rm CSM} = 5 \times 10^{11} A_{*} r^{-2}$ g cm$^{-3}$ where the radial position from the SN center (or the progenitor), $r$, is given in cm. We note that the radio properties are also related to another issue in the progenitor evolution. The shock velocity inferred from the radio data has been suggested to be one of the measures on the spatial extent of the progenitor's envelope, sometimes referred as `compact' (e.g., WR-like) or `extended' (e.g., RSG or YSG) \citep{chevalier2010}. However, a few recent examples have shown that this is not a perfect indicator of the progenitor's size. The examples include SN 2011dh, which was classified as a compact class from the radio data \citep{soderberg2012}, while the progenitor is identified as an YSG (see above) \citep[see also the case for SN IIb 2011hs: ][]{bufano2014}. Since the radio data provide constraints on the CSM density rather than the progenitor spatial size \citep[except for the very early phase data: ][]{maeda2013b}, this means that a relation between the CSM density (i.e., the mass loss rate and the mass loss velocity) and the progenitor's nature is yet to be clarified. This also provides strong motivation to study the CSM properties around SN 2011dh. X-rays provide additionally important information. For SN IIb 2011dh, X-ray observations have been performed in the first month after the explosion, by {\em Chandra}, {\em SWIFT/XRT} \citep{soderberg2012}, and by {\em XMM-Newton} \citep{campana2012,sasaki2012}. However, the interpretation of the X-ray emission mechanism is still controversy for SE-SNe \citep{chevalier2006,maeda2012}. {\em If the X-ray emission would be known} to be thermal, one can conclusively determine the CSM density \citep{fransson1996}, but a problem is that it has not been clarified if the X-ray is thermal or non-thermal for most SE-SNe except for SN IIb 1993J, which is among nearest {\em and} intrinsically brightest X-ray emitter as SE-SNe discovered so far. Generally the low photon statistics does not allow the spectral deconvolution into thermal and non-thermal components \citep[see e.g., ][]{chakraborti2013}, and it is also the case for SN IIb 2011dh \citep{soderberg2012,maeda2012}. For the typical CSM density of $A_{*} \sim 1$, thermal emission is not able to explain the observed X-ray flux level of typical SE-SNe when detected, and thus the inverse Compton (IC) scattering of SN photospheric photons has been sometimes invoked \citep{bjornsson2004,chevalier2006}. This is also the suggestion for SN 2011dh \citep{soderberg2012,maeda2012}, while there might be at least non-negligible thermal emission component as well \citep{sasaki2012}. In sum, problems in understanding the CSM density, the non-thermal electron acceleration mechanism, and the X-ray production mechanism are all coupled. In this paper, we analyze archival {\em Chandra} data of SN 2011dh at $\sim 500$ days. At this late epoch, the non-thermal processes are ineffective in creating X-ray photons, since the synchrotron emitting electrons should be in the cooling regime and there are not sufficient optical seed photons for the IC emission \citep[e.g., ][]{maeda2013a}. The late-time X-ray emission, if detected, must come from thermal emission either from a forward shock penetrating into the CSM or a reverse shock penetrating into the ejecta. By detecting the thermal X-ray emission from SN 2011dh, we are able to provide a conclusive determination of the mass loss rate. In \S 2, we describe data reduction and spectral analyses where we report on detection of strong late-time X-ray emission. In \S 3, we argue that the late-time X-ray emission is thermal emission from an adiabatic reverse shock penetrating into the outer envelope. In \S 4 (as complemented by discussion in Appendix), we discuss an origin of the early phase X-ray emission, in view of the natures of the reverse shock derived from the late-time emission. The paper is closed with conclusions and discussion in \S 5, where we discuss implications of our findings for the evolution toward the YSG progenitor and for non-thermal mechanism in the early phases. \begin{figure} \begin{center} \begin{minipage}[]{0.22\textwidth} \epsscale{0.95} \plotone{f1a.eps} \end{minipage} \begin{minipage}[]{0.22\textwidth} \epsscale{0.95} \plotone{f1b.eps} \end{minipage} \end{center} \caption {The {\em Chandra} X-ray image of SN 2011dh integrated in 467 - 498 days (right), as compared to the early phase image at 33 days (left). The energy band is 0.3 - 8 keV. The position of the detected source is consistent with the reported position of SN IIb 2011dh within a point spread function. The target spectra are extracted within the inner circular shown by a green circle, and the background is extracted within the annulus defined by the inner and outer circles. \label{fig1}} \end{figure} | We have found a strong X-ray emission from SN IIb 2011dh at $\sim 500$ days after the explosion. Such a late-time X-ray observation is quite rare for extragalactic SNe -- especially among SE-SNe, it has been done only for SN IIb 1993J \citep{immler2001} and SN Ic 1994I \citep{immler2002}. We have shown the thermal nature of the late-time X-ray emission, which must come from the reverse shock. The reverse shock must be in an adiabatic phase, as evidenced by various arguments. Therefore, we can provide a solid measurement of the CSM density as $A_{*} \sim 15$, which translates into the mass loss rate in the final $\sim 1,300$ years before the explosion as $\sim 3 \times 10^{-6} M_{\odot} (v_{w} / 20 \ {\rm km} \ {\rm s}^{-1})$. We have checked various sources of uncertainties in our estimate, and conclude that our derivation of the CSM density and the mass loss rate are robust. The discovery of the YSG progenitor of SN IIb 2011dh makes the progenitor evolution of this object highly interesting. The derived mass loss rate should be an extremely strong constraint on any evolution models, either a single or binary \citep{georgy2012,benvenuto2013}. We note that this is a rare case where the mass loss rate of a SE-SN progenitor has been conclusively determined in a way free from assumptions in radio and X-ray emission mechanisms. This is the first report to clarify many properties of the SN ejecta and the CSM through X-ray emissions for the YSG progenitor. In this respect, we find (1) high density CSM, (2) existence of the outer envelope of the progenitor at least $\sim 0.013 M_{\odot}$, (3) steep density gradient in the outer envelope. These are not expected from a WR progenitor, strengthening the case of the YSG progenitor and suggesting quite different properties in the progenitor/environment of a SE-SN from an extended progenitor as compared to a compact, WR-type star. The above properties are shared with a prototypical SN IIb 1993J. The CSM density is between `compact' SE-SNe \citep{chevalier2010} and SN 1993J with an RSG progenitor, linking the properties of the CSM along types of the progenitor star within SE-SNe. The steep outer density gradient ($n \sim 20$) is similar to what was found for SN IIb 1993J. A simple hydrodynamic model with either radiative or fully convective progenitor envelopes, $n \sim 10 - 12$, do not predict such steep density gradient. This will be a strong constraint on the nature of the envelope of the YSG (and RSG) progenitor. Also, we note that this finding is important in interpreting the very early UV/optical emission from SE-SNe with an extended progenitor. Frequently a fiducial value of $n \sim 10$ is assumed in analyzing such data \citep[e.g., ][]{rabinak2011}, but this would lead to significant errors if applied to SE-SNe from a progenitor with an extended envelope. The derived mass loss rate is indeed close to the `standard wind' mass-loss rate, in the final centuries toward the core-collapse, adopted in stellar evolution models of stars with $M_{\rm ZAMS} \sim 12 - 15 M_{\odot}$\citep{georgy2012}. Note that this mass range was well constrained for SN 2011dh through optical emission models \citep{bersten2012}. We emphasize that the CSM density distribution is close to that expected from the steady-state wind, with variation at the level of $\sim 30$\% within the last thousand years. Such a mass loss is unable to expel all the hydrogen-rich envelope within the life time of the progenitor star at the giant stage. Namely, our results suggest that in the final centuries the mass loss was just consistent with the YSG wind mass loss, and this steady mass loss cannot be the main driver of the whole hydrogen envelope stripping - this indicates that the mass loss rate was at some point much higher than that in the last thousand years before the explosion, supporting the binary interaction \citep{benvenuto2013} as the main driver of the mass loss to produce the YSG progenitor. Deriving the mass loss rate (and the CSM density) provides also a very strong constraint in understanding still-debated non-thermal/radiation physics in these wavelengths. With $A_{*} \sim 15$ as robustly determined, we can address the efficiency of the acceleration of non-thermal electrons producing the early-phase synchrotron radio emission as $\epsilon_{\rm e} \sim 0.01$ and the efficiency of the magnetic field amplification as $\epsilon_{B} \sim 5 \times 10^{-3}$ \citep{maeda2012}. By estimating the contribution of thermal emission from the reverse shock to the X-ray luminosity in the early phase, we conclude that it is very likely that the X-ray at 33 days was dominated by the same thermal component. The thermal emission from the reverse shock also accounts for about half of the X-ray luminosity at $\lsim 10$ days. Therefore, the soft persistent component in the early phase \citep{sasaki2012} should be identified as the reverse shock thermal emission. The early-phase X-ray spectra showed an additional hard component \citep{sasaki2012}. We find that it is unlikely that this component is the thermal emission (bremsstrahlung) from a forward shock -- such a solution requires a very large density jump by a factor of $\sim 15$ and this is inconsistent with radio data (see Appendix). An interesting alternative is the IC scattering of the SN photospheric photons \citep{bjornsson2004,soderberg2012,maeda2012}, and in this case the non-thermal electrons at low energy ($\gamma \lsim 50$) should have a flatter spectrum than those at higher energy. This may be consistent with the earlier suggestion that the IC scenario requires that these low energy electrons are a different component than the radio-synchrotron emitting electrons \citep{maeda2012}. While we suggest the IC scenario is more likely than the thermal emission from the forward shock, more detailed study will be necessary to resolve this issue. The shock wave has expanded to $\sim 8 \times 10^{16}$ cm at $\sim 500$ days and the swept ejecta mass is $M_{\rm RS} \sim 0.013 M_{\odot}$, which is either H-rich or He-rich. This indicates that the reverse shock should be still in the outer envelope of the progenitor star. Namely, we constrain the mass of the envelope in the YSG progenitor of SN IIb 2011dh to be $\gsim 0.013 M_{\odot}$. We note this is fully consistent with the early-phase optical spectral modeling, from which the hydrogen mass of $\sim 0.024 M_{\odot}$ in the outer layer has been derived \citep{arcavi2011}. Our estimate of the CSM density is not sensitive to particular interpretation. For example, we have shown that the derived CSM density is insensitive to the composition in the outer envelope. We also checked this issue, by performing the same analyses based on the one-component sub-solar model. There we have confirmed that all the arguments presented with the two-component model are qualitatively reproduced by the sub-solar model as well. The only difference, in a quantitative sense, is that the required CSM density is reduced by a factor of $\sim 2$. This stems from overall low temperature in the sub-solar model, lacking the high temperature component, leading to a larger cooling rate than in the two-component model. While we regard such a low metallicity model unlikely, we note that once an independent estimate of the metallicity is given, that will discriminate these two scenarios. In any case, this would not alter any of our conclusions. Finally, we note that, while our analyses are based on the assumption that the CSM density is represented by a smooth distribution ($\rho_{\rm CSM} \propto r^{-s}$), most of our conclusions are not influenced by this assumption. First of all, the derived relation between the density and the mass of the emitting region (i.e., equation 5) is basically independent from this assumption. Since the reverse shock position within the ejecta outer envelope is mostly determined by the CSM mass already encountered by the forward shock, the mass loss rate derived by this study can be regarded as the {\em average} mass loss in the final 1,300 years in the case with the huge CSM density variation (assuming $v_{\rm w} = 20$ km s$^{-1}$). Interpreting the X-ray emission as a thermal emission from the {\em forward} shock at the putative dense CSM will lead to a similar conclusion, since the density behind the forward shock is required to be similar to the reverse shock case, as constrained by the observed characteristic energy scale in the X-ray spectrum. Moreover, we emphasize that the interpretation based on the smooth CSM density distribution provides consistent results to various observed behaviors, and indeed the derived distribution is close to the steady-state mass loss. Therefore, it does not require the dense CSM shell, and as such we suggest that the smooth CSM distribution is very likely the case. | 14 | 3 | 1403.2455 |
1403 | 1403.0813_arXiv.txt | SONYC -- Substellar Objects in Nearby Young Clusters -- is a survey programme to investigate the frequency and properties of substellar objects in nearby star-forming regions. We present a new imaging and spectroscopic survey conducted in the young ($\sim1\,$Myr), nearby ($\sim200\,$pc) star-forming region Lupus$\,$3. Deep optical and near-infrared images were obtained with MOSAIC-II and NEWFIRM at the CTIO-4m telescope, covering $\sim1.4\,$deg$^2$ on the sky. The $i$-band completeness limit of 20.3$\,$mag is equivalent to $0.009-0.02$ \solm, for A$_V\leq5$. Photometry and $11-12\,$yr baseline proper motions were used to select candidate low-mass members of Lupus~3. We performed spectroscopic follow-up of 123 candidates, using VIMOS at the Very Large Telescope (VLT), and identify 7 probable members, among which 4 have spectral type later than M6.0 and $T_{\mathrm{eff}}\leq3000\,$K, i.e. are probably substellar in nature. Two of the new probable members of Lupus$\,$3 appear underluminous for their spectral class and exhibit emission line spectrum with strong H$_{\alpha}$ or forbidden lines associated with active accretion. We derive a relation between the spectral type and effective temperature: T$_{\mathrm{eff}}=(4120\pm175)-(172\pm26)\times\mathrm{SpT}$, where SpT refers to the M spectral subtype between 1 and 9. Combining our results with the previous works on Lupus$\,$3, we show that the spectral type distribution is consistent with that in other star forming regions, as well as is the derived star-to-BD ratio of $2.0-3.3$. We compile a census of all spectroscopically confirmed low-mass members with spectral type M0 or later. | \label{intro} SONYC - short for {\it Substellar Objects in Nearby Young Clusters} - is a comprehensive project aiming to provide a complete, unbiased census of substellar population down to a few Jupiter masses in young star forming regions. Studies of the substellar mass regime at young ages are crucial to understand the mass dependence in the formation and early evolution of stars and planets. Although the low-mass end of the Initial Mass Function (IMF) has been the subject of intensive investigation over more than a decade, and by various groups, its origin is still a matter of debate (e.g. \citealt{bonnell07, bastian10, jeffries12b}). The relative importance of several proposed processes (dynamical interactions, fragmentation, accretion, photoevaporation) responsible for the formation of brown dwarfs (BDs) is not yet clear. The SONYC survey is based on extremely deep optical- and near-infrared wide-field imaging, combined with the Two Micron All Sky Survey (2MASS) and $Spitzer$ photometry catalogs, which are correlated to create catalogs of substellar candidates and used to identify targets for extensive spectroscopic follow-up. In this work for the first time we also include a proper motion analysis, which greatly facilitates the candidate selection. Our observations are designed to reach mass limits well below 0.01$\,$M$_{\odot}$, and the main candidate selection method is based on the optical photometry. This way we ensure to obtain a realistic picture of the substellar population in each of the studied cluster, avoiding the biases introduced by the mid-infrared selection (only objects with disks), or methane-imaging (only T-dwarfs). So far we have published results for three regions: NGC 1333 \citep{scholz09,scholz12a,scholz12b}, $\rho$ Ophiuchi \citep{geers11,muzic12}, and Chamaeleon-I \citep{muzic11}. We have identified and characterized more than 50 new substellar objects, among them a handful of objects with masses close to, or below the Deuterium burning limit. Thanks to the SONYC survey and the efforts of other groups, the substellar IMF is now well characterized down to $\sim5 - 10 $M$_J$, and we find that the ratio of the number of stars with respect to brown dwarfs lies between 2 and 6. In NGC1333 we find that, down to $\sim5$M$_J$, the free-floating objects with planetary masses In this paper we present the SONYC campaign in the Lupus$\,$3 star forming region. The Lupus dark cloud complex is located in the Scorpius-Centaurus OB association and consists of several loosely connected dark clouds showing different levels of star-formation activity (see \citealt{comeron08} for a detailed overview). The main site of star formation within the complex is Lupus$\,$3, which contains one of the richest associations of T-Tauri stars \citep{schwartz77, krautter97}. The center of Lupus$\,$3 is dominated by the two most massive members of the entire complex, the two Herbig Ae/Be stars known as HR~5999 and HR~6000. More than half of the known Lupus$\,$3 members are found in the $0.3 \times 0.3$ pc$^2$ area surrounding the pair \citep{nakajima00, comeron08}. The low-mass (sub-)stellar content of Lupus$\,$3 was extensively investigated using the data from the Spitzer Space Telescope. % \citet{merin08} compiled the most complete census of stars and brown dwarfs in Lupus$\,$3 at the time, using the data from their survey, together with the results of previous surveys by \citet{nakajima00, comeron03,comeron08,lopez-marti05, allers06,allen07,chapman07,tachihara07,merin07,strauss92,gondoin06}. Spectroscopic follow-up in the optical (FLAMES/VLT) by \citet{mortier11} confirmed the effectiveness of MIR-excess selection from \citet{merin08}, with about 80\% of the 46 observed sources confirmed as members of Lupus$\,$3. Two wide-area photometric surveys in Lupus were conducted using the Wide Field Imager (WFI) at the La Silla 2.2-m telescope. \citet{lopez-marti05} identified 22 new low-mass member candidates in an area of 1.6 deg$^2$ in Lupus$\,$3, with about half of the candidates confirmed spectroscopically in surveys by \citet{allen07} and \citet{mortier11}. \citet{comeron09} surveyed an area of more than 6 deg$^2$ in the Lupus~1, 3, and 4 clouds and identified $\sim$70 new candidate members of Lupus$\,$3. About 50\% of the photometric sample was revealed to belong to a background population of giant stars in a follow-up spectroscopic study by \citet{comeron13}. The deepest survey so far in Lupus$\,$3 \citep{comeron11}, did not report any brown dwarfs, despite the sensitivity to objects with sub-Jupiter masses. However, the surveyed area of $\sim 7' \times 7'$ was very small compared to the total extent of the cluster, and thus contained only a minor fraction of the total population. The substellar population in Lupus$\,$3 remains therefore poorly constrained. The distance to the Lupus star forming region is still a matter of debate, with the distance determinations from different studies ranging from 140 pc to 240 pc \citep{comeron08}. A widely-used value of $140\pm20\,$ pc was derived by \citet{hughes93} using photometry and spectroscopy of F-G field stars in Lupus. Using an improved version of the same method, \citet{lombardi08} obtained a distance of $155\pm8\,$ pc, but also suggest that the apparent thickness of Lupus might be the result of different Lupus sub-clouds being at different distances. Distances obtained from the moving-cluster method \citep{makarov07} suggest a thickness of $\sim80\,$pc, and places Lupus$\,$3 at a distance of about 170 pc, or 25 pc farther than the center of the greater association. This spatial arrangement is in general agreement with the distances derived from Hipparcos parallaxes of individual stars located in different sub-clouds \citep{wichmann98, bertout99}. As pointed out by \citet{comeron08}, while a distance of 150 pc seems adequate for most of the clouds of the complex, a value of 200 pc is likely to be more appropriate for Lupus$\,$3. We therefore adopt a distance of 200 pc throughout the analysis presented in this paper. Adopting the distance of 200 pc, and evolutionary tracks of \citet{baraffe98}, \citet{comeron03} derive an age of 1-1.5 Myr for most of their observed members. In this work we present new observations of the Lupus$\,$3 cloud using the MOSAIC-II optical imager and NEWFIRM near-infrared imager at the CTIO-4m telescope. Observations and data reduction are explained in Section~\ref{Obs&DR}. Photometry, proper motions and criteria for candidate selection are presented in Section~\ref{candsel}, and the spectral analysis in Section~\ref{specanal}. The results are discussed in Section~\ref{discuss}. Finally, we summarize the main conclusions in Section~\ref{summary}. \begin{figure*} \centering \resizebox{14cm}{!}{\includegraphics{Lupus3_spatial.jpg}} \caption{{\it Upper panel:} Spatial distribution of the candidate sources in Lupus$\,$3. Photometric candidates are shown as black dots, those also selected by proper motions are shown in yellow. Diamonds mark previously known members, with the two red stars marking the brightest pair, HR5999/6000. Solid line shows the field-of-view of our optical and near-infrared survey, and the dotted line the extent of the WFI field used as the first epoch for the proper motions. The dashed lines indicate the region within which the VIMOS spectroscopic fields have been arranged. The shaded contours indicate A$_V$ from the extinction map by \citet{cambresy99}. {\it Lower panel:} Zoom into the region covered by spectroscopy. All the candidates observed in our follow-up are marked with squares, and the confirmed objects are represented by red diamonds. } \label{spatial} \end{figure*} \begin{figure*} \centering \resizebox{15cm}{!}{\includegraphics{Lupus3_IJcmd_2.jpg}} \caption{($i$,$i - J$) color-magnitude diagram. Open circles represent all the photometrically-selected candidates located inside the selection box (dashed lines), with those additionally selected on the basis of their proper motions shown as filled circles. Diamonds show SONYC spectroscopy follow-up targets, with spectroscopically confirmed VLMOs highlighted as red (identified in this work) and green (previously known members) stars. Atmosphere model isochrones for age 1 Myr are overplotted for DUSTY00 \citep{chabrier00} and BT-Settl \citep{allard11}. For clarity, the sources outside the candidate selection box are represented with contours, where the number on each contour represents the number of sources within a 0.2 mag $\times$ 0.2 mag bin. On the right-hand side of the figure we show 1$\sigma$ uncertainties of the photometry. The dotted line shows the completeness limit of our survey. } \label{IJCMD} \end{figure*} | \label{summary} In this work, we have presented deep optical and near-infrared images of the 1.4$\,$deg$^2$ area surrounding the two brightest members HR5999/6000 of the Lupus$\,$3 star forming region. From the optical+NIR photometry we selected 409 candidate VLM cluster members. Proper motion analysis, based on two epochs of imaging separated by 11-12 years, helped to narrow down the candidate selection to 59 high priority candidates. To confirm the membership of the selected candidates, we performed spectroscopic follow-up using VIMOS/VLT, in which we collected 123 spectra from the photometric selection box, including 32 from the high priority list. We confirm 7 candidates as probable members of Lupus$\,$3, among which 4 are later than M6.0 and with $T_{\mathrm{eff}}\leq 3000\,$K, i.e. are probably substellar in nature. Two of the sources identified as probable members of Lupus$\,$3 appear underluminous for their spectral class, similar to some previously known members exhibiting similar emission line spectra with strong H$_{\alpha}$ and several forbidden lines associated with active accretion. We derive a relation between the spectral types (from comparison to low-gravity objects in young regions) and $T_{\mathrm{eff}}$ (from BT-Settl models): $T_{\mathrm{eff}}$ = $4083 - 166$ $\times$ SpT, where SpT refers to the M spectral subtype between 1 and 9. The rms-error of the fit of 61 K. We derive a star-to-BD ratio of 2.0 - 3.3, consistent with the values observed in other star forming regions. Combining our results with previous work on Lupus$\,$3, we compile a table containing all spectroscopically confirmed low-mass objects with spectral type M0 or later, and show that the distribution of spectral types is in line with what is observed in other young star forming regions. \appendix | 14 | 3 | 1403.0813 |
1403 | 1403.1612.txt | Particles count rates at given Earth location and altitude result from the convolution of (i) the interstellar (IS) cosmic-ray fluxes outside the solar cavity, (ii) the time-dependent modulation of IS into Top-of-Atmosphere (TOA) fluxes, (iii) the rigidity cut-off (or geomagnetic transmission function) and grammage at the counter location, (iv) the atmosphere response to incoming TOA cosmic rays (shower development), and (v) the counter response to the various particles/energies in the shower. Count rates from neutron monitors or muon counters are therefore a proxy to solar activity. In this paper, we review all ingredients, discuss how their uncertainties impact count rate calculations, and how they translate into variation/uncertainties on the level of solar modulation $\phi$ (in the simple Force-Field approximation). The main uncertainty for neutron monitors is related to the yield function. However, many other effects have a significant impact, at the $5-10\%$ level on $\phi$ values. We find no clear ranking of the dominant effects, as some depend on the station position and/or the weather and/or the season. An abacus to translate any variation of count rates (for neutron and $\mu$ detectors) to a variation of the solar modulation $\phi$ is provided. | After the discovery of cosmic rays (CR) by Hess in 1912, ground-based CR detectors located at various latitudes, longitudes and altitudes, played a major role to determine the CR composition and spectrum (see \citealtads{2009AdSpR..44.1081S} for a historical perspective). From the 50's, networks of neutron monitors \citepads{2000SSRv...93...11S} and muon telescopes \citepads{2000SSRv...93..207D} were developed. They provide today one of the most valuable data to inspect time variations of the integrated CR flux in the $10-100$ GeV/n range. The formal link between these variations and the Sun activity was established in the mid-fifties, by means of a transport equation of CR fluxes in the solar cavity (\citeads{1965P&SS...13....9P}, \citealtads{1966ApJ...146..480J}). In the 80's, the effect of particle drift was shown to be responsible to a charge-sign dependent modulation \citepads{2013LRSP...10....3P}, following the Sun polarity cycle\footnote{A 11-yr average periodicity was established $\sim 250$ yrs ago from sunspot series \citep[see][for a review]{2007AdSpR..40..929V,2013LRSP...10....1U}. The now well observed 22-yr cycle (polarity reversal every 11 yrs) was first hinted at from magnetograph observations by \citetads{1961ApJ...133..572B}.}. However, the Force-Field approximation \citepads{1967ApJ...149L.115G,1968ApJ...154.1011G} has remained widely used thanks to its simplicity: this approximation, used in this work, has only one parameter $\phi(t)$. Several strategies have been developed for time series reconstruction of the modulation level $\phi(t)$, and/or CR TOA fluxes at any time (of interest for many applications): \begin{itemize} \item Using spacecraft measurements \citepads[e.g.,][]{2001JGR...10629979D,2011P&SS...59..355B,2013SoPh..284..599B}: it is the most direct approach, but the time coverage is limited to a few decades with a poor sampling; \item Comparison of calculated and measured count rates in ground-based detectors~\citepads{1999CzJPh..49.1743U,2002SoPh..207..389U,2005JGRA..11012108U,2011JGRA..116.2104U}: it covers a larger period (60~yrs), with a very good time resolution (a few minutes)\footnote{In the same spirit, the concentration of the cosmogenic radionuclide $^{10}$Be in ice cores \citepads{2003JGRA..108.1355W,2010JGRA..11500I20H} covers several thousands of years, but with a poor time resolution.}; \item Extracting relationships between the modulation level and solar activity proxies, based on empirical (\citealtads{1994AdSpR..14..749B,1996AdSpR..17....7B}; \citealtads{2006AdSpR..37.1727O,2010ITNS...57.3148O}) or semi-empirical (\citealt{R.A-1992}; \citealtads{1994AdSpR..14..759N,1996AdSpR..17...19N}; \citealtads{1997ITNS...44.2150T}; \citealtads{2007AdSpR..40..313N}; \citealtads{2013AdSpR..52.2112A}) approaches. \end{itemize} All these strategies provide a satisfactory description of CR fluxes, though some fare better than others (for comparisons, see \citealtads{2010AdSpR..45.1026B,2012JGRA..117.8109M,2013JGRA..118.1837Z,2013AdSpR..51..329M}). Note also that empirical methods are expected to provide effective and less meaningful values for $\phi$ \citepads{2006AdSpR..37.1727O}. In this paper, we focus on the second strategy, for a systematic study of the main uncertainties affecting the calculation of expected count rates in NM and muon detectors. This requires the description of the atmosphere and of the ground-based detector responses to incoming CRs \citep[e.g.,][]{2000SSRv...93..335C}. The various uncertainties, described in the \citetads{1974crvs.book.....D,2004ASSL..303.....D,2009crme.book.....D} textbooks, are generally discussed separately in research articles (uncertainty on the yield function, geomagnetic rigidity cutoff, seasonal effects\dots), and not propagated to the modulation parameter. For this reason, we believe it is useful to recap and gather them in a single study, re-assess which ones are the most important, and link these uncertainties to the expected level of variation/uncertainty they imply on the modulation level $\phi(t)$. The complementarity (different uncertainties and time coverage) of NM count rates and TOA CR flux measurements to obtain time-series of the solar modulation parameters is left to a second study\footnote{Recent CR instruments such as \href{http://pamela.roma2.infn.it}{PAMELA} and the \href{http://www.ams02.org}{AMS-02} on the International Space Station are or will be providing high-statistics fluxes on an unprecedented time frequency, which renders this comparison even more appealing.}. The paper is organised as follows: we start with a general presentation of the ingredients involved in the count rate calculations (Sect.~\ref{sec:count_rate}), and discuss a new fit for the IS fluxes (Sect.~\ref{sec:ref_is}). We then detail the calculation of the propagation in the atmosphere, providing a new yield function parametrisation (Sect.~\ref{sec:Yield}). Combining these inputs allows us to link the count rate variation with the solar modulation parameter, and to study the various sources of uncertainties (Sect.~\ref{sec:impact_factors}). The final ranking of the uncertainties in terms of both count rates and $\phi$ concludes this study (Sect.~\ref{sec:concl}). %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | \begin{enumerate}[a)] \item using a sigmoid function instead of a step function gives $\sim 4\%$ less count rates for NM64, and $\sim 2\%$ less for $\mu$ detectors. This effect is $R_c$ dependent, and maximal for large $R_c$ values (see bottom panel of Fig.~\ref{rate_err_rctransm}); \item even in the step-function approximation, the count rate variation is expected to change due to long-term or short term geomagnetic variations. We evaluate that over 50 years a typical decrease of $\sim 4\%$ for NM64 (and $\sim 1\%$ for $\mu$ detectors) can occur (see top panel of Fig.~\ref{rate_err_rctransm}). The level of the variation depends on the geomagnetic position and $R_c$; \item the use of the apparent cut-off rigidity of \citetads{1997JGR...10226919C} and \citetads{2008AdSpR..42..510D} accounting for obliquely incident particles (in the geomagnetic field) is found to have an impact of $\lesssim 2-4\%$ on the NM64 count rate and $\lesssim 1-2\%$ on $\mu$ detectors). As above, the effect depends on the geomagnetic position and $R_c$. \end{enumerate} \item We recap the various seasonal effects and their impact on count rates. First, muon detector data are dominated by temperature effects: the corresponding count rate variation is $\sim 8\%$ but corrections (which are seldom implemented in distributed data) are able to bring this variation down to $\sim 0.3\%$ \citepads{2011APh....34..401D}. Second, for NM64, all the following effects must be considered: \begin{enumerate}[a)] \item atmospheric pressure and temperature effects ($\sim 10\%$ level) are routinely corrected for in public data. The level of variation left after this correction is $\lesssim 0.5\%$ (pressure) and $\lesssim 1\%$ (temperature); \item water vapour is expected to lead to a $\sim 0.2-0.3\%$ effect; \item the effect of snow coverage in the surrounding of the detector is investigated by means of BSS measurements whose low energy spectrum is very sensitive to it. We obtain a $2-8\%$ seasonal variation for this effect (obviously strongly dependent on the climatic conditions at the station location), in agreement with direct measurements in NM stations. Recent efforts by \citetads{2011Ge&Ae..51..247K,2013JGRA..118.6852K} to provide real-time data corrected for this effect are an important step for the network of NMs around the world. \end{enumerate} \item Finally, some uncertainties are intrinsic to the detector itself, as thoroughly investigated by means of a calibrator \citepads{2010AdSpR..46.1394K}. These authors find a spread $\sim 4\%$ in their measurements attributed to local conditions, but it may be even larger for some stations. However, such effects, along with differences attributed, e.g., to the exact geometry of the detector \citepads{1964CaJPh..42.2443H}, are not expected to change in time, and thus are probably not as problematic as seasonal effects. \end{itemize} \begin{figure*}[!t] \begin{center} \includegraphics[width=0.995\columnwidth]{figs/ratevar2d_vs_rc_and_phi_nm64.pdf} \includegraphics[width=0.995\columnwidth]{figs/ratevar2d_vs_rc_and_phi_muon.pdf} \includegraphics[width=0.995\columnwidth]{figs/deltarate2d_vs_rc_and_phi_nm64.pdf} \includegraphics[width=0.995\columnwidth]{figs/deltarate2d_vs_rc_and_phi_muon.pdf} \end{center} \vspace{-4mm} \caption{Left panels are calculated for NM64 and right panels for $\mu$ detectors. {\bf Top panels:} Count rate relative variation $\Delta N/N_{\otimes}$ with respect to a reference count rate $N_{\otimes}=N(\phi=0.5,\,R_c=0)$. The relative variation (in \%) are shown as iso-contours in the plane $(R_c,\,\phi)$, with the 0\%-isocontour (passing through the reference point $\otimes$) in black solid line. {\bf Bottom panels:} scaling factor $f$ to infer the count rate relative variation $[\Delta N_{\delta\phi/\phi}/N](R_c,\,\phi)$ for any modulation relative change $\delta\phi/\phi$ (for NM64, slight differences in the contours arise for values of $\delta\phi/\phi>20\%$ and factor $f<-0.3$). For instance, for a NM64 detector at $R_c=6.5$~GV and a solar modulation period $\phi=1.2$~GV, the scaling factor is $f=-0.25$, which reads: an increase of 5\% (resp. 10\%) in the modulation level $\phi$ translates in a decrease of $0.25\times5\%=1.25\%$ (resp. 2.5\%) in the detector count rate, and vice versa.} \label{deltan_vs_phi} \vspace{-3mm} \end{figure*} \begin{table*}[!t] \caption{Impact of different effects/input ingredients ($1^{\rm st}$ and$2^{\rm nd}$ columns) on the relative count rate calculation for NMs and $\mu$-like detectors ($4^{\rm th}$ and $5^{\rm th}$ columns), and on the uncertainty associated to the derived modulation parameter value ($6^{\rm th}$ and $7^{\rm th}$ columns). The first two rows, in \textit{italic}, serve as a reference: they correspond to the maximum variation expected between periods of low and high solar activity, and low and high rigidity cut-off). The $3^{\rm rd}$ column provides the figure or section where the effect is discussed in the paper. See text for description.} \vspace{-3mm} \label{tab:summary_table} \begin{center} \begin{tabular}{cccccccc}\hline & & & & & & &\vspace{-3mm}\\ Ingredient & Effect &\!\!\!\!\!\!\!\!Fig./Sect.\!\!\!\!\!\!\!\!\!\!\!\!\!\!& \multicolumn{2}{c}{$\displaystyle\frac{\Delta N}{N}$}& \multicolumn{2}{c}{$\Delta\phi^\ast$ [MV]} & Comment \\ & & & NM & $\mu$ & NM & $\mu$ &\vspace{2mm}\\ \!\!\!\!\!{\em Solar modulation}\!\!\!\!\!& $\phi\in$~{\em [0.2,1.5]~GV} &{\em Fig.\ref{deltan_vs_phi}}&{\em [+15,-25]\%}&{\em [+5,-10]\%}&{\em -} &{\em -}& {\em w.r.t. $\mathit{\phi\!=\!0.5}$~GV}\\ \!\!\!{\em Cut-off rigidity}\!\!\!& $R_c\in$~{\em [0,10]~GV}& {\em Fig.\ref{deltan_vs_r}} & {\em [+10,-20]\%} & {\em [0,-5]\%} & {\em -} &{\em -}& {\em w.r.t. $\mathit{R_c\!=\!5}$~GV} \vspace{1.5mm}\\\hline % & & & & & & &\vspace{-2.5mm}\\ \multirow{3}{*}{TOA flux} & p and He CR data &Fig.\ref{rate_err_yieldjis}& $\pm2\%$ & $\pm2\%$ & $\pm66$ & $\pm140$& $(t,\,R_c,\,\phi)$-independent\\ & IS flux dispersion$^\P$ &Fig.\ref{rate_err_yieldjis}& $\pm6\%$ & $\pm8\%$ & $\pm200$ & $\pm570$& $\Downarrow$\\ & Heavy species &Fig.\ref{fraction_cr} & $\pm0.6\%$ & $\pm0.6\%$ & $\pm20$ & $\pm40$ & Global norm. factor$^\diamond$\vspace{3mm}\\ % \!\!\!Yield function\!\!\! & Dispersion &Fig.\ref{rate_err_yieldjis}&$\lesssim\pm4\%$& $<0.2\%$ &$\lesssim120$&$\lesssim14$&$(R_c,\,\phi)$ dependent\vspace{3mm}\\ % & Sigmoid$(R_c,x\!=\!+\frac{\sigma}{0.1})$&Fig.\ref{rate_err_rctransm}&-$2x\%$&-$0.5x\%$& +$66x$& +$35x$ &For $R_c\gtrsim5$~GV\\ Transfer &\!\!\!\!\!$H(R_c\!\!+\!\!\Delta R_c)$: $x\!\!=\!\!\frac{\!\!(\Delta R_c\!/\!R_c)\!\!}{0.05}$&Fig.\ref{rate_err_rctransm}&-$2x\%$& -$x\%$& +$66x$ & +$71x$& For $R_c\gtrsim5$~GV\\ function &- $R_c(t)$: $\frac{\Delta R}{R}\!\lesssim\!+0.2\%$/yr\!\!\!\!\!\!&\S\ref{sec:Rcut}&-0.4\%/yr\!\!\!\!\!\!\!\!\!&-0.1\%/yr\!\!\!\!\!\!\!\!\!&+13/yr\!\!\!\!\!&+7/yr\!\!\!& Depends on location\\ &\!\!\!\!\!\!- $R^{\rm eff}_c\!\rightarrow\!R^{\rm app}_c$: $\!+3\%$&\S\ref{sec:Rcut}& -1.2\%& -0.3\% &+40~ &+21~ & Depends on $R_c$\vspace{3mm}\\ % & Pressure & \S\ref{sec:pe} & $\pm0.2\%$ & $\pm0.2\%$ & $\pm6$ & $\pm14$ & After correction\\ Time-dep. & Temperature & \S\ref{sec:te} & $\pm0.5\%$ & $\pm4\%$ & $\pm15$ & $\pm290^\ddagger$& Not corrected\\ effects$^\dagger$ & Vapour water & \S\ref{sec:wv} & $\pm0.3\%$ & $\pm0.1\%$ & $\pm10$ & $\pm8$ & Not corrected\\ &\!\!\!\!\!\!\!\!Snow coverage ($T\!\!=\!\!1$~yr)\!\!\!\!\!\!\!& \S\ref{sec:se} & -7\% & - & +230 & - & Not corrected\vspace{3mm}\\ % & Diurnal variation & \S\ref{sec:calib} & 0.24\% & 0.24\% & 8 & 17 & NC$^\ddagger$ (24h period)\vspace{3mm}\\ % NM detector & Temperature & \S\ref{sec:calib} &+0.05\%/$^\circ$C\!\!\!\!\!\!& - & -1.5/$^\circ$C\!\!\!\!\!\! & - & $(t,\,R_c,\,\phi)$-independent\\ effects & $n$NM6 vs $m$NM64 & \S\ref{sec:calib} & few \% & - & $\sim 100$ & - & $\Downarrow$\\ & Surroundings (hut) & \S\ref{sec:calib} & few \% & - & $\sim 100$ & - & Global norm. factor$^\diamond$\\\hline \end{tabular}\\ \vspace{1mm} {\small $^\ast$ The variation $\Delta\phi$ of the modulation level is calculated for a detector at $R_c=5$~GV and $\phi=0.5$~GV:\\ refer to Fig.~\ref{deltan_vs_phi} to convert rate variations for any other $(R_c,\,\phi$) condition.\\ $^\P$ Very conservative estimate (some IS fluxes are based on old CR data).\\ $^\diamond$ Global normalisation factors can always be absorbed in the yield function normalisation.\\ $^\dagger$ Distributed data are either corrected or not corrected for these effects.\\ $^\ddagger$ After correction, $\Delta N/N\sim 0.3\%$ \citepads{2011APh....34..401D}, leading to $\Delta\phi\sim 22$~MV.} \end{center} \vspace{-6mm} \end{table*} \subsection{Abacus: count rates to solar modulation variations} To conclude, we propose a last figure and a table for a panoptic view of all the effects we have approached in this study. Actually, these plots provide a direct link between solar modulation level and count rate variations (and vice versa) for both NM64 and $\mu$ detectors. The top panels of Fig.~\ref{deltan_vs_phi} provide the relative count rate variation in the $R_c-\phi$ plane, with respect to a reference point $N_\otimes(R_c=0,\phi=0.5$~GV). In addition to providing a global view of the expected count rate variation between detectors at different $R_c$ and for different solar periods, it also gives a flavour of the precision required in order to be sensitive to changes in the $\phi$ parameter: the count rate variation over a full solar cycle is smaller for $\mu$ detectors than for NMs, but the latter are more sensitive to any uncertainty on $R_c$ (location in the geomagnetic field) than the former. The bottom panels of Fig.~\ref{deltan_vs_phi} go further in that direction, as they directly provide, for any value $(R_c,\,\phi)$, how much variation $\Delta N/N$ to expect in the count rates, whenever the solar modulation changes by $\delta\phi/\phi$. This abacus usage is two-folds: first, on short term variations, it can directly be used to extract $\delta\phi/\phi$ from count rate variation in NM (or $\mu$) data; second, it can be used to estimate how much uncertainty is propagated in $\phi$ from the various uncertainties calculated on count rates. This is what is gathered in Table~\ref{tab:summary_table}: for each input/effect discussed in the paper, we provide (in addition to the section/figure where it was dealt with) the typical uncertainty obtained on $\Delta N/N$, and the associated $\Delta\phi$ calculated for an NM64 or $\mu$ detector at $R_c=5$~GV and a solar modulation level of $\phi=500$~MV (using Fig.~\ref{deltan_vs_phi}). For $\phi$ calculations, the first thing to underline is that NM and $\mu$ detectors do not suffer the same amount of uncertainties, due to different sensitivities to the various effects explored. Moreover, there are no clear-cut ranking of these errors. Luckily, when interested in time series, time-independent normalisation effects can be absorbed in a normalisation factor \citepads[e.g.,][]{2011JGRA..116.2104U}: the latter accounts for differences in NM detector efficiency and their surroundings (last entries in table). The case of the IS fluxes (first entries in table) is peculiar, since different choices $i$ lead to an overall shift $\Delta\phi_i$ of the time series. For NMs, the main source of uncertainties are the seasonal snow effects (strength depending on position, some stations not affected), and the yield function dispersion (applicable for all stations). All other effects cannot be simply disregarded as they typically have a $5-10\%$ on $\phi$. For $\mu$ detectors, the main effect is that of the temperature variation, but after corrections, it is at the level of other uncertainties ($5-10\%$). Overall, $\mu$ detectors seem to suffer slightly less uncertainties than NM64, but of course the latter benefit from a much larger time and position coverage than the former. \subsection{Future works} The approach we have followed in this study could easily be extended to other types of ground-based measurements, by simply using the appropriate yield function for each type (e.g., $^{10}$Be production in ice cores, \citealtads{2010JGRA..11500I20H}; ionisation measurements in the atmosphere, \citealtads{2008SSRv..137..149B}; etc.). In any case, one of the main challenge of such approaches is to obtain an accurate yield function. In that respect, the efforts to improve that of NMs should be pursued, given their role in the history of solar activity monitoring. As underlined at the beginning of this study, our primary goal is to get time series of modulation parameters, taking advantage of the complementarity of NM count rates and TOA CR flux measurements. The above Table~\ref{tab:summary_table} provides a synthetic view of the difficulties. This table, and more importantly, the characterisation of the dependence of these uncertainties with $R_c$, `weather' conditions for the stations, etc., should help decide which stations to consider to minimise the uncertainties in the $\phi$ calculations. This is the aim of our next study. | 14 | 3 | 1403.1612 |
1403 | 1403.1212_arXiv.txt | We present a scenario where dark matter is in the form of dark atoms that can accomodate the experimentally observed excess of positrons in PAMELA and AMS-02 while being compatible with the constraints imposed on the gamma-ray flux from Fermi/LAT. This scenario assumes that the dominant component of dark matter is in the form of a bound state between a helium nucleus and a $-2$ particle and a small component is in the form of a WIMP-like dark atom compatible with direct searches in underground detectors. One of the constituents of this WIMP-like state is a $+2$ metastable particle with a mass of 1 TeV or slightly below that by decaying to $e^+e^+$, $\mu^+ \mu^+$ and $\tau^+ \tau^+$ produces the observed positron excess. These decays can naturally take place via GUT interactions. If it exists, such a metastable particle can be found in the next run of LHC. The model predicts also the ratio of leptons over baryons in the Universe to be close to -3. \\[.1cm] {\footnotesize \it Preprint: CP$^3$-Origins-2014-005 DNRF90 \& DIAS-2014-5.} | The possibility of dark matter being in the form of ``dark atoms'' has been studied extensively~\cite{Blin1,Blin2,shadow2,hodges,GH,BM,BDM,FV,MT,foot,SZ,arkani,kaplan,behbahani,kaplan2,cline,cline2,cyr,cyr2,dweinberg,cline3}. In this scenario new stable particles are bound by new dark forces (like mirror partners of ordinary particles bound by mirror electromagnetism~\cite{LeeYang,KOP,ZKrev,OkunRev,Paolo}). However, it turns out that even stable electrically charged particles can exist hidden in dark atoms, bound by ordinary Coulomb interactions (see \cite{Levels,Levels1,mpla,DMRev} and references therein). Stable particles with charge -1 (and corresponding antiparticles as tera-particles \cite{Glashow}) are excluded due to overproduction of anomalous isotopes. However, negatively doubly charged particles are not constrained by anomalous isotope searches as much as -1 charged particles~\cite{Fargion:2005xz}. There exist several types of particle models where heavy stable -2 charged species, $O^{--}$, are predicted: \begin{itemize} \item[(a)] AC-leptons, predicted as an extension of the 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 (WTC) \cite{KK,Sannino:2004qp,Hong:2004td,Dietrich:2005jn,Dietrich:2005wk,Gudnason:2006ug,Gudnason:2006yj}. \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 terrestrial matter. Therefore a cosmological scenario should provide a mechanism which suppresses anomalous helium. There are two possible mechanisms than can provide a suppression: \begin{itemize} \item[(i)] The abundance of anomalous helium in the Galaxy may be significant, but in terrestrial matter a recombination mechanism could suppress this abundance below experimental upper limits \cite{Khlopov:2006dk,FKS}. The existence of a new U(1) gauge symmetry, causing new Coulomb-like long range interactions between charged dark matter particles, is crucial for this mechanism. This leads inevitably to the existence of dark radiation in the form of hidden photons. \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,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 a charge asymmety of the $O^{--}$ which form the atom-like states with primordial helium \cite{mpla,I}. If new stable species belong to non-trivial representations of the SU(2) electroweak group, sphaleron transitions at high temperatures can provide the relation 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 formation in the Big Bang Nucleosynthesis (BBN), $^4He$ screens the $O^{--}$ charged particles in composite $(^4He^{++}O^{--})$ {\it $OHe$} ``atoms'' \cite{I}. In all the models of $OHe$, $O^{--}$ behaves either as a lepton or as a specific ``heavy quark cluster" with strongly suppressed hadronic interactions. Therefore $OHe$ interactions with matter are determined by the nuclear interactions 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 the $OHe$ Universe can explain many results of experimental searches for dark matter \cite{mpla}. Such a scenario is insensitive to the properties of $O^{--}$, since the main features of the $OHe$ dark atoms are determined by their nuclear interacting helium shell. In terrestrial matter such dark matter species are slowed down and cannot cause significant nuclear recoil in the underground detectors, making them elusive in direct WIMP search experiments (where detection is based on nuclear recoil) such as CDMS, XENON100 and LUX~\cite{CDMS,CDMS2,CDMS3,xenon,lux}. The positive results of DAMA and possibly CRESST and CoGeNT experiments \cite{DAMA,DAMA-review,Bernabei:2008yi,cresst,cogent} can find in this scenario a nontrivial explanation due to a low energy radiative capture of $OHe$ by intermediate mass nuclei~\cite{mpla,DMRev}. It has been also shown \cite{KK,KK2,unesco,iwara} that a two-component dark atom scenario is also possible. Along with the dominant $O^{--}$ abundance, a much smaller excess of positively doubly charged techniparticles can be created. These positively charged particles are hidden in WIMP-like atoms, being bound to $O^{--}$. In the framework of WTC such positively charged techniparticles can be metastable, with a dominant decay channel to a pair of positively charged leptons. In this paper we show that even a $10^{-6}$ fraction of such positively charged techniparticles with a mass of 1 TeV or less and a lifetime of~$10^{20} \s$, decaying to $e^+e^+$, $\mu^+ \mu^+$, and $\tau^+ \tau^+$ can explain the observed excess of cosmic ray positrons, being compatible with the observed gamma ray background. One should note that as it was shown in \cite{FKS,KK,I,KK2} (for a review see \cite{mpla,Khlopov:2006dk} and references therein) the case of -2 charged stable particles is significantly different from the case of stable or metastable particles with charge -1, avoiding severe constraints on charged particles from anomalous isotope searches and BBN due to their catalytic effects (see e.g. \cite{ATPTnew,CLLCMPnew,BBNPNPnew}). In the essence this difference comes from the fact that primordial He formed in BBN, captures -2 charged particles in neutral OHe states, while -1 charged particles are captured by He in +1 charged ions, which either (if stable) form anomalous isotopes of hydrogen, or (if long-lived, but metastable) catalyze processes of light element production and influence their abundance. Nuclear physics of OHe is in the course of development, but a qualitative analysis has shown \cite{unesco} that the OHe interactions with matter should not lead to overproduction of anomalous isotopes, while OHe catalytic effects in BBN can lead to primordial heavy element production, but not to overproduction of light elements. The paper is organized as follows: In section 2 we give a brief review of dark atoms made of stable charged techniparticles. In Section 3 we present the constraints and the predictions of the scenario with respect to the parameters of the Technicolor model we use as well as how the ratio of lepton over baryon number is deduced. In section 4 we show what GUT operators can implement the decay of the doubly charged particle to leptons. In section 5, we show how the scenario of decaying dark matter can be realized, and how it can explain the PAMELA and AMS-02 results while satisfying the Fermi/LAT constraints. We conclude in section 6. | Dark matter can potentially be in the form of neutral $OHe$ dark atoms made of stable heavy doubly charged particles and primordial He nuclei bound by ordinary Coulomb interactions. This scenario sheds new light on the nature of dark matter and offers a nontrivial solution for the puzzles of direct dark matter searches. It can be realized in the framework of Minimal Walking Technicolor, in which an exact relation between the dark matter density and baryon asymmetry can be naturally obtained predicting also the ratio of leptons over baryons in the Universe. In the context of this scenario a sparse component of WIMP-like dark atoms of charged techniparticles can also appear. Direct searches for WIMPs put severe constraints on the presence of this component. However we demonstrated in this paper that the existence of a metastable positively doubly charged techniparticle, forming this tiny subdominant WIMP-like dark atom component and satisfying the direct WIMP searches constraints, can play an important role in the indirect effects of dark matter. We found that decays of such positively charged constituents of WIMP-like dark atoms to the leptons $e^+e^+, \mu^+\mu^+,\tau^+ \tau^+$ can explain the observed excess of high energy cosmic ray positrons, while being compatible with the observed gamma ray background. These decays are naturally facilitated by GUT scale interactions. This scenario makes a prediction about the ratio of leptons over baryons in the Universe to be close to $-3$. The best fit of the data takes place for a mass of this doubly charged particle of 1 TeV or below making it accessible in the next run of LHC. {\bf Acknowledgements} C.K. is supported by the Danish National Research Foundation, Grant No. DNRF90. | 14 | 3 | 1403.1212 |
1403 | 1403.6604_arXiv.txt | {The large-scale magnetic field in the Sun varies with a period of approximately 22 years, although the amplitude of the cycle is subject to long-term modulation with recurrent phases of significantly reduced magnetic activity. It is believed that a hydromagnetic dynamo is responsible for producing this large-scale field, although this dynamo process is not well understood.} {Within the framework of mean-field dynamo theory, our aim is to investigate how competing mechanisms for poloidal field regeneration (namely a time delayed Babcock-Leighton surface $\alpha$-effect and an interface-type $\alpha$-effect), can lead to the modulation of magnetic activity in a deep-seated solar dynamo model.} {We solve the standard $\alpha\Omega$ dynamo equations in one spatial dimension, including source terms corresponding to both of the the competing $\alpha$-effects in the evolution equation for the poloidal field. This system is solved using two different methods. In addition to solving the one-dimensional partial differential equations directly, using numerical techniques, we also use a local approximation to reduce the governing equations to a set of coupled ordinary differential equations (ODEs), which are studied using a combination of analytical and numerical methods.} {In the ODE model, it is straightforward to find parameters such that a series of bifurcations can be identified as the time delay is increased, with the dynamo transitioning from periodic states to chaotic states via multiply periodic solutions. Similar transitions can be observed in the full model, with the chaotically modulated solutions exhibiting solar-like behaviour.} {Competing $\alpha$-effects could explain the observed modulation in the solar cycle.} | At the solar photosphere, bipolar active regions are formed when loops of magnetic flux rise to the surface from the base of the convection zone due to the action of magnetic buoyancy \citep[][]{PARKER3}. This implies that the properties of sunspot-bearing active regions can be used to deduce some of the features of the underlying large-scale magnetic field. It is well known \citep[see, for example,][]{STIX,CHARB,JONES} that zones of active region emergence follow a cyclic pattern with a period of approximately 11 years. At the beginning of each cycle, sunspots tend to be found at mid-latitudes, with zones of emergence drifting towards the equator as the cycle progresses. The underlying large-scale (predominantly azimuthal) magnetic field changes sign at the end of each cycle, giving a full magnetic period of approximately 22 years. However, the solar cycle is not strictly periodic. In particular, the peak amplitude (measured, for example, by the sunspot coverage) varies from one cycle to the next. Although this modulation does not usually disrupt the cycle, more extreme episodes of modulation have been recorded. For example, during a period known as the Maunder Minimum, very few sunspots were observed between approximately 1650 and 1720 \citep[][]{EDDY,RIBES}. However, sunspot records are not the only indicators of modulation. Due to the fact that the Sun's strong magnetic field protects the Earth from cosmic rays, the abundance of certain isotopes in the Earth's atmosphere is known to be anti-correlated with the solar cycle. Therefore, by analysing Beryllium-10 deposits in ice cores \citep[see, for example,][]{BE10} and Carbon-14 levels in tree rings \citep[see, for example,][]{C14} it is possible to deduce the history of the solar cycle. Such studies have indicated that cyclic activity did persist throughout the Maunder Minimum, but at a significantly reduced level \citep[][]{BEER}. Furthermore, it is clear that the Maunder Minimum is not exceptional -- the solar cycle has often been interrupted by recurrent ``Grand Minimum'' phases of significantly reduced magnetic activity. \par It is believed that the large-scale magnetic field in the solar interior is generated and maintained by a hydromagnetic dynamo. From a conceptual point of view, the large-scale field can usefully be decomposed into its toroidal (azimuthal) and poloidal (meridional) components -- a working dynamo requires mechanisms that allow the poloidal field to be regenerated from toroidal field and vice versa. It is widely accepted that differential rotation (usually referred to as the $\Omega$-effect in dynamo theory) is responsible for the generation of toroidal field from poloidal field. Surface observations indicate that equatorial regions rotate more rapidly than the poles, and helioseismological studies \citep[see, for example,][]{SCHOU} have shown that this rotation profile persists, approximately independently of radius, throughout most of the convection zone. At the base of the convection zone, a region of strong shear (the tachocline) couples the radiative zone, which rotates almost rigidly, to the differentially-rotating convective envelope. In most solar dynamo models, it is assumed that a significant fraction of the toroidal field is generated in the vicinity of the tachocline (where the $\Omega$-effect should be very efficient due to the presence of strong differential rotation). \par Although the $\Omega$-effect is well understood, the reverse process that generates poloidal field from toroidal field is a topic of some debate. In classical interface dynamo models \citep[see, for example,][]{PARKER,CHARBMAC} the poloidal field is regenerated at the base of the convection zone by the action of cyclonic convection upon toroidal magnetic field lines \citep[][]{PARKER2}. This process is usually referred to as the $\alpha$-effect. Strong toroidal fields will tend to inhibit (or quench) the operation of the $\alpha$-effect, so interface dynamo models are usually constructed in such a way that the $\alpha$-effect is restricted to the region just above the base of the convection zone, whilst the $\Omega$-effect operates just below the interface. The two layers are coupled by the effects of magnetic diffusion, as well as magnetic buoyancy and turbulent pumping \citep[see, for example,][]{TOBBRU}. Even with strong $\alpha$-quenching, it has been shown that an interface dynamo of this type can operate efficiently \citep[][]{CHARBMAC}. In Babcock-Leighton dynamo models \citep[][]{BL,LTON1}, the poloidal field is regenerated at the solar surface through the decay of active regions (which tend to emerge with a systematic tilt with respect to the azimuthal direction). This surface $\alpha$-effect can only contribute to the dynamo if there is some mechanism that is capable of transporting the resultant poloidal field to the tachocline. This could be achieved by diffusion or by pumping, but meridional flows also could play an important role in this respect. A polewards meridional flow is observed at the solar surface \citep[see, for example,][]{HATHAWAY} and, by mass conservation arguments, there must be a returning circulatory flow somewhere within the solar interior. A single-cell meridional circulation, with an equatorial flow at the base of the convection zone would couple the surface layers to the tachocline in an effective way, thus completing the dynamo loop. \par A complete model of the solar dynamo must be able to explain the observed modulation as well as the 22-year magnetic cycle. It has been shown that it is possible to induce modulation by introducing stochastic effects into Babcock-Leighton models \citep[][]{DIKCHARB,BUSHTOB}, as well as into models of interface type \citep[][]{OSSENDRIJVER2000}. However, fully deterministic models (with no random elements) can also produce modulated dynamo waves. \citet{WEISSCATJONES} and \citet{JWC} considered a simple system in which the dynamo was modelled using a set of coupled ordinary differential equations, which included the nonlinear interactions between the magnetic field and the flow. They found that it was possible to generate quasiperiodic and chaotically-modulated solutions in addition to standard periodic dynamo waves. More recent studies have shown that the full mean-field equations also exhibit significant modulation when dynamical nonlinearities are included in the governing equations \citep[see, for example,][]{TOBIAS96,BROOKE,BUSHBY}. An alternative approach was used by \cite{YOSHI} who demonstrated that modulation can arise if explicit time delays are built into the nonlinear terms in a simple system of model dynamo equations. A more sophisticated model was considered by \cite{JOUVE} who investigated the effects of magnetic buoyancy-induced time delays in the context of a two-dimensional Babcock-Leighton dynamo. By introducing time delays into the surface $\alpha$-effect term, they were able to demonstrate the existence of modulated cycles. They then went on to consider a simpler one-dimensional dynamo system in which the surface $\alpha$-effect term was represented by the inclusion of a time-delayed toroidal field (with a parameterised time delay that was dependent upon the magnetic field strength). They were able to demonstrate the existence of a sequence of bifurcations from periodic to chaotically modulated solutions as the time delay parameter was increased. \par The aim of this work is to investigate the competition between a deep-seated (interface) $\alpha$-effect and a surface $\alpha$-effect. Building on the approach described by \citet{JOUVE}, who did not include a deep-seated $\alpha$-effect, the influence of the surface $\alpha$-effect will be modelled using a time-delayed toroidal field. The use of a time delay is natural in this context: even if flux tubes rise rapidly to the surface, the time taken for the resultant poloidal field to be transported back to the tachocline will, in general, be non-negligible compared to the period of oscillation of the dynamo. Previous studies have investigated systems with competing $\alpha$-effects \citep[see, for example,][]{DIKPATI01,MASON,MANN}, but we believe that this is the first study to consider the effects of explicit time delays in a model of this type. The paper is structured as follows: In Section 2, we describe the full model and an idealised system of equations that can derived from it (based upon a local analysis). This is followed in Section 3 by an analysis of the stability of the idealised model and then in Section 4 by the corresponding numerical results. In Section 5, we describe some numerical calculations which demonstrate the existence of modulated solutions in the full one-dimensional model. Finally, in Section 6, we present our conclusions and discuss the relevance of our results to the solar dynamo. | In this paper, we have investigated the properties of an illustrative mean-field dynamo model which includes two competing $\alpha$-effects. The first of these is the standard deep-seated $\alpha$-effect, the second is due to a surface $\alpha$-effect (of Babcock-Leighton type). Following the approach described by \citet{JOUVE}, who did not consider competing $\alpha$-effects, the contribution from the surface $\alpha$-effect was modelled by assuming that it depends upon a time-delayed toroidal field (with a constant parameterised time delay $\tau$). Two different approaches were applied to this model. Initially, a local approximation was made to reduce the governing equations to a system of coupled ordinary differential equations. A linearised version of these equations was used to determine the dependence of the critical dynamo number upon $S$ (the magnitude of the surface $\alpha$-effect) and $\tau$. Generally, the larger the magnitude of $S$, the easier it becomes to excite the dynamo. However, there are some regions of parameter space in which the two competing $\alpha$-effects appear to impede each other, thus inhibiting the dynamo. Moving beyond linear theory, it was found that there are significant regions of parameter space in which the periodic solution becomes unstable with increasing $\tau$, leading to quasi-periodicity. This was verified numerically, where further increases in the time delay were shown to produce chaotically modulated states with phases of significantly reduced activity. The full PDE model was then investigated. Although modulation was found, this occurs in a different parameter regime to that predicted by the ODE model. This discrepancy could be model specific, although we expected to see some differences between the two models due to the fact that significant simplifications were made when deriving the set of coupled ODEs. Nevertheless, it was possible to find chaotically modulated solutions in the PDE model, and these solutions exhibit certain features that are (at least qualitatively) ``solar-like''. \par There are many possible areas of future work. In particular, more could be done to explore the robustness of the PDE model to variations in the boundary conditions and the non-linear quenching mechanisms. As has already been mentioned, preliminary calculations suggest that the adoption of different nonlinearities may make a very significant difference to the behaviour of the model. It may also be possible to improve the existing model by refining the way in which the time delay is implemented -- the current approach is simple and effective, but is derived by truncating a Taylor series expansion at lowest order. Retaining higher order terms may make a difference to the behaviour of the model. Moving beyond the one-dimensional Cartesian system, it would be natural to explore a two-dimensional version of this model in (axisymmetric) spherical geometry. This would open up the possibility of including a more realistic flow geometry (in both the meridional and azimuthal directions) as well as spatially dependent mean-field coefficients. Although this would still be within the framework of mean-field theory, a more realistic model would enable more detailed comparisons to be made between our results and the solar dynamo. | 14 | 3 | 1403.6604 |
1403 | 1403.4601_arXiv.txt | Understanding how distinct, near-spherical gas-free clusters of very young, massive stars shape out of vast, complex clouds of molecular hydrogen is one of the biggest challenges in astrophysics. A popular thought dictates that a single gas cloud fragments into many new-born stars which, in turn, energize and rapidly expel the residual gas to form a gas-free cluster. This study demonstrates that the above classical paradigm remarkably reproduces the well-observed central, young cluster (HD 97950) of the Galactic NGC 3603 star-forming region, in particular, its shape, internal motion and the mass distribution of stars, naturally and consistently follow from a single model calculation. Remarkably, the same parameters (star formation efficiency, gas expulsion time scale and delay) reproduce HD 97950 as were found to reproduce the Orion Nebula Cluster, Pleiades and R136. The present results thereby provide intriguing evidences of formation of star clusters through single-starburst events followed by significant residual gas expulsion. | Very young, massive star clusters (hereafter VYMCs), which are compact associations of stars of mass $\gtrsim10^4\Ms$ and up to a few Myr old\footnote{Such clusters are usually referred to as ``starburst clusters'' but here we prefer a more generic and origin-independent designation of such systems. They constitute a sub-category of Young Massive Star Clusters \citep[]{pz2010}.}, are typically found in locations of high star-formation activity (or ``starburst regions'') in our Milky Way and external galaxies. The classical notion implies that such systems form through a single episode of starburst. In this scenario, efficient cooling processes within a dense parent gas-cloud lead to the formation of a number of proto-stellar cores. Such an infant star cluster consists of massive main sequence (MS) and lower-mass pre-main-sequence (PMS) stars embedded within the residual gas. The radiation and winds from the MS and PMS stars inject energy into this gas until the latter becomes unbound and escapes the system. The associated rapid dilution of the potential well causes the system to expand which loses a fraction of its stars depending on its initial mass and concentration \citep[]{adm2000,pketl2001,bokr2002,bk2007,sb2013,pfkz2013}. The remaining gas-free system may eventually attain a state of virial equilibrium. Such a ``monolithic'' formation scenario has successfully explained the structure and kinematics of Galactic and extra-galactic VYMCs like the Orion Nebula Cluster (ONC) and R136 and intermediate-aged clusters like the Pleiades and Hyades \citep[]{pketl2001,sb2013}. In this study we incorporate the same properties and parameters (see below) of the gas expulsion process as those in the earlier studies \citep[]{pketl2001,sb2013} in our realistically detailed model calculations of the time-evolution of star clusters. We find good and consistent agreement between these model computations and detailed measurements of the central young cluster (HD 97950; hereafter HD97950) of the Galactic starburst region NGC 3603 which are based on infrared and optical observations with the ESO/VLT (Very Large Telescope) and the Hubble Space Telescope (HST), respectively \citep[]{hara2008,roch2010,pang2013}. In particular, we closely reproduce its observed (a) mass density profile \citep[]{hara2008}, (b) central velocity dispersion \citep[]{roch2010,pang2013}, (c) present-day stellar mass function (PDMF; \citealt[]{pang2013}), and as well (d) the spatial distribution of stars from the crude data \citep[]{pang2013}, simultaneously from one individual model calculation. This not only strongly points to a single-episode formation of HD97950 as envisaged previously \citep[]{stol2004,stol2006} but it also affirms the monolithic formation channel for VYMCs in general and suggests that such systems evolve according to universal principles. | \label{discuss} This study demonstrates that it is possible to reproduce the detailed observed features of the central cluster HD97950 of the NGC 3603 starburst region, within its age and mass constraints, from a monolithic formation standpoint and adopting the reasonable and widely used values of the parameters quantifying the gas dispersal phase \citep[]{pketl2001,sb2013}. It should be noted that the above results imply a \emph{natural} match of the computed models with the observed data \emph{without any scaling or fitting}. The above calculations have been performed without applying an external Galactic potential. As verified by comparing with calculations with a Solar-neighborhood-like field, the effect of the external field is only minor. This can be expected as the highly compact initial system remains well within its Roche lobe. Notably, the initial compact size ($r_h(0)\approx0.25$ pc) of the computed models is consistent with the observed widths of the compact gas-filaments that host the proto-clusters \citep[]{andr2011} and also with the mass-radius relation of embedded clusters \citep[]{mk2012}. In this work, we assume the ``classical'' SFE of $\epsilon\approx33$\% over the proto-cluster scale which is supported by both theoretical \citep[]{mnm2012} and observational \citep[]{lnl2003} studies. We do not take into account the variations of SFE over the proto-star/subcluster scale as suggested in recent theoretical studies \citep[]{giri2012,mok2012}. Such studies, based on hydrodynamic simulations, indicate much higher local SFE which would influence the gas expulsion. However, inclusion of proto-stellar outflows have been shown to substantially suppress the local SFE to $\lesssim 50$\% and is consistent with $\approx33$\% SFE over both local (stellar) and global (proto-cluster) scale \citep[]{PriceBate2010,mnm2012}. The analytic potential-lowering method applied here is widely used to this date for a good reason. The self-consistent treatment of gas removal requires three-dimensional radiation-hydrodynamic calculations which is formidable to date for the mass scale involved in this study. High-resolution (reaching the ``opacity-limit'') smoothed-particle-hydrodynamics (SPH) computations have been done so far forming stars in gas-spheres of only up to $\approx500\Ms$ \citep[]{kl1998,bate2004,bate2009,giri2011,giri2012,bate2012} but without any feedback and hence self-regulation mechanism. Radiation-magnetohydrodynamic (MHD) calculations incorporating feedback to the star-forming gas has also been carried out from proto-stellar scales \citep[]{mnm2012,bate2013} upto $\approx50\Ms$ gas spheres \citep[]{PriceBate2010}. While the latter studies provide insights into the self-regulation mechanisms in the star formation process, the processes that lead to the ultimate dispersal of the gas is still unclear. Therefore, the analytic potential lowering is the best that can be done now for the mass scales involved here. It should be noted, however, that the above monolithic cluster formation scenario has nevertheless successfully reproduced ONC, Pleiades and Hyades clusters previously \citep[]{pketl2001} and as well the R136 \citep[]{sb2013} and the HD97950 cluster as in here. This indicates that although deprived of the details of star-formation and radiation-hydrodynamic processes due to technological bottleneck, the adapted simplified gas-expulsion model and the choices of the parameters might still be an appropriate description of the overall gas expulsion process. Indeed, in a comparative study, \citet[]{gb2001} found that analytic gas-potential reduction has essentially the same effect on the cluster evolution as in an SPH computation (for a given SFE) where the gas is removed via shock heating. The present work thereby provides another pillar to the to-date widely used classical gas expulsion scenario. As such, the current state-of-the-art in theoretical research on star formation does not provide a definitive answer to the gas dispersal and cluster formation processes and their interrelations. Given this, \emph{the present work provides a plausible set of physical conditions that compellingly evolve to reproduce the detailed observed properties of HD97950, subject to its photometric and age constraints.} They comprise a \emph{solution} since the same computed model reproduces multiple observed properties of the cluster (it’s density profiles, velocity dispersion and stellar mass function). Remarkably, the same conditions ($\epsilon\approx0.33$, $v_g\approx10{\rm~km}{~s}^{-1}$, $\tau_g=r_h(0)/v_g$, $\tau_d\approx0.6$ Myr) that reproduced the known properties of the ONC and Pleiades \citep[]{pketl2001} also works for the much more massive R136 and HD97950 clusters (see Table~\ref{tab1}). The present work thereby serves as an intriguing evidence that these clusters might have formed monolithically with the chosen parameters being an appropriate description of the overall gas expulsion process. Several authors alternatively propose VYMCs to be formed via hierarchical merging of sub-clusters/structures \citep[]{fuji2012,sm2013,longm2014}. In particular, \citet[]{fuji2012} have demonstrated agreements with observed radial cumulative distributions of massive stars and merger products in HD97950. This work, although interesting, cannot be taken as a vivid counterexample because of notable drawbacks. Here, the representative age of HD97950 has been taken to be 2.5 Myr which is far too old compared to the best-fitting age of $\approx 1$ Myr with only a small age spread, as obtained from the cluster's PMS-MS CMD \citep[]{stol2004,pang2013}. In 1 Myr, the substructures hardly merge to form a single cluster (\cf, Fig.~2 of \citealt[]{fuji2012}). Furthermore, no attempt has been made to compare directly measurable quantities like the structure and kinematics of the merged cluster with observations. Hence, it still remains to be seen how well a hierarchical merging scenario can reproduce the structure and kinematics of HD97950 given its very young age. Based on the current theoretical literature, it is as such impossible to rule out a single-starburst scenario of the formation of smooth-structured VYMCs in preference to a hierarchical one, on fundamental grounds. \begin{figure} \includegraphics[width=9.2cm, angle=0]{fig5a.eps} \includegraphics[width=9.2cm, angle=0]{fig5b.eps} \includegraphics[width=9.2cm, angle=0]{fig5c.eps} \caption{The stellar mass function (of primaries) within the central $60''$($\approx 1.8$ pc) of the computed model with primordial binaries (model HD97950b) at $t\approx0.0$, 0.4 and 1.4 Myr (top, middle and bottom panels respectively). Beginning with the global canonical value of $\Gamma=-1.35$ ($\alpha_2=2.35$), the MF slope develops a break with time as a result of the switching to much tighter initial orbital periods for massive binaries (from $>5\Ms$; see Sec.~\ref{bincomp}). Upwards the break (at $\log(m/\Ms)\approx0.7$), the MF slope becomes increasingly shallow while it remains close to canonical for lower $m$. The corresponding slope at $t\approx1.4$ Myr (bottom), \viz, $\Gamma=-0.92\pm0.13$ (solid line), agrees closely with the corresponding measured one, \viz, $\Gamma=-0.88\pm0.15$ (dashed line; \citealt[]{pang2013}).} \label{fig:mfnow} \end{figure} \begin{figure} \includegraphics[width=9.2cm, height=6.3cm, angle=0]{fig6a.eps} \includegraphics[width=9.2cm, height=6.3cm, angle=0]{fig6b.eps} \caption{Radial variation of 1-dimensional velocity dispersion, $\sigma_{\rm 1d}$, for the computed model with primordial binaries (in presence of a tidal field) at $t=1.4$ Myr for stellar mass ranges $1.7\Ms\leq m\leq9.0\Ms$ (top panel) and $1.0\Ms\leq m\leq100.0\Ms$ (bottom panel). The overall increasing trend of $\sigma_{\rm 1d}$ with $R$ in the outer regions ($R\gtrsim40''$ in this case) indicates a recent gas expulsion. This can be tested in future by more accurate proper motion measurements (\eg, by \emph{Gaia}) in the outer regions of the HD97950 cluster. The tangential velocities of selected stars in \citet[]{pang2013} (for $R\lesssim60''$) do show an increasing trend with $R$. \label{fig:vprof} } \end{figure} Notably, the construction of spatial density profiles involves accumulation over finite-sized annuli and the heights over the profile depends on the choice of these annuli. However, note that in Fig.~\ref{fig:Mdensprof}, the angular annuli are chosen nearly the same as those used by \citet[]{hara2008} to obtain their observed profile. Also, in Fig.~\ref{fig:Ndensprof}, the same angular annuli are used to construct both of the number density profiles. Several earlier authors have considered HD97950 to be in complete virial equilibrium and estimated its dynamical mass to be $\approx20000\Ms$ from its observed velocity dispersion \citep[]{roch2010,pang2013}. We have computed initially $M_{ecl}(0)\approx20000\Ms$ models (with $\epsilon\approx33$\%) to find that such massive systems yield either too centrally-dense or too expanded profiles, \ie, no well-matching initial conditions could be found. The $M_{ecl}(0)\approx10000\Ms$ models presented here, that do reproduce the observations, are not in dynamical equilibrium at the epoch of matching ($t\approx1.4$ Myr), but are still expanding (the innermost regions are close to re-collapsing). A prediction from this study, therefore, is that the stars in the outer parts of HD97950 should have outgoing radial proper motions which can be tested by future observations. Fig.~\ref{fig:vprof} shows radial profiles of $\sigma_{\rm 1d}$ at $t=1.4$ Myr for $1.7\Ms\leq m\leq9.0\Ms$ and $1.0\Ms\leq m\leq100.0\Ms$ where $\sigma_{\rm 1d}$ tends to increase for $R\gtrsim40''(\approx1.2{\rm~pc})$. The latter feature can be attributed to the recent gas expulsion from the system resulting in its outer parts still expanding. Such a trend, which becomes more pronounced the closer the epoch of observation is to the gas expulsion, can be tested by future, more accurate determinations of stellar proper motions in the outer regions of HD97950, \eg, by \emph{Gaia}. Notably, the measured tangential velocities of selected stars in \citet[]{pang2013} do show an increasing trend with radial distance in the outer region (measured up to $R\approx60''$). The profiles in Fig.~\ref{fig:vprof} are obtained from a computation with primordial binaries in presence of a solar-neighborhood-like external Galactic potential. The above values of the parameters $\epsilon$, $\tau_d$ and $\tau_g$, quantifying the overall gas expulsion, fair well with several young clusters (see above) and hence can be considered well-representative. However, it would be of interest to consider how our conclusions can be affected by moderate alterations of these parameters, representing possible cluster-to-cluster variations. As shown in \citet[]{sb2013}, the effect of varying the delay time, $\tau_d$, is simply a time-shift in the cluster's violent expansion without substantially affecting the cluster's evolution for $t>\tau_d$. Hence, a delay of $\tau_d+\delta\tau_d$ in our computed models would reproduce HD97950 at $\approx t+\delta\tau_d$. As such, one can allow for $\delta\tau_d\approx\pm0.1$ Myr in the models computed here to obtain similarly good matchings. For $\epsilon<50$\%, as in the present case, the bound fraction of stars that re-virializes to form the remnant cluster \citep[]{pketl2001,sb2013} depends sensitively on $\tau_g$ \citep[]{bk2007}. However, the computed clusters here, which are still mostly expanding at the HD97950's age of $t\approx1$ Myr, are far from re-gaining dynamical equilibrium (see above; it takes $\approx3$ Myr to re-virialize an HD97950-like cluster as shown in \citealt[]{sb2013}) but still retains a compact spatial distribution. Therefore, at $t\approx1$ Myr, most of the cluster's primordial stellar population contributes to the computed radial profiles and the other quantities. Hence, the model-observation agreements, as presented here, would still mostly persist with a moderate change of $\tau_g$ ($v_g$) as long as the gas-removal is explosive (see Sec.~\ref{gdisp}). The computed values of $\sigma_{\rm 1d}$ can be expected to increase moderately with decreasing $\tau_g$ but this trend is likely to be suppressed by the substantial fluctuations in the values of $\sigma_{\rm 1d}$ (see Sec.~\ref{res}, Fig.~\ref{fig:v1d_pang}). With a substantially higher $\epsilon$, one would expect a $M_{ecl}(0)\approx20000\Ms$ cluster to be even more overdense/oversized and a matching model would be of $M_{ecl}(0)\lesssim10000\Ms$, \ie, below HD97950's photometric mass limit. The reproducibility of the NGC3603's young cluster as a function of the parameters $\epsilon$, $\tau_d$ and $\tau_g$ will be investigated in detail in a future project. An immediate improvement over this work would be to apply colour filters (\eg, V and I) to the stellar luminosities to calculate the density profiles and the mass functions for comparisons with observations even more accurately. | 14 | 3 | 1403.4601 |
1403 | 1403.4437_arXiv.txt | % Diversity of the X-ray observations of dwarf nova are still not fully understood. I review the X-ray spectral characteristics of dwarf novae during the quiescence in general explained by cooling flow models and the outburst spectra that show hard X-ray emission dominantly with few sources that reveal soft X-ray/EUV blackbody emission. The nature of aperiodic time variability of brightness of dwarf novae shows band limited noise, which can be adequately described in the framework of the model of propagating fluctuations. The frequency of the break (1-6 mHz) indicates inner disk truncation of the optically thick disk with a range of radii (3.0-10.0)$\times$10$^{9}$ cm. The $RXTE$ and optical (RTT150) data of SS Cyg in outburst and quiescence reveal that the inner disk radius moves towards the white dwarf and receeds as the outburst declines to quiescence. A preliminary analysis of SU UMa indicates a similar behaviour. In addition, I find that the outburst spectra of WZ Sge shows two component spectrum of only hard X-ray emission, one of which may be fitted with a power law suggesting thermal Comptonization occuring in the system. Cross-correlations between the simultaneous UV and X-ray light curves ($XMM-Newton$) of five DNe in quiescence show time lags in the X-rays of 96-181 sec consistent with travel time of matter from a truncated inner disk to the white dwarf surface. All this suggests that dwarf novae and other plausible nonmagnetic systems have truncated accretion disks indicating that the disks may be partially evaporated and the accretion may occur through hot (coronal) flows in the disk. | Dwarf novae (DNe) are a class of cataclysmic variables (CVs) which are interacting compact binaries in which a white dwarf (WD, the primary star) accretes matter and angular momentum from a main (or post-main) sequence star (the secondary) filling its Roche-lobe. The matter is transferred by means of an {\it accretion} disk that is assumed to reach all the way to the WD surface. Ongoing accretion at a low rate (quiescence) is interrupted every few weeks to months or sometimes with longer durations by intense accretion (outburst) of days to weeks where $\dot{\rm M}$ increases (see Warner 1995 for a review). The material in the inner disk of nonmagnetic CVs initially moving with Keplerian velocity dissipates its kinetic energy in order to accrete onto the slowly rotating WD creating a boundary layer (BL) (see Mauche 1997, Kuulkers et al. 2006 for an overview). Standard accretion disk theory predicts half of the accretion luminosity to originate from the disk in the optical and ultraviolet (UV) wavelengths and the other half to emerge from the boundary layer as X-ray and/or extreme UV (EUV)/soft X-ray emission which may be summerized as L$_{BL}$$\sim$L$_{disk}$=GM$_{WD}$$\dot M_{acc}$/2R$_{WD}$=L$_{acc}$/2 (Lynden-Bell $\&$ Pringle 1974, Godon et al. 1995). During low-mass accretion rates, $\dot M_{acc}$$<$10$^{-(9-9.5)}$M$_{\odot}$, the boundary layer is optically thin (Narayan $\&$ Popham 1993, Popham 1999) emitting mostly in the hard X-rays (kT$\sim$10$^{(7.5-8.5)}$ K). For higher accretion rates , $\dot M_{acc}$$\ge$10$^{-(9-9.5)}$M$_{\odot}$, the boundary layer is expected to be opticallly thick (Popham $\&$ Narayan 1995) emitting in the soft X-rays or EUV (kT$\sim$10$^{(5-5.6)}$ K). The transition between an optically thin and an optically thick boundary layer, also depends on the mass of the white dwarf (also rotation) and on the alpha viscosity parameter. | Studies of DNe broad-band noise characteristics in the X-rays (see also Balman \& Revnivtsev 2012) indicate that DNe have truncated accretion disks in quiescence detected in at least 8 systems in a range $\sim$(3.0-10.0)$\times$10$^9$\ cm including errors. The Magnetic CVs (MCVs) show rather smaller truncation radii (0.9-2.0)$\times$10$^9$\ cm (Revnivtsev et al. 2010, 2011). This can also explain the UV and X-ray delays in the outburst stage and the accretion may occur through hot (coronal) flows in the disk. Note that extended emission and winds are detected from DN in the outburst stage which may be an indication of the coronae/hot flows in these systems (e.g., Mauche 2004). Time delays detected in a range of 96-181 sec, are also consistent with matter propogation timescales onto the WD in a truncated nonmagnetic CV disk in quiescence. Balman \& Revnivtsev (2012) approximate an $\alpha$ of 0.1-0.3 for the inner regions of the DNe accretion disks in quiescence. In addition, the outburst spectra of WZ Sge shows two component spectrum of only hard X-ray emission, one of which may be fitted with a power law suggesting thermal Comptonization of the optically thick disk photons or scattering from an existing wind during the outburst. The spectral evolution and disapearance of the power law component after outburst may support this issue and that the accretion disk goes back to its quiescent truncated structure and Comptonization stops. A general picture of the accretion flow around a WD in quiescence thus might be somewhat similar to that of the black hole/neutron star accretors with an optically thick colder outer accretion disk and an optically thin hot flow in the inner regions where the truncation occurs (see Esin et al. 1997, Done et al. 2007). The appearance of a hot flow (e.g., ADAF-like) in the innermost regions of the accretion disk will differ from that of ordinary rotating Keplerian disk because it is no longer fully supported by rotation, but might have a significant radial velocity component with sub-Keplerian speeds. It is important to monitor DNe in the X-rays to measure the variability in the light curves in time together with the variations of the possible disk truncation and formation of plausible coronal (optically thin) and/or ADAF-like flows/regions on the disk in quiescence and outburst. \thanks SB thanks to M. Revnivtsev for his valuable collaboration on the timing analysis of DN. | 14 | 3 | 1403.4437 |
1403 | 1403.1354_arXiv.txt | We extend our approach of modeling spectral energy distribution (SED) and lightcurves of blazars to include external Compton (EC) emission due to inverse Compton scattering of an external anisotropic target radiation field. We describe the time-dependent impact of such seed photon fields on the evolution of multifrequency emission and spectral variability of blazars using a multi-zone time-dependent leptonic jet model, with radiation feedback, in the internal shock model scenario. We calculate accurate EC-scattered high-energy spectra produced by relativistic electrons throughout the Thomson and Klein-Nishina regimes. We explore the effects of varying the contribution of (1) a thermal Shakura-Sunyaev accretion disk, (2) a spherically symmetric shell of broad-line clouds, the broad line region (BLR), and (3) a hot infrared emitting dusty torus (DT), on the resultant seed photon fields. We let the system evolve to beyond the BLR and within the DT and study the manifestation of the varying target photon fields on the simulated SED and lightcurves of a typical blazar. The calculations of broadband spectra include effects of $\gamma-\gamma$ absorption as $\gamma$-rays propagate through the photon pool present inside the jet due to synchrotron and inverse Compton processes, but neglect $\gamma-\gamma$ absorption by the BLR and DT photon fields outside the jet. Thus, our account of $\gamma-\gamma$ absorption is a lower limit to this effect. Here, we focus on studying the impact of parameters relevant for EC processes on high-energy (HE) emission of blazars. | Introduction} Blazars are known for their highly variable broadband emission. They are characterized by a doubly humped spectral energy distribution (SED), attributed to non-thermal emission, and spectral variability. The SED and variability patterns can be used as key observational features to place constraints on the nature of the particle population, acceleration of particles, and the environment around the jet that is responsible for the observed emission. Conversely, incorporating the nature of the particle population and the jet environment, as accurately as possible, in modeling such observational features can enable us to reach a better agreement between theoretical and observational results. Thus, exploring the environment of a blazar jet in an anisotropic and time-dependent manner is important for connecting the pieces together and putting tighter constraints on the origin of $\gamma$-ray emission. Blazars, a combination of BL Lacertae (BL Lac) objects and flat spectrum radio-loud quasars (FSRQs), are divided into various subclasses depending on the location of the peak of the low-energy (synchrotron) SED component. The synchrotron peak lies in the infrared regime, with $\nu_{\rm s} \leq 10^{14}$ Hz, in low-synchrotron-peaked (LSP) blazars comprising FSRQs and low-frequency peaked BL Lac objects (LBLs). In the case of intermediate-synchrotron-peaked (ISP) blazars, consisting of LBLs and intermediate-frequency peaked BL Lacs (IBLs), the synchrotron peak lies in the optical - near-UV region with $10^{14} < \nu_{\rm s} \leq 10^{15}$ Hz. The synchrotron component of high-synchrotron-peaked (HSP) blazars, which include essentially all high-frequency-peaked BL Lac objects (HBLs), peaks in the X-rays at $\nu_{\rm s} > 10^{15}$ Hz \citep[]{ab2010, bm2012}. The high-energy (HE) component of blazars can be a result of inverse Compton (IC) scattering of synchrotron photons internal to the jet resulting in synchrotron self-Compton (SSC) emission \citep{bm1996}. It could also be due to upscattering of accretion-disk photons \citep[]{ds1993}, and/or photons initially from the accretion disk being scattered by the broad-line region (BLR) \citep[]{sbr1994, dss1997}, and/or seed photons from a surrounding dusty torus (DT) \citep[]{ka1999, bl2000}. In the case of HBLs, the HE component is usually well reproduced with a synchrotron/SSC leptonic jet model \citep[e.g.,][]{fdb2008, aj2012}, whereas an additional external Compton (EC) component is almost always required to fit the high-energy spectra of FSRQs, LBLs, and IBLs \citep[e.g.,][]{cg1999, co2010}. Detailed numerical calculations for Compton scattering processes have been carried out for many specific models of blazar jet emission that involve their environment. \citet[]{ds1993, ds2002} have calculated Compton scattering of target photons in the Thomson regime from an optically thick and geometrically thin, thermal accretion disk based on the model of \cite{ss1973}. Quasi-isotropic seed photon fields due to BLR or DT have also been considered to obtain Compton-scattered high-energy spectra in the Thomson limit by several authors \citep[]{sbr1994, dss1997, bl2000}. On the other hand, extensive calculations involving anisotropic accretion-disk and BLR seed photon fields have been considered as well \citep[]{bms1997, bb2000, br2004, kt2005}. Anisotropic radiation fields of the disk, the BLR, and the DT have been studied previously by \cite{dp2003}, but primarily in the context of $\gamma-\gamma$ interaction of these photons with the GeV and TeV photons produced in the jet. Anisotropic treatment of BLR and DT photons, focussing on jet emission and rapid non-thermal flares, was carried out by \cite{sm2005}. These authors studied parameters describing the properties of BLR and DT that govern the interplay between the dominance of SSC and EC emission and their subsequent impact on SEDs, as well as relative time delays between light curves at different frequencies. For the purposes of their study, they used an integrated intensity - instead of considering line and continuum intensities separately - of the incident emission from the BLR. The emitting plasma was assumed to be located at parsec scales and the evolution of HE emission at sub-pc distances was ignored. Recently, anisotropic treatment of disk and BLR target radiation fields has been considered by \cite{de2009}. The authors have calculated accurate $\gamma$-ray spectra due to inverse Comptonization of such seed photon fields throughout the Thomson and Klein-Nishina (KN) regimes to model FSRQ blazars, although in a one-zone scenario. Also, these authors evolve the system to only sub-pc distances along the jet axis, limiting themselves to locations within the BLR. In addition to this, one-zone leptonic jet models were recently shown \citep{brm2009} to have severe limitations in attempts to reproduce very high energy (VHE) flares, such as that of 3C~279 detected in 2006 \citep{al2008}. In a more recent approach, \cite{ma2013} has considered an anisotropic seed photon field of the DT to calculate the resultant EC component of HE emission from blazar jets, in a turbulent extreme multi-zone scenario. While the $\gamma$-ray spectra are calculated throughout the Thomson and KN regimes, the energy loss rates are limited to only the Thomson regime. For the problem that work addresses, the system is located \textit{beyond} the BLR, at parsec-scale distances from the central engine. Here, we extend the previous approach of \citet[][hereafter Paper 1]{jb2011}, which calculated the synchrotron and SSC emission from blazar jets, to address some of the limitations of the models mentioned above. We use a fully time-dependent, 1-D multi-zone with radiation feedback, leptonic jet model in the internal shock scenario, shortened to the MUlti ZOne Radiation Feedback, MUZORF, model. We evolve the system from sub-pc to pc scale distances along the jet axis. We consider anisotropic target radiation fields to calculate the HE spectra resulting from EC scattering processes. The entire spectrum is calculated throughout the Thomson and KN regimes, thereby making it applicable to all classes of blazars. We include the attenuation of jet $\gamma$-rays through $\gamma-\gamma$ absorption (described in Paper 1) due to the presence of target radiation fields inside the jet, in a self-consistent manner. The generalized approach of our model lets us account for the constantly changing contribution of each of the seed photon field sources in producing HE emission in a self-consistent and time-dependent manner. This is especially relevant for understanding the origin of $\gamma$-ray emission from blazar jets. In a number of previous analyses, the region within the BLR has been considered the most favorable location for $\gamma$-ray emission, with a range limited to between 0.01 and 0.3 pc \cite[]{ds1994, bl1995, gm1996}. The reason behind this is the short intra-day variability timescales observed in some $\gamma$-ray flares, which indicated on the basis of light crossing timescales that the emission region is small and hence not be too far away from the central engine \cite[]{gm1996, gt2009}. At the same time, the emission region cannot be too close to the central engine without violating constraints placed by the $\gamma-\gamma$ absorption process \cite[]{gm1996, lb2006}. As a result, an emission region location closer to the BLR was considered the most favorable position due to the strong dependence of the scattered flux on the level of boosting and the energy of incoming photons \citep{sbr1994}. Contrary to the above scenario, recent observations have shown coincidences of $\gamma$-ray outbursts with radio events on pc scales \cite[e.g.,][]{lt2012, jo2013}. This seems to suggest a cospatial origin of radio and $\gamma$-ray events located at such distances. As a result, some authors conclude that the $\gamma$-ray emitting region could also lie outside of the BLR \cite[]{sm2005, lvt2005}. Thus, in order to understand the origin of $\gamma$-ray emission, it is important to let the system being modeled evolve to beyond the BLR and into the DT, and to include its contribution to the production of $\gamma$-ray emission. Here, we focus our attention toward understanding the dependence of $\gamma$-ray emission on the combination of various intrinsic physical parameters. We explore this aspect by including various components of seed photon fields in order to obtain a complete picture of their contribution in producing $\gamma$-ray emission and understand their effects on the dynamic evolution of SEDs and spectral variability patterns. In \S\ref{method}, we describe our EC framework of including anisotropic seed photon fields from the accretion-disk, the BLR, and the DT. We lay out the expressions used to calculate accurate Compton-scattered $\gamma$-ray spectra resulting from the seed photon fields and the corresponding electron energy loss rates throughout the Thomson and KN regimes. In \S\ref{paramstud}, we describe our baseline model, its simulated results, and the relevant physical input parameters that we use in the study. In \S\ref{outcome}, we present our results of the parameter study and discuss the effects of varying the input parameters on the simulated SED and lightcurves. We discuss and summarize our findings in \S\ref{disco}. Throughout this paper, we refer to $\alpha$ as the energy spectral index such that flux density, $F_{\nu}, \propto \nu^{- \alpha}$; the unprimed quantities refer to the rest frame of the AGN (lab frame), primed quantities to the comoving frame of the emitting plasma, and starred quantities to the observer's frame; the dimensionless photon energy is denoted by $\epsilon = \frac{h\nu}{m_{\rm e}c^{2}}$. Appendix \ref{loseqns} delineates the details of line-of-sight calculations for the BLR line and diffuse continuum emission used in obtaining the intensity of incoming BLR photons. | 14 | 3 | 1403.1354 |
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1403 | 1403.2741_arXiv.txt | A major impediment to understanding star formation in massive star forming regions (MSFRs) is the absence of a reliable stellar chronometer to unravel their complex star formation histories. We present a new estimation of stellar ages using a new method that employs near-infrared (NIR) and X-ray photometry, $Age_{JX}$. Stellar masses are derived from X-ray luminosities using the $L_X - M$ relation from the Taurus cloud. $J$-band luminosities are compared to mass-dependent pre-main-sequence evolutionary models to estimate ages. $Age_{JX}$ is sensitive to a wide range of evolutionary stages, from disk-bearing stars embedded in a cloud to widely dispersed older pre-main sequence stars. The MYStIX (Massive Young Star-Forming Complex Study in Infrared and X-ray) project characterizes 20 OB-dominated MSFRs using X-ray, mid-infrared, and NIR catalogs. The $Age_{JX}$ method has been applied to 5525 out of 31,784 MYStIX Probable Complex Members. We provide a homogeneous set of median ages for over a hundred subclusters in 15 MSFRs; median subcluster ages range between $0.5$~Myr and $5$~Myr. The important science result is the discovery of age gradients across MYStIX regions. The wide MSFR age distribution appears as spatially segregated structures with different ages. The $Age_{JX}$ ages are youngest in obscured locations in molecular clouds, intermediate in revealed stellar clusters, and oldest in distributed populations. The NIR color index $J-H$, a surrogate measure of extinction, can serve as an approximate age predictor for young embedded clusters. | \label{introduction_section} Infrared (IR) and millimeter astronomy, together with astrophysical theory, have made great progress in understanding the formation of single stars in small molecular cloud cores on $\leq 0.1$~pc scales. However, most stars do not form in such simple environments, arising rather in turbulent giant molecular clouds where different modes of star formation are seen: rich OB-dominated stellar clusters, smaller clusters in dense IR Dark Clouds and filaments, and stellar groups on the peripheries of expanding H\,{\sc{ii}} regions. It is not clearly known how cloud cores and filaments form, whether clusters form during a single free-fall time or over millions of years, or whether triggering by H\,{\sc{ii}} regions or supernova remnants play minor or major roles. Combined infrared and X-ray surveys, such as the {\it Chandra Carina Complex Project} \citep[CCCP;][]{Townsley2011}, have the advantage of locating both recently formed disk-bearing (proto)stars and older pre-main sequence (PMS) stars. These studies reveal spatially distinct clusters that appear to have different ages, and often a population of older widely distributed stars dispersed from earlier generations of star formation. Our group is now engaged in the Massive Young Star-Forming Complex Study in Infrared and X-ray \citep[MYStIX;][]{Feigelson2013} that examines 20 star forming complexes. MYStIX fields include a range of small and large cloud structures producing sparse and rich stellar clusters. A major impediment to understanding star formation in massive star forming regions (MSFRs) has been the absence of a reliable stellar chronometer so that the star formation history in a region can be reconstructed. While embedded protostars with Class 0 and Class I infrared spectral energy distributions (IR SEDs) clearly represent the most recent episodes of star formation activity, the sequence of past activity is difficult to unravel from IR SEDs alone because dusty protoplanetary disks have a wide range of longevities. Approximately 10\% of stars lose their disks in $\leq 0.5$~Myr, 10\% retain their disks for $\geq 6$~Myr, and the remaining 80\% have disk longevities smoothly distributed between 0.5 and 6~Myr \citep{Hernandez2008, Mamajek2009}. The disk fraction commonly used to estimate the age of a young stellar sample \citep[][ and citations thereof]{Haisch2001}, roughly the ratio \#(Class~I + Class~II)/\#(Class~I + Class~II + Class~III), therefore can only give a crude measure of stellar ages. Other stellar properties once proposed as chronometers, such as the depletion of photospheric lithium or the evolution of stellar rotation, do not show simple scalings with stellar age during the PMS phase \citep{Piau02, Eggenberger12, Bouvier13}. The best hope for an accurate stellar age estimator during the PMS evolutionary phase has been location on the Hertzsprung-Russell diagram (HRD) or photometric color-magnitude diagram (CMD), where stars of a fixed mass are predicted to progress along distinct tracks defined by convective interior models. Unfortunately, many problems are now recognized to affect a star's location in HRDs: photometric variability from accretion and magnetic activity, different accretional histories, binarity, extinction uncertainty, veiling from accretion, scattering and absorption by disks, stellar interiors model uncertainty, and distance uncertainty \citep[e.g.,][]{Baraffe2009, Naylor2009, Soderblom2010, Jeffries2011, Jeffries2012}. These issues have recently been reviewed by \citet{Preibisch2012} who emphasizes that `severe misinterpretations, gross overestimates of the age spread, and ill-based conclusions about the star formation history' can result from insufficient treatment of these effects. Preibisch also emphasizes the importance of disentangling the distinct subgroups that are often present in a star forming region in order to establish the `temporal sequence of several discrete star formation episodes'. In \S\S \ref{concept_section} and \ref{understanding_method_section}, we describe a new PMS stellar age estimator, $Age_{JX}$, that depends on near-IR stellar photospheric emission and on hard X-ray luminosity which arises mostly from enhanced magnetic flaring of PMS stars. Near-IR and X-ray photometric measurements are provided in surveys such as CCCP and MYStIX where they are used in the selection of MSFR members. $Age_{JX}$ is applied only to carefully chosen subsamples of members with appropriate near-IR and X-ray data, and is based on an empirical X-ray/mass relation calibrated to well-studied Taurus PMS stars \citep{Telleschi2007} and to theoretical evolutionary tracks calculated by \citet{Siess2000}. This age estimate mitigates some (although not all) of the above problems encountered in HRD- and CMD-based chronometry (\S\ref{some_advantage_subsection}). Individual star $Age_{JX}$ values for the MYStIX MSFRs range from $<1$ to $\sim 5$~Myr (Table~\ref{tbl_individual_ages}). As with other methods, we find a range of PMS $Age_{JX}$ values in large-scale MSFRs that could be considered to be a general wide spread in ages (\S\ref{reddening_effect_section}-\ref{subregions_subclusters_section}). But, in accord with the advice of \citet{Preibisch2012}, we recompute ages averaged over spatially distinct subclusters (\S\ref{subregions_subclusters_section}) that were recently obtained by \citet{Kuhn2014c} for MYStIX MSFRs. In that study, a statistical `finite mixture model' of isothermal ellipsoids was fit to the sky distribution of probable cloud members using maximum likelihood techniques. Each star can then be assigned as a member of a specific subcluster or as a member of a distributed unclustered population. The result of computing $Age_{JX}$ averaged over these subclusters is often remarkable: the full MSFR age distribution spanning several million years now appears as spatially segregated structures with narrowly constrained ages. Furthermore, we often see that the subclusters form spatial-absorption gradients consistent with astronomically reasonable patterns of star formation histories, with older lightly absorbed structures on one side, younger heavily absorbed structures on the other side, and intermediate-age structures in the middle (\S\ref{age_gradients_section}). Widely distributed populations are nearly always older. The progression of star formation within a MSFR can be traced, and patterns of star formation histories can be compared between MSFRs. We provide a homogeneous set of averaged $Age_{JX}$ for a hundred subclusters and regions of interest in MYStIX MSFRs (Table~\ref{tbl_cluster_ages}). | \label{conclusion_section} Researchers have historically used a number of different methods for dating young stellar clusters: isochrone fitting on the HRD or on photometric CMD, protoplanetary disk fraction, lithium depletion, and other procedures \citep{Mamajek2009b, Soderblom2010, Jeffries2012}. The placement of PMS stars on the HRD or CMD with theoretical evolutionary isochrones is probably the most popular method in the literature for dating young ($\la 10$~Myr) clusters but the individual stellar ages are highly uncertain. The distributions of stellar ages within rich clusters and across large-scale MSFRs often show wide age spreads; it has been unclear whether this is entirely due to poor individual age estimates or to real extended star formation histories. The problem of age estimation in MSFRs is particularly challenging, as reliable censuses of stellar members are hard to obtain and complex star formation histories are likely to be present. The MYStIX project has produced rich samples in 20 OB-dominated young star forming regions at distances $<4$~kpc using sensitive X-ray, NIR, and mid-infrared photometric catalogs \citep{Feigelson2013}. \citet{Kuhn2014c} has segmented the spatial distribution of MYStIX Probable Cluster Members into dozens of distinct subclusters. Some subclusters are deeply embedded in cloud material and have mostly disk-bearing stars, while others have low obscuration and mostly disk-free stars. The MYStIX sample thus provides a valuable opportunity to characterize star formation histories in MSFRs where optical or NIR spectroscopy is as yet unavailable. In this work, we develop a new PMS stellar age estimator, $Age_{JX}$, based on $J$-band stellar photospheric emission and on X-ray emission from coronal magnetic flaring. Stellar masses are directly derived from absorption-corrected X-ray luminosities using the $L_X - M$ relation from the Taurus cloud. These masses are combined with $J$-band magnitudes, corrected for source extinction and distance, for comparison with PMS evolutionary models \citep{Siess2000} to estimate ages. We provide a homogeneous set of median $Age_{JX}$ for over a hundred subclusters identified by \citet[][Table~\ref{tbl_cluster_ages}]{Kuhn2014c}. The method has a number of advantages over traditional methods of age estimation for young stellar clusters (\S\ref{some_advantage_subsection}). The main conclusions of our study are as follows: \begin{enumerate} \item We demonstrate that the $L_X-M$ relation, rather than the $L_X/L_{bol}$ ratio, remains slowly varying during the early PMS phase (\S \ref{lx_mass_pms_section}). \item The $Age_{JX}$ stellar ages are anticorrelated with the source extinction $A_V$ and the $J-H$ color. At high absorptions, the $Age_{JX} - A_V$ relationship is consistent with that of \citet{Ybarra2013} that was independently derived using disk fraction techniques. The $Age_{JX}$ vs. $J-H$ trend can serve as an approximate age predictor for young, highly reddened clusters, provided that the bulk of the reddening is local to the clusters (\S \ref{reddening_effect_section}). \item To overcome the dispersion of highly uncertain individual ages, we compute median ages (and their confidence intervals) of stellar samples within subclusters defined by the companion study of \citet{Kuhn2014c} (\S\S \ref{subregions_subclusters_section} and \ref{age_gradients_section}). We find narrowly constrained ages for these spatially distinct structures. Thus, we establish that (at least some) of the apparent age spread in MSFRs is real and can be attributed to clusters formed at different times. \item Spatial gradients in subcluster ages are often seen, consistent with astronomically reasonable patterns of star formation histories with older lightly absorbed structures on one side, younger heavily absorbed structures on the other side, and intermediate-age structures in between. Some regions do not show coherent star formation patterns. However, widely distributed populations nearly always have older ages than the principal clusters. The progression of star formation within a MSFR can be traced over spatial scales of several parsecs and time scales of several million years. Patterns of star formation histories can be compared between MSFRs. We caution that uncertainties in distance to MYStIX MSFRs can cause systematic shifts in the age scales (\S \ref{some_advantage_subsection}). \end{enumerate} | 14 | 3 | 1403.2741 |
1403 | 1403.5027_arXiv.txt | \noindent We explore the possibility of explaining a gamma-ray excess in the Galactic Center, originally pointed out by Hooper, collaborators, and other groups, in an effective field theory framework. We assume that dark matter annihilation is mediated by particles heavy enough to be integrated out, and that such particles couple to all quark families. We calculate the effective coupling required to explain the annihilation signal in the Galactic Center, and compare with bounds from direct detection, collider searches, and the requirement that the dark matter particle make up the appropriate fraction of the universal energy budget. We find that only a very small set of operators can explain the gamma-ray excess while being consistent with other constraints. Specifically, for scalar dark matter the viable options are one scalar-type coupling to quarks and one interaction with gluons, while for fermionic (Dirac) dark matter the viable options are two scalar-type dimension-7 operators or a dimension-6 vector-type operator. In all cases, future searches with the Large Hadron Collider should probe the relevant operators' effective energy scale, while all viable interactions should escape direct detection experiments. | \label{sec:intro} The gamma-ray flux from the Galactic Center region contains an excess, at energies in the 1-10 GeV range, over standard choices for the astrophysical background, as noted in 2009 by Goodenough and Hooper \cite{Goodenough:2009gk}, reiterated in 2010 by the same authors \cite{Hooper:2010mq}, then in 2011 by Hooper and Linden \cite{Hooper:2011ti}, and recently again by Hooper and collaborators \cite{newhooperon} and by other groups \cite{Abazajian:2012pn, Macias:2013vya, Abazajian:2014fta}. The Fermi-LAT Collaboration has also searched the Galactic Center \cite{fermilatgc}, reaching less striking conclusions. The excess has been tentatively associated with the pair-annihilation of a weakly-interacting massive particle (WIMP) in the inner Galaxy. For natural reasons, we indicate the tentative dark matter (DM) particle whose annihilation would produce said excess as the ``Hooperon''. If, indeed, astrophysical backgrounds cannot reproduce the observed excess, such a discovery would be an exciting breakthrough in the ongoing attempt to understand the nature of DM in the Universe. Pinpointing the implications of this scenario for particle physics model building might be perhaps premature, but could eventually become an exercise of the utmost importance to shed light on physics beyond the Standard Model (SM). Generally, to learn something about the implications of a DM model one requires a complete understanding of its interactions with SM particles. In some cases, however, it is possible to model such interactions using a relatively small number of possible operators, so long as the force-mediating particles are much heavier than the DM particles, and can thus be effectively integrated out. The effective field theory (EFT) framework necessary to perform such an analysis has been thoroughly explored in the literature (see e.g. \cite{everyone, Goodman:2010ku, Profumo:2013hqa}). Here, we apply EFT techniques to model a DM ``Hooperon'' particle that could produce the observed Galactic Center excess while being consistent with collider searches, direct DM searches, and producing an acceptable universal DM density in the early Universe. Indirect detection bounds are also relevant and are on the verge of probing the excess discussed here \cite{dwarfsbounds,GCbounds}. The present study is organized as follows: In the next section, we succinctly review the EFT framework; In Sec.~\ref{sec:bounds} we summarize current constraints on the relevant EFT operators from direct detection, relic density considerations, and collider searches; In Sec.~\ref{sec:signal} we calculate the required interaction strengths to reproduce annihilation cross sections consistent with the Galactic Center excess, and compare those values with current bounds; Finally, we conclude in Sec.~\ref{sec:conc}. | \label{sec:conc} We investigated the possibility that the Galactic Center excess in gamma rays could be explained by annihilation of a DM particle whose interactions with the SM are described by higher-dimensional effective operators considering direct detection, collider and abundance arguments. We explored a complete operator basis of all interactions of DM with hadronic matter within the heavy-mediator limit, for both scalar and fermionic (Dirac) DM. We found that there is a set of phenomenologically viable operators capable of producing the right pair-annihilation cross section today to explain the gamma-ray excess. For a scalar DM two options are viable: a scalar-type coupling to quarks, C2, and an operator with gluonic couplings, C6, while for Dirac DM the scalar-type operators D2 and D4 are most promising. Only one vector-type operator, D7, can explain consistently the gamma-ray excess, and only if its couplings are isospin-blind. If the couplings are opposite to up- and down-type quarks the LHC has imposed strict bounds on this operator as well. We emphasize that the relevant scales explored in this work might be lower than cut off scale which the EFT approximation is valid, and referred to stringent and robust direct detection constraints when applicable. We believe that these models provide an important benchmark which can be aimed for in DM searches. Ongoing work on LHC DM searches will continue to improve LHC bounds on effective theories of DM, and ongoing work in understanding the direct detection signatures due to operators which lead to suppressed scattering \cite{Fitzpatrick:2012ib} will help to shed light on the future detectability of the relevant effective operators with increasingly sensitive direct detection experiments. The effective Hooperon is a viable and testable scenario. | 14 | 3 | 1403.5027 |
1403 | 1403.1432_arXiv.txt | We present spectroscopic X-ray data of two candidate ultra-compact X-ray binaries: 4U~0614+091 and 4U~1543$-$624. We confirm the presence of a broad O VIII Ly$\alpha$ reflection line (at $\approx18\ \AA$) using {\it XMM-Newton} and {\it Chandra} observations obtained in 2012 and 2013. The donor star in these sources is carbon-oxygen or oxygen-neon-magnesium white dwarf. Hence, the accretion disc is enriched with oxygen which makes the O VIII Ly$\alpha$ line particularly strong. We also confirm the presence of a strong absorption edge at $\approx14$ \AA\ so far interpreted in the literature as due to absorption by neutral neon in the circumstellar and interstellar medium. However, the abundance required to obtain a good fit to this edge is $\approx3-4$ times solar, posing a problem for this interpretation. Furthermore, modeling the X-ray reflection off a carbon and oxygen enriched, hydrogen and helium poor disc with models assuming solar composition likely biases several of the best-fit parameters. In order to describe the X-ray reflection spectra self-consistently we modify the currently available {\sc xillver} reflection model. We present initial grids that can be used to model X-ray reflection spectra in UCXBs with carbon-oxygen-rich (and hydrogen and helium poor) accretion disc. We find that the new reflection model provides a better overall description of the reflection spectra of 4U~0614+091 and 4U~1543$-$624 than the reflection models that assume solar abundances. Besides a strong O VIII Ly$\alpha$ line the new reflection model also shows a strong O VIII K-edge (at $14.23$ \AA). We find that the absorption edge at $\approx 14$ \AA\ present in the data can be described by a O VIII K-edge formed due to reflection in the accretion disc and a Ne I K-edge originating mostly (if not entirely) in the interstellar medium, mitigating the problem of the apparent very high neon abundance. Additionally, based on the spectral properties of 4U~1543$-$624 we consider a scenario in which this source is accreting near the Eddington limit. | X-ray reflection is a common phenomenon in accreting X-ray sources. X-rays originating from the neutron star (NS), in flares above an accretion disc or at the bottom of a jet irradiate the disc. Some of this X-ray radiation is then reflected off the disc leading to an X-ray continuum spectrum and emission lines formed in the fluorescent and recombination process. The shape of the reflected continuum spectrum is determined by photoelectric absorption which is dominant at lower energies, electron scattering which is dominant at higher energies, and the continuum emission from the reflecting material itself. If the reflection spectrum originates close to the compact object relativistic effects such as gravitational redshift and the relativistic Doppler effect broaden the reflection signatures in a characteristic way \citep{Fabian1989}. By modeling the relativistically broadened reflection spectra one can in principle infer the inner radius of the accretion disc. The signatures of X-ray reflection such as the relativistically broadened Fe K-shell emission line have been observed in the X-ray spectra of a number of accreting compact objects e.g. Active Galactic Nuclei \citep[e.g.][]{Tanaka1995} and X-ray binaries \citep[][]{Miller2007,Cackett2009}. \\ The most frequently used reflection models: {\sc reflionx} \citep{Ross2005} and {\sc xillver} \citep{Garcia2010,Garcia2013} are best suited for spectra of AGN or X-ray binaries in the low (luminosity), hard (spectrum) state. These models assume a power-law incident spectrum and solar abundances in the accretion disc (with the sole exception of the Fe abundance which can vary). X-ray reflection in X-ray binaries in the high (luminosity), soft (spectrum) state was considered by \citet{Ross2007}. In this model the disc is no longer assumed to be cold like it is assumed in e.g. {\sc reflionx} but its emission is described by a black body with a variable temperature. A chemical composition of the accretion disc that differs from solar has not yet been considered when calculating the reflection spectra (with the said exception of Fe). Given that there is a growing number of sources showing X-ray reflection signatures from a disc with non-solar composition, there is need for reflection models which would be able to describe these reflection features self-consistently. \\ In this study we consider two sources: 4U~0614+091 and 4U~1543$-$624 that likely belong to the subclass of low-mass X-ray binaries (LMXBs) called ultra-compact X-ray binaries (UCXBs). UCXBs consist of a white dwarf donor star that transfers mass to a NS or a black hole (BH) accretor. They have an orbital period of $\lessapprox 80$ min. The accretor in the case of 4U~0614+091 is a NS given the presence of type I X-ray bursts \citep{Swank1978,Brandt1992}. In the case of 4U~1543$-$624 the nature of the accretor is uncertain (it could either be a NS or a BH). For the two sources studied in this paper tentative periods have been suggested: $P_{\rm orb}\approx 50$ min for 4U~0614+091 (Shahbaz et al. 2008) and $P_{\rm orb}\approx 18$ min for 4U~1543$-$624 \citep{Wang2004}. The optical spectra of both of these sources show emission lines of carbon and oxygen indicating that the donor star is either a carbon-oxygen (CO) or an oxygen-neon-magnesium (ONeMg) white dwarf \citep{Nelemans2004}. Furthermore, \citet{Madej2010} and \citet{Madej2011} have discovered a relativistically broadened O VIII Ly$\alpha$ reflection line in the X-ray spectra of these two sources. \\ The aim of this study is to investigate the potential changes in the properties of the reflected emission when considering CO-rich discs instead of disc material consisting of elements with solar abundances. The first sample of reflection models assuming CO-rich reflecting material are compared to archival as well as new observations of 4U~0614+091 and 4U~1543$-$624 obtained using the {\it XMM-Newton} and {\it Chandra} satellites. \begin{figure*} \vspace{-3mm} \includegraphics[width=0.8\textwidth]{plot_xillver.ps} \caption{Two {\sc xillver$_{\rm CO}$} test models (grid 1, see Table 1) with $kT^{\rm ref}_{\rm bb}$=0.1 keV (black curves) and $kT^{\rm ref}_{\rm bb}$=0.2 keV (red curves) and the $Frac=Flux_{\rm pl}/Flux_{\rm bb}=1$. The extremely narrow emission lines could be smeared out depending on the instrument resolution. {\it Top, left panel}: X-ray spectra, note the strong O VIII Ly$\alpha$ emission line at $\approx0.7$ keV in the case of a $kT^{\rm ref}_{\rm bb}=0.1$ keV. {\it Top, right panel}: temperature of the reflecting material as a function of the Thomson optical depth. {\it Bottom, left panel}: the incident power-law spectrum and the black body spectrum (included as a boundary condition at the bottom of the illuminated atmosphere). In order to keep the $Frac$ parameter the same also the luminosity of the incident power-law has increased for the case where $kT^{\rm ref}_{\rm bb}=0.2$ keV. {\it Bottom, right panel}: O ion fractions as a function of the Thomson optical depth. The O VIII ion fraction is plotted as a solid line and O IX ion fraction is plotted as a dashed line (see Sec. 2.1.1 for more detailed description). } \label{fig:xillver_test} \vspace{-3mm} \end{figure*} | We have taken the first steps to adapt the {\sc xillver} model \citep{Garcia2010,Garcia2013} to the case of reflection in UCXBs with CO-rich accretion discs. We have increased the abundances of all the elements except He and allow the abundance of C and O to vary during the fit. Additionally, we have considered a cutoff power-law incident spectrum with a cutoff energy in the range $2-10$ keV. As expected, the new reflection model ({\sc xillver$_{\rm CO}$}) shows stronger C and O emission lines. On the other hand the strength of the other emission lines e.g. Ne, Si, S, Fe is decreased. In contrast to the {\sc xillver} model which assumes the reflecting material to be cold, the reflecting material in the new {\sc xillver$_{\rm CO}$} model has a higher temperature which provides a better physical description of the gas in the accretion disc in X-ray binaries. \\ We have tested the modified reflection model {\sc xillver$_{\rm CO}$} using archival and new {\it XMM-Newton} spectra of 4U~0614+091 and {\it Chandra} spectra of 4U~1543$-$624. In the case of 4U~0614+091 we find that the new reflection model can describe the reflection signatures in the spectra and it indicates an overabundance of C and O of $A_{\rm C\&O}\approx 100-500$ with respect to the solar photospheric value of \citet{Lodders2003}. We note, however, that the current {\sc xillver$_{\rm CO}$} model is preliminary as it has a limited number of grid points. We have found that some of the parameters of this model e.g. $kT^{\rm ref}_{\rm bb}$, $Frac$, $A_{\rm C\&O}$ settle on the upper or lower limit values in the current grid which suggests that more grid points will need to be calculated before robust conclusions about the various parameter values can be drawn. Additionally, given that there are no significant C emission lines visible in the RGS and LETGS spectra of the studied sources, it is difficult to establish whether the overabundance of C is also required in order to describe the reflection spectra. We note that the {\sc xillver$_{\rm CO}$} model absorbed by the ISM also does not show significant C emission lines for the best-fit $A_{\rm C\&O}$ abundance. Hence, it is possible that the C emission is not strong enough if the C is overabundant.\\ Considering the parameters of the relativistic broadening obtained when fitting the {\sc xillver$_{\rm CO}$} model to observations 1, 2 and 3 of 4U~0614+091, we find that the value of the inner radius of the accretion disc is close to the innermost stable circular orbit (ISCO) $R_{\rm in}\approx 6\ GM/c^2$ in both low/hard and high/soft states. This result contradicts the standard accretion disc model \citep{Esin1997} in which the accretion disc is truncated further from the ISCO in the low/hard state. It needs to be mentioned, however, that it is possible that the properties of the reflecting material assumed in the {\sc xillver$_{\rm CO}$} (e.g. the abundances of elements or ionization structure of the disc) still deviate from the properties of the accretion disc in the UCXB 4U~0614+091 which prevent us from obtaining a self-consistent description of the data. Additionally, as mentioned above the caveat about the incomplete grid applies here as well. \\ \subsection{Ne I K-edge/O VIII K-edge and O VIII Ly$\alpha$ line} We confirm the presence of a strong absorption edge at $\approx 14$ \AA\ in the {\it XMM-Newton} spectra of 4U~0614+091 obtained in 2013 and LETGS spectra of 4U~1543$-$624 obtained in 2012 found before by \citet{Juett2001}, \citet{Madej2010}, \citet{Madej2011}, \citet{Schulz2010}. Our results suggest that the absorption edge at $\approx 14$ \AA\ can be partly described by the Ne I K-edge present in the neutral absorption model and partly by the O VIII K-edge present in the {\sc xillver$_{\rm CO}$} model. The possible presence of an O VIII K-edge or radiative recombination continuum (RRC) in the X-ray spectra of 4U~0614+091 has been suggested before by \citet{Schulz2010}. We find that the abundance of Ne decreases when the {\sc xillver$_{\rm CO}$} model is used. The best-fit parameters suggest, however, that the overabundance of Ne with respect to the solar photospheric value of \citet{Lodders2003} is still required. It is important to stress that the value of the Ne abundance can be uncertain, given the difficulty in measuring the abundance of this element in the Sun. We measure $A_{\rm Ne}\approx1.6^{+0.1}_{-0.2}$ with respect to the solar photospheric value of \citet{Lodders2003} when fitting the RGS data (see Sec. 4.3.1). However, the solar photospheric abundance of Ne in \citet{Lodders2003} is for example lower by a factor of $\approx 1.7$ and $\approx1.9$ with respect to the solar abundance of Ne in \citet{Anders1989} and \citet{Anders1982}, respectively. Hence, it is possible that the measured $A_{\rm Ne}$ is much closer to (or even consistent with) the abundance of Ne in the ISM. Additionally, we note that part of the Ne in ISM is thought to be in the ionized form \citep[e.g. Ne II, Ne III, see ][]{Juett2006} which is currently not taken into account in the {\sc tbnew} model. As a result the measurement of the neutral Ne abundance could be affected to some degree. \subsection{4U~1543$-$624: UCXB accreting near \\ the Eddington limit ?} We have found that the combination of a cutoff power-law with a cutoff energy of $\approx 5$ keV and a black body with a temperature of $\approx 0.4$ keV can describe the LETGS data of 4U~1543$-$624. These characteristics of the spectrum of 4U~1543$-$624 appear similar to the characteristics of spectra of ultra-luminous X-ray sources when observed in the ultraluminous state \citep{Gladstone2009}. Accretion near the Eddington limit can happen in UCXBs when the system comes into contact \citep{vanHaaften2012}. Theory predicts that this stage happens in BH or NS UCXB for orbital periods of around $8-11$ min depending on the type of the accretor and the influence of the thermal pressure on the donor star radius (Lennart van Haaften, private communication). This limit on the orbital period is around half of the period observed so far in the lightcurve of 4U~1543$-$624 \citep[$P_{\rm orb}=18\ {\rm min}$,][]{Wang2004}. The unabsorbed flux of the UCXB 4U~1543$-$624 measured using the LETGS spectra is $1.3\pm0.2\times10^{-9}\ {\rm erg\ cm^{-2}\ s^{-1}}$ ($0.1-10$ keV energy range). Assuming that this source is accreting near the Eddington limit $L_{Edd}\approx1-2\times10^{38}\ {\rm erg\ s^{-1}}$, we estimate the distance to the source to be $d\approx30-40\ {\rm kpc}$. Given the source Galactic coordinates ($l=322^{\circ}, b=-6^{\circ}$) and the estimated distance, 4U~1543$-$624 could be in the outskirts of the Galaxy and possibly belong to a halo population of LMXBs, implying that this source was kicked out of our Galaxy at formation. \\ We note that if 4U~1543$-$624 is indeed accreting near the Eddington limit, the $R_{\rm in}$ parameter measured using the relativistically broadened reflection model might not represent the true inner radius of the accretion disc. An optically thick corona present above the inner part of the accretion disc can modify the reflection spectrum (through absorption and multiple scattering events) in a way that will destroy the characteristic, inner-disc signatures. \subsection{Assumptions in the reflection models \& Future prospects} We have taken the first steps into describing the X-ray reflection spectra of UCXBs with CO-rich disc. However, apart from the issues already mentioned before in this section such as the limited number of grid points in the new model there are still questions remaining on the details of the X-ray reflection process and accretion disc physics in UCXBs. Those questions will need to be addressed before we are able to provide robust constraints on e.g. the geometry of the accretion disc or the abundances of elements in the accretion disc using X-ray reflection signatures. \\ For example, we have no robust observational constraints on the hydrogen number density in the inner regions of the accretion discs in UCXBs. Therefore, we have assumed in this analysis a hydrogen number density of $n_{\rm H}=10^{17}$ cm$^{-3}$. However, the ionization state of the reflecting material and hence the shape of the reflection spectrum depends on this parameter. Given the quality of the current data, introducing the density as a free parameters would probably reveal many degenerate solutions for the same dataset. However, future missions e.g. {\it Astro-H}, {\it Athena+} with higher spectral resolution and effective area than {\it XMM-Newton} and {\it Chandra} will reveal more details of the X-ray reflection signatures observed in the spectra and potentially be able to constrain the density of the reflecting material along with the currently considered parameters.\\ Another concern is the underabundance of H in the accretion disc of UCXBs. In our analysis we mimic the underabundance of H by increasing the abundance of all the other elements. Given the uncertainly on the exact abundance of H in UCXBs and the capabilities of the currently available reflection models we fixed the abundance of all the elements except He, C and O at ten times the solar value. However, future reflection models should be able to provide reflection spectra for the abundances of H even lower with respect to the other elements than currently assumed. | 14 | 3 | 1403.1432 |
1403 | 1403.6354_arXiv.txt | We develop an analytic model of intermittent, three-dimensional, strong, reduced magnetohydrodynamic (RMHD) turbulence with zero cross helicity. We take the fluctuation amplitudes to have a log-Poisson distribution and incorporate into the model a new phenomenology of scale-dependent dynamic alignment between the Els\"asser variables~$\displaystyle \bm{z}^\pm $. We find that the structure function $\displaystyle \langle |\Delta \bm{z}^\pm _\lambda|^n\rangle$ scales as $\displaystyle \lambda^{1-\beta^n}$, where $\displaystyle \Delta \bm{z}^\pm_\lambda$ is the variation in~$\displaystyle \bm{z}^\pm$ across a distance~$\displaystyle \lambda$ perpendicular to the magnetic field. We calculate the value of~$\beta$ to be~$\simeq 0.69$ based on our assumptions that the energy cascade rate is independent of~$\displaystyle \lambda$ within the inertial range, that the most intense coherent structures are two-dimensional with a volume filling factor~$\propto \lambda$, and that most of the cascade power arises from interactions between exceptionally intense fluctuations and much weaker fluctuations. Two consequences of this structure-function scaling are that the total-energy power spectrum is $\displaystyle \propto k_\perp^{-1.52}$ and that the kurtosis of the fluctuations is~$\displaystyle \propto\lambda^{-0.27}$. Our model resolves the problem that alignment angles defined in different ways exhibit different scalings. Specifically, we find that the energy-weighted average angle between the velocity and magnetic-field fluctuations is $\displaystyle \propto \lambda^{0.21}$, the energy-weighted average angle between~$\displaystyle \Delta \bm{z}^+$ and $\displaystyle \Delta \bm{z}^-$ is~$\displaystyle \propto \lambda^{0.10}$, and the average angle between~$\displaystyle \Delta \bm{z}^+$ and $\displaystyle \Delta \bm{z}^-$ without energy weighting is~$\propto [\ln(L/\lambda)]^{-1/2}$ when $L/\lambda \gg 1$, where $L$ is the outer scale. These scalings appear to be consistent with numerous results from direct numerical simulations. | \label{sec:intro} \vspace{0.0cm} Plasma turbulence plays an important role in many astrophysical systems, including accretion flows around black holes, intracluster plasmas in clusters of galaxies, and outflows from stars, including the solar wind. In many of these systems, the energetically dominant component of the turbulence is non-compressive and can be modeled, at least in an approximate way, within the framework of incompressible magnetohydrodynamics (MHD). In incompressible MHD, velocity and magnetic-field fluctuations ($\displaystyle \delta \bm{v}$ and $\displaystyle \delta \bm{B}$) propagate either parallel or anti-parallel to the local background magnetic field~$\displaystyle \bm{B}_{\rm loc}$, and nonlinear interactions occur only between counter-propagating fluctuations~\citep{iroshnikov63,kraichnan65}. As a consequence, the energy cascade is anisotropic, producing small-scale structures or ``eddies'' that satisfy $\lambda \ll l$, where $l$ ($\lambda$) is the correlation length of an eddy parallel (perpendicular) to~$\bm{B}_{\rm loc}$~\citep{shebalin83,goldreich95,ng96,galtier00,cho00,maron01}. When $\lambda \ll l$, the components of $\delta \bm{v}$ and $\delta \bm{B}$ perpendicular to~$\bm{B}_{\rm loc}$ evolve independently of the components parallel to~$\bm{B}_{\rm loc}$ and are well described by reduced MHD (RMHD)~\citep{kadomtsev74,strauss76}. When $\delta B \ll B_{\rm loc}$ and $\rho_{\rm p} \ll \lambda \ll l$, where $\rho_{\rm p}$ is the proton gyroradius, RMHD is a rigorous limit of gyrokinetics and is valid for both collisional and collisionless plasmas~\citep{schekochihin09}. In this paper, we propose a phenomenological theory of RMHD turbulence that goes beyond scaling theories for spectra~\citep{iroshnikov63,kraichnan65,goldreich95,boldyrev06} and allows us to make predictions concerning the scale dependence of arbitrary-order structure functions and the relative orientation of the turbulent magnetic field and velocity. A new feature of this theory is that it accounts, within one model, for both intermittency and scale-dependent dynamic alignment (SDDA). The concept of SDDA was introduced by \cite{boldyrev05,boldyrev06}, who argued that the angle $\phi_\lambda$ between~$\delta \bm{v}_\lambda$ and $\delta \bm{B}_\lambda$ decreases with decreasing~$\lambda$, where $\delta \bm{v}_{\lambda}$ and $\delta \bm{B}_\lambda$ are the fluctuations in the velocity and magnetic field at perpendicular scale~$\lambda$. As $\phi_\lambda$ decreases, nonlinear interactions in RMHD weaken, causing the power spectrum of the fluctuation energy to flatten relative to models that neglect~SDDA. Intermittency is the phenomenon in which the fluctuation energy is concentrated into an increasingly small fraction of the volume as~$\lambda \rightarrow 0$. Intermittency has been measured in hydrodynamic turbulence~\citep[e.g.,][]{benzi93}, solar-wind turbulence~\citep{burlaga91,horbury97,sorriso99,forman03,bruno07,wan12b,osman12,perri12,osman14}, numerical simulations of MHD turbulence and RMHD turbulence~\citep{muller00,maron01,muller03,beresnyak06,mininni09,imazio13}, and hybrid-Vlasov and particle-in-cell simulations of plasma turbulence~\citep{greco12,servidio12,wan12,karimabadi13,wu13}. A number of theoretical models have been introduced to describe intermittency, including the log-normal model~\citep{kolmogorov62,gurvich67}, the ``constant-$\beta$'' model~\citep{frisch78}, and multi-fractal models in which the fluctuation amplitudes scale differently on different subsets of the volume that have different fractal dimensions~\citep{parisi85,paladin87}. One such multi-fractal model, based on a log-Poisson probability distribution function for the local dissipation rate, was developed by \cite{she94} \citep[see also][]{dubrulle94}. She \& Leveque's~(1994) approach served as the basis for several previous studies of intermittency in both compressible and incompressible MHD turbulence~\citep{grauer94,politano95,muller00,boldyrev02a,boldyrev02b}. We draw upon ideas from the She-Leveque model to construct an analytic model of strong RMHD turbulence that incorporates a new phenomenology of SDDA. We present this model in Section~\ref{sec:theory}. In Section~\ref{sec:comp}, we compare our model with previously published numerical simulations, and in Section~\ref{sec:discussion} we discuss our results and the relation between our work and previous turbulence models. \vspace{0.0cm} | \label{sec:conclusion} \vspace{0.0cm} We have constructed an analytic model of intermittent, three-dimensional, strong RMHD turbulence that incorporates a new phenomenology of scale-dependent dynamic alignment. We restrict our analysis to the case of ``globally balanced'' turbulence, in which the cross helicity is zero. There are three main assumptions in our model. First, we take the fluctuation amplitudes to have a scale-dependent, log-Poisson PDF. In Section~\ref{sec:statistical}, we describe how this assumption can be motivated by treating a fluctuation's evolution as a random, quantized, multiplicative process, as in the work of \cite{she95}. Second, we assume that the most intense $\delta z^\pm_\lambda$ fluctuations are two-dimensional current/vorticity sheets with a volume filling factor~$\propto \lambda$. Third, we assume that the turbulence obeys a refined similarity hypothesis (Equation~(\ref{eq:eps1})) that includes the effect of dynamic alignment. We argue that the largest contribution to the average $z^+$ cascade power at any inertial-range scale~$\lambda$ comes from regions in which $\delta z^+_\lambda \gg \delta z^-_\lambda$ and $\delta z^+_\lambda \gg \delta z^\ast_\lambda$, where $\delta z^\ast_\lambda$ is the typical (median) fluctuation amplitude at scale~$\lambda$. We then develop an approximate theory describing how a large-amplitude, coherent $\delta z^+_\lambda$ structure interacts with a much weaker $z^-$ fluctuation. We show that during such an interaction, the $z^-$ fluctuation cascades rapidly to smaller scales without a reduction in amplitude and rotates into alignment with the coherent $\delta z^+_\lambda$~structure. By accounting for these effects, we compute the average $z^+$ cascade power using the assumed log-Poisson PDF of~$\delta z^\pm_\lambda$. This log-Poisson PDF has two free parameters, $\mu$ and~$\beta$ (see Equations~(\ref{eq:z1}) and (\ref{eq:chi1})). Our assumption that the most intense fluctuations form two-dimensional structures with a filling factor~$\propto \lambda$ determines~$\mu$ up to an additive constant~$A$, which affects neither the power-law scalings in our model nor the fact that~$\theta^\ast_\lambda$ (Equation~(\ref{eq:thetaast})) decreases logarithmically as~$\lambda \rightarrow 0$. The condition that the average cascade power is independent of~$\lambda$ then determines~$\beta$. Once we have determined $\mu$ and~$\beta$, we compute the scalings of the $z^\pm$ power spectrum, higher-order structure functions, and three different average alignment angles. Given the assumptions stated above, the scalings in our model do not depend upon free parameters and agree reasonably well with previously published numerical results. There are a number of ways in which our model could be improved. As presented, our model can approximate a broad distribution of outer-scale fluctuation amplitudes through the parameter~$A$, but the outer-scale distribution is then forced to be log-Poisson. A more realistic approach might be to allow the quantity $\overline{ \delta z}$ (Equation~(\ref{eq:z1})) to be random with a distribution that could be adjusted so as to model different forcing mechanisms in forced turbulence or different initial conditions in decaying turbulence. Our finding that $\tau_{{\rm nl},\lambda}^\pm$ is an increasing function of~$\delta z^\pm_\lambda$ at each scale suggests that, at least for some dissipation mechanisms such as Laplacian viscosity and resistivity, the dissipation scale is an increasing function of fluctuation amplitude. This would mean that the unusually intense fluctuations that make the dominant contribution to the power spectrum begin dissipating at a larger scale than the fluctuations that fill most of the volume. A useful direction for future research would be to develop this idea further by exploring the consequences of intermittency for the transition between the inertial and dissipation ranges within the framework of our analytic model. It would also be useful to extend our model to allow for nonzero cross helicity in order to investigate how intermittency affects strong ``imbalanced'' RMHD turbulence. Finally, inhomogeneity of the background plasma can fundamentally alter RMHD turbulence by causing the non-WKB reflection of Alfv\'en waves~\citep{heinemann80}. This linear coupling between counter-propagating Alfv\'en waves occurs in the solar atmosphere and solar wind~\citep{dmitruk02,cranmer05,verdini07,chandran09c} and can modify the power spectrum and energy-cascade timescales in solar-wind turbulence~\citep{velli89,verdini12,perez13}. Extending our model to account for background inhomogeneity and non-WKB wave reflection would be helpful for understanding intermittent turbulence in the inner heliosphere. | 14 | 3 | 1403.6354 |
1403 | 1403.3092_arXiv.txt | The mass of the lenses giving rise to Galactic microlensing events can be constrained by measuring the relative lens-source proper motion and lens flux. The flux of the lens can be separated from that of the source, companions to the source, and unrelated nearby stars with high-resolution images taken when the lens and source are spatially resolved. For typical ground-based adaptive optics (AO) or space-based observations, this requires either inordinately long time baselines or high relative proper motions. We provide a list of microlensing events toward the Galactic Bulge with high relative lens-source proper motion that are therefore good candidates for constraining the lens mass with future high-resolution imaging. We investigate all events from 2004 -- 2013 that display detectable finite-source effects, a feature that allows us to measure the proper motion. In total, we present 20 events with $\mu \gtrsim$ 8 mas yr$^{-1}$. Of these, 14 were culled from previous analyses while 6 are new, including OGLE-2004-BLG-368, MOA-2005-BLG-36, OGLE-2012-BLG-0211, OGLE-2012-BLG-0456, MOA-2012-BLG-532, and MOA-2013-BLG-029. In $\lesssim$12 years from the time of each event the lens and source of each event will be sufficiently separated for ground-based telescopes with AO systems or space telescopes to resolve each component and further characterize the lens system. Furthermore, for the most recent events, comparison of the lens flux estimates from images taken immediately to those estimated from images taken when the lens and source are resolved can be used to empirically check the robustness of the single-epoch method currently being used to estimate lens masses for many events. | Gravitational microlensing provides a useful tool for characterizing Galactic objects in a way that is unbiased by their brightness. One way to characterize a lens system is to determine its physical parameters by simultaneously measuring the microlens parallax {\pie} and Einstein radius {\thetae}. The Einstein ring represents the image of the lensed star (source) in the case of exact lens-source alignment, and its radius is commonly used as an angular scale in lensing phenomena. The Einstein radius is related to the physical parameters of the lens system by ${\thetae} = (\kappa M {\pirel})^{1/2}$, where $M$ is the total mass of the lens system, $\kappa = 4G/(c^{2} {\rm AU})$, ${\pirel} = {\rm AU}(D_{l}^{-1} - D_{s}^{-1})$, and $D_{l}$ and $D_{s}$ are the distances to the lens and source, respectively. The magnitude of {\pie} corresponds to the relative lens-source parallax, {\pirel}, normalized to {\thetae}. However, it is generally difficult to measure either {\pie} or {\thetae}, making it all the more unlikely to be able to measure both quantities simultaneously. As a result, determining the physical quantities of the lens system via this method has been possible only for a small fraction of lensing events. Another way to characterize a lens is to directly detect the light from the lens system. In general, direct lens detections are difficult because the typical separations between the lens and source at the time of the event or soon after are on the order of milli-arcseconds (mas), precluding resolution of the lens from the source. Nevertheless, in principle, measuring the flux of the lens and the source separately can be done without resolving the two systems. The flux of the source star can be ``de-blended'' from the combined blend flux from all unresolved objects, associated or otherwise, by fitting a microlensing model to the ground-based light curve, since only the source star is magnified during the event. A high-resolution image of the target then resolves out all unrelated stars with a high probability. Finally, converting the source flux derived from the ground-based data to the photometric system of the high-resolution data, typically taken in the near-infrared (NIR), and subtracting it from the flux of the target measured in the high-resolution data yields a measurement of excess flux that is not due to the source star itself or to stars with angular separations from the source that are greater than the resolution of the high-resolution data. Assuming this excess flux is due solely to the lens, the lens mass can then be estimated by combining the resulting mass-distance relation with a mass-luminosity relation and measuring the relative lens-source proper motion (e.g., \citealt{bennett06,bennett07}). This ``single-epoch'' method has been used to estimate the mass of a number of the hosts of planetary microlensing events (e.g., \citealt{sumi10,batista11}). However, there exists the possibility that some or all of the excess flux arises from companions to either the lens or the source, or even (with much lower probability) flux from unrelated stars that are blended with the source even in the higher-resolution images. Furthermore, ground-based $H-$band data exist for only a subset of microlensing events. For these cases, aligning the ground-based data from which the source flux is derived to the photometric system of the high-resolution data requires a large number of bright and isolated stars, which is difficult to achieve in the ground-based data due to the crowded fields and in the high-resolution image due to the small field of view. In the case that no ground-based $H-$band data exist, the source flux derived from the ground-based light curve must be transformed from the $I-$band optical flux measurements taken during the event to the NIR flux measurements of the high-resolution data. Thus, there are significant uncertainties present, whether in the alignment between the two photometric systems or the flux transformation from optical to NIR. It is possible to circumvent many of the potential difficulties of the single-epoch method by simply waiting until the source and lens have separated sufficiently that they are resolved in a high-resolution image \citep{han03}. Then, under the assumption that any potential companions to the lens are sufficiently dim, the flux of the lens can be measured directly, thereby eliminating any potential contamination from unrelated stars or companions to the source as well as any need to calibrate the photometry to the ground-based light curve data. However, most microlensing events toward the Galactic Bulge have sufficiently small relative lens-source proper motions such that resolving the lens and source requires one to wait an inordinately long time. In this paper, we present a catalog of lensing events toward the Galactic Bulge discovered over the period 2004 -- 2013 with sufficiently high relative proper motions of $\mu \gtrsim 8$ mas yr$^{-1}$ to allow for resolution of the lens and source and thus a direct measurement of the lens flux within $\lesssim$12 years, given the $\sim$0.1$\arcsec$ resolution of ground-based Adaptive Optics (AO) systems and space telescopes. A subset of these events currently retain modest lens-source separations of $\lesssim$20 mas and thus are unresolvable with present high-resolution systems. For these events, taking an immediate NIR high-resolution image will allow for a single-epoch estimate of the lens flux. By comparing this measurement to the direct measurement obtained later with an additional high-resolution image taken when the lens and source have separated enough to be resolved, it will be possible to directly and empirically test for contamination from companions to the source or from unrelated stars and to uncover any systematic errors involved in the flux alignment or transformation. Our catalog includes 20 total events with relative lens-source proper motions of $\mu \gtrsim$ 8 mas yr$^{-1}$. Of these, 14 have previously published values of $\mu$ while we present the analysis for 6 new events: OGLE-2004-BLG-368, MOA-2005-BLG-36, OGLE-2012-BLG-0211, OGLE-2012-BLG-0456, MOA-2012-BLG-532, and MOA-2013-BLG-029. The outline of the paper is as follows. In Section \ref{sec:events} we discuss the selection criteria for the events and describe the data for each new event. In Section \ref{sec:analysis} we explain the analysis procedure of the events. We then combine the results of our analysis with a thorough literature search to provide a catalog of lensing events with high relative proper motion in Section \ref{sec:results}. We also specifically address the current and future prospects of directly imaging the lenses in our catalog. In Section \ref{sec:conclusions} we present our conclusions. | \label{sec:conclusions} Here we have presented a catalog of microlensing events over the period 2004 -- 2013 with high proper motion, $\mu \gtrsim 8$ mas yr$^{-1}$. In the next $\lesssim$12 years, each of the events in our catalog will be sufficiently separated for direct imaging of the lens system. There are several ground-based telescopes with AO systems operating at or near the diffraction limit and space telescopes that have or will have the angular resolution necessary to resolve the lens from the source on that time scale. Such observations will further characterize each lens system and provide valuable insight into Galactic structure. We also urge for immediate high-resolution images to be taken of the six events whose lens and source are currently $\lesssim$20 mas. This includes MOA-2011-BLG-040 ($\sim$24 mas), OGLE-2012-BLG-0211 ($\sim$24), OGLE-2012-BLG-0456 ($\sim$24), MOA-2012-BLG-532 ($\sim$18), MOA-2013-BLG-029 ($\sim$8.4), and MOA-2013-BLG-220 ($\sim$12.5). For these events, the lens and source will still be unresolved in high-resolution images taken in the very near future. It will then be possible to compare the estimate of the lens flux obtained using this initial image with that obtained directly using a future high-resolution image taken when the source and lens are resolved. This comparison will provide an empirical check on the robustness of the single-epoch method currently being used to estimate lens masses for many events. | 14 | 3 | 1403.3092 |
1403 | 1403.3086_arXiv.txt | We present integral field unit (IFU) observations covering the [O~{\sc iii}]$\lambda\lambda4959,5007$ and H$\beta$ emission lines of sixteen $z<0.2$ type~2 active galactic nuclei (AGN). Our targets are selected from a well-constrained parent sample of $\approx24,000$ AGN so that we can place our observations into the context of the overall AGN population. Our targets are radio-quiet with star formation rates ($\lesssim$[10--100]\,\Msolyr) that are consistent with normal star-forming galaxies. We decouple the kinematics of galaxy dynamics and mergers from outflows. We find high-velocity ionised gas (velocity widths $\approx600$--1500\kms; maximum velocities $\le1700$\,km\,s$^{-1}$) with observed spatial extents of $\gtrsim$\,(6--16)\,kpc in all targets and observe signatures of spherical outflows and bi-polar superbubbles. We show that our targets are representative of $z<0.2$, luminous (i.e., $L_{{\rm [O~III]}}>10^{41.7}$\,erg\,s$^{-1}$) type~2 AGN and that ionised outflows are not only common but also in $\ge$70\% (3$\sigma$ confidence) of cases, they are extended over kiloparsec scales. Our study demonstrates that galaxy-wide energetic outflows are not confined to the most extreme star-forming galaxies or radio-luminous AGN; however, there may be a higher incidence of the most extreme outflow velocities in quasars hosted in ultra-luminous infrared galaxies. Both star formation and AGN activity appear to be energetically viable to drive the outflows and we find no definitive evidence that favours one process over the other. Although highly uncertain, we derive mass outflow rates (typically $\approx$10$\times$ the SFRs), kinetic energies ($\approx0.5$--10\% of $L_{{\rm AGN}}$) and momentum rates (typically $\gtrsim10$--$20\times L_{{\rm AGN}}/c$) consistent with theoretical models that predict AGN-driven outflows play a significant role in shaping the evolution of galaxies. | One of the most remarkable discoveries in modern astronomy is that all massive galaxies host supermassive black holes (BHs) with masses that are proportional to that of their host galaxy spheroid (e.g., \citealt{Kormendy95}; \citealt{Magorrian98}; \citealt{Tremaine02}; \citealt{Gultekin09}; see \citealt{Kormendy13} for a review). These BHs primarily grow through mass accretion that is visible as active galactic nuclei (AGN) in the centre of galaxies. Theoretical models of galaxy formation have found it necessary to implement AGN ``feedback'' processes, during which AGN activity injects energy into the gas in the larger-scale environment, in order to reproduce the properties of local massive galaxies, intracluster gas and the intergalactic medium (e.g., BH-mass--spheroid mass relationship; the sharp cut-off in the galaxy luminosity function; colour bi-modality; metal enchrichment; X-ray temperature--luminosity relationship; e.g., \citealt{Silk98}; \citealt{Churazov05}; \citealt{Bower06}; \citealt{Hopkins06}; \citealt{McCarthy10}; \citealt{Gaspari11}). Placing observational constraints on how AGN activity couples to the gas in galaxies and halos, and where these processes are most prevalent, is an important area of ongoing research (for reviews see: \citealt{Cattaneo09}; \citealt{Alexander12}; \citealt{Fabian12}; \citealt{McNamara12}). Several of the most successful galaxy formation models invoke a dramatic form of energy injection (sometimes called the ``quasar mode'' or ``starburst mode'') where AGN drive galaxy-wide (i.e., $\gtrsim0.1$--10\,kpc) energetic outflows that expel gas from their host galaxies and consequently this shuts down future BH growth and star formation and/or enriches the larger-scale environment with metals (e.g., \citealt{Silk98}; \citealt{Fabian99}; \citealt{Benson03}; \citealt{King03}; \citealt{Granato04}; \citealt{DiMatteo05}; \citealt{Springel05}; \citealt{Hopkins06,Hopkins08a}; \citealt{Booth10}; \citealt{DeBuhr12}). This is in contrast to the ``maintenance mode'' (or ``hot-halo'') feedback where radio jets, launched by AGN, control the level of cooling of the hot gas in the most massive halos (see \citealt{Bower12} and Harrison~2014\nocite{Harrison14a} for a discussion on the two modes). While there is little doubt that star formation processes (e.g., stellar winds and supernovae) drive galaxy-wide outflows (e.g., \citealt{Heckman90}; \citealt{Lehnert96}; \citealt{Swinbank09}; \citealt{Genzel11}; \citealt{Newman12b}; \citealt{Bradshaw13}; see review in \citealt{Veilleux05}) and are an integral part of galaxy evolution (e.g., \citealt{DallaVecchia08}; Hopkins et~al. 2013a\nocite{Hopkins13a}), it is believed that AGN activity is required to drive the highest velocity outflows and are particularly important for the evolution of the most massive galaxies (e.g., \citealt{Benson03}; \citealt{McCarthy11}; Hopkins et~al. 2013b\nocite{Hopkins13b}; \citealt{Zubovas14}). X-ray and ultraviolet spectroscopy has shown that a high-fraction, and potentially all, of high-accretion rate AGN drive high-velocity outflows ($v\approx0.1c$) close to their accretion disks (i.e., on sub-parsec scales; e.g., \citealt{Blustin03}; \citealt{Reeves03}; \citealt{Ganguly08}; \citealt{Tombesi10}; \citealt{Gofford11}). However, are AGN capable of driving outflows over galaxy scales as is required by galaxy formation models? A diagnostic that is commonly used to identify outflowing gas over large scales is broad (i.e., exceeding that expected from galaxy dynamics), asymmetric and high-velocity [O~{\sc iii}]$\lambda$5007 emission-line profiles. This emission line traces the kinematics of the ionised gas; however, we briefly note that outflowing gas has been observed in multiple gas phases in some AGN (e.g., \citealt{Rupke05b}; \citealt{Martin05}; \citealt{Fischer10}; \citealt{Feruglio10}; \citealt{Alatalo11}; \citealt{Veilleux13}; \citealt{Cimatti13}; \citealt{Rupke13}). As a forbidden transition the [O~{\sc iii}]$\lambda$5007 emission line cannot be produced in the high-density sub-parsec scales of the AGN broad-line region (BLR) making it a good tracer of the kinematics in the narrow-line region (NLR) and can be observed over parsecs to tens of kiloparsecs (e.g., \citealt{Wampler75}; \citealt{Wilson85}; \citealt{Boroson85}; \citealt{Stockton87}; \citealt{Osterbrock89}). The [O~{\sc iii}]$\lambda$5007 emission line has long been used to identify outflowing ionised gas in small samples of local and low redshift AGN (e.g., \citealt{Weedman70}; \citealt{Stockton76}; \citealt{Veron81}; \citealt{Heckman81,Heckman84}; \citealt{Feldman82}; \citealt{Vrtilek85b}; \citealt{Whittle85,Whittle88}; \citealt{Veilleux91,Veilleux95}; \citealt{Boroson92}; \citealt{Nelson96}); however, the small sample sizes makes it difficult to know how representative these observations are. More recently, large systemic spectroscopic surveys (e.g., the Sloan Digital Sky Survey [SDSS]; \citealt{York00}) have enabled the study of NLR kinematics in hundreds to tens of thousands of AGN (e.g., \citealt{Boroson05}; \citealt{Greene05a}; \citealt{Komossa08}; \citealt{Zhang11}; \citealt{Wang11}; \citealt{Mullaney13}) that can constrain both the ubiquity of these outflow features and study them as a function of key AGN properties. \cite{Mullaney13} used the SDSS spectroscopic database to study the one-dimensional kinematic properties of [O~{\sc iii}]$\lambda$5007 by performing multi-component fitting to the optical emission-line profiles of $\approx24,000$, $z<0.4$ optically selected AGN. They showed that $\approx17$\% of the AGN have emission-line profiles that indicate their ionised gas kinematics are {\em dominated} by outflows and a considerably larger fraction are likely to host ionised outflows at lower levels. The fraction of AGN with ionised gas kinematics dominated by outflows increases to $\gtrsim40$\% for the more radio-luminous AGN (i.e., those with $L_{\rm{1.4 GHz}}>10^{23}$\,W\,Hz$^{-1}$), in contrast, when taking into account intrinsic correlations, this fraction shows little dependence on [O~{\sc iii}] luminosity or Eddington ratio (\citealt{Mullaney13}). In agreement with smaller studies (e.g., \citealt{Heckman81}; \citealt{Whittle85,Whittle92}; \citealt{Gelderman94}; \citealt{Nelson96}; \citealt{Nesvadba11}; \citealt{Kim13}; see also \citealt{Greene05a}), this result shows that ionised outflows are most common in AGN that have moderate-to-high radio luminosities. However, while insightful, the origin of the radio emission is often unknown, particularly at the moderate radio luminosities (i.e., $L_{\rm{1.4 GHz}}\approx10^{23}$--10$^{24}$\,W\,Hz$^{-1}$) where AGN cores, radio jets, shocks and high-levels of star formation could all contribute (e.g., \citealt{DelMoro13}; \citealt{Condon13}; \citealt{Zakamska14}). It is therefore vital to measure star-formation rates (SFRs) and properly investigate the origin of the radio emission in the sources that host these outflows to properly interpret these results. The one-dimensional spectra discussed above provide no insight into the spatial extent or structure of the outflows, for this, we must appeal to spatially resolved spectroscopy. Both longslit and integral-field unit (IFU) observations of AGN, over a large range of redshifts, have identified disturbed and high-velocity ionised gas over kiloparsec scales (e.g., \citealt{McCarthy96}; \citealt{VillarMartin99}; \citealt{Colina99}; \citealt{Holt08}; \citealt{Nesvadba07a,Nesvadba08}; Lipari et~al. 2009a,b\nocite{Lipari09a,Lipari09b}; \citealt{Fu09}; \citealt{Alexander10}; \citealt{Humphrey10}; \citealt{Greene11}; \citealt{Rupke11,Rupke13}; \citealt{Harrison12a}; \citealt{Westmoquette12}; \citealt{CanoDiaz12}; \citealt{Husemann13}; \citealt{Liu13a,Liu13b}; \citealt{ForsterSchreiber13}). Several of these studies have revealed considerable masses of outflowing ionised gas with velocities higher than the galaxy escape velocities, in apparent agreement with basic predictions from galaxy formation models. However, a key limitation of these studies is that it is often difficult to place these observations into the context of the overall AGN and galaxy populations as the samples are small, inhomogeneous and/or only represent the most extreme AGN or star-forming systems in the Universe. In this paper we are interested in measuring the prevalence, properties, and the potential impact of galaxy-wide energetic outflows. We present Gemini (South) Multi-Object Spectrograph (GMOS; \citealt{AllingtonSmith02}) IFU observations of sixteen $0.08<z<0.2$ type~2 AGN drawn from the parent sample of \cite{Mullaney13}. Importantly, this means that we can place our IFU observations into the context of the overall AGN population. So that we can properly interpret our results we perform SED fitting and analyse the available radio data to measure SFRs, AGN luminosities and to search for evidence of radio jets in all of our targets. In Section~\ref{Sec:obs}, we give details of the IFU observations, data reduction and SED fitting. In Section~\ref{Sec:Analysis} we provide details of our analysis of the ionised gas kinematics and in Section~\ref{Sec:Results} we present our results. In Section~\ref{Sec:Discussion} we discuss our results and their implication for understanding galaxy evolution and in Section~\ref{Sec:conclusions} we give our conclusions. We provide background information and a discussion of the results on individual sources in Appendix~\ref{Sec:appA}. We have adopted $H_0 = 71$\kms\,Mpc$^{-1}$, $\Omega_{\rm{M}} = 0.27$ and $\Omega_{\Lambda} = 0.73$ throughout; in this cosmology, 1 arcsecond corresponds to $\approx$\,1.5--3.3\,kpc for the redshift range of our sample ($z=0.08$--0.2). | 14 | 3 | 1403.3086 |
|
1403 | 1403.1083_arXiv.txt | We show the results of microsecond resolution radio data analysis focused on flux measurements of single pulses of PSR~B1133+16. The data were recorded at 4.85~GHz and 8.35~GHz with 0.5~GHz and 1.1~GHz bandwidth, respectively, using Radio Telescope Effelsberg (MPIfR). The most important conclusion of the analysis is, that the strongest single pulse emission at 4.85~GHz and 8.35~GHz contributes almost exclusively to the trailing part of the leading component of the pulsar mean profile, whereas studies at lower frequencies report that the contribution is spread almost uniformly covering all phases of the pulsar mean profile. We also estimate the radio emission heights to be around 1\%--2\% of the light cylinder radius which is in agreement with previous studies. Additionally these observations allowed us to add two more measurements of the flux density to the PSR B1133+16 broadband radio spectrum covering frequencies from 16.7 MHz up to 32 GHz. We fit two different models to the spectrum: the broken power law and the spectrum based on flicker noise model, which represents the spectrum in a simpler but similarly accurate way. | Pulsar radio emission is still not explained in details but it is believed to originate close to the pulsar surface (\citealt{krzeszowski2009} and references therein). The analysis of 62 mean profiles of 23 pulsars at different frequencies regarding aberration and retardation effects yielded the estimation of the emission height being below 1500 km, but most probably being of the order of 500 km above the pulsar surface. The advance in theoretical understanding of the emission mechanism and conditions in the magnetosphere have been driven mainly by the observations. Recently many observations have been concentrated on observing single pulses which carry detailed information about the physics of radio emission. Analysis of single pulse high resolution time series can show a few interesting properties i.e. giant and bright pulses, subpulse drift, nulling, microstructure, etc. Observations of single pulses for most pulsars may be carried out mainly at lower frequencies because pulsars are weak radio sources at high frequencies which is clearly seen in their steep spectra. The pulsar spectrum in general can be described by a power law $S\propto\nu^\alpha$, where $\alpha$ is the spectral index. The average spectral index for 266 pulsars for frequency spread from 0.4 to 23~GHz is $\alpha= -1.8$ \citep{mkk+00}. \begin{table} \caption{Basic properties of PSR B1133+16 \citep{Brisken2002, manchester2005}.} \label{tab:props} \begin{tabular}{ll} \hline BNAME &B1133+16\\ JNAME &J1136+1551\\ $P$ &1.188 s\\ $\dot{P}$ &$3.73 \times 10^{-15}$ s~s$^{-1}$\\ RA ~~(J2000) &11$^\mathrm{h}$36$^\mathrm{m}$03$^\mathrm{s}$\\ DEC (J2000) &15$^\circ$51'04''\\ DM &4.86 pc~cm$^{-3}$\\ RM &1.1 rad~m$^2$\\ Age &$5.04 \times 10^6$ Yr\\ Distance &350 $\pm$ 20 pc\\ Proper motion &375 mas yr$^{-1}$\\ Transverse velocity &$631^{+38}_{-35}$~km~s$^{-1}$\\ $B_\mathrm{surf}$ &$2.13 \times 10^{12}$ G\\ $\dot{E}$ &$8.8 \times 10^{31}~\mathrm{erg}~\mathrm{s}^{-1}$\\ \hline \end{tabular} \end{table} PSR B1133+16 is a nearby middle--aged pulsar with one of the highest proper motion, and thus, one with the highest transverse velocity \citep{Brisken2002}. Basic properties of PSR B1133+16 are gathered in Table~\ref{tab:props}. Its faint optical counterpart (B=28.1 $\pm$ 0.3 mag) was firstly detected by \cite{Zharikov2008}. Recently \cite{Zharikov2013} detected the optical candidate of the pulsar counterpart on the GTC and VLT images that is consistent with the radio coordinates corrected for its proper motion. This source was also detected in X-rays by \cite{Kargaltsev2006} using the \textit{Chandra} satellite with the flux of (0.8 $\pm$ 0.2) $\times 10^{-14}$ ergs~cm$^{-2}$~s$^{-1}$ in the 0.5--8.0~keV range. For the X--rays fit the assumed hydrogen column density was $n_\mathrm{H} = 1.5 \times 10^{20}$ cm$^{-2}$. Low value of $n_\mathrm{H}$ and no $H_\alpha$ Balmer bow shock imply a low density of ambient matter around the pulsar. This pulsar has not been detected by the \textit{Fermi} satellite. In this paper in Sec. \ref{sec:observations} we describe observational parameters and technical issues about the recorded data. Sec. \ref{sec:single.pulses} covers analysis of the data. We present two different approaches for data analysis: mean profiles composed of pulses that their flux fall into specific intensity range as well as the phase position and flux of single pulses that are stronger than 20$\sigma$. In Sec. \ref{sec:emission.heights} we discuss radio emission height estimations, while in Sec. \ref{sec:spectrum} we present the radio spectrum of PSR B1133+16 and discuss different spectrum models. We conclude our results in Sec. \ref{sec:conclusions}. | \label{sec:conclusions} The analysis of PSR B1133+16 single pulses is a process that needs a certain amount of care. First off all, it is important to take into account different observational and technical effects that can affect recorded data, especially with very high time resolution. The effects that are presented in this paper play huge role and alter the data significantly. Understanding of such effects and their influence on the data recording process is important for proper data reduction. Some of the effects are not visible in the mean profiles but only in single pulses data. The mean profiles of PSR B1133+16 at 4.85~GHz and 8.35~GHz consist of two components (Fig.~\ref{fig:mean.profiles}). The second component is emitted almost exclusively by low intensity individual pulses. On the other hand, the first component is seen in single pulses regardless of their intensity except for the cases when it is not present at all. However, lower intensity emission contributes mostly to the leading part of the first component whereas higher intensity single pulses contribute mainly to its trailing part (Fig.~\ref{fig:intensity.ranges}) which was also reported by \citet{maron2013}. The results of analysis of 4.85~GHz and 8.35~GHz data are consistent with previous studies by \citep{nowakowski96} at 430~MHz but studies of B0329+54 \citep{mitra2007} show entirely opposite behaviour without full explanation. This inconsistency requires further studies of other pulsars single pulses to explain this effect. We show, in contradiction to studies at lower frequencies by \citet{kss+11}, who report an almost uniform spread of single pulse maxima, that the maximum emission of B1133+16 single pulses at 4.85 and 8.35~GHz contributes almost exclusively to the trailing part of the leading component of the mean profile. Our result is consistent with the behaviour at 341, 626, 1412 and 4850~HMz mentioned by \citet{kra+03} and extends the studies up to 8.35 GHz. Radio emission arises close to the pulsar surface at the distances of around 65 stellar radii at frequencies of 4.85~GHz and 8.35~GHz. Weaker emission, which contributes to the leading part of the leading components, comes in earlier phases which suggests that originates in magnetosphere further from the pulsar surface than more energetic emission. Our calculations shows, that the difference of the emission heights for stronger and weaker emission is of the order of a few stellar radii, amounting to a change of 1\%--2\% of the emission height, which is consistent with previous estimations \citep{krzeszowski2009}. There are 60 mean flux measurements in the literature of PSR B1133+16 that are known to us. They span a very wide radio frequency range from 16.7~MHz up to 32~GHz. To reproduce the spectrum we fitted two different models: the broken power--law model and one based on flicker noise model of pulsar radio emission \citep{ljk+08}. Surprisingly, the model proposed by \citet{ljk+08} is not widely used in the literature although it reproduces the pulsar spectrum comparably well to the power-law model. Future high time resolution observations might be useful to verify nano--pulse emission model. Due to the fact, that the pulsar radio emission is weaker at higher frequencies, giant pulses are the ones that can allow us to study closely their structure. In our both data samples there is roughly one per cent of bright pulses that are at least ten times stronger than mean flux and their microstructure is clearly visible. | 14 | 3 | 1403.1083 |
1403 | 1403.3753_arXiv.txt | We present variational Monte Carlo calculations of the neutron matter equation of state using chiral nuclear forces. The ground-state wavefunction of neutron matter, containing non-perturbative many-body correlations, is obtained from auxiliary-field quantum Monte Carlo simulations of up to about 340 neutrons interacting on a $10^3$ discretized lattice. The evolution Hamiltonian is chosen to be attractive and spin-independent in order to avoid the fermion sign problem and is constructed to best reproduce broad features of the chiral nuclear force. This is facilitated by choosing a lattice spacing of 1.5 fm, corresponding to a momentum-space cutoff of $\Lambda = 414$ MeV/c, a resolution scale at which strongly repulsive features of nuclear two-body forces are suppressed. Differences between the evolution potential and the full chiral nuclear interaction (Entem and Machleidt $\Lambda = 414$ MeV) are then treated perturbatively. Our results for the equation of state are compared to previous quantum Monte Carlo simulations which employed chiral two-body forces at next-to-next-to-leading order (N2LO). In addition we include the effects of three-body forces at N2LO, which provide important repulsion at densities higher than 0.02 fm$^{-3}$, as well as two-body forces at N3LO. | 14 | 3 | 1403.3753 |
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1403 | 1403.3279_arXiv.txt | Astroparticle experiments have provided a long list of achievements both for particle physics and astrophysics. Many of these experiments require to be protected from the background produced by cosmic rays in the atmosphere. The main options for such protection are to build detectors deep under ground (mines, tunnels) or in the deep sea or antarctic ice. In this proceeding we review the main results shown in the RICAP 2013 conference related with these kind of experiments and the prospects for the future. | \label{intro} Astroparticle experiments have provided richful physics results for many years. In this conference have seen many examples of the most recent advances and the prospects for the future, showing that this yield is growing and will continue giving us answers (and new questions) during the the following years. In this paper I will report on the results presented which are related with one three following topics: dark matter, under-ground experiments and neutrinos. | \label{conclusions} There are many frontier physics topics for which it is necessary to be protected from the atmospheric muon cosmic ray background and therefore to install detectors deep under-ground, under-sea or under-ice. These hot topics include the search for dark matter, the Majorana nature of neutrinos, the solar neutrinos or other astrophysical neutrino sources. We have seen in this conference important advances in several of these topics. In dark matter searches, for intance, we are in the very thrilling situation in which there are hints of first detection together with exclusion limits which in their simplest interpretation seem incompatible with such signals. In the mean time, we are likely attending the dawn of the neutrino astronomy. The events observed by IceCube are very difficult to explain from the background due to atmospheric muon and neutrinos, not only for the observed rates at high energies, but also for the ratio between showers and tracks, the observed spectral index or the distribution of the vertices in the detector. This first evidence of cosmic neutrinos should be confirmed by more data from IceCube and by the future detectors like KM3NeT and Baikal-GVD, already successfully deploying the first prototype lines. In paralel, we have also seen the potential of the multi-messenger approach, which allow to increase the sensitivy using the information of other channels like gamma rays or gravitational waves. Next years will provide us many more information (and probably puzzles) with all the data we will gather under the surface. | 14 | 3 | 1403.3279 |
1403 | 1403.3938_arXiv.txt | {We present the VIMOS Ultra Deep Survey (VUDS), a spectroscopic redshift survey of $\sim$10\,000 very faint galaxies to study the major phase of galaxy assembly $2<z\simeq6$. The survey covers 1 deg$^2$ in 3 separate fields: COSMOS, ECDFS and VVDS-02h, with targets selection based on an {\it inclusive} combination of photometric redshifts and color properties. Spectra covering $3650<\lambda<9350$\AA ~are obtained with VIMOS on the ESO-VLT with integration times of 14h. Here we present the survey strategy, the target selection, the data processing, as well as the redshift measurement process, emphasizing the specific methods adapted to this high redshift range. The spectra quality and redshift reliability are discussed, and we derive a completeness in redshift measurement of 91\%, or 74\% for the most reliable measurements, down to $i_{AB}=25$, and measurements are performed all the way down to $i_{AB}=27$. The redshift distribution of the main sample peaks at $z=3-4$ and extends over a large redshift range mainly in $2 < z < 6$. At $3<z<5$, the galaxies cover a large range of luminosities $-23< M_{NUV} < -20.5$, stellar mass $10^9$M$_{\sun}< M_* < 10^{11}$M$_{\sun}$, and star formation rates $1$ M$_{\sun}$/yr$< SFR < 10^3$M$_{\sun}$/yr. We discuss the spectral properties of galaxies using individual as well as stacked spectra. The comparison between spectroscopic and photometric redshifts as well as color selection demonstrate the effectiveness of our selection scheme. With $\sim6000$ galaxies with reliable spectroscopic redshifts in $2<z<6$ expected when complete, this survey is the largest at these redshifts and offers the opportunity for unprecedented studies of the star-forming galaxy population and its distribution in large scale structures during the major phase of galaxy assembly. } | The study of the first billion years of galaxy evolution is one of the key frontiers of modern cosmology. The current theoretical paradigm rests on the hierarchical build-up of dark matter halos in a $\Lambda$CDM cosmology (see e.g. Mo, van den Bosch \& White 2010). Galaxies formed in these deep potential wells, are expected to transform primordial gas into stars from the initial reservoir and to be fed from new accreted gas. As dark matter halos merge, galaxies in them are also expected to merge, these events deeply transforming the dynamics and overall star, gas, and dark matter content of the merging galaxies into the newly formed one. Along with these processes super-massive black holes are expected to form at the bottom of the potential wells. Complex processes are invoked to regulate the growth of galaxies, including supernovae or AGN feedback possibly quenching star formation, or the role of different environments impacting the way galaxies are nurtured. All these processes are combined in increasingly sophisticated numerical models producing galaxy simulations in representative volumes (e.g. Springel et al. 2008) coupled to semi-analytic description of galaxy evolution (e.g. Guo et al. 2011). These simulations face the challenge to reproduce both the internal physics in complex galaxy systems and the general volume-averaged properties of large galaxy populations as a function of cosmic time, and need better constraints from observations to be thoroughly tested. Impressive progress has been made on the observational front over the past two decades in an attempt to test and detail a galaxy formation and evolution scenario from robust measurements. A key element driving observational progress is the need to cover all major phases of galaxy evolution from the early formation and galaxy assembly to today, a formidable endeavor. Deep galaxy surveys have florished to conduct this exploration. The latest few billion years have been extensively mapped by surveys like the 2dFGRS (Colless et al. 2001) and then the various stages of the SDSS (Abazajian et al. 2009), setting a firm observational reference for status of galaxies after more than 13 billion years of evolution. At larger redshifts deep surveys are providing a complex picture with strong evolutionary features like the build-up along cosmic time of stellar mass in galaxies of different types or the star formation history. Different types of imaging and spectroscopic surveys are playing a complementary role, the deepest studies being performed in photometry and augmented with photometric redshifts (e.g. Ilbert et al. 2006), and with spectroscopic surveys bringing accurate redshifts and spectro-photometry, spectral features properties, as well as internal velocity information. The contribution of space observations in combination with ground-based surveys has been key to provide morphological information (Rix et al. 2004, Koekemoer et al. 2007, Koekemoer et al. 2011) and access to photometric bands invisible from the ground either in the far UV or in the mid to far infrared. Despite this progress, the exploration of early phases of galaxy formation and evolution is still largely incomplete. We do not know which objects ignited first at the end of the dark ages, when and how the universe was reionised, when and how the first massive galaxies formed, and the importance of any physical connection between galaxy and black-hole formation and growth. Star-forming galaxies are providing key information to understand how galaxies grow with time, in particular enabling to measure fundamental quantities such as the cosmic star formation history (e.g. Lilly et al. 1996, Madau et al. 1996, Tresse et al. 2007, Bouwens et al. 2009, Cucciati et al. 2012, Madau \& Dickinson 2014), and the history of stellar mass assembly (e.g. Arnouts et al. 2007, Ilbert et al. 2013). At redshifts $z\sim1$ the pioneering CFRS survey (Lilly et al. 1995, Le F\`evre et al. 1995) was followed by more extensive galaxy redshift surveys covering larger volumes like the DEEP2 (Davis et al. 2003), the VVDS (Le F\`evre et al. 2005a, Le F\`evre et al. 2013), zCOSMOS (Lilly et al., 2007), now reaching the 100\,000 redshift mark at $z\sim1$ with VIPERS (Guzzo et al. 2014). These surveys have brought a wealth of quantitative and accurate measurements now reaching large enough areas of a few tens of square degrees and volumes of $\sim 5 \times 10^7 h^{-3}$Mpc$^3$ so that the most fundamental statistical quantities describing the galaxy population like the luminosity function (LF), the mass function (MF) or the correlation function (CF) are becoming very accurate and less affected by the cosmic variance related to the small fields of earlier studies. At higher redshifts ($z >\sim 2$), the rapid progress has been driven by the effectiveness in selecting high redshift galaxies and, most importantly, by the impressive gains in sensitivity and efficiency provided by high multiplex multi-slit spectrographs like LRIS (Oke et al. 1995) and DEIMOS (Faber et al. 2003) on the Keck telescope, FORS (Appenzeller et al. 1998) and VIMOS (Le F\`evre et al. 2003) on the VLT. The effectiveness of the Lyman-break galaxies (LBG) selection has provided the capability to find large numbers of galaxies at $z>2.5$, and is continuing to be the single-most used technique to select galaxies at the highest possible redshifts (e.g. Steidel et al. 2003, Bouwens et al. 2009, Ellis et al. 2013). It is supplemented by narrow band imaging techniques isolating Lyman-$\alpha$ emitters (LAE; Taniguchi et al. 2005, Shimasaku et al. 2006; Ouchi et al. 2008), highly efficiency when reaching sufficiently deep to dig deeper into the LAE luminosity function at increasingly high redshifts. In addition to these pre-selection techniques, deep purely magnitude selected surveys were conducted in order to probe a large population mostly free of pre-selection biases. The largest to date probing redshifts $z>1.5$ is the $i-band$ magnitude selected VVDS survey covering up to $z\sim6.5$ with $>1000$ galaxies with redshifts $z\geq2$ (Le F\`evre et al. 2005a, 2013). Other magnitude selected surveys have attempted using redder bands to alleviate selecting only galaxies with strong rest-frame UV continuum. The K20 survey used K-band magnitude selection down to $K=20$ to identify extremely red objects (EROs) at $z\sim1.5-2$ which turned out to be dust obscured star forming galaxie or old passive early-type galaxies (Cimatti et al. 2002). In a following work the GMASS survey selected objects on the basis of Spitzer NIR photometry with $m_{4.5\mu m} < 23.0$ (AB) coupled to photometric redshift $z_{phot}>1.4$ to identify a few hundred galaxies with $1.5<z<3$, including 13 red passive galaxies (Cimatti et al. 2008). However, and despite these attempts, the approach using pure magnitude selection is costly in observing time when going much beyond redshift $z\sim2$ or so. Performing a complete galaxy census is a basic {\it Astronomy} input necessary for any subsequent {\it astrophysical} analysis. While at redshifts $\sim1$ this census is mostly complete down to stellar mass $10^8$M$_{\sun}$, it is not yet the case at redshifts $z>1$ for several reasons. First, the color selection schemes applied to photometric samples to extract the high redshift populations, while efficient to identify galaxies, are affected by significant incompleteness, loosing some fraction of the population at the selected redshift, or by contamination from galaxies at other redshifts. While the latter can be dealt with by obtaining spectroscopic redshifts, the former remains a serious difficulty especially at faint magnitude and at the highest redshifts. Unfortunately the level of incompleteness strongly depends on the photometric filters used for imaging, the depth of the observations, as well as the image quality, which requires a case by case study involving source simulations complicated by the need to make apriori hypotheses on the properties of galaxies one is trying to establish. It was realised that color-color selection like the LBG technique at $z>2.5$ or like the BzK working at $z\simeq2$ (Daddi et al. 2004) would miss a part of the general galaxy population in their selection process (Le F\`evre et al. 2005b), therefore making the galaxy census incomplete (Le F\`evre et al. 2014). The consequences of this maybe far-reaching, as incompleteness in counts leads to underestimates in luminosity density, star formation rates, as well as mass density, just to cite these few. An important aspect of on-going and future studies is to revisit galaxy counts as a function of redshift making sure that no significant population is missing and that no significant bias is introduced in deriving astrophysical quantities. A key element is then the availability of large samples of galaxies with a well defined and well controlled selection function. Spectroscopic redshift surveys play a key role as they provide samples with confirmed redshifts. Photometric redshift surveys are widely used and have now reached an impressive accuracy. However the level of 'catastrophic failures' when photometric redshifts disagree with their training set of spectroscopic redshifts, even if low at a few percent (Ilbert et al. 2013), could still produce large unknowns because of the shape of the N(z) of flux limited samples. An error of 1\% at the peak, $z\simeq1$, of the N(z) of a $i_{AB}=25$ sample could spread galaxies with wrong photometric redshifts to higher redshifts, e.g. at $z\sim3$ where the projected galaxy density is less than 10 times the N(z) at peak, which could then represent several tens of percent of uncertainty. One recent example is the difficulty to distinguish $z\sim5$ very massive objects from lower redshifts $z\sim2$ galaxies of lower mass (Wiklind et al. 2008, Caputi et al. 2012). Obtaining a spectroscopic redshift therefore remains a fundamental measurement. Unfortunately, the total number of galaxies spectroscopically confirmed today by all these surveys at $z>2$ is still limited. Published LBG samples reach $\sim$2000 redshifts at $z\sim3$, only $\sim150$ at $z\sim4$, and a few tens of galaxies beyond that (e.g. Steidel et al., 2003, Vanzella et al., 2009, Bielby et al. 2013). Samples of LAE galaxies selected with narrow band techniques and confirmed in spectroscopy reach a few hundred objects beyond $z=3$ (Ouchi et al. 2008, Kashikawa, 2011). The VVDS has assembled $\sim$35\,000 galaxies with spectroscopic redshifts down to $i_{AB}=24.75$, but the high redshift tail at $z > 2$ contains about 1000 galaxies (Le F\`evre et al. 2014). In front of the difficulty of obtaining large samples of spectroscopicaly confirmed galaxies, many surveys use samples defined solely on the basis of photometric color selection techniques like LBG or other simple color cuts, relying on completeness and contamination estimates difficult to control. To overcome the uncertainties linked to small existing spectroscopic samples, and to understand the biases and limitations of photometry-based studies in their ability to provide a complete census of the galaxy population, extremely deep spectroscopic surveys over large volumes are needed. Here we present VUDS, the VIMOS Ultra Deep Survey of $\sim$10\,000 galaxies specifically designed to study the early phases of galaxy evolution $2 < z < 6+$. The VUDS sample contains an unprecedented number of galaxies with secure spectroscopic redshifts at this epoch, obtained in three different fields: COSMOS, ECDFS and VVDS-02h. The survey design including the target selection is based mainly on photometric redshifts, as presented in Section \ref{overview}. The VIMOS multi-slit spectroscopic observations, the data reduction and redshift measurement scheme are discussed in Section \ref{observations}. Properties of the VUDS sample are presented in Section \ref{sample}, including the redshift distribution of the sample, the distribution of intrinsic properties like stellar mass and star formation rate, and the average spectra properties based on high signal-to-noise stacked spectra. After comparing to other surveys in Section \ref{comparison}, we summarize our results in Section \ref{summary} and conclude on the usefulness of the VUDS to study the early phases of galaxy assembly with unprecedented accuracy. \\ \\ All magnitudes are given in the AB system unless specified, and we use a Cosmology with $\Omega_M=0.3$, $\Omega_{\Lambda}=0.7$ and $h=0.7$. | \label{summary} The VIMOS Ultra Deep Survey (VUDS) is a deep spectroscopic redshift survey aiming to study the early phases of galaxy assembly at $2<z<6.5$ from a sample of $\sim10\,000$ galaxies observed with the VIMOS multi-slit spectrograph at the ESO-VLT. The survey target selection is based on photometric redshifts derived from extensive multi-wavelength data, combined to color and color-color selection as well as magnitude-selection. Most of the sample is limited down to $i_{AB}=25$, but galaxies are observed as faint as $i_{AB}=27$. The combination of a wide wavelength coverage from 3650\AA ~to 9350\AA, and exposure times of $\simeq14$h, lead to a spectroscopic success rate in redshift measurement of about 91\% (74\% for spectroscopic reliability flags 2 to 9) down to $i_{AB}=25$. The comparison of photometric redshifts to the VUDS spectroscopic redshifts shows that the VUDS strategy minimizes the loss of galaxy populations compared to more restrictive selection criteria. We report on the general properties of the sample based on $\sim80$\% of the data which has already been processed. The redshift distribution of the current sample at $z \geq 2$ peaks at a mean $z=3$, and extends beyond $z=6$. A secondary sample at $z<2$ is the result of the selection function and provides interesting very low intrinsic luminosity galaxies. The average spectral properties of galaxies with $z>2$ are discussed based on high S/N stacks of VUDS spectra in several increasing redshift bins. Galaxies with and without Ly$\alpha$ in emission are found at any redshift, but it is found that the fraction of galaxies with Ly$\alpha$ in emission increases with redshift. This is quantified in the accompaning paper by Cassata et al. (submitted). Using stacked spectra we find that there is observed flux below the 912\AA ~Lyman limit in all the redshift ranges explored, the origin of this is being investigated and will be the subject of future papers. Following an early measurement of the merger rate at $z\sim3$ (Tasca et al. 2013), several papers presenting results from the VUDS are submitted together with this paper (Cucciati et al. submitted, Cassata et al. submitted, Lemaux et al. submitted, Amorin et al. submitted). A number of other analyses are in progress. VUDS is the first deep spectroscopic survey covering such a large redshift range with such a large sample of galaxies with confirmed spectroscopic redshifts. It is ideally suited for detailed studies of the galaxy population at early times $2<z<6$. When the full data set will be completed, it is foreseen to make future VUDS data releases publicly available. | 14 | 3 | 1403.3938 |
1403 | 1403.3615_arXiv.txt | We study a hundred of galaxies from the spectroscopic Sloan Digital Sky Survey with individual detections in the Far-Infrared Herschel PACS bands (100 or 160 $\mu$m) and in the GALEX Far-UltraViolet band up to z$\sim$0.4 in the COSMOS and Lockman Hole fields. The galaxies are divided into 4 spectral and 4 morphological types. For the star forming and unclassifiable galaxies we calculate dust extinctions from the UV slope, the H$\alpha$/H$\beta$ ratio and the $L_{\rm IR}/L_{\rm UV}$ ratio. There is a tight correlation between the dust extinction and both $L_{\rm IR}$ and metallicity. We calculate SFR$_{total}$ and compare it with other SFR estimates (H$\alpha$, UV, SDSS) finding a very good agreement between them with smaller dispersions than typical SFR uncertainties. We study the effect of mass and metallicity, finding that it is only significant at high masses for SFR$_{H\alpha}$. For the AGN and composite galaxies we find a tight correlation between SFR and L$_{IR}$ ($\sigma\sim$0.29), while the dispersion in the SFR - L$_{UV}$ relation is larger ($\sigma\sim$0.57). The galaxies follow the prescriptions of the Fundamental Plane in the M-Z-SFR space. | One of the main aspects to understand galaxy formation and evolution focuses on the mass assembly of galaxies at different epochs. There are three galaxy properties which are fundamental when studying these processes and that are strongly interrelated to each other: the galaxy stellar mass (M), the metallicity (Z) and the star formation rate (SFR). The existence of a main sequence, MS, in the M-SFR relation (and its evolution with redshift) has been widely demonstrated \citep{Brinchmann04, Noeske2007, Elbaz2007, Daddi2007, Rodighiero2010}, as well as the mass-metallicity M-Z relation \citep{Tremonti2004, Lara09b}. Moreover, in the last few years it has been shown that the M-SFR and M-Z relations for star-forming galaxies are particular cases of a more general relationship, defined as the Fundamental Plane (FP) by \citet{Laralopez2010b} or the Fundamental Metallicity Relation (FMR) by \citet{Mannucci2010}. In a more recent work, \citet{LaraLopez2013} refine the parameters of the original representation of this plane, where the mass is a function of the metallicity and the SFR (M=f(Z,SFR)). This suggests that the stellar mass can be calculated as a linear combination of the rate at which a galaxy is currently forming stars (SFR) plus a measure of the star formation history, represented by the metallicity (corresponding to the amount of gas reprocessed by past stellar generations). It is of special relevance to have accurate measurements of these three galaxy parameters. The stellar masses can be estimated by fitting the photometric spectral energy distributions (SEDs) or the spectral features, while for the metallicity derivation it is strictly necessary to have robust emission line measurements from spectroscopy. The SFR, however, can be measured through different indicators in a wide wavelength range. Among the many SFR indicators we have X-rays, tracing X-ray binary emission (e.g. \citealt{Ranalli2003}); the UV, where the recently formed massive stars emit the bulk of their energy \citep{Schmitt2006,Rosa-Gonzalez2007}; optical wavelengths, from the recombination lines emission of the young massive population \citep{Kewley2002}; the mid-IR and FIR, since a significant fraction of the UV light of a galaxy is absorbed by the interstellar dust and reemitted in the infrared (e.g., \citealt{Kennicutt1998}, \citealt{Calzetti2005}, \citealt{Calzetti2007}, \citealt{Alonso-Herrero2006}, \citealt{Calzetti2009}, \citealt{Calzetti2010}); or the radio wavelengths, which traces supernova activity (e.g. \citealt{Yun2001}). One of the major problems when using the SFR indicators in the UV or optical regimes is the absorption of a great part of the light emitted by the young population by the dust surrounding the star-forming regions. Previous works combine optical and infrared observations to derive attenuation-corrected H${\alpha}$ and UV continuum luminosities of galaxies (e.g. \citealt{Gordon2000}, \citealt{Inoue2001}, \citealt{Hirashita2001}, \citealt{Bell2003}, \citealt{Hirashita2003}, \citealt{Iglesias-Paramo2006}, \citealt{Cortese2008}, \citealt{Kennicutt2009}, \citealt{Wuyts2011a}). The main advantage of the SFR from the FIR emission is that it is not affected by the dust extinction. However, in the pre-\textit{Herschel} era, the \LIR~estimates had to rely on the detection of IR emission at 24 or 70 $\mu$m from the \textit{Spitzer} data, meaning that the emission at longer wavelengths had to be extrapolated. Due to the different physical mechanisms and assumptions made to estimate the SFR at different wavelengths it is of great importance to see how these SFRs indicators compare to each other and which galaxy properties have a more important impact on their agreement/disagreement. Testing the validity of the SFRs indicators with complete samples of galaxies at low-z is fundamental to extend their validity at higher redshifts where the available data is usually scarcer. In a recent work, \cite{DominguezSanchez2012} studied a \textit{Herschel} selected sample at z $\leq$ 0.46 with \Ha~emission from the zCOSMOS survey \citep{Lilly2007,Lilly2009} and compared the SFR from the FIR with the SFR from the dust corrected \Ha~emission. We found an excellent agreement between the SFR indicators, except for very metal rich/poor galaxies. In this work we extend this analysis and compare various SFR indicators with each other (UV, \Ha, IR, SDSS, total). We combine for the first time the deep IR data from the latest PEP (PACS Evolutionary Probe, \citealt{Lutz2011}) \textit{Herschel} public data release with the extensive and already processed ancillary data from the Sloan Digital Sky Survey - Data release 7 (SDSS-DR7, \citealt{Abazajian2009}) in the COSMOS and Lockman Hole fields. The SDSS-DR7 data include masses, metallicities, emission line fluxes (e.g., \Ha~and \Hb) and SFRs for $\sim 10^6$ galaxies up to $z \sim 0.6$. We also use The Galaxy Evolution Explorer satellite (\textit{GALEX}, \citealt{Martin2005}) data in the far and near Ultra-Violet (FUV, NUV). The \textit{Herschel Space Telescope} has performed the deepest surveys in the FIR bands, which sample the IR peak of the galaxy spectra, helping to derive accurate \LIR~values. Using \LIR~and \Luv~ we derive SFR$_{total}$=SFR$_{UV}$+SFR$_{FIR}$, i.e., the obscured (SFR from the FIR) plus the unobscured SFR (SFR from the UV uncorrected for dust extinction). We also derive SFRs from the FUV and \Ha~fluxes, using two extinction correction approaches (from the observed \Ha/\Hb~ratio and the UV slope). We compare the predictions of the different SFRs estimators and study how the galaxy stellar mass and metallicity affect the comparison. We also investigate the relation between SFR and \LIR~and \Luv~ for a sample of AGN and composite galaxies. We locate the FIR counterparts of the SDSS galaxies in the M-Z-SFR space. This paper is organised as follows. In Section 2 we explain the sample selection and the data from the SDSS-DR7, \textit{Herschel} and \textit{GALEX} surveys. In section 3 we study the differences between the whole SDSS sample and the FIR detected sample. In Section 4 we derive luminosities at different wavelengths (\Luv, \LHa, \LIR, $L_{100}$ and $L_{160}$). Then we separately study the SF and unclassifiable galaxies in Section 5 (where we calculate dust extinctions and SFRs and compare the results), while in Section 6 we study the AGN and composite galaxies and the correlation of their \LIR~and \Luv~with SFR. In Section 7 we study the location of our sample of galaxies in the M-Z-SFR space. Finally in Section 8 we summarise our results and highlight the most important conclusions. Throughout this paper we use a standard cosmology ($\Omega_{m}=0.3,\Omega_{\Lambda}=0.7$), with $H_{0}=70$ km s$^{-1}$ Mpc$ ^{-1}$. The stellar masses are given in units of solar masses (M$_{\odot}$), and both the SFRs and the stellar masses assume a \cite{Kroupa2001} IMF. The SFR estimates from the SDSS-DR7 are derived with the Kroupa IMF. We convert the SFR derived with the \cite{Kennicutt1998} (K98 herafter) recipes, which assume a \citet{Salpeter1955} IMF into Kroupa IMF, by dividing the Salpeter SFR values by 1.5 \citep{Brinchmann04}. | In this work we have analysed a hundred SDSS galaxies at z $<$ 0.4 with a FIR counterpart in the PACS bands (100 and 160 $\mu$m) from the PEP survey carried out with the \textit{Herschel Space Telescope} in the COSMOS and Lockman Hole fields. From this sample, we have isolated 105 robust counterparts in FIR, with FUV emission from the GALEX data, and such galaxies constitute our main sample. We have divided this set into different spectral (SF, AGN, composites and unclassifiable galaxies) and morphological (E, S0, Sab, Scd) types. We have made use of extensive ancillary data from the SDSS, which includes masses, SFRs, metallicities or emission lines. We have derived different SFR indicators and compared them to study their validity and limitations. We have also placed our FIR SDSS counterparts in the Fundamental plane formed by the M-Z-SFR space. Our main conclusions are: \begin{itemize} \item The percentage of SF galaxies increases from 41.5$\%$ for the whole SDSS sample to 55.2$\%$ when galaxies are detected both in the FIR and FUV bands. The unclassifiable galaxies decrease from 43.5 to 11.4$\%$. With respect to the morphological classification, the number of Scd galaxies increases from 22.5 to 31.7$\%$, while the percentage of E decreases from 20.2$\%$ to 6.9$\%$. This selection effects are expected, as FIR and FUV emission are closely related with the star-formation (which mainly occurs in SF and late type galaxies). On the other hand, the unclassifiable galaxies present weak emission lines (i.e., low star-formation) and thus no significant FIR and FUV emission. \item The distribution of redshifts and masses is not strongly affected by the FIR/FUV detection. However, the FIR counterparts of the SDSS galaxies show larger SFRs (and therefore sSFRs) and slightly larger metallicities than the whole SDSS sample. \item $L_{100}$ seems to be a very good approximation of the total \LIR, with a slope in the \LIR-$L_{100}$~ relation m=0.99 and a dispersion $\sigma$=0.07. \item We derive dust extinction values from two different methods: from the observed \Ha/\Hb~ ratio, E(B-V)$_{R}$, towards the emission lines, and from the UV slope, E(B-V)$_{\beta}$, towards the continuum. We compare these extinction values with E(B-V)$_{IRX}$, from the \LIR/\Luv~ratio. E(B-V)$_{IRX}$ and E(B-V)$_{\beta}$ correlate very well, with small dispersion $\sigma$=0.06. The dispersion is larger for the E(B-V)$_{IRX}$-E(B-V)$_{R}$ comparison, $\sigma$=0.12. We derive a conversion factor between E(B-V)$_{UV}$ and E(B-V)$_{R}$ of 0.44, in excellent agreement with that from \citet{Calzetti2000}. \item We find a tight correlation between the E(B-V) and the \LIR, for the three studied methods of dust attenuation (dust extinction increasing with \LIR). The correlation between E(B-V) and metallicity is also significant (metal rich galaxies have higher dust extinctions). The relation between the E(B-V) and the stellar mass shows a too large dispersion to derive any significant correlation, specially at large masses (log M $>$ 10 \Msun). \item We have derived SFR$_{total}$ as the sum of the obscured (SFR$_{IR}$) and the unobscured (SFR$_{UV}$ without extinction correction) SFRs. The SFR$_{IR}$ represents more than 75 $\%$ of the SFR$_{total}$ for galaxies with log \LIR~$>$ 10 \Lsun~ and more than 90$\%$ for galaxies with log \LIR~$>$ 10.7 \Lsun. However, caution must be taken when deriving the SFR from the FIR emission only for low \LIR~galaxies, as the unobscured contribution may account for $\sim$ 50$\%$ of the total SFR. \item We have compared the SFR$_{total}$ with the one derived by the MPA-JHU group, SFR$_{SDSS}$, for the SF and unclassifiable galaxies. The agreement between the two SFRs for the SF sample is excellent with a slope in the SFR$_{total}$-SFR$_{SDSS}$ relation m=1.05 and a dispersion $\sigma$=0.20, smaller than typical SFR$_{SDSS}$ uncertainties for the SF sample ($\sim$ 0.28). \item SFR$_{total}$ and SFR$_{UV}$ or SFR$_{H\alpha}$ are also in a very good agreement for the SF sample, with slopes and dispersions $m_{UV}$=1.16, $\sigma_{UV}$=0.28, $m_{H\alpha}$=1.11, $\sigma_{H\alpha}$=0.43. The zero point in the SFR$_{total}$-SFR$_{H\alpha}$ relation is a=0.28, which causes that $\sim$ 84 $\%$ of the galaxies have SFR$_{total}$ $>$ SFR$_{H\alpha}$. There are $\sim$ 45$\%$ galaxies with SFR$_{total}$ $<$ SFR$_{UV}$, which indicates that there may be problems related to the dust extinction values derived from the UV. \item The relations obtained for the late type galaxies sample are very similar to those for the SF sample while for the unclassified and ETGs the relations show significant dispersions (larger than typical SFRs uncertainties). \item We have studied the dependence of the SFR comparison with the galaxy stellar mass and the metallicity. While the mass does not seem to affect the comparison of SFR$_{total}$ with SFR$_{SDSS}$ or SFR$_{UV}$, we find a significant difference between SFR$_{total}$ and SFR$_{H\alpha}$ for high galaxy stellar masses ($\sim$ 1 dex for log M $>$ 11 \Msun). The effect of the metallicity seems to be less important ($\sim$ 0.7 dex), but the number of galaxies with accurate metallicity values is small (32 high S/N SF galaxies) and they are mostly low metallicity galaxies (log (O/H) + 12 $<$ 9.3). \item We have studied the SFR$_{SDSS}$-\LIR~ and -\Luv~ relations for the AGN and composite galaxies. The dispersion of the SFR$_{SDSS}$-\Luv~ relation is too large ($\sigma$=0.57) to derive any recipe, but SFR seems to correlate very well with the \LIR~for both AGN and composite galaxies ($\sigma$=0.29). \item The SF sample of FIR SDSS counterparts seems to follow the MS relation obtained for the whole SDSS sample (m$_{FIR}$=0.79; m$_{SDSS}$=0.77, B04); while the AGNs, composites and unclassifiable galaxies always show lower sSFRs and are located below the MS. The best-fitting slope for the late type galaxies is larger (m$_{late}$=0.92) and shows an offset in the zero point due to the presence of late type galaxies with low sSFR. The majority of E and S0 galaxies lie below the MS. \item We have located the FIR counterparts of the SDSS galaxies in the fundamental plane formed by M-SFR-Z and confirmed that they follow the prescriptions derived by LL2013. \end{itemize} | 14 | 3 | 1403.3615 |
1403 | 1403.4729_arXiv.txt | The BICEP2 collaboration has recently reported a large tensor fluctuation in the cosmic microwave background, which suggests chaotic inflation models. In this letter, we reconsider the chaotic inflation model in the supergravity. We introduce a non-holomorphic shift-symmetry breaking parameter, which we expect to exist in general, and discuss its effect on the inflaton dynamics. We show that the model predicts a sizable deviation from the original chaotic inflation model and the predicted tensor fluctuation can lie between the BICEP2 result and the upper bound given by the Planck experiment with a small shift-symmetry breaking parameter. The model is characterized by only two parameters, which yields predictability and testability in future experiments. | Cosmic inflation~\cite{Guth:1980zm} is a natural scenario which not only solves the flatness and the horizon problem, but also explains the large scale structure of the universe and the fluctuation of the cosmic microwave background (CMB) radiation. Precise observations of the CMB~\cite{Hinshaw:2012aka,Ade:2013zuv,Ade:2013uln} begin to reveal nature of inflation. Recently, the BICEP2 collaboration has reported a large tensor fraction, $r= {\cal O}(0.1)$~\cite{Ade:2014xna}, which favors chaotic inflation models~\cite{Linde:1983gd}. Chaotic inflation models have been studied in the literature, especially in the context of the supergravity theory (SUGRA). In this letter, we reconsider chaotic inflation models in the SUGRA. In SUGRA chaotic inflation models, the shift-symmetry proposed in Ref.~\cite{Kawasaki:2000yn} is a crucial assumption. In order to obtain non-zero potential energy, the shift symmetry must be explicitly broken. In Ref.~\cite{Kawasaki:2000yn} the shift-symmetry breaking is introduced in the superpotential. However, it would be natural to consider that the Kahler potential also has shift-symmetry breaking terms. The shift-symmetry breaking in the Kahler potential is discussed in Refs.~\cite{Kallosh:2010ug,Li:2013nfa}, and it is shown that the prediction of the model deviates from that of Ref.~\cite{Kawasaki:2000yn} significantly. However, it is not clear whether the model possesses predictability. Higher dimensional terms in the Kahler potential may change inflaton dynamics due to large inflaton field value during inflation, once the shift-symmetry breaking is introduced. In this letter, we propose to treat the shift-symmetry breaking in a systematic way by introducing a non-holomorphic shift-symmetry breaking spurion ${\cal E}$ and discussing its effect on the inflaton dynamics. We restrict our attention to the range of the breaking parameter where higher dimensional terms are negligible for the inflaton dynamics and the model possesses predictability and testability. We show that the prediction for the spectral index and the tensor fraction can lie between the results of the Planck and the BICEP2 experiments with a small shift-symmetry breaking parameter. We also show that future observations of the CMB can quantify the reheating temperature of the universe within a factor of ${\cal O}(10)$. This letter is organized as follows. In the next section, we review the SUGRA chaotic inflation model. In Sec.~\ref{sec:breaking}, we introduce a non-holomorphic shift-symmetry breaking parameter ${\cal E}$ and discuss how the prediction on the spectral index and the tensor fraction is modified by the shift-symmetry breaking. We estimate the range of the shift symmetry breaking where higher dimensional terms in the Kahler potential are negligible, and show the prediction of the model within the range. The last section is devoted to discussion and conclusions. | In this letter, we have reconsidered chaotic inflation models in the SUGRA. We have introduced a non-holomorphic shift-symmetry breaking parameter ${\cal E}$ and discussed its effect on the inflaton dynamics. We have clarified the range of ${\cal E}$ where higher dimensional terms are negligible for the inflaton dynamics and the model possesses predictability and testability. We have shown that the prediction for the spectral index $n_s$ and the tensor fraction $r$ are given by $n_s\sim 0.96$ and $r = 0.11 - 0.18$. The prediction can lie between the results of the Planck and the BICEP2 experiments. It is interesting that future experiments will measure $n_s$ and $r$ accurately and reveal the structure of the shift-symmetry breaking in the inflaton sector. We have also shown that future observations of the CMB can quantify the reheating temperature of the universe within a factor of ${\cal O}(10)$, as long as ${\cal E}$ is in the range we have clarified. Let us comment on the magnitude of the shift symmetry breaking. We have introduced two shift-symmetry breaking parameters, $m$ and ${\cal E}$. The magnitude of the curvature perturbation indicates that $m\sim 10^{-5}$ and the consistency with the observed spectral index and the tensor fraction suggests that $|{\cal E}| \sim 10^{-3}$. Therefore, the two shift-symmetry breaking parameters are different by order of magnitudes. Note that the $m$ is a holomorphic parameter while ${\cal E}$ is a non-holomorphic one, and hence they may have different origins. We hope that a more fundamental theory explains the origin of the shift-symmetry breaking. Finally, let us briefly consider a model without the $Z_2$ symmetry. In this case, the Kahler potential is expanded as \begin{eqnarray} K = c (\Phi + \Phi^*) + \frac{1}{2} (\Phi + \Phi^*)^2 - i \frac{{\cal E}'}{\sqrt{2}} (\Phi - \Phi^*) - \frac{\kappa'}{2}\left(\frac{{\cal E}'}{\sqrt{2}}\right)^2 (\Phi - \Phi^*)^2 + \cdots, \end{eqnarray} and the scalar potential of the inflaton is given by \begin{eqnarray} V (\phi) = {\rm exp}\left( {\cal E}' \phi + \frac{\kappa'}{2}{\cal E}'^2 \phi^2 + \cdots \right) \frac{1}{2} m^2 \phi^2. \end{eqnarray} We can clarify the predictability of the model, that is, insensitivity to $\kappa'$, as we have done in this letter. It can be shown that the model possesses the predictability as long as \begin{eqnarray} |{\cal E}'| < 10^{-2.2}. \end{eqnarray} The prediction on $n_s$ and $r$ for $|{\cal E}'| < 10^{-2.2}$ is shown in Figure~\ref{fig:tilt-fraction-2}. Here, we have assumed that the inflaton field value is positive during the inflation. \begin{figure}[tb] \begin{center} \includegraphics[width=.8\linewidth]{tilt-fraction-2.pdf} \end{center} \caption{\sl \small The prediction on the spectral index $n_s$ and the tensor fraction $r$ for the model without the $Z_2$ symmetry. We also show the constraint from the Planck and the BICEP2 experiments. } \label{fig:tilt-fraction-2} \end{figure} | 14 | 3 | 1403.4729 |
1403 | 1403.7200_arXiv.txt | X-rays are particularly suited to probe the physics of extreme objects. However, despite the enormous improvements of X-ray Astronomy in imaging, spectroscopy and timing, polarimetry remains largely unexplored. We propose the photoelectric polarimeter Gas Pixel Detector (GPD) as an instrument candidate to fill the gap of more than thirty years of lack of measurements. The GPD, in the focus of a telescope, will increase the sensitivity of orders of magnitude. Moreover, since it can measure the energy, the position, the arrival time and the polarization angle of every single photon, allows to perform polarimetry of subsets of data singled out from the spectrum, the light curve or the image of source. The GPD has an intrinsic very fine imaging capability and in this work we report on the calibration campaign carried out in 2012 at the PANTER X-ray test facility of the Max-Planck-Institut f\"ur extraterrestrische Physik of Garching (Germany) in which, for the first time, we coupled it to a JET-X optics module with a focal length of 3.5~m and an angular resolution of 18~arcsec at 4.5~keV. This configuration was proposed in 2012 aboard the X-ray Imaging Polarimetry Explorer (XIPE) in response to the ESA call for a small mission. We derived the imaging and polarimetric performance for extended sources like Pulsar Wind Nebulae and Supernova Remnants as case studies for the XIPE configuration, discussing also possible improvements by coupling the detector with advanced optics, having finer angular resolution and larger effective area, to study with more details extended objects. | X-ray Astronomy obtained important results so far by using imaging, spectroscopy and timing. New observational techniques are required to refine theoretical models and remove degeneracies by adding new observational parameters. X-ray polarimetry would allow for introducing the degree and the angle of polarization that relate closely to the emission mechanism and to the source geometry. However, despite the enormous improvements in X-ray Astronomy, polarimetry remained largely unexplored. The first detection of polarized X-rays from an astrophysical source was obtained for the \object[M1]{Crab nebula} in the 1971, by means of a sounding rocket experiment \citep{Novick1972}. The result was later confirmed and the polarization was precisely measured with a degree of ($19.2\% \pm 1.0\%$) at 2.6 keV and ($19.5\% \pm 2.8\%$) at 5.2 keV \citep{Weisskopf1976, Weisskopf1978} by the polarimeter on board the OSO-8 satellite. This result ultimately proved the synchrotron origin of the X-ray emission of the nebula and remains, still today, the only precise non-zero one since the `70s, while upper limits were measured by \cite{Hughes1984}. The need to measure the polarization of high energy emission of many other sources remains urgent, and new technological solutions are available today. So far only few measurements of non-imaging polarimetry, with moreover a low significance, where performed. Non-imaging polarimetry averages the polarization of all subsystems within the field of view. For extended sources this can result in a substantial spoiling of the physical information. The reduction of measured polarization arises from the cancellation of the polarization vectors coming from regions with a different polarization state. This is crucial for extended sources such as the PWNe and the SNRs. X-ray polarimeters with imaging capability would allow to overcome this problem and to obtain polarization maps of extended sources. Moreover, imaging is useful to increase the signal to noise ratio for polarimetry by developing analysis strategies aimed to reduce the contamination of the emission due to source regions or emission components different from the ones of interest (see in particular the study of the pulsar signal in PWNe of Sect.~\ref{sec:PWN}). This improvement is possible only with a detector having simultaneously both polarimetric and imaging capabilities. A combination of an imager detector and a non-imaging polarimeter would not be adequate to this aim. The Gas Pixel Detector (GPD) \citep{Costa2001, Bellazzini2003} exploits the photoelectric effect to perform polarimetry and it is able also to make simultaneously spectral and timing measurements. The tracks of photoelectrons are produced in gas with a charge content proportional to the photon energy. From their initial emission direction the polarimetric measurement is derived, while the image is obtained as a map of the photoionization locations. The GPD is the most advanced 2-D imaging polarimeter with high polarimetric sensitivity and spatial resolution with respect to other instruments. For example CCDs were considered to perform polarimetry \citep{Tsunemi1992, Buschhorn1994} by exploiting the border effect among neighbour pixels to detect photoelectron polarization. However, this technique is heavily limited by systematics due to the small range of photoelectrons in silicon (only $\simeq$1.5$\mu$m at 10 keV) with respect to the pixel size. An other technique by \cite{Sakurai2004} exploits CCDs to detect the UV scintillation images of photoelectron tracks in a Capillary Gas Proportional Counter. Photoelectric effect in gas is exploited also in TPCs for GEMS \citep{Black2010}. A high quantum efficiency is obtained at expense of imaging, while its 1-D imaging capability is very much blurred by inclined penetration (see Sect.~\ref{sec:GPD}), due to focusing, in astronomical implementation \citep{Jahoda2010} with a consequent much larger background. At higher energies the Compton scattering polarimeter by \cite{Hayashida2012} has some imaging capability, with an angular resolution of few arcmin. In this case the spatial resolution depends basically on the width (few millimetres) of the scattering rods. The intrinsic imaging capability of the GPD was already studied by \cite{Soffitta2013a} who measured the \textit{spatial resolution} of the detector alone (with a narrow parallel X-ray beam). In our work we study the performance of a GPD combined with an X-ray telescope and compare them with predictions. From simulation studies \citep{Fabiani2009, Lazzarotto2010} we expect that the GPD, if coupled with an X-ray optical module with an intrinsic angular response in the range of a fraction of arcminute, should allow for imaging without a significant loss of performance with respect to the intrinsic angular resolution of the telescope. In this work we report about this, by proving it experimentally for the first time. Even if this paper is focused on the analysis of the imaging properties, we discuss, briefly, also the relationship between polarization and grazing incidence reflection. This is useful to clarify what is the expected limit of spurious polarization induced by optics and why we have no concern about the feasibility of polarimetry by means of the GPD coupled with X-ray telescopes. The GPD was placed at the focal plane of the Flight Module No. 2 (FM2) of the JET-X telescope \citep{Citterio1994, Spiga2013}. We will show the results of the measurement campaign performed at the PANTER X-ray test facility carried out between the 27$^{th}$ of November and the 1$^{st}$ of December 2012. The JET-X telescope (see Tab.\ref{tab:jetx} for the characteristics) was originally built for the former SPECTRUM-X GAMMA mission. \begin{deluxetable*}{c|c} \tabletypesize{\scriptsize} \tablecaption{JET-X telescope characteristics \citep{Spiga2013}.\label{tab:jetx}} \tablewidth{0pt} \startdata \hline \hline Configuration & Wolter-I\\ Focal length & 3500 mm \\ Diameter at entrance pupil (outer shell) & 300 mm \\ Diameter at entrance pupil (inner shell) & 191.1 mm \\ On-axis incid. angle at the intersection plane (outer shell) & 0.60$^\circ$ \\ On-axis incid. angle at the intersection plane (inner shell) & 0.39$^\circ$ \\ Mirror length (parabolic + hyperbolic) & 2$\times$300 mm \\ Reflecting surface material & Gold \\ no. of shells & 12 \\ Eff. area at 1.5 keV & 147 cm$^2$ \\ Eff. area at 8 keV & 53 cm$^2$ \\ FOV -- GPD+Telescope & 14.7~arcmin~$\times$~14.7~arcmin \\ \hline \enddata \end{deluxetable*} Finally, we show the simulated response for two kinds of extended sources, namely PWNe and Shell-like SNRs. The discussion is addressed with particular emphasis with respect to the detector configuration proposed on board the small pathfinder mission XIPE (X-ray Imaging Polarimetry Explorer) \citep{Soffitta2013b} which was presented, but not selected, to the ESA call of 2012 for a small mission to be launched in 2017. Two GPDs, effective in the 2--10 keV energy band, were meant to be coupled with two JET-X optics modules to perform polarimetry of astrophysical sources. In Sect.~\ref{sec:GPD} the GPD polarimeter and the main properties of the JET-X telescope are introduced. In Sect.~\ref{sec:GPDFocalPlane} the arrangement of the experimental set-up is explained. In Sect.~\ref{sec:onaxis} the on-axis angular resolution is studied, while the off-axis angular resolution is treated in Sect.~\ref{sec:offaxis}. In Sect.~\ref{sec:opticspolarization} we discuss briefly the effects on the polarization of grazing incidence reflection of X-ray in the optics. In Sect.~\ref{sec:astroobservations} the implications in terms of observational targets are discussed. | The GPD is a photoelectric polarimeter with an intrinsic imaging capability that makes it suitable to be used as a focal plane instrument to do the image of the source while performing polarimetry in the X-rays. So far, the imaging performance of the GPD coupled with X-ray telescopes were only studied by means of Montecarlo simulations and for the first time we demonstrated its feasibility by means of experimental measurements. We measured the Point Spread Function of the GPD placed at the focus of the JET-X X-ray telescope in the 2-10 keV energy band, both for on-axis and off-axis radiation beams. This detector/optics system is the configuration proposed for the pathfinder mission XIPE as response to the ESA small mission Call of 2012. We measured the angular resolution in terms of HEW, that is 22.7~arcsec at 2.98 keV, 23.2~arcsec at 4.51 keV and 28.9~arcsec at 8.05 keV for on axis radiation. In this work we showed that a detector/optics configuration typical of a pathfinder mission is able to obtain important results, opening the field of imaging polarimetry also in the X-rays. PWNe and SNRs were considered as case studies and the relation between the polarized emission, the source geometry and the magnetic field configuration were analysed. We demonstrated experimentally that the image quality of the optical system given by the GPD coupled to an X-ray mirror module depends mainly on the telescope intrinsic PSF. We showed that even with a small mission like XIPE bright SNRs and PWNe can be studied by imaging polarimetry. The availability of optics with a better angular resolution (few arcseconds) and a large effective area would allow GPD to obtain even more detailed images while performing sensitive polarimetry for many extended sources. Therefore, with IXO/ATHENA-like optics smaller and fainter features of a larger population of sources would be accessible. | 14 | 3 | 1403.7200 |
1403 | 1403.7807_arXiv.txt | Solar filament are commonly thought to be supported in magnetic dips, in particular, of magnetic flux ropes (FRs). In this Letter, from the observed photospheric vector magnetogram, we implement a nonlinear force-free field (NLFFF) extrapolation of a coronal magnetic FR that supports a large-scale intermediate filament between an active region and a weak polarity region. This result is the first in that current NLFFF extrapolations with presence of FRs are limited to relatively small-scale filaments that are close to sunspots and along main polarity inversion line (PIL) with strong transverse field and magnetic shear, and the existence of a FR is usually predictable. In contrast, the present filament lies along the weak-field region (photospheric field strength $\lesssim 100$~G), where the PIL is very fragmented due to small parasitic polarities on both side of the PIL and the transverse field has a low value of signal-to-noise ratio. Thus it represents a far more difficult challenge to extrapolate a large-scale FR in such case. We demonstrate that our CESE--MHD--NLFFF code is competent for the challenge. The numerically reproduced magnetic dips of the extrapolated FR match observations of the filament and its barbs very well, which supports strongly the FR-dip model for filaments. The filament is stably sustained because the FR is weakly twisted and strongly confined by the overlying closed arcades. | \label{sec:intro} Filaments are thin structures consisting of cool, dense plasma suspended in the tenuous hot corona. They lie above polarity inversion lines (PILs) on the photosphere, and are formed in filament channels where the chromospheric fibrils are aligned with the PIL \citep{Gaizauskas1997}. Filaments can be found inside active regions (ARs, ``AR filaments''), at the border of ARs (``intermediate filaments''), and on the quiet Sun (``quiescent filaments''). Most filaments eventually erupt and lead to coronal mass ejections, which are major drivers of space weather \citep{Schmieder2012}. The magnetic field plays a primary role in all the coronal processes including filament formation because the plasma $\beta$ is low. It is now commonly accepted that filament plasma is supported in magnetic dips, in particular, within twisted flux ropes (FRs) \citep[e.g.,][]{Rust1994SoPh, Chae2001ApJ, Ballegooijen2004}. Obtaining the three-dimensional (3D) magnetic field that supports filaments is key to understanding their structure, stability and eruption. Unfortunately, the 3D coronal field is very difficult to measure directly. People thus seek numerical models to construct the coronal field involved with filaments. The first numerical model was developed by \citet{Aulanierfilament1} using linear force-free field extrapolation from line-of-sight (LoS) magnetogram, which is proven to be a powerful tool to simulate filaments and related small structures like filament barbs \citep{Aulanierfilament1, Aulanierfilament2, Aulanier2000, Aulanier2002, Dudik2008, Dudik2012}. \citet{Ballegooijen2004} then developed a nonlinear force-free field (NLFFF) model, the FR insertion method, in which a FR with its axis following the targeted filament channel is first inserted into a potential field environment and the system is then relaxed to a force-free equilibrium. The FR insertion method has been applied in modeling various filaments \citep[e.g.,][]{Ballegooijen2004, Bobra2008, Su2011, Su2012,Savcheva2012}. NLFFF extrapolation from the photospheric vector magnetograms (VMs) is a more general method to reconstruct the coronal field \citep{Sakurai1981, Wu1990, Amari1997, Wiegelmann2012solar}, regardless of the presence of filaments. Unlike the FR insertion method, the NLFFF extrapolation can reconstruct a FR naturally and in a straightforward way, if it exists. Recent studies have reported many examples with FR extrapolated from different VMs using various NLFFF codes \citep[e.g.,][]{Yan2001, Regnier2004, Canou2010, Cheng2010, Guo2010, Jing2010, Jiang2013NLFFF}. Some studies also show that the extrapolated FR structure matches the related filament, e.g., good spatial correlation of the FR dips with the filament channels \citep{Canou2010,Guo2010}. We note that all these works are limited to relatively small-scale filaments (length within tens of megameters) whose channel is close to sunspots and along the main PIL with strong transverse field and magnetic shear. In such cases, the existence of a FR is usually predictable by inspecting the VM. The sheared PIL even with bald patches \footnote{Places on PIL where the transverse field is so strongly sheared that it points from the negative to the positive polarity, which is opposite to a potential field case \citep{Titov1999}.} usually indicates the presence of a coronal FR \citep{Titov1999, Aulanier2010}. On the contrary, many large-scale filaments (length up to hundreds of megameters), like the intermediate and quiescent filaments, usually lie above the photospheric weak-field region (field strength $\lesssim 100$G), where the PIL is very fragmented due to small parasitic polarities on both side and the transverse field is very noisy. As a result, the signature of FR, e.g., unbroken PIL with strong magnetic shear, is difficult to be observed directly from the magnetogram. Thus in such conditions it is far more of a challenge to extrapolate a large-scale FR from currently observed VM. In this Letter, we report such an challenging NLFFF extrapolation which recovers a large-scale FR that supports an intermediate filament with a length up to 300~Mm. The extrapolation is based solely on the SDO/HMI VM without any other observation or artificial input. It reproduces a coronal FR matching the filament recorded by AIA and H$\alpha$ strikingly well, which strongly supports the FR dip model for filaments. We further study why the FR can keep its stability from eruption. \begin{figure*}[htbp] \centering \includegraphics[width=0.8\textwidth]{fig1} \caption{The background shows a large AIA-304 image at 22:00~UT on September 6, overlaid with contours of $\pm 100$~G (red/blue) for LoS photospheric magnetic field. The white arc is the solar limb. The dashed box denotes the FoV of the VM used for extrapolation (see \Fig~2). The inserted panels are observations of the filament channels in AIA-304 and the H$\alpha$ filament at different times.} \label{fig1} \end{figure*} \begin{figure*}[htbp] \centering \includegraphics[width=0.8\textwidth]{fig2} \caption{HMI VM at 22:00~UT on September 6. Its FoV is outlined by the dashed box in \Fig~1. For left to right are vertical field strength, horizontal field strength, and azimuthal angle, respectively. The map shown here is Gaussian-smoothed from the original data with FWHM of 10~arcsec. The photospheric PILs are shown by the black contours in the left panel. The white thick lines overlaid on the maps are the PIL derived from the potential field extrapolated to a height of 10~arcsec, and only the part along the filament channel is shown.} \label{fig2} \end{figure*} | \label{sec:colu} This Letter reports an NLFFF extrapolation of solar coronal field that holds a large-scale filament from photospheric VM, which provides concrete evidence for filaments being supported by magnetic FR. Although the presence of a large-scale FR can hardly be predicted from the noisy VM, our CESE-MHD-NLFFF extrapolation code is able to overcome this difficulty to extract the key information of the coronal field from the magnetogram. We have also examined the robustness of the extrapolation by using VMs at two other different times nearby (not shown here), which also reproduce the similar FR structure. A detailed comparison with multiple observations, including those of stereoscopic viewpoints, demonstrates that the filament structure is well reproduced by the extrapolation. The FR that supports the filament is very stable because it is weakly twisted and strongly confined by the overlying closed arcades. Its eruption is likely triggered by a nearby AR eruption, which awaits further investigation. | 14 | 3 | 1403.7807 |
1403 | 1403.5669_arXiv.txt | The Rastall's theory is a modification of General Relativity touching one of the cornestone of gravity theory: the conservation laws. In Rastall's theory, the energy-momentum tensor is not conserved anymore, depending now on the gradient of the Ricci curvature. In this sense, this theory can be seen as a classical implementation of quantum effects in a curved background space-time. We exploit this structure in order to reproduce some results of an effective theory of quantum loop cosmology. Later, we propose a model for the dark sector of the universe. In this case, the corresponding $\Lambda$CDM model appears as the only model consistent with observational data. | The Rastall's proposal is based on the following observation: The conservation laws are tested effectively only on flat spacetime \cite{rastall}. Hence, in a non-flat spacetime, there is room for generalizations of the usual conservation laws. The choice made by Rastall to introduce a modification of the usual conservation law is: \begin{eqnarray} \label{cons} {T^{\mu\nu}}_{;\mu} = 0 \quad \Rightarrow \quad {T^{\mu\nu}}_{;\mu} = \frac{(\lambda -1)}{16\pi G}R^{;\nu}. \end{eqnarray} with $\lambda$ a constant. When $\lambda = 1$, General Relativity, with its usual conservation law, is recovered. \par Today, a new motivation to Rastall's theory can be given. Since the violation of the usual conservation law is related to the gradient of the curvature scalar, such violation can be associated to quantum effects in curved spacetime. In fact, the expression (\ref{cons}), in a particular form, appears in the so called gravitational anomaly, a anomaly in the usual gravitational equations due to quantum effects \cite{anomaly}. According to this remark, the Rastall's theory can be considered as a classical, effective implementation of quantum effects due to fundamental fields living in a curved spacetime. \par The field equations of the Rastall's theory can be written alternatively as follows: \begin{eqnarray} R_{\mu\nu} - \frac{1}{2}g_{\mu\nu}R &=& 8\pi G\biggr\{T_{\mu\nu} - \frac{\gamma - 1}{2}g_{\mu\nu}T\biggl\},\\ {T^{\mu\nu}}_{;\mu} &=& \frac{\gamma - 1}{2}T^{;\nu}. \end{eqnarray} In these equations $\gamma$ is a dimensionless parameter related to the previous one, $\lambda$, by the relation, \begin{equation} \gamma = \frac{2 - 3\lambda}{1 - 2\lambda}. \end{equation} When $\gamma = 1$, General Relativity, with the usual conservation law, is recovered. \par Is this theory well justified theoretically? First of all, there is in principle no Lagrangian formulation leading to the Rastall's field equation. But this is strictly true only at the context of the Riemannian geometry. Possible exterions (e.g., Weyl geometry) may cure this possible drawback. Moreover an introduction of an "external field" may lead to a Lagrangian formulation even in the context of the Riemannian geometry. \par In fact, consider the following Lagrangian: \begin{eqnarray} L = \sqrt{-g}\biggr\{R + \phi(R - \Lambda)\biggl\} + L_m, \end{eqnarray} where $L_m$ is the usual matter Lagrangian, $\Lambda$ is an external constant field, and $\phi$ is a Lagrange multiplier. Such modification of the usual Hilbert Lagrangian of General Relativity in order to obtain the Rastall's field equation. The parameter $\gamma$ is connected with $\Lambda$. \par Moreover, besides the initial motivations for the Rastall's theory, there are others possible connections with effective quantum models. As discussed in references \cite{anomaly,birrell}, quantum effects in curved spacetime lead to violations of the usual conservation law of the energy-momentum tensor. This phenomena is usually known as {\it gravitational anomalies}. Such gravitational anomaly may be very important, for example, in the computation of the Hawking radiation of a black hole \cite{glauber}. \par For example, consider the effective two-dimensional metric of a black hole: \begin{eqnarray} \label{met1} ds^2 = f(r)dt^2-f(r)^{-1}dr^2\quad. \end{eqnarray} The absence of ingoing models leads to \cite{wil}, \begin{eqnarray} \label{eang}\nabla_\mu T^\mu_{\nu}(r_H)\equiv\Xi_\nu(r)\equiv\frac{1}{\sqrt{-g}}\,\partial_\mu N^\mu_\nu\quad, \end{eqnarray} where, \begin{eqnarray} \label{ano} N^\mu_\nu=\frac{1}{96\pi}\,\epsilon^{\beta\mu}\partial_\alpha\Gamma^\alpha_{\nu\beta}\rightarrow N^r_t=\frac{1}{192\pi}\bigg(ff''+f'\,^2\bigg)\quad, \end{eqnarray} with $\epsilon^{01}=\epsilon^{tr}=1$ and ($'$) is a derivative with respect to $r$. \par Such interpretation of the Rastall's theory as an classical, effective implementation of quantum effects leads to many interesting applications, as we will see in what follows. First, we will use the Rastall's theory to implement a classical version of the Loop quantum cosmology effective equation, showing that singularity-free solutions emerge in that classical context. Second, we will apply Rastall's theory to the description of the present universe. It will be shown that the corresponding $\Lambda$CDM model is unique when a confrontation with observations is made. | Rastall's theory is a possible mechanism to introduce in a gravity theory quantum effects in a classical, effective approach. Rastall's theory touches one of the cornerstone of the General Relativity theory: the conservation of the energy-momentum tensor. Now, the divergence of the energy-momentum tensor is proportional to the gradient of the Ricci curvature scalar. \par One natural application of the Rastall's theory, viewed as an effective theory implementing classically the general features of quantum effects, is the primordial universe. With a suitable equation of state, it is possible to reproduce in the context of the Rastall's cosmology, the general behaviour predicted by Loop quantum cosmology. Singularity-free cosmological models are obtained. Some types of soft singularities are also displayed. \par In applying the Rastall's cosmology to the present universe, and using a model containing pressureless matter and a cosmological term, it is possible to keep the success of the $\Lambda$CDM model, but introducing new features at non-linear perturbation level. This {\it modified $\Lambda$CDM model} is, in some sense, unique, since any deviation of this configuration, by generalizing the equation of the state of the dark energy component, leads to a complete disagreement with observations, mainly with respect to the power spectrum data. \newline \vspace{0.5cm} \newline {\bf Acknowledgements:} We thank FAPES (Brazil) and CNPq (Brazil) for partial financial support. J.C.F acknowledges also support by “FONDECYT-Concurso incentivo a la Cooperaci´on Internacional” No. 1130628. | 14 | 3 | 1403.5669 |
1403 | 1403.0799_arXiv.txt | Gravitational instabilities play a primary role in shaping the clumpy structure and powering the star formation activity of gas-rich high-redshift galaxies. Here we analyse the stability of such systems, focusing on the size and mass ranges of unstable regions in the disc. Our analysis takes into account the mass-size and linewidth-size scaling relations observed in molecular gas, originally discovered by Larson. We show that such relations can have a strong impact on the size and mass of star-forming clumps, as well as on the stability properties of the disc at all observable scales, making the classical Toomre parameter a highly unreliable indicator of gravitational instability. For instance, a disc with $Q=1$ can be far from marginal instability, while a disc with $Q\ll1$ can be marginally unstable. Our work raises an important caveat: if clumpy discs at high redshift have scale-dependent surface densities and velocity dispersions, as implied by the observed clump scaling relations, then we cannot thoroughly understand their stability and star formation properties unless we perform multi-scale observations. This will soon be possible thanks to dedicated ALMA surveys, which will explore the physical properties of super-giant molecular clouds at the peak of cosmic star formation and beyond. | Today it is well established that the majority of the stellar mass observed in galaxies formed at high redshift, and in particular that the mean cosmological star-formation-rate density peaks at redshift 1--3 (e.g., Hopkins \& Beacom 2006). Recent semi-empirical studies (e.g., Behroozi et al.\ 2013) have allowed understanding how the peak redshift of individual galaxies depends, on average, on the mass of the host dark matter halo, with today's $L_{*}$ galaxy population forming stars at peak efficiency around $z=1\mbox{--}2$. Understanding the complex behaviour of galaxy assembly is a daunting task for galaxy formation theory (e.g., Weinmann et al.\ 2011), and highlights the need to build robust models that are capable of predicting how star formation proceeds in high-redshift galaxies. The morphology and star formation properties of massive high-redshift galaxies are very different from those of present-day quiescent spirals and ellipticals. Extended clumpy irregular discs with kpc-sized star-forming clumps as massive as $M\sim10^{7}\mbox{--}10^{9}\,\mbox{M}_{\odot}$ are observed in the Hubble Ultra Deep Field (UDF; e.g., Elmegreen et al.\ 2007, 2009), a population that is rare today. Multi-wavelength observational evidence (e.g., Elmegreen \& Elmegreen 2006; Shapiro et al.\ 2008; Tacconi et al.\ 2010) suggests that clumps generally form in gas-rich spiral discs rather than in mergers, although the latter scenario cannot be completely ruled out (e.g., Overzier et al.\ 2008). Numerical work by Bournaud et al.\ (2007) and Elmegreen et al.\ (2008) demonstrated that internal disc fragmentation can reproduce many of the observables of clumpy high-redshift galaxies. Using high-resolution hydrodynamical simulations in a fully cosmological framework, Agertz et al.\ (2009b) demonstrated that super-massive clumps are a natural outcome of fragmenting massive gas-rich discs, formed from multi-phase cosmological accretion (see also Ceverino et al.\ 2010). The size and mass of such clumps can be predicted using simple arguments, if one assumes that the disc is marginally unstable according to Toomre's stability criterion (e.g., Noguchi 1998, 1999; Dekel et al.\ 2009; Genzel et al.\ 2011). This assumption makes sense because current dynamical models of high-redshift star-forming galaxies suggest that their discs are driven by self-regulation processes, which keep them close to marginal instability (e.g., Noguchi 1998, 1999; Agertz et al.\ 2009a; Dekel et al.\ 2009; Burkert et al.\ 2010; Krumholz \& Burkert 2010; Cacciato et al.\ 2012; Forbes et al.\ 2012, 2014). If $Q\equiv\kappa\sigma/\pi G\Sigma=1$, then there is a single unstable wavelength, $\lambda=2\sigma^{2}/G\Sigma$, and the associated mass is $M\sim\Sigma\lambda^{2}=4\sigma^{4}/G^{2}\Sigma$. In the gas disc of the Milky Way, these quantities are comparable to the maximum size and mass of giant molecular clouds, i.e.\ $\lambda\sim100\,\mbox{pc}$ and $M\sim10^{6}\,\mbox{M}_{\odot}$ (see, e.g., Glazebrook 2013). In high-redshift discs, both the surface density $\Sigma$ and the velocity dispersion $\sigma$ of molecular gas are typically one order of magnitude larger than in the Milky Way (see again Glazebrook 2013). As $\lambda\propto\sigma^{2}/\Sigma$ and $M\propto\sigma^{4}/\Sigma$, we get $\lambda\sim1\,\mbox{kpc}$ and $M\sim10^{9}\,\mbox{M}_{\odot}$, which are the typical clump size and mass. Genzel et al.\ (2011) and Wisnioski et al.\ (2012) showed that the clumps are located in regions of the disc where $Q\la1$. This provides further evidence that in clumpy discs at high redshift there is a strong link between star formation and gravitational instability. In spite of its predictive power, such a scenario neglects an important aspect of the problem: in clumpy discs, the surface density and velocity dispersion depend on the size of the region over which they are measured (Romeo et al.\ 2010; Hoffmann \& Romeo 2012), contrary to what is generally assumed (see, e.g., Glazebrook 2013). In fact, there is mounting evidence that molecular gas is characterized by mass-size and linewidth-size scaling relations: \begin{equation} \Sigma\propto\ell^{a}, \;\;\;\;\;\mathrm{i.e.}\ M\propto\ell^{2+a}, \end{equation} \begin{equation} \sigma\propto\ell^{b}, \end{equation} where $\Sigma$ and $M$ are the mass column density and the mass of the clump, $\sigma$ is its 1D velocity dispersion, and $\ell$ is the clump size. \begin{enumerate} \item The most compelling evidence of such a link comes from observations of molecular clouds in the Milky Way and nearby galaxies (see, e.g., Hennebelle \& Falgarone 2012, and references therein; Donovan Meyer et al.\ 2013; Kauffmann et al.\ 2013; Kritsuk et al.\ 2013; Kruijssen \& Longmore 2013). These observations show that both Galactic and extragalactic molecular clouds are fairly well described by the so-called `\emph{Larson's scaling laws}', $a=0$ and $b=\frac{1}{2}$ (Larson 1981; Solomon et al.\ 1987), although the uncertainties are still significant: $-0.8\la a\la0.7$ (Beaumont et al.\ 2012), and $0.2\la b\la1.1$ (Shetty et al.\ 2012). \item Similar scaling exponents are found in high-resolution simulations of molecular clouds and supersonic turbulence (see, e.g., Hennebelle \& Falgarone 2012, and references therein; Beaumont et al.\ 2013; Federrath 2013; Kritsuk et al.\ 2013; Bertram et al.\ 2014; Fujimoto et al.\ 2014; Ward et al.\ 2014). The latter simulations show that $a$ depends not only on the Mach number of the gas, but also on turbulence forcing (Federrath et al.\ 2009, 2010; Federrath 2013) and self-gravity (Collins et al.\ 2012; Kritsuk et al.\ 2013). In contrast, at high Mach numbers, $b$ is approximately constant and close to 0.5 (see again Federrath 2013; Kritsuk et al.\ 2013). \item Larson-type scaling relations have recently been observed, for the first time, in the dense star-forming clumps of a high-redshift galaxy: the strongly lensed sub-millimetre galaxy SMM J2135--0102 at $z=2.32$, also known as the cosmic eyelash (Swinbank et al.\ 2011). Although this is the only detection of super-giant molecular clouds at high redshift, it will soon be followed by many such observations, which will exploit the unprecedented resolution and sensitivity of the Atacama Large Millimeter/submillimeter Array (ALMA) for exploring the physical properties of molecular gas at $z\ga2$ (see, e.g., Glazebrook 2013). \end{enumerate} Romeo et al.\ (2010) explored the gravitational instability of clumpy gas discs, and showed that the mass-size and linewidth-size scaling relations of the clumps can have a strong impact on disc instability. For instance, they can excite three main instability regimes, two of which have no classical counterpart. Hoffmann \& Romeo (2012) generalized this result to two-component discs of clumpy gas and old stars, and analysed the stability of spirals from The H\,\textsc{i} Nearby Galaxy Survey (THINGS). In this paper, we investigate the gravitational instability of clumpy discs at high redshift, focusing on the size and mass ranges of unstable regions (see Sect.\ 2). We begin by spelling out the assumptions of our stability analysis and summarizing the results of Romeo et al.\ (2010), which are fundamental to a proper understanding of this paper (see Sect.\ 2.1). Next, we discuss the effects of varying the clump scaling relations across the observed ranges of $a$ and $b$, and illustrate how the spatial resolution affects the \emph{inferred} stability properties of the disc, if the observed $\Sigma$ and $\sigma$ are scale-dependent (see Sect.\ 2.2). This is a complex aspect of the problem, which should be taken into account when analysing the stability of high-redshift star-forming galaxies. Last but not least, we discuss the properties of discs close to marginal instability (see Sect.\ 2.3). As pointed out above, this is the condition generally assumed for estimating the typical size and mass of the clumps. The disc scale height is expected to play a significant role in this scenario, since it is the scale at which galactic turbulence undergoes a transition from 3D to 2D (e.g., Bournaud et al.\ 2010), and this may be accompanied by a break in the clump scaling relations. We discuss this aspect of the problem in Sect.\ 3. The conclusions of our paper are drawn in Sect.\ 4. | We have investigated the gravitational instability of clumpy disc galaxies, focusing on the size and mass ranges of unstable regions. Multi-frequency observations of both the gas and the stellar contents (e.g., Elmegreen \& Elmegreen 2006; Shapiro et al.\ 2008; Tacconi et al.\ 2010) have established that such galaxies are ubiquitous at high redshift. Furthermore, the majority of stars in the Universe are known to form at $z>1$ (e.g., Hopkins \& Beacom 2006). Thus it is crucial to understand the properties of unstable star-forming gas at this epoch of galaxy evolution. Clumpy discs at high redshift are dynamically similar to gas discs with scale-dependent surface density and velocity dispersion, i.e.\ $\Sigma\propto\ell^{a}$ and $\sigma\propto\ell^{b}$, where $\ell$ is the clump size. Taking these `turbulent' scaling relations into account, and extending the traditional Toomre stability analysis as in Romeo et al.\ (2010), a wide variety of non-classical stability properties arise. We have illustrated this scenario for the whole observed range spanned by the clump scaling relations, which is centred around Larson's scaling laws $(a,b)=(0,\frac{1}{2})$, and for a range of spatial resolution scales typical of current high-redshift surveys. Our key results and a few eloquent examples are summarized below. \begin{enumerate} \item The scale-dependence of the surface density and velocity dispersion plays a \emph{crucial} role in determining the size and mass ranges of unstable regions. For example, in the rings and outer discs of SINS/zC-SINF galaxies at $z\sim2$ (Genzel et al.\ 2014), where the spatial resolution scale is close to the inferred 2D Jeans length, small variations in the logarithmic slope of $\Sigma(\ell)$ can lead to dramatic, order-of-magnitude, changes in the mass of the most unstable clumps. For the same observed surface density and velocity dispersion, logarithmic slopes of $\sigma(\ell)$ steeper than $b\approx0.4$ and flatter than $b\approx0.5$ ($a=0$) lead to complete disc stability. This illustrates the dynamical complexity introduced by the clump scaling relations. \item Variations in the logarithmic slopes of $\Sigma(\ell)$ and $\sigma(\ell)$ can drive significant changes in the stability properties of the disc at \emph{all} scales. For example, a clumpy disc can be marginally stable even if the classical Toomre parameter $Q_{0}\ll1$. In the case of Larson's scaling laws, the disc is always stable, however small $Q_{0}$ is, if the inferred 2D Jeans length $L_{\mathrm{J}0}$ is larger than the spatial resolution scale $\ell_{0}$. \item For discs with $Q_{0}=1$, we have payed special attention to $b\approx0.5$ and $-0.1\la a\la0.6$, since this range encompasses the most representative values of $a$ and $b$ found in observational (e.g., Larson 1981; Solomon et al.\ 1987) and theoretical (e.g., Federrath 2013; Kritsuk et al.\ 2013) works on supersonic turbulence. In spite of being marginally stable in the classical sense, such discs can be anywhere from highly stable to unstable, depending on the value of $a$. In fact, as $a$ approaches 0.6 while $L_{\mathrm{J}0}=\ell_{0}$, all observable scales $\ell\ga2L_{\mathrm{J}0}$ become unstable, \emph{even though} the disc is close to the stability threshold ($\overline{Q}_{0}\approx1.1$). \end{enumerate} Points (i)--(iii) illustrate the peculiar stability regimes possessed by discs with scale-dependent surface densities and velocity dispersions, and why it is important to take such regimes into account when predicting the size and mass of star-forming clumps in high-redshift galaxies. Note also that our work raises an important caveat: as the interstellar medium (ISM) is characterized by scale-dependent surface densities and velocity dispersions, we cannot thoroughly understand its global stability properties unless we carry out \emph{multi-scale} observations. This will soon be possible thanks to dedicated ALMA surveys, which will explore the physical properties of super-giant molecular clouds at the peak of cosmic star formation and beyond. Our work provides a new set of tools for exploring galactic star formation. In the ISM, there exist different sources of turbulence driving, such as large-scale gravitational stirring and stellar feedback (e.g., Mac Low \& Klessen 2004; Agertz et al.\ 2009a), and it is still unclear how they affect the ISM at various scales. Understanding the origin and evolution of $a$ and $b$, and how they vary with galactic environment, is a daunting task for numerical simulations, given the vast dynamical range involved in the star-forming ISM: from scales $\ell\la0.1\,\mbox{pc}$ to scales $\ell\sim10\,\mbox{kpc}$. Preliminary results from numerical simulations of entire galactic discs (Agertz, Romeo \& Grisdale, in preparation) show that large-scale gravitational stirring and stellar feedback can generate markedly different scaling properties in both $\Sigma(\ell)$ and $\sigma(\ell)$. This is a promising and novel avenue for constraining the role of stellar feedback in galaxy evolution, a topic that we will address further in future work. | 14 | 3 | 1403.0799 |
1403 | 1403.2049_arXiv.txt | { In this paper, we scrutinize very closely the cosmology in the proxy theory to massive gravity obtained in Phys. Rev. D84 (2011) 043503. This proxy theory was constructed by covariantizing the decoupling limit Lagrangian of massive gravity and represents a subclass of Horndeski scalar-tensor theory. Thus, this covariantization unifies two important classes of modified gravity theories, namely massive gravity and Horndeski theories. We go beyond the regime which was studied in Phys. Rev. D84 (2011) 043503 and show that the theory does not admit any homogeneous and isotropic self-accelerated solutions. We illustrate that the only attractor solution is flat Minkowski solution, hence this theory is less appealing as a dark energy model. We also show that the absence of de Sitter solutions is tightly related to the presence of shift symmetry breaking interactions. } | Whether the law of gravitation at cosmological distances can be described by general relativity or not will provide us a rich information of dark energy, which is responsible for the present accelerated expansion of the universe. One of such a candidate for alternative theories of gravity is massive gravity, originally proposed by Fierz and Pauli \cite{FP:1939aa}. They introduced a mass term in the linearized theory of general relativity in the context of Lorentz invariant theory. Unfortunately once Fierz-Pauli massive gravity is extended to a nonlinear theory the sixth degree of freedom called Boulware-Deser ghost appears \cite{Boulware:1972aa}. This problem was recently solved by de Rham and Gabadadze by adding higher order potentials, making the sixth degree of freedom removed \cite{deRham:2010ik}. It turned out that this infinite potential can be resummed by introducing the tensor, which has a square-root structure \cite{Rham:2011aa}, and this theory is now referred as de Rham-Gabadadze-Tolley (dRGT) massive gravity, which has been shown to be technically natural \cite{deRham:2012ew, deRham:2013qqa}. Since the inception of the dRGT theory there has been a flurry of investigations related to the self-accelerating solutions in the full theory. In dRGT theory, the universe can not be of a flat or closed Friedmann-Robertson-Walker (FRW) form \cite{PhysRevD.84.124046}, nonetheless, open FRW universe is still allowed \cite{Gumrukcuoglu:2011aa}. In this solution, the mass term exactly behaves as the cosmological constant, which allows a self-accelerating universe. However, the perturbations suffer from the instabilities, and the kinetic terms in the scalar and vector sector vanishes, which signals the strong coupling at a certain scale \cite{Gumrukcuoglu:2012aa,Gumrukcuoglu:2012ab,Felice:2012aa}. On the other hand, in Ref. \cite{Rham:2011ab} it has been shown, that there are exact de Sitter solutions in the decoupling limit theory, which is only valid within a certain region in the universe. This solution however suffers from ghost instabilities of the vector modes unfortunately \cite{Rham:2011ab, Koyama:2011wx, Gabadadze:2013aa}. In any case, it is very interesting that the mass of the graviton can drive an accelerated expansion of the universe \cite{Leon:2013qh}. As an alternative of massive gravity, one can covariantize the decoupling limit theory \cite{deRham:2011by}, and this "proxy theory" is not a massive gravity theory any longer but rather a non-minimally coupled subclass of Horndeski scalar-tensor theory \cite{Horndeski:1974aa}. Horndeski theory is the scalar-tensor theory, whose equations of motion remain second order differential equations, while the Lagrangian contains second derivatives with respect to space-time. It has been shown that Horndeski theory is equivalent with generalized Galileon theory \cite{Deffayet:2011aa}, which is the general extension of the Galileon theory \cite{Nicolis:2009aa}, and these theory contains four arbitrary functions in the Lagrangian\footnote{See Ref. \cite{Tasinato:2014eka, Heisenberg:2014rta} for the generalized vector Galileons.}. In the proxy theory these arbitrary functions can be automatically determined by covariantization, and it shares the same decoupling limit with dRGT massive gravity. In Ref. \cite{deRham:2011by}, the authors found a self-accelerating solution in a given approximated regime driven by the scalar field , which originally represents the helicity-0 mode in massive gravity. In contrast to the pure Galileon models, generalized Galileons do not impose the Galileon symmetry. The naive covariantization of the Galileon interactions on non-flat backgrounds breaks the Galileon symmetry explicitly, however one can successfully generalize the Galileon interactions to maximally symmetric backgrounds while keeping the corresponding symmetries \cite{Burrage:2011bt}. Inspired by these Horndeski scalar-tensor interactions, one can in a similar way construct the most general vector-tensor interactions with non-minimal couplings with only second order equations of motion \cite{Tasinato:2014eka, Heisenberg:2014rta}. The cosmology of these theories has been explored in \cite{Jimenez:2013qsa}. In the present paper, we study the cosmological evolution in the proxy theory in more detail beyond the approximations used in \cite{deRham:2011by} and show the absence of de Sitter attractor solutions which renders the theory not suitable as dark energy model. In Sec.2, we briefly review dRGT massive gravity and the derivation of the proxy theory. In Sec.3, we first investigate the de Sitter solution, then we study the dynamical system of cosmological solutions by using phase analysis. In Sec 4, we summarize our results. Throughout the paper, we use units in which the speed of light and the Planck constant are unity, $c=\hbar=1$, and $M_{\rm Pl}$ is the reduced Planck mass related to Newton's constant by $M_{\rm Pl}=1/\sqrt{8 \pi G}$. We follow the metric signature convention $(-,+,+,+)$. Some contractions of rank-2 tensors are denoted by ${\cal K}^{\mu}_{~\mu}=[{\cal K}]$,~ ${\cal K}^{\mu}_{~\nu}{\cal K}^{\nu}_{~\mu}=[{\cal K}^2]$,~ ${\cal K}^{\mu}_{~\alpha}{\cal K}^{\alpha}_{~\beta}{\cal K}^{\beta}_{~\mu}=[{\cal K}^3]$, and so on. | In this paper, we studied the cosmological dynamics of the proxy theory. For homogeneous and isotropic universe, there is de Sitter solution found in \cite{deRham:2011by}; however we show that this solution can be realized during only transient regime and can not be an attractor. In order to realize this transient de Sitter regime, we need fine-tuning of the initial conditions of the scalar field, thus the homogeneous and isotropic universe in the proxy theory can not be an alternative theory for dark energy model. Instead, the space-time approaches Minkowski space-time or type II singularity at the end, depending on initial conditions. In the proxy theory, the constant shift symmetry, $\pi \to \pi +c$, is broken by $\pi R$ interaction term while the decoupling limit theory in massive gravity satisfies this symmetry. If the theory satisfies the shift symmetry, then the field equation of the scalar field obeys $\ddot\phi+3H\dot\phi=0$, where $\dot \phi$ depends on models. In this case this equation can be easily solved, which gives ${\dot \phi} \propto a^{-3}$. This means that whatever the $\dot\phi$ is this variable will be diluted in the future, signaling an attractor solution. Furthermore, thanks to the shift symmetry, $\dot \phi$ only depends on $\dot \pi$, and $\pi$ never comes in any equations of motion, which means ${\dot \pi}={\rm const}$ could be an attractor solution with the wide range of initial conditions. This can be applied to the most general second order scalar-tensor theory which satisfies the shift symmetry. However, it should be noted that this is just a sufficient condition to have de Sitter attractor, not a neccessary condition. The shift symmetry breaking example in the Galileon theory can be found in \cite{Silva:2009aa,Kobayashi:2010ab}, there exist (quasi-) de Sitter attractor solution in these models. In addition, the case of massive scalar fields is an exception. One should note that there is exact de Sitter solution in the decoupling limit theory of massive gravity. Since the proxy theory shares the same decoupling limit with massive gravity, there should be exact de Sitter attractor solution within the patch enclosed by a sphere of radius, whose domain is order of the current horizon scale $H_0^{-1}$. This approximate solution should be connected to inhomogeneous or anisotropic solutions in the proxy theory in a similar way as it is the case in massive gravity itself. However, this would rely on the successful implementation of the Vainshtein mechanism \cite{PhysRevD.84.124046}. It would be interesting to study this kind of inhomogeneous and/or anisotropic solutions in a future work. | 14 | 3 | 1403.2049 |
1403 | 1403.3490_arXiv.txt | Based on SDSS and South Galactic Cap of u-band Sky Survey (SCUSS) early data, we use star counts method to estimate the Galactic structure parameters in an intermediate latitude with 10,180 main-sequence (MS) stars in absolute magnitude interval of $4 \leq M_r \leq 13$. We divide the absolute magnitude into five intervals:$4 \leq M_r < 5$, $5 \leq M_r < 6$, $6 \leq M_r < 8$, $8 \leq M_r < 10$, $10 \leq M_r \leq 13$, and estimate the Galactic structure parameters in each absolute magnitude interval to explore their possible variation with the absolute magnitude. Our study shows the parameters depend on absolute magnitude. For the thin disk, the intrinsic faint MS stars have large local space density and they tend to stay close to the Galactic plane. A plausible explanation is that faint MS stars with long lifetime experience long gravitational interaction time result in a short scaleheight. However, for the thick disk, the parameters show a complex trend with absolute magnitude, which may imply the complicated original of the thick disk. For the halo, the intrinsic faint MS stars have large local density and small axial ratio, which indicate a flattened inner halo and a more spherical outer halo. | The star counts method has been used to study the structure of our galaxy by generations of astronomers. This useful method can provide a measurement of the density distribution of the stellar component of the Galaxy. Since \citet{Bahcall1980} fitted observations with two components Galactic model, namely disk and halo, and then improved by \citet{Gilmore1983} through introducing a third component, namely the thick disk, the star counts method has provided a picture of the standard Galactic model. As the great improvement on data collection over the years, more and more researchers try to refine both the estimation of standard Galactic model parameters and the standard Galactic model to explain the currently available observations well. The better we know the Galactic structure, the better we understand the formation and evolution of the Galaxy. A simple way to study Galactic structure is assuming a global smooth structure model of our Galaxy, which is based on the modern of physics, and also on the pioneering work of \citet{Eggen1962}, who argued that the galaxy formed from a relatively rapid ($\thicksim$$10^8$ years) radial collapse of the protogalactic cloud through studying the correlation between ultraviolet excess, metal abundance, angular momentum and the eccentricity of the Galactic orbit. The astronomers have done lots of works on the smooth structure parameters. We can find elaborate list of the structural parameters of the Milky Way in Table 1 of \citet{Chang2011}. We relist these parameters in Table~\ref{table_ref} to conveniently find their improvements. Previous works show the estimation of the parameters suffer from the degeneracy \citep{Robin00,Chen01,Siegel02,Phleps05,Juric08,Bilir08,Chang2011}, which could influence the confidence of the final results, but it is unconvinced to take it as a main reason to explain the spread of values in Table~\ref{table_ref}. We should notice the sample stars used in researchers' works are different. These parameters depend on both the properties and locations of the sample stars, such as absolute magnitude \citep{Bilir06,Bilir2006,Karaali07}, Galactic longitude \citep{Du06,Cabrera-Lavers07,AK2007, Bilir08,Yaz10,Chang2011} and Galactic latitude \citep{Du06, AK07}, thus these parameters that they derived could different from each other. The explanation of these dependence is still a topic of debate. The issue is not to give a explanation that could interpret the observation well, but rather to test whether the theory of the Galaxy is reasonable and therefore to answer a given question related to the Galactic formation and evolution. Recent researches give us more information about the structure of the Galaxy: the Galactic structure is not as smooth as we thought. Many more substructures have been discovered, such as Sagittarius dwarf tidal stream \citep[e.g.][]{Majewski2003}, Monoceros stream \citep[e.g.][]{Newberg2002}, and Virgo overdensity \citep{Juric08}. The researches show the Milky Way is a complex and dynamic structure that is still being shaped by the merging of neighboring smaller galaxies, and the dynamic process may plays a crucial role in establishing the galactic structure. These are reviewed in detail by \citet{Ivezic12} and references therein. Needless to say, the presence of irregular structure make the research about Galactic structure more complex. We are interested in the global smooth structure of the Galaxy. It can help us to define the irregular structure, and then allow us to study the formation and evolution of our Galaxy. In this paper, based on SDSS and the South Galactic Cap u-band Sky Survey early data (SCUSS), we attempt to study the structure parameters' possible variations with absolute magnitude. We present density function in Sect. 2. The introduction of SDSS and SCUSS and the reduction of the data would be described in Sect. 3. Sect. 4 provide the method we used to determine the parameters. Finally, our main results are discussed and summarized in Sect. 5. \begin{table*} \begin{center} \begin{tabular}{lllllllll} \hline $H_{z1}$ & $H_{r1}$ & $f_2$ & $H_{z2}$ & $H_{r2}$ & $f_h$ & Re(S) & $\kappa$ & Reference \\ (pc) & (kpc) & & (kpc) & (kpc) & & (kpc) & & \\ \hline 310-325 & - & 0.0125-0.025 & 1.92-2.39 & - & - & - &- & Yoshii (1982) \\ 300 & - & 0.02 & 1.45 & - & - & - & - & Gilmore $\&$ Reid (1983) \\ 325 & - & 0.02 & 1.3 & - & 0.002 & 3 & 0.85 & Gilmore (1984) \\ 280 & - & 0.0028 & 1.9 & - & 0.0012 & - & - & Tritton $\&$ Morton (1984) \\ 125-475 & - & 0.016 & 1.18 - 2.21 & - & 0.0013 & 3.1* & 0.8 & Robin $\&$ Creze (1986) \\ 300 & - & 0.02 & 1 & - & 0.001 & - & 0.85 & del Rio $\&$ Fenkart (1987) \\ 285 & - & 0.015 & 1.3 - 1.5 & - & 0.002 & 2.36 & Flat & Fenkart $\&$ Karaali (1987) \\ 325 & - & 0.0224 & 0.95 & - & 0.001 & 2.9 & 0.9 & Yoshii et al. (1987) \\ 249 & - & 0.041 & 1 & - & 0.002 & 3 & 0.85 & Kuijken $\&$ Gilmore (1989) \\ 350 & 3.8 & 0.019 & 0.9 & 3.8 & 0.0011 & 2.7 & 0.84 & Yamagata $\&$ Yoshii (1992) \\ 290 & - & - & 0.86 & - & - & 4 & - & von Hippel $\&$ Bothun (1993) \\ 325 & - & 0.020-0.025 & 1.6-1.4 & - & 0.0015 & 2.67 & 0.8 & Reid $\&$ Majewski (1993) \\ 325 & 3.2 & 0.019 & 0.98 & 4.3 & 0.0024 & 3.3 & 0.48 & Larsen (1996) \\ 250-270 & 2.5 & 0.056 & 0.76 & 2.8 & 0.0015 & 2.44 - 2.75* & 0.60 - 0.85 & Robin et al. (1996, 2000) \\ 260 & 2.3 & 0.74 & 0.76 & 3 & - & - & - & Ojha et al. (1996) \\ 290 & 4 & 0.059 & 0.91 & 3 & 0.0005 & 2.69 & 0.84 & Buser et al. (1998, 1999) \\ 240 & - & 0.061 & 0.79 & - & - & - & - & Ojha et al. (1999) \\ 280/267 & - & 0.02 & 1.26/1.29 & - & - & 2.99* & 0.63 & Phleps et al. (2000) \\ 330 & 2.25 & 0.065 - 0.13 & 0.58 - 0.75 & 3.5 & 0.0013 & - & 0.55 & Chen et al. (2001) \\ - & 2.8 & 3.5 & 0.86 & 3.7 & - & - & - & Ojha (2001) \\ 280(350) & 2 - 2.5 & 0.06 - 0.10 & 0.7 - 1.0 (0.9 - 1.2) & 3 - 4 & 0.0015 & - & 0.50 - 0.70 & Siegel et al. (2002) \\ 285 & 1.97 & - & - & - & - & - & - & Lopez-Corredoira et al. (2002) \\ - & 3.5 & 0.02-0.03 & 0.9 & 4.7 & 0.002-0.003 & 4.3 & 0.5-0.6 & Larsen $\&$ Humphreys (2003) \\ 320 & - & 0.07 & 0.64 & - & 0.00125 & - & 0.6 & Du et al. (2003) \\ 265-495 & - & 0.052-0.098 & 0.805-0.970 & - & 0.0002-0.0015 & - & 0.6-0.8 & Karaali et al. (2004) \\ 268 & 2.1 & 0.11 & 1.06 & 3.04 & - & - & - & Cabrera-Lavers et al. (2005) \\ 300 & - & 0.04-0.10 & 0.9 & - & - & 3/2.5* & 1/0.6 & Phleps et al. (2005) \\ 220 & 1.9 & - & - & - & - & - & - & Bilir et al. (2006b) \\ 160-360 & - & 0.033-0.076 & 0.84-0.87 & - & 0.0004-0.0006 & - & 0.06-0.08 & Bilir et al. (2006a) \\ 301/259 & - & 0.087/0.055 & 0.58/0.93 & - & 0.001 & - & 0.74 & Bilir et al. (2006c) \\ 220-320 & - & 0.01-0.07 & 0.6-1.1 & - & 0.00125 & - & $>$0.4 & \citet{Du06} \\ 206/198 & - & 0.16/0.10 & 0.49/0.58 & - & - & - & 0.45 & Ak et al. (2007a) \\ 140-269 & - & 0.062-0.145 & 0.80-1.16 & - & - & - & - & Cabrera-Lavers et al. (2007) \\ 220-360 & 1.65-2.52 & 0.027-0.099 & 0.62-1.03 & 2.3-4.0 & 0.0001-0.0022 & - & 0.25-0.85 & Karaali et al. (2007) \\ 167-200 & - & 0.055-0.151 & 0.55-0.72 & - & 0.0007-0.0019 & - & 0.53-0.76 & Bilir et al. (2008) \\ 245(300) & 2.15(2.6) & 0.13(0.12) & 0.743(0.900) & 3.261(3.600) & 0.0051 & 2.77* & 0.64 & \citet{Juric08} \\ 325-369 & 1.00-1.68 & 0.0640-0.0659 & 0.860-0.952 & 2.65-5.49 & 0.0033-0.0039 & - & 0.489-0.654 & Yaz $\&$ Karaali (2010) \\ 360 & 3.7 & 0.07 & 1.02 & 5 & 0.002 & 2.6* & 0.55 & Chang et al. (2011) \\ 103-350 & - & 0.083-0.165 & 0.525-0.675 & - & 0.0005-0.0065 & - & 0.20-0.84 & this work \\ \hline \end{tabular} \end{center} \caption{Galactic model parameters tabulated base on the Table 1 in \citet{Chang2011}. $H_z$ and $H_r$ mean scaleheight and scalelengh, respectively. And suffix 1 and 2 denote thin disk and thick disk, repectively. $f_2$ and $f_h$ are the local stellar density ratio of the thick-to-thin disk and halo-to-thin disk, respectively. The parentheses are the corrected values for binarism. The asterisk denotes the power-law index replacing Re, which is commonly known as the de Vaucouleurs radius, and $\kappa$ is axial ratio. } \label{table_ref} \end{table*} | We estimate the Galactic model parameters using $\chi^2$ method in absolute magnitude intervals: $4 \leq M_r < 5$, $5 \leq M_r < 6$, $6 \leq M_r < 8$, $8 \leq M_r < 10$, $10 \leq M_r \leq 13$, with a unique density law for each population individually of 7.07 $deg^2$ field locate at $50 ^\circ \leq l \leq 55^\circ$, $-46 ^\circ \leq b\leq -44^\circ $, to explore their possible variation with absolute magnitude from SDSS and SCUSS photometry. We also estimate the parameters in absolute magnitude interval $4 \leq M_r \leq 13$ in order to compare with other researchers' works. Our results show the parameters are absolute magnitude dependent (see Table~\ref{tab2}). This is a possible reason to explain why different researchers obtained different parameters. \subsection{Thin disk} The local density of the thin disk increases and the scaleheight decreases, with the sample stars get fainter. The local density $n_1(R_\circleddot,z_\circleddot)$ varies from $1.25\times10^6$ $star/kpc^3$ in absolute magnitude interval $5 \leq M_r < 6$ to $2.32\times10^7$ $star/kpc^3$ in absolute magnitude interval $10 \leq M_r \leq 13$, and the scaleheight changes from 350 $pc$ to 103 $pc$ in the corresponding absolute magnitude interval. The results reveal a phenomenon for thin disk: the intrinsic faint MS stars tend to stay closer to the Galactic plane, and have larger local density. In order to compare with other researchers' results, we derive the scaleheight of the thin disk in absolute magnitude interval $4 \leq M_r \leq 13$, that is 205 $pc$, which is a little smaller than most previous work (see Table~\ref{table_ref}). As these parameters depend on absolute magnitude, we can conclude that the small scaleheight of thin disk in absolute magnitude interval $4 \leq M_r \leq 13$ is caused by the contribution of the faint MS stars in our sample. It is obvious that different types of MS stars could have experienced different dynamic process, thus could have different specific distributions. In other words, the specific distribution of MS stars could dependent on absolute magnitude. The quantitative calculation to explain this phenomenon may need the chemical, kinematic and dynamic information, the star formation history and star formation rate may also be included. Here we just give a possible qualitative explanation: bright MS stars have large mass and short lifetime which lead the local density to be relatively smaller than that of the faint one , thus the gravitational interaction time they experienced are short, this may make the bright MS stars stay a little further in mean than the fainter one, so the scaleheight $h_1$ of the bright MS stars would be larger. \subsection{Thick disk} The final results of the thick disk suffer from the degeneracy between the two halo parameters: $\kappa$ and $n_3$, see Figure~\ref{fig_halocontour}. It raise the uncertainty of the relationship between thick disk parameters and absolute magnitude. \textbf{As our local density of thin disk $n_1$ changes with absolute magnitude, it is significative to compare the value of $n_2$ with other works rather than $n_2/n_1$. The values of $h_2$ we obtained are a little smaller than previous works, thus the value of $n_2$ (normalized at $(R_\circleddot,z_\circleddot)$) should be a little larger. As shown in Table~\ref{tab2}, $h_2$ roughly fall into the range of $520 < h_2 (pc) < 680$, simultaneously, $n_2$ roughly fall into the range of $2.0\times10^5 < n_2 (star/kpc^3) < 1.4\times10^6$. The results show the parameters of thick disk are non-monotonic change with absolute magnitude, which may imply the origin of the thick disk is complicated.} \subsection{Halo} As shown in Table~\ref{tab2}, the halo parameters $\kappa$ and $n_3$ are monotonic change with absolute magnitude. The axial ratio $\kappa$ varies from 0.84 in bright interval to 0.20 in faint interval which indicates a flattened inner halo and a more spherical halo. The trend of $n_3$ shows intrinsic faint MS stars have large local density. In general, the values of $\chi^2$ for halo-fitting are larger than disk-fitting in all intervals. Figure~\ref{fig_plot} can reflect the large deviation of halo-fitting. The deviation may come from the influence of the uncompleted data, especially at large distance. As seen in Figure~\ref{fig_halocontour}, the degeneracy exists between the halo parameters $n_3$ and $\kappa$. So the confidence of the halo parameters monotonic change with absolute magnitude will be influenced. We do not know what cause the halo parameters degenerate, but the most likely reason is the unsatisfactory of the halo component fitting: whether the observation is not good enough (e.g. the incomplete data in halo component) or the halo model is not precisely describe the Galaxy. In summary, we fit the observations with three components Galactic model to estimate the Galactic structure parameters to explore their possible variations with absolute magnitude. Our results show the parameters of the thin disk and halo are monotonic change with absolute magnitude, but the thick disk is non-monotonic. The explanation of this phenomenon need more work. Also, we still need a new method to estimate the Galactic structure parameters in order to break the degeneracy. As greatly improve on data collection and more and more researchers work on these field in recent years, we believe these questions can be solved in the near future. We still have a long way to know our galaxy well. | 14 | 3 | 1403.3490 |
1403 | 1403.3804_arXiv.txt | The statistics of black holes and their masses strongly suggests that their mass distribution has a cutoff towards lower masses near $3 \times 10^{6}$ M$_{\odot}$. This is consistent with a classical formation mechanism from the agglomeration of the first massive stars in the universe. However, when the masses of the stars approach $10^{6}$ M$_{\odot}$, the stars become unstable and collapse, possibly forming the first generation of cosmological black holes. Here we speculate that the claimed detection of an isotropic radio background may constitute evidence of the formation of these first supermassive black holes, since their data are compatible in spectrum and intensity with synchrotron emission from the remnants. The model proposed fulfills all observational conditions for the background, in terms of single-source strength, number of sources, far-infrared and gamma-ray emission. The observed high energy neutrino flux is consistent with our calculations in flux and spectrum. The proposal described in this paper may also explain the early formation and growth of massive bulge-less disk galaxies as derived from the massive, gaseous shell formed during the explosion prior to the formation of a supermassive black hole. | Black holes are ubiquitous in the centers of early Hubble type galaxies, and their mass distribution shows evidence for a low mass cut-off near $3 \cdot 10^{6} \, M_{\odot}$ (Greene \etal 2006, 2008; Greene \& Ho 2007a, b; Caramete \& Biermann 2010). Massive stars readily turn into black holes (Heger \etal 2003, 2005; Woosley \etal 2002), and the agglomeration of massive stars (Spitzer 1969; Sanders 1970; Quinlan \& Shapiro 1990, Portegies Zwart \etal 2004, 2007, 2010; McMillan \etal 2007), perhaps aided by a gravo-thermal collapse (Spitzer 1969, 1987), can turn them into yet more massive stars. However, their powerful winds counteract the increase in mass, and the maximal mass which can be reached from merging is limited to a few hundred solar masses (Yungelson \etal 2008, Crowther \etal 2010). On the other hand, the wind is driven by radiation interacting with metal ions, thus for zero metal stars there is no such wind, and no mass loss (Heger \etal 2003, Woosley \etal 2002). It follows that massive, zero-metal stars can indeed reach very high masses in agglomeration. As the instability refers to that part of the star which is in hydrostatic equilibrium, and the further outer layers of such a rapidly growing supermassive star may still be relaxing towards hydrostatic equilibrium, the fall-back may increase the mass of the final black hole beyond the mass of the hydrostatic mass fraction of the star itself. In such a way the initial black hole mass can possibly become larger than the stellar instability threshold. An alternate picture is that of the growth of a budding super-massive star by direct accretion of collapsing material (e.g. (e.g., Begelman \etal 2006, Bonoli \etal 2012). The formation of this super-massive star and then a black hole may be aided by the potential well of a surrounding dark matter clump (see, e.g., Munyaneza \& Biermann 2005, 2006 % ; Destri \etal 2012 % ). Massive stars are dominated by radiation pressure, and with increasing zero-age main sequence mass their effective adiabatic gas index approaches the unstable limit of 4/3 (Chandrasekhar 1939). Subtle effects of general relativity push the stars over the limit, causing stars to blow up at a mass approaching $10^{6} \, M_{\odot}$ (Appenzeller \& Fricke 1972a, b). It follows that the agglomeration of zero-metal stars can readily produce super-massive black holes. % It is possible in some cosmological models (e.g., Biermann \& Kusenko (2006)) for star formation to occur early ($z \ge 20$), paving the way to an early black hole formation. Whalen \etal (2013) have recently discussed the possibility of supernovae explosions of massive population III stars which could lead to the formation of the first generation of black holes. In such cases, massive black holes can grow rather rapidly, by merging with other black holes or by accretion (e.g., Wang \& Biermann 1998, Silk \& Rees 1998, Gergely \& Biermann 2009). In this way the existence of extremely massive black holes at high redshift can be understood (e.g. Mortlock \etal 2011). The formation of super-massive black holes at very high redshift may have consequences in galaxy evolution and cosmology (e.g. Kormendy \etal 2010, 2011, Kormendy \& Bender 2011, Conselice \etal 2011, Biermann \& Harms 2012, 2013a, b, Buitrago \etal 2013, Conselice \etal 2013). This line of reasoning suggests that it is worth exploring % possible observational tests in order to confirm or refute such a picture. The recent claim of the detection of an isotropic radio background unrelated to any known population of galaxies (Kogut \etal 2011, Fixsen \etal 2011, Seiffert \etal 2011, Condon \etal 2012, Formengo \etal 2014; but see also Subrahmanyan \& Cowsik 2013)) raises the possibility that the explosions of these super-massive stars which give rise to the first generation of super-massive black holes might be considered as being similar to a supernova explosion with the concomitant acceleration of particles producing radio emission. We work out the strength of this radio emission. We do not question here whether this is the only possible explanation, as clearly the thick cosmic ray disk in our own Galaxy can also produce strong background emission (Sun \etal 2008; Everett \etal 2010), but it seems unlikely to be able to give the spectrum derived by Fixsen \etal (2011). However, the strength of the spectral constraint depends on the error estimate. As a test of our model, we consider the conditions derived by Condon \etal (2012), and show that their conditions can all be met; these include the strength of each source, the number of sources, and the absence or weakness of far-infrared emission. The flux density determined for the background, which was not explained by known source populations by Condon \etal (2012), can be determined using the spectrum obtained by Fixsen \etal (2011). The observed flux density at 3.02 GHz corresponds at 1 GHz to $10^{-18.5} \, {\rm erg \, cm^{-2} \, s^{-1} \, Hz^{-1} \, sr^{-1}}$ with an error of about 16 percent. Earlier attempts to interpret this radio background were done by Singal et al. (2010), Meiksin \& Whalen (2013), and Holder (2014). They noted some of the same difficulties emphasized by Condon et al. (2012). | We have derived the radio background due to the formation of the first population of super-massive black holes. Their production is assumed to lead to radio remnants, quite similar to normal supernova remnants, just scaled up. The prediction presented here falls within the IceCube sensitivity (Schukraft \etal 2013) and therefore provides a possible explanation of the recently announced IceCube neutrino excess (Aartsen \etal 2013). The model also leads to a possible explanation of the observed flux density and spectrum of the gamma-ray background. The model obeys all known radio observational constraints, including single source strength, total number, lack of far-infrared emission, and radio spectrum. Adopting cautious values for the parameters of the model suggests that the formation redshift may be quite large, consistent with % a very early epoch of star formation. Matching both the neutrino and radio background gives rather strong constraints on the factor $\eta_{\rm B} ^{0.8} \eta_{\rm CR,e}/\eta_{\rm CR}$, since the ratio of these two backgrounds depends only weakly on the explosion energy and redshift. If upon further exploration the radio background is lowered in its flux density (e.g., see Subrahmanyan \& Cowsik 2013), then the explanation of the neutrino background stands, and this model would then predict a radio background at a level depending on this factor, relative to the neutrino background. Interestingly, the scenario also has the potential to solve the formation riddle of large massive galaxies, which to all appearances never merged (see Kormendy \etal 2010). The massive shells formed by the explosion of the super-massive stars give the right order of magnitude both for the mass and the space density to form such galaxies. A direct consequence is that a large number of galaxies never merged. The redshift in this scenario could possibly be determined using the absorption spectra of hydrogen molecules, H$_{2}$, HD$^{+}$, H$_{2}^{+}$ or H$_{3}^{+}$, although the observations could be challenging. If the interpretation can be confirmed, it would demonstrate the formation of the first super-massive black holes in the universe. The radio emission is non-thermal and together with the recent detection of a high energy neutrino background constitutes evidence for the first cosmic ray population in the universe. | 14 | 3 | 1403.3804 |
1403 | 1403.5755_arXiv.txt | The absolute--magnitude distributions of seven supernova types are presented. The data used here were primarily taken from the Asiago Supernova Catalogue, but were supplemented with additional data. We accounted for both foreground and host--galaxy extinction. A bootstrap method is used to correct the samples for Malmquist bias. Separately, we generate volume--limited samples, restricted to events within 100 Mpc. We find that the superluminous events (M$_B < -21$) make up only about 0.1\% of all supernovae in the bias--corrected sample. The subluminous events (M$_B > -15$) make up about 3\%. The normal Ia distribution was the brightest with a mean absolute blue magnitude of $-$19.25. The IIP distribution was the dimmest at $-$16.75. | About two decades ago, a study was carried out on the absolute--magnitude distributions of supernovae (SNe) separated by types \citep[][hereafter MB90]{miller90}. This study was based on data taken from the Asiago Supernova Catalogue \citep[the ASC;][]{barbon89}, which at the time had 687 SNe. A decade later \citet[][hereafter R02]{richardson02} carried out a similar study; however, by that time the ASC had grown to 1910 SNe. Currently, the ASC has over 6100 SNe. A histogram of the number of supernovae versus the year of discovery is shown in Figure 1. In this graph we can see a drastic increase in the discovery rate of SNe over the last 20 years. Also shown in this figure are labels for important milestones in astronomy that had an effect on the discovery rate. The years for the studies MB90 and R02 are shown for context. The inset shows the same data on a log scale. Here we plot log of N+1 so that years with only one SN discovery can be shown. The large spike in the number of SNe discovered over the last decade is primarily due to the number of large searches that have taken place during that time. Many of these searches were motivated by the discovery of the accelerating universe \citep{riess98,perlmutter99}. Because of this drastic increase in the overall number of SNe, it is time to update the absolute--magnitude distributions for the various SN types. Absolute--magnitude distributions are important in determining SN rates and providing information useful for constraining progenitor and explosion models. They are also useful for the planning of future SN searches. The main push for the large SN searches is the discovery of type Ia SNe, which are used as cosmological distance indicators. As a by--product of this, more SNe of all types are being discovered, although follow up is not necessarily done. Many studies have looked at SNe~Ia in great detail and through some form of light--curve standardization have reduced the dispersion to about 0.12 \citep{bailey09}. This type of standardization is not possible with our large data set and is not carried out here. Most of the data used in this study are taken from the ASC. However, a number of other sources are used as well. The data used in this study, and their sources, are discussed in Section~2. The different types of analysis are discussed in Section~3. The results of our analysis are presented in Section~4 along with comparisons to other studies. A summary is given in Section~5. | In this study, we collected basic data for as many SNe as possible, mainly from the Asiago Supernova Catalogue. From that basic data, we calculated absolute magnitudes for each SN in a self--consistent manner. Most of the apparent magnitudes were reported in the B filter. If B--filter peak magnitudes were not available, we approximated the K--corrections using spectra from SNe of the same type. We also estimated the host--galaxy extinctions for as many individual SNe as possible and used a model to assign values for the rest. In order to account for Malmquist bias, we used a bootstrap method to generate bias--corrected distributions for seven SN types. These distributions were then compared to volume--limited distributions. We found that the distribution with the brightest peak mean absolute magnitude (Ia) was 0.72 magnitudes brighter than the next brightest (IIn). The SNe Ib and Ic have similar mean absolute magnitudes (M$_B = -$17.45, $-$17.66 respectively). Their average is roughly 1.7 magnitudes dimmer than that of the SNe~Ia. Type IIb SNe were roughly half a magnitude dimmer than the other stripped--envelope SNe (Ib \& Ic). The dimmest type in the study is the SN~IIP distribution which is 1.23 magnitudes dimmer than the SNe~IIL distribution. The mean absolute magnitude for the dimmest distribution (IIP) is 2.5 magnitudes dimmer than that for the the brightest (Ia). We found that volume--limited samples (limited to $\mu < 35$) produce good approximations of the unbiased absolute--magnitude distributions for all SN types. For most of the distributions, the bias--correction process didn't go far beyond $\mu = 35$ and so we would expect them to be similar. However, the SN~Ia and IIn samples went out to $\mu = 38$ and 39 respectively. The SN~Ia distributions were nearly identical and the SN~IIn distributions differed by only 0.09 mag in mean and by only 0.12 in standard deviation. As expected, the SNe~Ia have the smallest standard deviation and are highly concentrated within the $2\sigma$ limits. This distribution also has the smallest fraction of extreme SNe (2.6\%), not counting the small sample of SN~IIb with 0\%. When comparing the results of the current distributions to those of R02, we found that the two biggest differences were for the SNe~IIL and IIn distributions. The large number of superluminous SNe~IIn discovered during the last decade played a role in this difference. We should also note that the SN~Ia distribution is somewhat brighter than that of R02 and this is most likely due to the fact that most of the dimmest SNe~Ia suffered from significant host--galaxy extinction. Once that was taken into account, the mean absolute--magnitude became brighter. We also compared our results to those of the Lick Observatory Supernova Search (LOSS) as reported in L11. After adjusting their values for a different H$_0$, different filter (R to B) and making an adjustment for host--galaxy extinction, we found that most of their distributions are relatively close to ours (within the, sometimes considerable, error bars of the two distributions). However, there is one noticeable exception, the SN IIn distribution. As mentioned above, our sample is considerably larger (containing all of the LOSS sample) and simply has more bright SNe~IIn than the LOSS sample alone. There are nine superluminous SNe (M$_B < -21$) when considering all types. All of these were discovered during the last decade and six of them are SNe~IIn. From the bias--corrected sample representing SNe of any type, we see that superluminous SNe make up only about 0.1\% of all SNe. There are seven SNe in this study that are considered to be subluminous (M$_B > -15$). We see that approximately 3\% of all SNe in the bias--corrected sample are subluminous. We would like to thank David Branch for his helpful suggestions. This work was partially supported by the National Science Foundation under Grant No. CHE-1005026. This work was also supported by the U.S. Nuclear Regulatory Commission under Grant No. 23N811. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. | 14 | 3 | 1403.5755 |
1403 | 1403.5563_arXiv.txt | We investigate whether the satellite luminosity function (LF) of primary galaxies identified in the Sloan Digital Sky Survey (SDSS) depends on whether the host galaxy is in a filament or not. Isolated primary galaxies are identified in the SDSS spectroscopic sample while potential satellites (that are up to 4 magnitudes fainter than their hosts) are searched for in the much deeper photometric sample. Filaments are constructed from the galaxy distribution by the ``Bisous'' process. Isolated primary galaxies are divided into two subsamples: those in filaments and those not in filaments. We examine the stacked mean satellite LF of both the filament and non-filament sample and find that, on average, the satellite LFs of galaxies in filaments is significantly higher than those of galaxies not in filaments. The filamentary environment can increases the abundance of the brightest satellites ($M_\mathrm{sat.} < M_\mathrm{prim.} + 2.0$), by a factor of $\sim 2$ compared with non-filament isolated galaxies. This result is independent of primary galaxy magnitude although the satellite LF of galaxies in the faintest magnitude bin, is too noisy to determine if such a dependence exists. Since our filaments are extracted from a spectroscopic flux-limited sample, we consider the possibility that the difference in satellite LF is due to a redshift, colour or environmental bias, finding these to be insufficient to explain our result. The dependence of the satellite LF on the cosmic web suggests that the filamentary environment may have a strong effect on the efficiency of galaxy formation. | The \lcdm~model predicts that structure forms in a hierarchical manner. The first objects to collapse and virialize at high redshift are small dark matter haloes that later merge to form larger objects. Small haloes that host satellite or dwarf galaxies, can often survive the violent process associated with halo mergers for many Giga-years providing important information about galaxy formation, the population of subhaloes, and even the nature of dark matter. Moreover these dark matter structures form an intricate pattern on the megaparsec scale, known as ``cosmic web'' \citep{bon96}, consisting of regions termed voids, filaments, sheets and knots. The cosmic web is a direct consequence of the gravitational instabilities that emerge out of the primordial density field. The presence of such cosmic pattern has been confirmed observationally by the distribution of the galaxies from the large surveys such as the 2dF Galaxy Redshift Survey \citep[2dFGRS;][]{col01, col03}, the Sloan Digital Sky Survey \citep[SDSS;][]{yok00,teg04}, the Two Micron All Sky Survey \citep[2MASS;][]{huc05}, Galaxy And Mass Assembly \citep[GAMA;][]{alp04} and the CosmicFlow-2 survey of peculiar velocities \citep{Tully:14}. A number of studies have examined how the cosmic web can affect specific halo properties such as abundance, shape, or assembly history \citep{ara07,hah07a,hah07b,lib12,lib13b,cau13}. Correlations have also been found between halo spin and the principle axis of filaments (their spine) and walls (their normal) that they are embedded in \citep{alt06, ara07, hah07a, hah07b, zha09, lib13, ara13, dub14}. Although more difficult to quantify owing to degeneracies and inherent biases, similar studies have been conducted in observational samples \citep{jon10,tem13a,tem13b,zha13}. At $z=3.1$ \cite{mats04} found that the spatial distribution of galaxies within Ly$\alpha$ haloes trace the underlying large-scale filamentary structure of the universe. The morphology and spatial extent of Ly$\alpha$ haloes depend on the environment as well \citep{mats11, mats12}. Tying galaxy or halo properties to the environment is not a new idea and many studies have quantified the importance of such relations for galaxy formation and evolution \citep[e.g.][]{dre80, kau04, bla05b, Tempel:11}. On the larger scales of the cosmic web, the properties of galaxies depend on the cosmic filament it inhabits \citep[e.g.][]{mur11,jon10} or on the supercluster environment \citep[e.g.][]{Lietzen:12,Einasto:14}. However, to date no observational studies have examined the effect of the cosmic web environment on satellite galaxies. Analyzing satellite systems of external galaxies is challenging, because typically only several satellites are detected per primary galaxy. Furthermore the real space position of a satellite with respect to its primary is uncertain. Owing to the advent of large galaxy surveys, a statistically robust estimate of the satellite luminosity function (LF) has become possible \citep[e.g.][]{my11,tal11,wang12}. In this study, we investigate if the satellite LF depends on whether the host galaxy is located within a filament. We divide isolated primary galaxies into two categories: those within galaxy filaments defined by the ``Bisous process'' as in \cite{tem14} and those not in these filaments. Each subsample is then divided into three bins according to the primary's magnitude. The resulting subsamples enable us to study the satellite LFs of isolated primary galaxies in filaments and compare it with isolated primary galaxies that are not in filaments. Such an analysis will quantify how the filament environment affects galaxy formation. Throughout the paper we assume a fiducial ${\rm\Lambda CDM}$ cosmological model with $\Omega_{\rm M}=0.3$, $\Omega_{\Lambda}=0.7$ and $H_0 = 70\ $km~s$^{-1}$Mpc$^{-1}$. | We have examined the luminosity function (LF) of satellites close to isolated primary galaxies in the SDSS. Isolated primary galaxies have been split by filamentary environment into two camps: those primaries in and not in filaments. Background galaxies are subtracted statistically and the mean satellite LF is computed by stacking all centrals of a given absolute $r$ band magnitude. Our results indicate that primaries in filaments have more satellites than those that are not found in filamentary environments. This is most evident for the brightest satellites but is true for satellites up to 4 magnitudes fainter of their host. Except for the faintest magnitude bin (where the signal-to-noise is too low to judge), in-filament primaries have a factor of $\sim 1.5$--$2$ more bright satellites than primaries not in filaments. This slightly varies with redshift possibly because the contamination in the sample of galaxies not in filaments increases with the redshift. We have examined if the difference in the satellite LF is due to other controlling factors including differences in colour, density, redshift distribution or variance of small sample sizes. None of these control factors can explain the result found here. The differences exhibited are statistically significant and robust and thus reflect an inherent difference in the satellite population of galaxies in and not in filaments. We have also examined the mean radial distribution of satellites brighter than a $r$-band magnitude of $-20$ in Appendix~\ref{appendix:profile} and found no significant difference in the projected spatial distribution of satellites in or not in filaments in spite of the obvious difference between the total numbers of satellites. \begin{figure} \plotone{fig8.eps} \caption{The mean satellite LFs of in-filaments primaries (solid lines) and a control sample of ``twins'' chosen from the not-in-filaments sample (dashed lines) according to Equation~2, color coded by magnitude. The ratio of the satellite LFs of in-filaments and not-in-filaments are shown in the small lower panel.} \label{fig:lf_control} \end{figure} Previous numerical studies have shown that the abundance of dark matter subhaloes may depend on the environment: haloes in filaments may have more subhaloes than those in other cosmic web environments \citep[Cautun et al. in preparation;][]{ney14}. The properties of subhaloes in filaments can be also different from those in the other environments and this may affect galaxy formation since the satellite LF is a result of the gas physics that regulate star formation in small haloes. Our results thus suggest that galaxy formation itself may be more efficient in subhaloes that are born in and accreted through filaments. The filamentary environment may be a crucial component of galaxy formation on such small scales. | 14 | 3 | 1403.5563 |
1403 | 1403.6169_arXiv.txt | We consider a low-dimensional model of convection in a horizontally magnetized layer of a viscous fluid heated from below. We analyze in detail the stability of hydromagnetic convection for a wide range of two control parameters. Namely, when changing the initially applied temperature difference or magnetic field strength, one can see transitions from regular to irregular long-term behavior of the system, switching between chaotic, periodic, and equilibrium asymptotic solutions. It is worth noting that owing to the induced magnetic field a transition to hyperchaotic dynamics is possible for some parameters of the model. We also reveal new features of the generalized Lorenz model, including both type I and III intermittency. | 14 | 3 | 1403.6169 |
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1403 | 1403.6675_arXiv.txt | The nuclear radio emission of low-luminosity active galactic nuclei (LLAGN) is often associated with unresolved cores. In this paper we show that most LLAGN present extended jet radio emission when observed with sufficient angular resolution and sensitivity. They are thus able to power, at least, parsec-scale radio jets. To increase the detection rate of jets in LLAGN, we analyze subarcsecond resolution data of three low-ionization nuclear emission regions. This yields the detection of extended jet-like radio structures in NGC\,1097 and NGC\,2911 and the first resolved parsec-scale jet of NGC\,4594 (Sombrero). The three sources belong to a sample of nearby LLAGN for which high-spatial-resolution spectral energy distribution of their core emission is available. This allows us to study their accretion rate and jet power ($Q_\mathrm{jet}$) without drawing on (most) of the ad hoc assumptions usually considered in large statistical surveys. We find that those LLAGN with large-scale radio jets ($>$100 pc) have $Q_\mathrm{jet} > 10^{42}$ erg s$^{-1}$, while the lowest $Q_\mathrm{jet}$ correspond to those LLAGN with parsec-scale ($\leq$100 pc) jets. The $Q_\mathrm{jet}$ is at least as large as the radiated bolometric luminosity for all LLAGN, in agreement with previous statistical studies. Our detection of parsec-scale jets in individual objects further shows that the kinematic jet contribution is equally important in large- or parsec-scale objects. We also find that the Eddington-scaled accretion rate is still highly sub-Eddington ($<10^{-4}$) when adding the $Q_\mathrm{jet}$ to the total emitted luminosity (radiated plus kinetic). This indicates that LLAGN are not only inefficient radiators but that they also accrete inefficiently or are very efficient advectors. | \label{intro} The presence of an accreting engine in the active galactic nuclei (AGN) with the lowest luminosities (LLAGN; $L_\mathrm{bol}\leq10^{42}$ erg s$^{-1}$) is a matter of debate (see review by \citealt{2008ARA&A..46..475H}). LLAGN are found at the center of nearly one-third of all local galaxies and encompass low-luminosity Seyferts, low-ionization nuclear emission regions (LINERs; \citealt{1980A&A....87..152H}), and transition nuclei (which have spectral properties intermediate between those of LINERs and HII regions). LINERs are characterized by very strong low-ionization optical forbidden lines whose source of ionization has been debated for decades. The main scenarios include a low-luminosity version of bright AGN, shock heating (e.g., \citealt{1995ApJ...455..468D}), or stellar photoionization (e.g., \citealt{1985MNRAS.213..841T}; \citealt{1992ApJ...397L..79F}). The spectral energy distribution (SED) of LLAGN is very prominent in the radio band and lacks the canonical optical-UV ``big blue bump'' of AGN (indicative of the presence of an optically thick, geometrically thin, accretion disk). A combination of an optically thin advection-dominated accretion flow (ADAF; \citealt{1994ApJ...428L..13N}; see \citealt{2014arXiv1401.0586Y} for a review), an outer truncated disk (\citealt{1999ApJ...525L..89Q}), and a radio jet (\citealt{1995A&A...293..665F,1999A&A...342...49F}; \citealt{1996ApJ...464L..67F}) was proposed to model the SED of LLAGN (e.g., \citealt{1999ApJ...516..672H}; \citealt{2001ApJ...562L.133U}; \citealt{2006ApJ...643..652N,2012Sci...338.1445N,2014MNRAS.438.2804N}; \citealt{2010ApJS..187..135E}; \citealt{2011ApJ...726...87Y}; \citealt{2012JPhCS.372a2006F}; although see also \citealt{2013ApJ...777..164M} and references therein). However, the ADAF model alone is not able to satisfactorily reproduce the observed radio luminosities of LLAGN (e.g., \citealt{2003ApJ...588..175F}; \citealt{2005MNRAS.360..119D}a; \citealt{2005ApJ...621..130W}; \citealt{2007ApJ...669...96W}; \citealt{2008ApJ...681..905M}) and predicts an inverted radio spectral index ($\alpha\sim$0.4\footnote{The spectral index $\alpha$ is defined from $S_{\nu} \propto \nu^{\alpha}$, where \textit{S} is the flux density at frequency $\nu$.}; \citealt{1998ApJ...499..198Y}) much higher than the one typically observed in LLAGN spectra ($\alpha\sim -0.2$ to 0.2; e.g., \citealt{2000ApJ...542..197F}; \citealt{2001ApJ...559L..87N}; \citealt{2005MNRAS.363..692D}b; \citealt{2011AJ....142..167D}). Although this does not rule out the ADAF model, it argues for the necessity of a parsec-scale (pc-scale) radio jet to model the SEDs of these sources. Despite the results showing that the SED of LLAGN is compatible with jet-dominated emission, observational evidence of radio jets in these sources still remains elusive. Several studies have aimed at detecting jet radio emission in LLAGN, finding that they are usually associated with compact radio sources where most of the radio emission is highly concentrated (e.g., \citealt{2000ApJ...542..197F}; \citealt{2000ApJS..129...93F}; \citealt{2001ApJ...562L.133U}; \citealt{2002A&A...385..425F}; \citealt{2004ApJ...603...42A}; \citealt{2005MNRAS.363..692D}b; \citealt{2005A&A...435..521N}; \citealt{2006A&A...451...71F}). Sub-parsec-scale radio emission was detected in 98\% LLAGN and AGN of the Palomar Sample observed with the VLBA\footnote{Very Long Baseline Array of the National Radio Astronomy Observatory (NRAO).} at 5 GHz (e.g., \citealt{2005A&A...435..521N}). The characteristic variability of LINERs, sometimes on timescales of months (e.g., \citealt{2002A&A...392...53N}; \citealt{2005ApJ...627..674A}; \citealt{2005ApJ...625..699M}; \citealt{2007MNRAS.377.1696M}; \citealt{2011A&A...527A.142G}), agrees with the confinement of the radio emission to compact cores. Most of these sources present a flat or slightly inverted radio spectrum and non-thermal brightness temperatures $T_\mathrm{B} > 10^{5}$ K footprint of a relativistic jet (\citealt{1979ApJ...232...34B}), indicating thus the presence of an AGN. However, only few of them show jet-like outflows or slightly resolved core emission at pc-scales (e.g., \citealt{2000ApJ...542..197F}; \citealt{2002A&A...392...53N}; \citealt{2005A&A...435..521N}; \citealt{2002A&A...385..425F}a; \citealt{2004A&A...418..429F}; \citealt{2007A&A...464..553K}; \citealt{2013ApJ...779....6H}), while larger kpc-scale radio jets in LLAGN is even scarcer (e.g., Cen A, \citealt{1992ApJ...395..444C}; M87, \citealt{2005AJ....130.1389L}; NGC\,1052, \citealt{1984ApJ...284..531W}; \citealt{2004A&A...420..467K}). The question that arises is: do radio observations fail to detect radio jets in LLAGN due to insufficient angular resolution and sensitivity, or do most LLAGN show compact cores due to insufficient energy to launch a jet at parsec or larger scales? It is clear that the answer must come from high-resolution and -sensitivity radio campaigns reaching sub-arcsecond resolutions, which are the only ones plausibly able to detect pc-scale jets in these sources. In this paper we analyze subarcsecond archival VLA\footnote{Very Large Array, now officially known as the Karl G. Jansky VLA, of the NRAO.} and VLBA data so far unpublished of a sample of nearby LLAGN studied in the IR at high spatial resolution by \cite{2010MNRAS.402..724P} and \cite{2010MNRAS.402..879R}, from now on PR2010. The results yield the first resolved radio structure of the pc-scale jet in NGC\,4594 (or Sombrero) and the detection of jet-like extended emission in NGC\,1097 and NGC\,2911. We also compile and derive the jet power, bolometric luminosity, and Eddington accretion rates of the eight sources included in the sample of PR2010, which allows us to perform a study of the nuclear energetics of LLAGN devoid of most of the assumptions usually adopted in large statistical surveys. The objects and the radio data analyzed are presented in Section~\ref{data}, while the results obtained are shown in Section~\ref{results} and discussed in Section~\ref{discussion}. Final conclusions are summarized in Section~\ref{conclusions}. \begin{table*} \centering \begin{threeparttable} \caption{Targets and Radio Archival Data \label{table1}} \begin{tabular}{lcccccc} \tableline \tableline Name & Array & Obs. date & Frequency & Resolution & Peak & rms \\ & & & (GHz) & (mas) & (mJy beam$^{-1}$) & (mJy beam$^{-1}$) \\ (1) & (2) & (3) & (4) & (5) & (6) & (7) \\ \hline NGC\,1097 & VLBA & 2010 May & 8.4 & 6.34 $\times$ 1.66 & 0.9 & 0.1 \\ NGC\,2911 & VLA-A & 1986 Mar & 4.9 & 0$\arcsec$.72 $\times$ 0$\arcsec$.35 & 70.1 & 0.2 \\ & & 1986 Mar & 14.9 & 0$\arcsec$.21 $\times$ 0$\arcsec$.11 & 21.5 & 0.6 \\ NGC\,4594 & VLBA & 2005 Jun & 8.6 & 4.43 $\times$ 0.97 & 68.0 & 0.4 \\ & & 2011 Mar & 15.3 & 2.11 $\times$ 0.60 & 73.5 & 0.3 \\ & & 2011 Mar & 23.8 & 1.13 $\times$ 0.62 & 8.9 & 0.3 \\ \tableline \end{tabular} \begin{tablenotes} \small \item {\bf Notes.}~(1) Source name; (2) array configuration; (3) observing date; (4) observing frequency; (5) beam size; (6) peak flux density; (7) off-source rms noise. \end{tablenotes} \end{threeparttable} \end{table*} | \label{conclusions} LLAGN were suggested to be typically associated with unresolved radio sources even at mas scales; however, we find that 67\% of the LLAGN with subarcsecond radio emission of the Palomar Sample present extended pc-scale jets or slightly resolved core emission when observed with sufficient angular resolution and sensitivity. Most LLAGN are thus powerful enough to eject radio jets at least at pc scales. We also report the detection of extended jets in the LLAGN/LINERs NGC\,1097, NGC\,2911, and NGC\,4594 based on the analysis of archival VLA and VLBA data. For NGC\,4594, the pc-scale jet is resolved, for the first time, into several components. With these detections, all the sources (eight out of eight) in the sample of the brightest nearby LLAGN in the southern hemisphere of PR2010 present parsec-scale or larger radio jets. The extended detected structures of the LLAGN in the Palomar and the PR2010 samples, together with the derived high brightness temperatures and radio spectral indices, indicate a clear association of the radio emission with AGN. This reinforces the long debated scenario in which LLAGN are also powered by accretion onto a massive BH as AGN but have a different accretion mode (i.e., are radiatively inefficient), and supports those models in which the subarcsecond SED of these sources is dominated by synchrotron jet emission. The pc-scale jet emission is found to be aligned with the large (100 pc to kpc) jet emission in 10 out of 13 sources for which both a small- and large-scale jet is detected, indicating that the nuclear engine of these sources is powerful enough to collimate the jets from pc to kpc scales. The misalignment between the small and large-scale jet emission for the other sources can be explained by jet interactions with the ISM. The high-spatial-resolution core SED available for the LLAGN in the PR2010 sample together with their jet power, either directly measured from X-ray cavities or derived from their core radio luminosity, allows us to derive the Eddington-scaled accretion rate and Eddington ratio of these sources with minimum ad-hoc assumptions. For the eight sources in the PR2010 sample we obtain that: \begin{itemize} \item A general trend is observed between jet power and jet size: those powerful radio galaxies with the largest radio jets (Cen A, M87, NGC\,1052) are also the ones with largest values of jet power, while the lowest values correspond to those LLAGN whose largest observed jet size is $\leq$ 100 pc. All the sources with jet power $>$ 10$^{42}$ erg s$^{-1}$ do show large-scale ($>$ 100 pc) radio jets. \item The jet power is comparable to or larger than the bolometric luminosity for all LLAGN, in agreement with statistical studies of larger LLAGN samples. Our study of individual sources reveals the presence of radio jets when these sources are observed with enough resolution and sensitivity, indicating that the jet kinetic power contributes at least as much as the radiative output regardless of the jet size. The nuclear energetics of LLAGN is thus dominated by the jet kinematic output. \item The Eddington ratio ranges $10^{-6}$-$10^{-4}$ and is thus highly sub-Eddington. When including the jet power to the total emitted luminosity (radiated plus kinetic), the Eddington-scaled accretion rate increases nearly one order of magnitude but is still $<10^{-4}$ for nearly all sources. These results indicate that LLAGN are not only unequivocally inefficient radiators, the main channel for releasing their energy being the jet, but also either inefficient accretors or very efficient advectors. Even with the current best estimate of kinetic power and bolometric luminosity, they remain highly sub-Eddington. \end{itemize} | 14 | 3 | 1403.6675 |
1403 | 1403.6396_arXiv.txt | The CMB map provided by the {\it Planck} project constrains the value of the ratio of tensor-to-scalar perturbations, namely $r$, to be smaller than $0.11$ (95 \% CL). This bound rules out the simplest models of inflation. However, recent data from BICEP2 is in strong tension with this constrain, as it finds a value $r=0.20^{+0.07}_{-0.05}$ with $r=0$ disfavored at $7.0 \sigma$, which allows these simplest inflationary models to survive. The remarkable fact is that, even though the BICEP2 experiment was conceived to search for evidence of inflation, its experimental data matches correctly theoretical results coming from the matter bounce scenario (the alternative model to the inflationary paradigm). More precisely, most bouncing cosmologies do not pass {\it Planck's} constrains due to the smallness of the value of the tensor/scalar ratio $r\leq 0.11$, but with new BICEP2 data some of them fit well with experimental data. This is the case with the matter bounce scenario in the teleparallel version of Loop Quantum Cosmology. | The latest {\it Planck} temperature data for cosmic inflation constrains the spectral index for scalar perturbations to be $n_s=0.9603\pm 0.0073$, ruling out exact scale invariance with over $5\sigma$ confidence, and establishes an upper bound for tensor/scalar ratio given by $r\leq 0.11$ (95 \% CL) \cite{Ade}. Such data shrinks the set of allowed simplest inflationary models: power law potentials in chaotic inflation {\cite{linde}}, exponential potential models {\cite{Lucchin}, inverse power law potentials \cite{Barrow}, are disfavored because they do not provide a good fit to {\it Planck's} data \cite{Ade, Steinhardt}. In fact, this data set prefers a subclass of inflationary models with {\it plateau-like} inflation potentials (see for example \cite{Olive}) and $R^2$ gravity \cite{Odintsov}. On the other hand, recent results from the BICEP2 experiment \cite{Kovac}, designed to look for the signal of gravitational waves in the $B$-mode power spectrum, lead to the same constrain for the spectral index, but constrain the ratio of tensor-to-scalar perturbations to be $r=0.20^{+0.07}_{-0.05}$ with $r=0$ disfavored at $7.0 \sigma$ (see figure $13$ of \cite{Kovac} to compare {\it Planck's} with BICEP2 data). This higher value of $r$ extends the set of compatible inflationary models, allowing back some of the simplest inflationary models cited above. Dealing with the matter bounce scenario, the alternative to the inflationary paradigm (see \cite{brandenberger} for a report about bouncing cosmologies), one encounters a similar problem when one tries to match {\it Planck's} data with theoretical results: theoretical results provide, in general, values of $r$ higher than $0.11$ and, then, to sort out this problem some very complicated mechanism has to be introduced to enhance the power spectrum of scalar perturbations \cite{Cai}, reducing the ratio $r$ enough to achieve the bound $0.11$. However, {in this work we will show that} the higher value of $r$ provided by BICEP2 allows the viability of some bouncing models. {This is the main goal of the paper.} {As a matter of fact, we will deal with the matter bounce scenario in Loop Quantum Cosmology (LQC) which, when one only takes into account holonomy corrections, provides the simplest bounce. More precisely, it is well known that LQC contains two kind of corrections: holonomy corrections and inverse-volume effects. When one deals with the flat Friedmann-Lema{\^\i}tre-Robertson-Walker (FLRW) geometry, holonomy corrections always lead to a big bounce (see for instance \cite{singh}), however this could not happen when one considers inverse-volume effects. For example, when the universe is filled by a field under the action of a non-negative potential (to guarantee a positive energy density), one will obtain a non bouncing universe because the Hubble parameter never vanishes (see equations (5) and (8) of \cite{Bojowald}). That is the reason why, in this paper, we will do not take into account inverse-volume corrections.} { On the other hand, for the flat FLRW geometry, it has been recently showed in \cite{Haro1} that holonomy corrected LQC can be formulated as a particular example of teleparallel $F(T)$ gravity, where $T$ is the so-called {\it torsion scalar} whose value in the flat FLRW spacetime is equal to $-6H^2$. This new formulation of LQC with holonomy corrections has been named {\it teleparallel LQC} and only coincides with the standard holonomy corrected LQC in the FLRW geometry. Dealing with cosmological perturbations both formulations lead to different perturbation equations and, of course, to different results. The reason of this difference is that in holonomy corrected LQC, working in the Hamiltonian framework, the corresponding perturbation equations are obtained replacing the Ashtekar connection by a suitable sinus function in the classical Hamiltonian and inserting in it counter-terms to preserve the algebra of constrains \cite{Barrau,Cailleteau}). In constrast to holonomy corrected LQC, the perturbation equations in teleparallel LQC are directly obtained, in the Lagrangian framework, from the well-known perturbation equations in teleparallel $F(T)$ gravity \cite{Cai1}. In fact, it has been shown in \cite{Haro} that for scalar perturbations both formulations lead to the same kind of results, the difference appears when one deals with tensor perturbations, because in teleparallel LQC the equation of perturbations \cite{Haro} is a {\it regular} equation, but in holonomy corrected LQC the corresponding equation \cite{Cailleteau} has two singular points (at the beginning and end of the super-inflationary phase). This difference is what leads to completely different results.} {To show that, we deal with the matter bounce scenario in LQC, where the universe is filled by only a scalar field whose potential is the simplest one leading, at early times, to a matter domination in the contracting phase. In this case, the conservation equation is a second order differential equation (a Klein-Gordon equation). Each orbit, i.e. each solution of this differential equation, depicts a different matter dominated universe at early times. We will see that one of these orbits can be calculated analytically (the orbit that depicts a matter dominated universe for all time), but the other ones have to be calculated numerically. Then, for all of these orbits } we will calculate {analytically and numerically}, the { corresponding} tensor/scalar ratio {for adiabatic perturbations (we only considers one matter field, meaning there are not entropy perturbations)} coming from holonomy corrected and teleparallel LQC, and we will check that in the case of teleparallel LQC there are {orbits} leading to theoretical results that match correctly with BICEP2 data, and there are other {orbits} that provide theoretical results that fit well with {\it Planck's} data. {On the other hand, we will also show numerically that holonomy corrected LQC, provides theoretical results, that only match correctly with {\it Planck's} data.} The units used in the {paper} are $\hbar=c=8\pi G=1$. | In this work we have studied cosmological perturbations produced by one matter scalar field (adiabiatic perturbations) in the context of holonomy corrected and teleparallel LQC. We have explained that, in the flat FLRW geometry, both formulations coincide, but dealing with perturbations they provide different results. This is basically due to tensor perturbations, that satisfy completely different equations depending on the formulation used. Our results show that holonomy corrected LQC only fits well with {\it Planck's} data due to the low value of the tensor/scalar ratio provided by this theory. Then, since {\it Planck's} and BICEP2 data are in strong tension (they provide completely different experimental data for the ratio of tensor to scalar perturbations), holonomy corrected LQC is only viable if the correct experimental data are the ones given by the {\it Planck} project. However, teleparallel LQC provides theoretical results that match with {\it Planck's} data and others that fit well with BICEP2. Then, whichever between the {\it Planck} or BICEP2 experimental results are most accurate (at this moment, there is not an answer), teleparallel LQC always has a set of solutions whose theoretical results match correctly with the accurate experimental data. \vspace{0.5cm} \centerline{\bf Aknowledgements} \vspace{0.5cm} We thank Professor Sergei D. Odintsov for correspondence and useful comments. This investigation has been supported in part by MINECO (Spain), project MTM2011-27739-C04-01, MTM2012-38122-C03-01, and by AGAUR (Generalitat de Ca\-ta\-lu\-nya), contracts 2009SGR 345 and 994. | 14 | 3 | 1403.6396 |
1403 | 1403.1263_arXiv.txt | Counter-rotating galaxies host two components rotating in opposite directions with respect to each other. The kinematic and morphological properties of lenticulars and spirals hosting counter-rotating components are reviewed. Statistics of the counter-rotating galaxies and analysis of their stellar populations provide constraints on the formation scenarios which include both environmental and internal processes. | Counter-rotating galaxies belong to the class of multi-spin galaxies. They are characterized by the presence of two components that are observed rotating in opposite directions with respect to each other. Before their discovery, counter-rotating galaxies were considered from a theorethical point of view and dismissed as elegant curiosities \citep[see][for an historical perspective on the investigation of galactic kinematics]{Rubin1994b}. This belief changed when \citet{Galletta1987} measured the kinematics of the ionized-gas and stellar components of the early-type barred galaxy NGC~4564 and showed they are rotating in opposite directions around the same rotation axis. In the same period, the first counter-rotating elliptical \citep[NGC~5898;][]{Bettoni1984, Bertola+Bettoni1988} was found in Padua too. As more and more data became available, the presence of counter-rotating components was detected in tens of galaxies along all the Hubble sequence, from ellipticals to irregulars. Previous reviews about counter-rotation are those by \citet{Galletta1996} and \citet{Bertola+Corsini1999}, while \citet{Corsini+Bertola1998} listed all the counter-rotating galaxies known at the time. This paper focuses on the counter-rotation in lenticulars and spirals. | After they were discovered three decades ago, counter-rotating galaxies still represent a challenging subject for both theorists and observers. Although the broad picture of the formation of counter-rotating galaxies is in place, we still miss many details. A few issues should be attacked first in the near future to make a step forward in our understanding of counter-rotation in disk galaxies. They include: a deep imaging survey to look for the fingerprints of accretion and merging events in the environment of counter-rotating galaxies at very low levels of surface brightness; the analysis of a complete sample of spiral galaxies to drive unbiased conclusions about the frequency of the different kinds of counter-rotation; the derivation of the stellar LOSVD from high (spectral and spatial) resolution data obtained with wide-field integral-field units to look for yet undetected retrograde stars; and the extensive measurement of the stellar populations of the prograde and retrograde components in counter-rotating galaxies to test the predictions of the different formation scenarios. | 14 | 3 | 1403.1263 |
1403 | 1403.1780_arXiv.txt | Observing sites at the East-Antarctic plateau are considered to provide exceptional conditions for astronomy. The aim of this work is to assess its potential for detecting transiting extrasolar planets through a comparison and combination of photometric data from Antarctica with time series from a midlatitude site. During 2010, the two small aperture telescopes ASTEP\,400 (Dome~C) and BEST\,II (Chile) together performed an observing campaign of two target fields and the transiting planet \mbox{WASP-18b}. For the latter, a bright star, Dome~C appears to yield an advantageous signal-to-noise ratio. For field surveys, both Dome~C and Chile appear to be of comparable photometric quality. However, within two weeks, observations at Dome~C yield a transit detection efficiency that typically requires a whole observing season in Chile. For the first time, data from Antarctica and Chile have been combined to extent the observational duty cycle. This approach is both feasible in practice and favorable for transit search, as it increases the detection yield by 12--18\%. | To obtain a better understanding of our Universe, improved observing conditions have constantly been sought by astronomers. While the limiting factors can be many and diverse, the selection of an observing site particularly impacts the quality of the astronomical data recorded. Among the most important constraints are the fraction of clear skies, the level of atmospheric seeing and scintillation, the accessibility of a wide spectral range, and a low sky brightness due to emission, scattered light, and light pollution. Optimal conditions are generally achieved high above the atmosphere, i.e., using air-borne or space observatories. However, such projects are limited by extensive costs, and technical considerations impose further constraints. Therefore, the search for excellent observational sites on ground is being pursued with unwaned interest. Over the last few decades, high-altitude sites such as in the Chilean Atacama desert or the mountain tops of Hawaii have generally been recognized to provide the best observing conditions for large ground-based observatories. In recent times, Antarctica is expected to provide a number of advantages for astronomy \citep[see, e.g.,][]{Indermuehle2005,Saunders2009,Burton2010,Fossat2010}, and therefore sites at the East Antarctic plateau such as Dome~C are currently being considered for future large-scale observatories \citep{Burton2005,Cui2010,Ichikawa2010,Abe2013a}. In particular, time series observations are considered to benefit from a high duty cycle and low photometric noise. These are of key importance for detecting and characterizing transiting extrasolar planets. First, more and/or smaller planets are expected to be found at Dome~C due to an increased photometric precision \citep{Rauer2010a}. Two conditions are considered important in this respect: Less systematic noise due to stable environmental conditions (in particular, the lack of day/night temperature variations; \citealt{Pont2005a}), and less scintillation noise due to a low level of atmospheric turbulence. The latter is expected to be 2--4~times smaller at Dome~C compared to temperate sites \citep{Kenyon2006b}. Second, the Antarctic winter allows for an almost continuous time series to be obtained \citep{Caldwell2004,Pont2005a}, although the total amount of usable dark time is not increased compared to midlatitude sites \citep{Kenyon2006a}. However, observations can cover large planetary orbits better than temperate sites with diurnal interruptions. \citet{Rauer2008} showed that planets with orbital periods of up to two weeks are covered well within one observing season at Dome C; in contrast, a similar performance with midlatitude sites could only be achieved if three stations were combined into a network. While \citet{Rauer2008} relied on observing times modeled from weather data and astronomical dark time, \citet{Crouzet2010} obtained statistics directly from the 10\,cm Antarctic Search for Transiting ExoPlanets (ASTEP)~South telescope. They estimated the transit yield and compared it to an analogous instrumental setup at La Silla: The ASTEP~South 2008 campaign is expected to yield a number of planet detections comparable to a modeled observing season at La Silla. However, if the ASTEP~South observations were extended over the whole winter season, the expected yield would be larger at Dome~C. In addition, \citet{Abe2013} recently obtained an unprecedented ground-based duty cycle while monitoring the transiting planet WASP-19b from Dome~C. While these previous studies indicate an advantage for transit search at Dome~C, this still needs to be confirmed on the basis of extensive photometric data. For example, the study of \citet{Kenyon2006b} derived the scintillation noise from measurements of atmospheric turbulence profiles above Dome~C; however, this will only yield an advantage if it forms the dominant component in the noise budget for bright stars. \citet{Crouzet2010} used observing statistics from Dome~C, but modeled the photometric quality in Antarctica and Chile from instrument characteristics. This study aims to address two open questions: First, whether transiting planets can be better photometrically characterized or detected from Antarctica in comparison to midlatitude sites. Second, if a transit survey from Antarctica together with a midlatitude site is feasible and promising in practice. In order to obtain a first comparison based on photometric data, a coordinated survey has been performed both from Dome~C and Chile. The paper is outlined as follows. It first introduces the instruments used (Section~\ref{sec:telescopes}), the observations (Section~\ref{sec:obs}), and their analysis (Section~\ref{sec:analysis}). Section~\ref{sec:phot} presents the scientific results regarding the photometric quality (i.e., limiting the radius of detectable transiting planets), while Section~\ref{sec:dc} focuses on the observational phase coverage (limiting the orbital period found by transit surveys). Section~\ref{sec:summary} summarizes the paper. | \label{sec:summary} This study presents first joint observations from Antarctica and Chile: The known transiting planet WASP-18b and two target fields were monitored together with ASTEP and BEST\,II in 2010. For the two target fields, ASTEP measurements span 25 nights and include 94{,}965 stars, while BEST\,II obtained 224{,}552 light curves during 18 nights; joint observations are available for 58{,}822 stars. Their comparison aims at a first quantitative evaluation of the potential for transit search at Dome~C that is solely based on real photometric time series. Particular attention was paid to the photometric precision and observational phase coverage, which are both expected to yield advantageous conditions for transit searches in Antarctica. For a single transit of WASP-18b, ASTEP yields an unbinned, out-of-transit standard deviation of 1.9\,mmag. The data from Antarctica thus show a smaller noise level compared to 2.6\,mmag achieved with BEST\,II. However, the difference is not in the order of a factor~2--4 as expected from a smaller scintillation noise \citep{Kenyon2006b}, indicating that unfiltered stellar variability and/or systematic effects still contribute significantly to the noise budget of either light curve. An analysis of the two large data sets from both telescopes showed that their photometric quality as well is excellent, reaching a precision of 2--4\,mmag for bright stars from both Antarctica and Chile over each observing campaign. However, the photometric precision is very similar: A simulation shows that BEST\,II and ASTEP overall yield a well comparable detection yield for a range of planetary radii. An advantage to find significantly smaller planets from Antarctica is thus not evident from these first data. Whether the limiting noise component is an inherent site characteristic or can be further decreased will become more apparent while the experience with the ASTEP\,400 system, its unique environment, and the data reduction grows through continued operation. In this respect, it would also be most interesting to directly compare the photometric quality with another Antarctic site, e.g., with data from one of the Antarctic Survey Telescopes \citep[AST3;][]{Li2013} at Dome~A. In contrast to the photometric quality, the long polar night yields a clear advantage for transit search in Antarctica. Within two weeks of observations, ASTEP yields a detection for planets at short periods that can only be achieved during a whole season from Chile. For the first time, light curves from Antarctica and Chile were combined in order to extend the observational duty cycle. A case-by-case comparison showed that the photometric systems of ASTEP and BEST\,II compare well and such a combination is practicable to the precision needed for transit search. If BEST\,II observations are added, the yield increases by 12--18\%. Furthermore, a case study has shown that a similar relative increase is even encountered if the duty cycle of ASTEP was extended further, and could be increased up to 26--33\% if BEST\,II observations could be obtained in parallel to each night with observations from Antarctica. Thus, such network observations can indeed increase the detection yield significantly compared to time series obtained from Antarctica alone. Note that a key advantage of the combination Chile-Dome~C is its large longitudinal separation (166$\degr$), and slightly less favorable conditions are expected from combining a Chilean site with another Antarctic observatory such as Dome~A~(148$\degr$) or Dome~F~(110$\degr$). | 14 | 3 | 1403.1780 |
1403 | 1403.2260_arXiv.txt | { The highest energies of solar energetic nucleons detected in space or through gamma-ray emission in the solar atmosphere are in the GeV range. Where and how the particles are accelerated is still controversial. } { We search for observational information on the location and nature of the acceleration region(s) by comparing the timing of relativistic protons detected on Earth and radiative signatures in the solar atmosphere during the particularly well-observed 2005 Jan 20 event. } { This investigation focusses on the post-impulsive flare phase, where a second peak was observed in the relativistic proton time profile by neutron monitors. This time profile is compared in detail with UV imaging and radio spectrography over a broad frequency band from the low corona to interplanetary space. } {It is shown that the late relativistic proton release to interplanetary space was accompanied by a distinct new episode of energy release and electron acceleration in the corona traced by the radio emission and by brightenings of UV kernels. These signatures are interpreted in terms of magnetic restructuring in the corona after the CME passage. } {We attribute the delayed relativistic proton acceleration to magnetic reconnection and possibly to turbulence in large-scale coronal loops. While Type II radio emission was observed in the high corona, no evidence of a temporal relationship with the relativistic proton acceleration was found.} | \label{Sec_Intro} On certain occasions transient energetic particle fluxes from the Sun, called solar energetic particle (SEP) events, may comprise relativistic nucleons at energies up to several GeV or even tens of GeV. Upon impinging on the Earth's atmosphere, these particles trigger nuclear cascades that produce secondaries detectable by ground-based neutron monitors and muon telescopes. These particular SEP events are therefore called ground level enhancements (or ground level events; GLEs). The rarity of GLEs - only 72 were detected since 1942, including a small event on 2014 Jan 06 - clearly shows that GeV energies are extreme in solar events. Understanding their origin is therefore one of the more challenging tasks in the research on solar eruptive activity. A comprehensive summary of GLE observations was given by \cite{Lop-06}, and the review by \cite{Car-62} is still very informative. GLEs are produced in conjunction with intense flares and extremely fast coronal mass ejections \citep[CMEs;][]{Blv:al-10,Gop:al-12}. The acceleration mechanisms are thought to be related to the flare, which usually means magnetic reconnection, or to the shock wave generated by the CME. Which of the two possibilities, the flare or the shock wave, is actually the accelerator is hard to say on observational grounds. Physical relationships between the particles detected on Earth and in dynamical processes in the solar corona can in principle be inferred from comparing the arrival times of the SEPs and radiative signatures of flares and CMEs. This is especially possible when the SEP time profile has an intrinsic structure. Evidence of successive distinguishable SEP releases within a given event has been reported from particle observations in the MeV to tens-of-MeV range \citep[e.g.,][]{Koc:al-07} and more often in GLEs, because the scattering mean free path in the interplanetary medium is larger at relativistic energies \citep{Dro-00}. It is indeed well established that GLEs often have a double-peaked structure, with an initial fast rise and an anisotropic particle population - called the `prompt component' - followed by a more gradual and less anisotropic `delayed component' \citep[see review in ][chap. 7.3]{Mir-01}. \cite{McC:al-12} conclude that the sequence of an anisotropic impulsive peak, and a less anisotropic gradual peak occurring 7-15~min later is a common occurrence when the parent active region is magnetically connected to the Earth, while the absence of the impulsive peak is typical of poorly connected activity near to or east of the central solar meridian or well beyond the western limb. A prominent case illustrating this double-peaked structure is the GLE of 2005 Jan 20. It displayed a well-defined, rapidly rising time profile at the beginning and a distinct second peak a few minutes later. Evidence that the first release was related to particle acceleration in the flaring active region in the low corona was brought by different publications \citep[e.g.,][]{Sim-06,Sim-07, Kuz:al-08,Grc:al-08,McC:al-08,Msn:al-09}. In the present paper we complete the investigation of \cite{Msn:al-09}\footnote{\cite{Msn:al-09} will be cited as Paper~1 in the following.} through a detailed comparison between the second peak of the relativistic proton time profile derived from neutron monitor measurements with high-quality radio and UV observations. This article is structured as follows. Section~\ref{Sec_Obs} introduces the observations (\ref{Sec_Obs_inst}) and describes the time profile of relativistic protons detected at Earth (\ref{Sec_Obs_p}). We conclude that two successive proton releases near the Sun can be identified in the GLE observation. The evolution of radio, hard X-ray, and gamma-ray emissions is tentatively separated into different acceleration episodes in Sect.~\ref{Sec_Obs_eps}, and the relation with relativistic protons before and during the first release is briefly discussed in Sect.~\ref{Sec_Obs_pr1} with reference to Paper~1 and to the evolution of the dm-m-wave dynamic spectrum studied by \cite{Brt:al-10}. A more detailed analysis of the radio and UV emission accompanying the second relativistic proton release (Sect.~\ref{Sec_Obs_pr2}) leads us to suggest that the particles are accelerated during dynamical processes in the magnetically stressed corona after the CME passage. The findings are summarised in Sect.~\ref{Sec_Obs_sum} and discussed in Sect.~\ref{Sec_Disc} with respect to previous work on the origin of relativistic solar nucleons. \begin{figure} \centering \includegraphics[width=9cm]{Klein_al_23783_fg1a.eps} \includegraphics[width=9cm]{Klein_al_23783_fg1b.eps} \caption[]{X-ray and radio emission and the relativistic proton profile of the 2005 Jan 20 event. From bottom to top: (a) soft X-rays $\lambda=0.1-0.8$~nm (dark line) and 0.05-0.4~nm (light; red in the colour plot of the online version); (b) microwaves (dark line 2.7 GHz, light - red in the colour display - a combination of 17 GHz (NoRP) before and 15.4~GHz (LEAR) after 06:55 UT); (c) dynamic radio spectrum at dm-m waves (ARTEMIS-IV; inverse colour scale; 1~s integration time); (d) decametre-kilometre wave radio emission (Wind/WAVES; inverse colour scale; 1~min integration); (e) proton flux time history at 2~GV (dark curve) and 5~GV (light curve; red in the online version) rigidity (kinetic energy 1.27 and 4.15~GeV, respectivelyy), time axis shifted back by 216~s. The intervals delimited by vertical lines and numbered 0 to 6 are different episodes of particle acceleration, as discussed in the text. } \label{Fig_ovw} \end{figure} | Results for the relativistic SEP event of 2005 Jan 20 presented in Paper~1 and the present work are consistent with the basic scheme of successive relativistic particle releases during a GLE devised by \citeauthor{Mir-01} (\citeyear{Mir-01}, and references therein), \cite{Vas:al-06}, and \citeauthor{McC:al-08} (\citeyear{McC:al-08}, \citeyear{McC:al-12}). Our observations relate the acceleration processes to evolving large-scale magnetic structures in the corona and therefore to the CME development. This does not mean that the CME shock is the accelerator. To the extent that the timing relationship between relativistic protons on Earth and the radio and UV emission in the solar corona discloses a physical relationship, rather than mere coincidence, the observations suggest that relativistic particle acceleration throughout the impulsive and post-impulsive phase is related to the magnetic restructuring of the corona during and after the liftoff of the CME. There is no evidence of a temporal connection between the relativistic proton acceleration and the evolution of a coronal shock as traced by the Type~II radio emission. | 14 | 3 | 1403.2260 |
1403 | 1403.2783_arXiv.txt | I present a comprehensive review of the evolution of galaxy structure in the universe from the first galaxies we can currently observe at $z \sim 6$ down to galaxies we see in the local universe. I further address how these changes reveal galaxy formation processes that only galaxy structural analyses can provide. This review is pedagogical and begins with a detailed discussion of the major methods in which galaxies are studied morphologically and structurally. This includes the well-established visual method for morphology; S{\'e}rsic fitting to measure galaxy sizes and surface brightness profile shapes; non-parametric structural methods including the concentration ($C$), asymmetry ($A$), clumpiness ($S$) (CAS) method, the Gini/M$_{20}$ parameters, as well as newer structural indices. Included is a discussion of how these structural indices measure fundamental properties of galaxies such as their scale, star formation rate, and ongoing merger activity. Extensive observational results are shown demonstrating how broad galaxy morphologies and structures change with time up to $z \sim 3$, from small, compact and peculiar systems in the distant universe to the formation of the Hubble sequence dominated by spirals and ellipticals we find today. This review further addresses how structural methods accurately identify galaxies in mergers, and allow measurements of the merger history out to $z \sim 3$. The properties and evolution of internal structures of galaxies are depicted, such as bulges, disks, bars, and at $z > 1$ large star forming clumps. The structure and morphologies of host galaxies of active galactic nuclei and starbursts/sub-mm galaxies are described, along with how morphological galaxy quenching occurs. Furthermore, the role of environment in producing structure in galaxies over cosmic time is treated. Alongside the evolution of general structure, I also delineate how galaxy sizes change with time, with measured sizes up to a factor of 2-5 smaller at high redshift at a given stellar mass. This review concludes with a discussion of how the evolving trends in sizes, structures, and morphologies reveal the formation mechanisms behind galaxies which provides a new and unique way to test theories of galaxy formation. | Galaxy structure is one of the fundamental ways in which galaxy properties are described and by which galaxy evolution is inferred. There is a long history of the development of this idea, which began with the earliest observations of galaxies, and continues up to the modern day as one of the major ways we study galaxies. This review gives a detailed description of the progress made up to late-2013 in using galaxy structure to understand galaxy formation and evolution. It is meant to be used as a primer for obtaining basic information from galaxy structures, including how they are measured and applied through cosmic time. The introduction to this review first gives an outline of the basic events in the history of galaxy morphology and structure analyses, while the second part of the introduction describes how galaxy structure fits into the general picture of galaxy formation. I also give a detailed description of the goals of this review at the end of the introduction. \subsection{Historical Background} Galaxy morphology has a long history, one that even predates the time we knew galaxies were extragalactic. When objects which today we call galaxies were first observed what clearly distinguished them from stars was their resolved structure. Since this time, structure and morphology has remained one of the most common ways galaxies are described and studied. Initially this involved visual impressions of galaxy forms. This has now been expanded to include quantitative methods to measure galaxy structures all the way back to the earliest galaxies we can currently see. The first published descriptions of galaxy structure and morphology predates the telescopic era. For example, the Andromeda nebula was described as a 'small cloud' by the Persian astronomer Abd al-Rahman al-Sufi in the 10th century, (Kepple \& Sanner 1998). The study of galaxies remained descriptive until the late 20th century, although more and more detail was resolved as technology improved. As a result, for about 150 years the science of galaxies was necessarily restricted to cataloging and general descriptions of structure, with notable achievements by Messier and William and John Herschel who located galaxies or `nebula' by their resolve structure as seen by eye. Even before photography revolutionized the study of galaxies some observers such as William Parsons, the 3rd Earl of Rosse, noted that the nebulae have a spiral morphology and first used this term to describe galaxies, most notably and famously in the case of M51. It was however the advent of photography that astronomers could in earnest begin to study the morphologies and structures of external galaxies. The most notable early schemes were developed by Wolf (1908), and Lundmark (1926), among others. This ultimately led to what is today called the Hubble classification which was published in essentially its modern form in Hubble (1926), with the final 'Hubble Tuning Fork' established in Hubble (1936) and Sandage (1961). The basic Hubble sequence (Figure~1) consists of two main types of galaxies, ellipticals and spirals, with a further division of spirals into those with bars and those without bars. Hubble, and the astronomers who followed him, could classify most nearby bright galaxies in terms of this system. \begin{figure}% \centerline{\psfig{figure=con.fig1.ps,height=14pc}} \caption{A modern form of the Hubble sequence showing the sequence of ellipticals and S0s, and the `tuning fork' in spirals. The elliptical sequence is determined by the overall shape of the galaxy, while spiral classifications are divided into different types (a-c) depending upon how wound-up spiral arms are, how large the bulge relative to the disk is, and how smooth the spiral arms in the spirals arm. The tuning-fork is the differential between spirals with and without bars. Also shown is the extension of this sequence to dwarf spheroidal galaxies and irregular galaxies, both of which are lower mass systems (Kormendy \& Bender 2012). } \label{fig1} \end{figure} The development of morphological classification methods continued into the 20th century, with newer methodologies based solely on visual impressions. For example, de Vaucouleurs (1959) developed a revised version of the Hubble sequence which included criteria such as as bars, rings and other internal features that were prominent on photographic plates of galaxies. Likewise, van den Bergh (1960, 1976), and later Elmegreen \& Elmegreen (1987) developed a system to classify galaxies based on the form of spiral arms, and the apparent clumpiness of the light in these arms. While it is important to classify galaxies visually, and all systems have some use, as all features should be explained by physics, it is not obvious which structural features of galaxies are fundamental to their formation history. Ultimately morphology and structure needs to prove to be useful for understanding galaxies, as there is now extensive use of photometric and spectroscopic methods permitting measurements of perhaps more fundamental measures of stellar populations and dust/gas properties in galaxies. Along these lines, at roughly the same time as progressively complicated classification systems were developed, astronomers such as Holmberg (1958) established that physical properties of nearby galaxies correlate with morphology in a broad context. Holmberg (1958) found that ellipticals are typically massive and red, and show little star formation, while spirals are less massive, bluer and have evidence for ongoing star formation. This quantitatively expands into other physical parameters as well (e.g., Roberts 1963; Roberts \& Haynes 1994; Conselice 2006a; Allen et al. 2006). It is also well known that this segregation of morphology in the local universe provides an important clue for understanding the physics of galaxy formation, especially as local environment is found to strongly correlate with a galaxy's morphology (e.g., Dressler 1984; see \S 4.7). A revolution in morphological and structurally studies came about with the advent of photometric photometry, and especially the later use of Charged Coupled Devices (CCD), which made detailed quantitative measurements of light distributions in galaxies possible. The first major contribution from this type of work was by de Vaucouleurs (1948) who used photometry to show that the light profiles of what we would identify today as massive ellipticals all follow roughly the same fundamental light distribution, known as the de Vaucouleurs profile. This was later expanded by others, most notably S{\'e}rsic (1963), who demonstrated that a more general form of light distributions matched galaxy light profiles with disks having exponential light profiles, while the light distribution within massive ellipticals generally following the de Vaucouleurs profile. This has led to a huge industry in measuring the light profiles of galaxies in the nearby and distant universe which continues today (\S 2.2). During the 1970s and 1980s the study of galaxy structure expanded to include the decomposition of galaxy light into bulge and disk profiles (e.g., Kormendy 1977) as well as features such as bars, rings and lenses (e.g., Kormendy 1979; de Vaucouleurs et al. 1993). The three dimensional structure of disk galaxies was investigated (e.g., van der Kruit \& Searle 1982), as well as detailed studies of bulges and disk in spiral systems (e.g., de Jong 1996; Peletier \& Balcells 1996). We also now know there is a great diversity in elliptical galaxy internal structures (e.g., Kormendy et al. 2009). Similar investigations demonstrated that secular evolution within disks can provide an explanation for how bars, rings and lenses can form (e.g., Kormendy 1979; Combes \& Sanders 1981). These effects, not driven by hierarchical galaxy formation, are also likely responsible for the formation of pseudo-bulges and may drive the formation of central massive black holes (e.g., Kormendy \& Kennicutt 2004; Sellwood 2013). While there is a large amount of work done on the structures and morphologies of galaxies in the nearby universe (e.g., see Kormendy et al. 2009; Buta 2013), it is difficult to investigate more than the very basics of structure and morphology when studying distant galaxies. This is due to the fact that current technology does not allow us to resolve these distant galaxies in the same detail as we can for closer systems. As such, this review will concentrate on the features and properties of galaxy structure which we can measure in distant galaxies, and how this reveals how galaxy evolution and formation occurs. The result of this is that one of the areas where galaxy structure and morphology has made its biggest impact is its ability to measure fundamental properties of distant galaxies that we can compare with nearby galaxies to determine evolution. There are extensive methods for studying galaxy evolution which galaxy structure analyses are becoming an essential aspect of, and providing unique information on, the history and physics of galaxy assembly, which I detail in this review. \subsection{Galaxy Structure within the Context of Galaxy Formation} We know that there is significant evolution in galaxies over time as the stellar mass density of galaxies evolves rapidly at $1 < z < 3$, with about half of all stellar mass formed by $z = 1$ (e.g., Bundy et al. 2005; Mortlock et al. 2011). We also know that there is a vast diversity of star formation histories for individual galaxies, and that the integrated star formation rate density in the universe's history peaks at $z \sim 2.5$, and declines at higher and lower redshifts (e.g., Shapley 2011; Madau \& Dickinson 2014, this volume). However, it is not clear from these observations what are/were the driving forces creating galaxies. Theory offers several approaches for understanding how galaxies form which detailed studies are starting to probe. We now believe that galaxy formation can happen in a number of ways. This includes: in-situ star formation in a collapsed galaxy, major and minor mergers, and gas accretion from the intergalactic medium. Galaxy structure and morphology are perhaps the best ways to trace these processes, as I discuss in this review. Another major question I address in this review is how do the structures and morphologies of galaxies change through cosmic time. Major issues that this topic allows us to address include: the formation history of the Hubble Sequence; whether galaxies form 'in-side-out' or 'out-side-in'?; how long does a galaxy retain its morphology?; is morphology a invariant quantity in a galaxy over a long cosmic time-span?, and furthermore what relative role does star formation and merging play in galaxy formation? Galaxy structure and morphology has made a significant impact on these questions largely due to the Hubble Space Telescope (HST) and its various 'Deep Field' campaigns starting in the mid-1990s, finding thousands of galaxies at redshifts $z >1$ within these images. This is complemented by extensive imaging and spectroscopy for nearby galaxies carried out by surveys such as the Sloan Digital Sky Survey (SDSS) and the Millennium Galaxy Catalog (e.g., Shen et al. 2003; De Propris et al. 2007). Combining these surveys makes it possible to study in detail the structures of distant galaxies, and to compare these with structures at different redshifts. This has led to a renaissance in the analysis of galaxy structure, including parametric fitting using S{\'e}rsic profiles, and the development of non-parametric measurements of galaxy structure that have allowed us to use galaxy morphology/structure as a tool for deciphering how galaxy assembly occurs over cosmic time. We are in fact now able to resolve galaxies back to redshifts $z = 8$ with imaging from space, and recently as well with adaptive optics from the ground (e.g., Conselice \& Arnold 2009; Carrasco et al. 2010; Akiyama et al. 2008). This reveals that galaxy structure is significantly different in the early universe from what it is today, and that there is a progression from the highest redshifts, where galaxies are small, peculiar, and undergoing high star formation rates to the relative quiescent galaxies that we find in the nearby universe. How this change occurs, and what it implies for galaxy evolution, is another focus of this review. Another ultimate goal is to describe the methods for measuring galaxy structure and morphology for nearby galaxies up to the most distant ones we can see. I also discuss how galaxy structure correlates with physical properties of galaxies, such as their star formation rate, merging and their overall scale. I then provide a description of the observed structural evolution of galaxies, and a discussion of what this implies for the driving mechanisms behind galaxy formation using the calibrated methods. The amount of information we have about the structures and properties of galaxies declines as one goes to higher redshift systems, and issues that arise due to observational bias must be dealt with. I therefore also discuss systematics that can be addressed through imaging simulations to determine the real evolution of the morphologies and structures of galaxies. I finish this review with a discussion of future uses of galaxy structure/morphology, including the potential with the advent of JWST and Euclid. This review is structured as follows. In \S 2, I describe the analysis methods used for measuring the morphologies and structures of galaxies. In \S 3 I describe how structures and morphologies reveal fundamental galaxy properties and evolutionary processes, while \S 4 describes the observed evolution of the structures of galaxies through cosmic time. I finish this review with a description of how galaxy structure and evolution is becoming an important aspect for understanding the underlying theory of galaxy formation and cosmology in \S 5 and give a summary and future outlook in \S 6. | I present here a review of galaxy structure and morphology studies in the galaxy population through cosmic time from $z = 8$ until today. The approach taken in this review is largely observational with a limited amount of interpretation, although I do show where galaxy structure and morphology can test galaxy formation and even cosmological models in a new, largely unexplored way. As of January 2014, the major conclusions concerning galaxy structure and its evolution, as discussed in this article, can be summarized as: \vspace{0.2cm} \noindent I. Galaxy structure and morphology is the longest studied observational feature of galaxies. In this review galaxy morphology is the apparent classification based on visual inspection, while structure is a way to quantify the light distributions in galaxies. In many ways morphology is still a descriptive science, and visual efforts continue to provide useful information in the form of large-scale projects to classify many galaxies, as in the Galaxy Zoo effort. The Hubble sequence has and likely will remain the major paradigm in which we consider galaxy morphology, although this system does not 'work' at high redshifts where most galaxies cannot be classified into a single Hubble type (\S 4.1). \vspace{0.2cm} \noindent II. Using the Hubble scheme the evolution of three broad classes of galaxies are now classified accurately out to $z \sim 3$ -- namely ellipticals, spirals and peculiars. The relative abundance of these galaxies has been measured as a function of redshift out to these early times. What we find is that the peculiar galaxies dominate the galaxy population at $z \sim 2.5 - 3$, with a relative fraction of at least 70\%. Galaxies which are elliptical and spiral like in appearance (but not necessarily in physical properties, see \S 4.1) become progressively more common at lower redshifts. The number densities of these two normal galaxies together equals that of the peculiars by $z \sim 1.4$ (Mortlock et al. 2013). \vspace{0.2cm} \noindent III. Since galaxy morphology by visual estimates is limited in its ability to derive the physics behind galaxy formation and by its nature is not quantitative, the use of parametric (\S 2.2) and non-parametric (\S 2.3) methods are essential for deriving in a quantitative way how galaxies are evolving. These quantitative indices also correlate to some degree with the present and past star formation history and properties of a galaxy. More work is needed to establish these relations with more certainty, but it appears that the S{\'e}rsic index and concentration correlate with the scale or mass of a galaxy, the clumpiness index with the star formation, and the asymmetry parameter with ongoing merging activity (\S 3). \vspace{0.2cm} \noindent IV. The merger history is now know from applying structural analyses to galaxy images in deep Hubble Space Telescope surveys such as the Hubble Deep Field (\S 4). The result of this is that the galaxy merger fraction increases with redshift at all stellar mass and luminosity selections. This increase can be fit well by a power-law $(1+z)^{m}$ up to $z \sim 3$, although at higher redshifts the structurally derived merger fraction may plateau (Conselice \& Arnold 2009). Using numerical/hydrodynamical simulations the time-scales for these mergers can be calculated, and thus merger fractions can be converted into merger rates (\S 4.3.2). The merger rate allows for the calculation of the number of mergers galaxies at various masses undergo, as well as the amount of stellar mass which is added to galaxies due to the merger process. The result of this is that it appears that up to half of the stellar mass in modern massive galaxies were formed in mergers between $1 < z < 3$, although at $z < 1$ dry mergers are likely more responsible for the further formation of these galaxies. \vspace{0.2cm} \noindent V. The resolved structures of galaxies also allows us to measure the internal features of galaxies and how they are assembling. There is some controversy over the formation history bulges, disks and bars, although many of these are likely formed by secular processes produced internally by disk dynamical evolution. This is an area where significant progress could be made in the next few years. The most up to date results suggest that the bar fraction for spiral galaxies at $z < 1$ depends upon the stellar mass of the galaxy. The most massive galaxies have a similar bar fraction at $z \sim 0.8$ as they do today, yet lower mass and bluer disk galaxies have a significantly lower bar fraction than similarly low mass nearby disks. This mirrors the evolution of the Hubble sequence itself where more massive galaxies settle into normal ellipticals and disks before lower mass galaxies. Spiral structure is a difficult problem and while some examples exist at high redshift, even at $z > 2$, the general onset of when disks form spirals is almost totally unconstrained by observations. \vspace{0.2cm} \noindent VI. Resolved imaging also permits us to measure the spectral energy distributions and colors, of galaxy components and individual pixels of different galaxies. This is another area where more work needs to be performed, but it appears that bulges of spirals tend to be older than their disks at high redshift, but there are examples of many ellipticals which have blue cores and central star formation (\S 4.1.1). Pixel-pixel analyses show that galaxies have a mixed star formation history, and that the inner parts of galaxies are often older than their outer parts. Pixel studies also show that the clumps seen in distant star forming galaxies are composed of young stellar populations, and thus must have recently formed or regenerate themselves. \vspace{0.2cm} \noindent VII. Perhaps the most popular (at present) problem in galaxy structural evolution is the apparent compactness in size of galaxies at high redshifts. The observations show that massive galaxies at $z > 1$ have sizes which are a factor of 2-5 smaller than similar massive galaxies in today's universe. This result has been studied in many different ways, and the sizes of a stellar mass selected sample of galaxies increases gradually as a function of redshift with a power-law increase $\sim (1+z)^{\beta}$, where $\beta$ varies from $-0.8$ to $-1.5$ depending upon whether the selected samples are disk-like or elliptical-like (\S 4.2). Results to date suggest that these galaxies are building up their outer parts over time to become larger systems, rather than adding mass to their centers. This process is unlikely driven by star formation, and theory suggests that this formation is produced by minor merger events (\S 4.2). In summary, we have learned much about galaxy morphology and structure over the past 15 years. There are however many open questions still remaining on all aspects of using structure to determine evolution. More work needs to be done in tying galaxy structure to underlying physics, both through empirical work and in simulations. Furthermore, the time-scales for structural features such as mergers and large clump survival are critical to better understand. Broad morphological features will remain important over the next decades as telescopes such as JWST, Euclid, LSST, the Dark Energy Survey, amongst others, will all resolve many more galaxies than we can currently study, and at higher redshifts. This opens up entirely new possibilities, and the with careful thoughtful planning a new revolution in galaxy structure may be upon us soon. A review such as this is written with help from many people. In particular I thank Alice Mortlock, Asa Bluck, Fernando Buitrago, Jamie Ownsworth, and Ken Duncan for illuminating conversations and collaboration on these topics over the past few years. I also thank Jennifer Lotz, Fernando Buitrago, Stijn Wuyts, Alice Mortlock for kindly providing figures. I personally thank the STFC and the Leverhulme Trust in the UK, as well as the NSF and NASA in the USA for financial support. | 14 | 3 | 1403.2783 |
1403 | 1403.5676_arXiv.txt | Recently it has been proposed that dark matter axions from the galactic halo can produce a small Shapiro step-like signal in Josephson junctions whose Josephson frequency resonates with the axion mass [C. Beck, PRL 111, 231801 (2013)]. Here we show that the axion field equations in a voltage-driven Josephson junction environment allow for a nontrivial solution where the axion-induced electrical current manifests itself as an oscillating supercurrent. The linear change of phase associated with this nontrivial solution implies the formal existence of a large magnetic field in a tiny surface area of the weak link region of the junction which makes incoming axions decay into microwave photons. We derive a condition for the design of Josephson junction experiments so that they can act as optimum axion detectors. Four independent recent experiments are discussed in this context. The observed Shapiro step anomalies of all four experiments consistently point towards an axion mass of $(110\pm 2) \; \mu$eV. This mass value is compatible with the recent BICEP2 results and implies that Peccei-Quinn symmetry breaking was taking place after inflation. | About 95\% of the energy contents of the universe appears to be of unknown origin, in the form of dark matter and dark energy. While there is a lot of astrophysical evidence for the existence of dark matter and dark energy, a deeper understanding of the physical nature of these main ingredients of the universe is still lacking. Clearly it is important to design new experiments on earth that could have the potential to unravel some of the unknown physics underlying dark matter and dark energy. At the particle physics level, there are two main candidates what dark matter could be. These are WIMPS (weakly interacting massive particles) \cite{bertone} and axions \cite{duffy, peccei, wilczek-old, sikivie-old}. WIMPS are motivated by supersymmetry, whereas axions are motivated by the solution of the strong CP problem in QCD. Various experimental searches to detect WIMPS \cite{lab1,bernabei} and axion-like particles \cite{lab2,lab3,lsw1,admx1} on the earth are currently going on. Very recently, there have been a couple of new suggestions how one could possibly detect dark matter axions in laboratory experiments on the earth \cite{graham,prl2013,sikivie2014}. All these proposals have in common that they are based on relatively small devices and that they suggest to look for small oscillating electric currents induced by axion flow, with a frequency given by the axion mass. Proposal 1 \cite{graham} is based on a technique similar to nuclear magnetic resonance (NMRI), known from medical imaging. Proposal 2 \cite{prl2013} is based on resonance effects in Josephson junctions. Proposal 3 \cite{sikivie2014} suggests to use LC circuits cooled down to mK temperatures. Further interesting proposals are based on topological magnetic insulators \cite{zhang} and atomic systems \cite{roberts}. In this paper we present a detailed calculation describing the physics of proposal 2, starting from the field equations of axion electrodynamics in a Josephson environment. In contrast to axions in vacuum, in a Josephson junction the axion has the possibility to induce electric supercurrents, rather than just ordinary currents. Our main result presented in this paper is that, besides the trivial solution where the axion passes through the Josephson junction without interaction, there is a nontrivial solution to the axion field equations due to these supercurrents. We show that the nontrivial solution implies the existence of a huge (formal) axion-flow generated magnetic field in a tiny surface area of the weak-link region of the junction, which makes incoming axions decay into microwave photons. The axion flow from the galactic halo through the junction then leads to a small measurable excess current of Cooper pairs, for which we will derive a concrete formula. The experimental consequence of this are Shapiro steps \cite{shapiro,tinkham} generated by axion flow, which are small but observable provided certain conditions on the design of the Josephson junction are satisfied. We will derive these conditions explicitly. An experiment by Hoffmann et al. based on S/N/S Josephson junctions \cite{hoffmann}, discussed in detail in \cite{prl2013}, provided evidence for an axion mass of 110 $\mu eV$ and an axionic dark matter density of about 0.05 GeV/$cm^3$ if interpreted in this way. Here we will discuss the results of four different experiments \cite{hoffmann,golikova,he,bae}. In all four cases small Shapiro step-like anomalies have been observed that, if interpreted within our theory, point towards an axion mass of $m_ac^2=(110\pm 2) \mu$eV. The predicted axion mass value has profound cosmological implications. If this value is confirmed by further experiments, it means that the Peccei-Quinn symmetry breaking took place {\em after} inflation \cite{visinelli2}. Employing the recent results of \cite{visinelli2,bicep} our result implies that the fractional contribution $\alpha^{dec}$ to the cosmic axion density from decays of axionic strings and walls is $\alpha^{dec}=0.66 \pm 0.05$. This paper is organized as follows: In section 2 we write down the axion field equations in a Josephson junction. The nontrivial solution, where the axion-induced electric current manifests itself as a supercurrent within the junction, is discussed in section 3. The physical interpretation of this solution is further worked out in section 4. In section 5 we present a short calculation how S/N/S Josephson junctions should be designed in order to serve as optimum axion detectors. Section 6 discusses some experimental candidate signals seen in various Josephson experiments that could possibly be associated with the nontrivial solution of section 3. Section 7 compares our mass estimate from Josephson resonances with cosmological and astrophysical bounds on the axion mass. Finally, our concluding remarks are given in section 8. \ | In this paper we have presented a detailed derivation why axions can generate small measurable electric currents in Josephson junctions. We started from the field equations of axion electrodynamics, plus the assumption that the axion-induced electric current can manifest itself as a supercurrent in a Josephson junction. We found a nontrivial solution of the axion field equations in the weak-link region of the junction, for which the phase grows linearly in time. This was interpreted in terms of axions decaying via Josephson radiation and triggering at the same time additional Cooper pair flow through the junction. A huge formal magnetic field appeared in our calculations, making the axions decay if they are still present at the surface of the junction, but the final result for the axion-generated additional critical current $I_c^a$ as given by eq.~(\ref{ic1}) is actually independent of the precise value of this formal $\vec{B}$-field, as it drops out of the equations. Overall the effect of the galactic axion background is small but measurable and the decaying axions produce a small Shapiro step-like feature when the axion mass resonates with the Josephson frequency. We derived concrete formulas for the additional critical current $I_c^a$ generated by axion flow, and derived conditions for different types of Josephson junctions to act as optimum axion detectors. The measured voltage where axion-generated Shapiro steps occur can be used to estimate the axion mass $m_a$, and their intensity can be used to estimate the axionic dark matter density $\rho_a$ near the earth \cite{prl2013}. We discussed peculiarities in the Shapiro step patterns measured by four different experimental groups for very different types of Josephson junctions. All four experiments point towards an axion mass of 110 $\mu$eV. Further systematic measurements should still be performed to test whether these candidate signals are really due to axions. The axion mass value of 0.11meV to which the various Josephson experiments point to is compatible with current astrophysical and cosmological bounds on the axion mass. | 14 | 3 | 1403.5676 |
1403 | 1403.0862_arXiv.txt | Impact craters exist on various solid objects in the planetary system. A simplified analogy of the process of their formation is here analyzed by standard solid state physics and the so called dynamic quantized fracture mechanics. An expression which links the crater volume to the parameters of the impactor and the target is obtained within the two approaches. For low impactor energy, this expression is of the same mathematical form as the one resulting from recent experiments.It is shown that the formation of an impact crater is possible even without heating of the target, if the critical stress in the target satisfies certain conditions. The critical value of the stress needed for the occurence of a fracture is calculated for three craters: two terrestrial and one lunar crater. The approach presented here uses only measurable material parameters, and is therefore more realistic than the treatement of the same problem using the cohesive energy of materials. | The existence of some craters on the surface of the Earth is due to impacts of small bodies in the planet. A list of such craters is avaliable at the web site http://www.passc.net/EarthImpactDatabase/index.html. The study of these craters, and the constraints which their existence places on the impactors has become a separate field of research in planetary science (for example [1]). Two probably best known examples of such events are the impacts which have led to the formation of the Barringer crater in Arizona, and the Tunguska event of 1908. What can be concluded about the impactors by combining astronomical data with results of solid state physics? It was recently shown ( [2] and related work) that by using basic principles of condensed matter physics, it becomes possible to derive an expression for the product $\rho_{1} r_{1}^{3} v_{1}^{2}$, where the three symbols denote the mass density, radius and speed of the impactor. In that calculation, the notion of cohesion energy of a solid was used. This quantity is defined as the energy needed to transform a sample of a solid into a gas of widely separated atoms. It is not easy to measure experimentally, and it is not related to the practical strength of solids, which is regulated by their resistance to fracture [3]. It this letter we investigate a simplified analogy of the formation of an impact crater. The object we investigate is a hole of given dimensions, resulting from the impact of an external object in a material with known parameters. Using the dimensions of the "hole" as the final result of the impact and parameters of the target, what can be concluded about the impactor? The aim of this paper is to derive an expression for $\rho_{1} r_{1}^{3} v_{1}^{2}$ by using the notion of stress instead of the notion of cohesive energy. The stress is defined as the ratio of the force applied on a body to the cross section of the surface of a body normal to the direction of the force [4]. In order for a crater to form in the target as a result of an impact, stress must become sufficiently high so as to allow the formation of a fracture in the material of the target. The following section contains a brief reminder of the required definitions and previous results, while the third part is devoted to the calculations. A discussion of the results and the conclusions are presented in the fourth and fifth sections. The approach used in the present work is based on general principles of solid state physics, and it can be applied to any material. The main tool used in derivation of scaling theories, discussed for example in [6],is dimensional analysis. See also [5]. According to the scaling theory, one of the factors on which the volume of an impact crater depends is the gravitational acceleration at the surface of the target. In the section devoted to discussion it will be shown that the calculation reported here leads to the same result. This paper uses solid state physics throughout. This means that it is assumed that the material of the target does not melt in the impacts,which implies small kinetic energies of the impactors. Various aspects of heating in mutual collisions of solids has recently been discussed ; examples are [7], [8]. Apart from fundamental interest, studies of impact craters and the projectiles which have made them have a very "practical" motivation. Impacts into the Earth have been occuring throughout the history of our planet,and will occur again. An impact, if sufficiently energetic, could have serious consequences for the region where it occurs, or the planet as a whole, so predicting the outcome of such events is highly important. At the time of this writing, the last example of such an event is the impact of a small asteroid designated $2014AA$ into the Atlantic on January 2,2014. | In this letter we have discussed to some extent the process analogous to the formation of impact craters on the surfaces of the solid object in the planetary system. This problem is of high practical importance, because impacts of small bodies into the Earth have been occuring and will occur again,as testified by the impact of asteroid $2014AA$ in the Atlantic on January 2, 2014. The approach to the problem discussed here, differs from the one used in [2] in the physics used. In [2] the calculation was performed using the cohesion energy (of the material of the target). The treatement used here is a distinct advantage because it uses only measurable quantities-various parameters of the target and the impactor. The values of the critical stress needed for the occurence of a fracture were calculated for three impact craters: two terrestrial and one lunar,and the values obtained differ by one order of magnitude. This difference can be ascribed to two factors: varying quality of initial data, and the real physical difference of the materials on the three crater sites. Note that the value of $\sigma_{C}$ of a material depends on the chemical composition. In the calculation for the Barringer crater it was assumed that the most abundant mineral there is $SiO_{2}$. Changing this assumption would change the values of $\sigma_{C}$ and $C_{V}$. It will be attempted to improve the results of the approach discussed in this letter by taking into account more physical details of the process of formation of impact craters. Note that within the $DQFM$ theory the value of the critical stress needed for a material to fracture, and therefore a crater to form, depends also on the geometrical parameters of cracks existing in the material. This conclusion potentially opens the possibility of applying the $DQFM$ in terrestrial laboratory work. Details will be discussed in the future. | 14 | 3 | 1403.0862 |
1403 | 1403.5506_arXiv.txt | {} {This paper introduces a kinetic code that simulates gamma-ray burst (GRB) afterglow emission from the external forward shock and presents examples of some of its applications. One interesting research topic discussed in the paper is the high-energy radiation produced by Compton scattering of the prompt GRB photons against the shock-accelerated electrons. The difference between the forward shock emission in a wind-type and a constant-density medium is also studied, and the emission due to Maxwellian electron injection is compared to the case with pure power-law electrons.} {The code calculates the time-evolving photon and electron distributions in the emission region by solving the relativistic kinetic equations for each particle species. For the first time, the full relativistic equations for synchrotron emission/absorption, Compton scattering, and pair production/annihilation were applied to model the forward shock emission. The synchrotron self-absorption thermalization mechanism, which shapes the low-energy end of the electron distribution, was also included in the electron equation.} {The simulation results indicate that inverse Compton scattering of the prompt GRB photons can produce a luminous $\gtrsim \textrm{TeV}$ emission component, even when pair production in the emission region is taken into account. This very high-energy radiation may be observable in low-redshift GRBs. The test simulations also show that the low-energy end of a pure power-law distribution of electrons can thermalize owing to synchrotron self-absorption in a wind-type environment, but without an observable impact on the radiation spectrum. Moreover, a flattening in the forward shock X-ray light curve may be expected when the electron injection function is assumed to be purely Maxwellian instead of a power law. The flux during such a flattening is likely to be lower than the {\it Swift}/XRT sensitivity in the case of a constant-density external medium, but a wind environment may result in a higher flux during the shallow decay.} {} | Gamma-ray burst (GRB) afterglows are produced by relativistic electrons radiating mainly via the synchrotron and inverse Compton mechanisms. According to the standard afterglow model, the electrons are accelerated to highly relativistic energies at two shock fronts, the forward shock and the reverse shock, which are the result of the interaction between the relativistic jet from the GRB central engine and the surrounding medium \citep[for reviews, see, e.g.,][]{2004RvMP...76.1143P,2006RPPh...69.2259M}. The earliest afterglow models invoke pure synchrotron radiation from the forward shock in a constant-density interstellar medium (ISM) or a wind-type environment and yield analytic time-evolving synchrotron spectra of the decelerating blast wave \citep{1998ApJ...497L..17S,2000ApJ...536..195C,2002ApJ...568..820G}. The role of inverse Compton scattering of the synchrotron photons has also been investigated, typically relying on an approximate treatment of the scattering process because no analytic solution for the inverse Compton spectrum is available \citep{1998ApJ...501..772P,1999ApJ...512..699C, 2000ApJ...543...66P,2001ApJ...548..787S}. The GRB observations by the {\it Swift} and {\it Fermi} satellites have revealed some surprising features in the afterglow and prompt light curves, resulting in a need to improve the models for GRB emission. For example, the {\it Fermi}/LAT telescope has observed $> 100\:\textrm{MeV}$ emission from several GRBs. In the literature, the high-energy radiation has been attributed to the prompt emission \citep[e.g.,][]{2009ApJ...706L.138A}, to pure synchrotron radiation from the external shock \citep{2009ApJ...706L..33G,2010MNRAS.403..926G,2010MNRAS.409..226K}, to a combination of external synchrotron photons and synchrotron self-Compton emission \citep{2013ApJ...771L..13T,2013ApJ...771L..33W,2013ApJ...773L..20L,2013ApJ...776...95F}, and to a superposition of the prompt and afterglow emission \citep{2011MNRAS.415...77M}. Another possibility is that some of the high-energy emission stems from prompt photons being Compton scattered to higher energies by the afterglow-emitting electrons \citep{2005ApJ...618L..13B,2012ApJ...753..178H,2013arXiv1307.2663B,2013ApJ...776...95F}. Owing to the {\it Swift} observations, it has been discovered that a typical X-ray light curve begins with a phase of steeply decaying flux, which is often followed by a shallow decay segment. The late-time afterglow, on the other hand, can often be explained by the standard synchrotron model \citep{2006MNRAS.366L..13G,2006MNRAS.369..197F}. Energy injection to the blast wave is currently the most popular explanation for the shallow decay phase observed both in X-ray and optical afterglows \citep{2006MNRAS.366L..13G,2006MNRAS.369..197F, 2006ApJ...642..389N,2006ApJ...642..354Z,2011MNRAS.414.3537P,2012ApJ...758...27L}. Other models introduced to explain the shallow decay phase include the evolution of microphysical parameters \citep{2006MNRAS.369.2059P,2006MNRAS.370.1946G}, emission due to an outflow ejected before the prompt GRB \citep{2009ApJ...690L.118Y, 2012MNRAS.422..393B}, an off-axis viewing angle of the jet \citep{2006ApJ...641L...5E}, dust scattering of X-rays \citep{2007ApJ...660.1319S}, late prompt emission \citep{2007ApJ...658L..75G,2011ApJ...732...77M}, and an adiabatic evolution of the shock following a radiative phase \citep{2007ApJ...664..384D}. It has also been suggested that the main contribution to the afterglow could come from a long-lived reverse shock, which may also explain the shallow decay phase in the X-ray afterglows \citep{2007ApJ...665L..93U,2007MNRAS.381..732G}. Models aiming to explain all the different slopes seen in the light curves have also been presented, including accretion of different layers of the progenitor star \citep{2008Sci...321..376K} and the curvature effect that is usually only invoked to explain the early steep X-ray decay \citep{2008ApJ...683..900Q}. The evolution of the GRB blast wave is described well by the self-similar solution by \citet{1976PhFl...19.1130B}, which is valid in the deceleration phase while the blast is still highly relativistic. The evolution in the late non-relativistic phase is given by the Sedov-Taylor solution \citep{1959sdmm.book.....S,1950RSPSA.201..159T}. A mechanical model of the relativistic blast ensuring mass, energy, and momentum conservation has been presented by \citet{2006ApJ...651L...1B}, and it is nearly identical to the Blandford-McKee solution at late times after the shock has started to decelerate. However, the mechanical model gives an accurate description of the blast also before the deceleration time, while earlier models unphysically assume an equal pressure at the forward and reverse shock. The evolution of the shell in the mildly relativistic phase can be found by means of hydrodynamic simulations, which can then be coupled to a radiation code to find the radiation spectrum from the shock. Such simulations can also be applied to calculate the afterglow emission for an observer with an off-axis viewing angle \citep{2010ApJ...722..235V}. Results of one- and two-dimensional hydrodynamic simulations of the blast wave have been presented by \citet{1999ApJ...513..669K}, \citet{2007MNRAS.376.1189M}, \citet{2009A&A...494..879M}, and \citet{2010ApJ...716.1028R} but without discussing the radiation mechanism of the afterglow. A calculation of the synchrotron radiation from the blast has been coupled to the hydrodynamic simulations of \citet{2002MNRAS.332..144D}, \citet{2009ApJ...698.1261Z}, \citet{2010MNRAS.403..300V,2011MNRAS.410.2016V}, and \citet{2011ApJ...738L..23W}. However, none of these works calculate the afterglow component due to Compton scattering, which is expected to appear at high energies. Simulations including an accurate treatment of both synchrotron and Compton processes, as well as pair production, have been presented by \citet{2009A&A...507..599P} (PM09), who use the solution of \citet{1976PhFl...19.1130B} to evaluate the evolution of the emitting shell. The code developed by PM09 is similar to the one presented in this paper, but it does not account for the electron heating due to synchrotron self-absorption. For the first time, we present simulations of afterglow emission from the forward shock with a relativistic kinetic code that treats synchrotron emission and absorption, Compton scattering, and electron-positron pair production/annihilation in a self-consistent way. The kinetic equations determining the time evolution of the electron and photon distributions are solved simultaneously at each timestep. We also consider electron heating due to synchrotron self-absorption, which shapes the electron distribution at low energies. Our treatment accounts for the fact that electrons injected at different times also have different cooling histories. For example, the magnetic field that determines the synchrotron cooling rate evolves while the electrons are cooling. It follows that there are no sharp cooling breaks in the electron distribution \citep{2014ApJ...780...82U}. The current version of the code applies a one-zone model of the emission region. It does not account for the different locations of the particles behind the shock and assumes a constant magnetic field throughout the shell. A more accurate treatment of synchrotron emission would require a model of the magnetic field structure behind the shock. Also, knowledge of the spatial photon and electron distributions is required for an exact calculation of Compton scattering. As an example of the applications of the code, we report the results of a simulation where the afterglow-emitting electrons interact with an external source of photons roughly corresponding to prompt GRB emission. The shocked electrons are expected to upscatter a small fraction of the prompt photons to GeV$-$TeV energies as long as the prompt emission overlaps with the shocked electrons \citep{2005ApJ...618L..13B,2005ApJ...629..334F}. Some of the high-energy photons then produce pairs with the prompt MeV photons, which in turn are able to scatter radiation to higher energies. In addition, we compare the forward shock emission in a wind environment with the emission in a constant-density ISM. The results indicate that a power-law electron distribution can thermalize at low energies thanks to synchrotron self-absorption heating in a wind medium with a typical density structure expected from the surroundings of a Wolf-Rayet star. Along with the ambient density, the importance of thermalization is mainly determined by the fraction of shock-generated energy given to the magnetic field. Our simulations imply that the thermalized electrons are unlikely to produce an observable signature in the afterglow spectrum. In our final example, we study the difference between the forward shock radiation due to Maxwellian and power-law electron injection. The standard afterglow model assumes that the injection function is a pure power law, even though a large fraction of the shock-generated energy goes to a thermal population of electrons. We find that pure Maxwellian injection can lead to a flattening in the X-ray light curve. The flux during this phase is found to be very low compared to {\it Swift}/XRT detections for a constant-density ISM, but detectable flux levels during the shallow decay may be achieved in a wind-type environment. | In this paper, we have studied the modeling of GRB forward shock emission for the first time with a code that treats all the radiation processes self-consistently. The main advantages of our code are the inclusion of synchrotron self-absorption heating and Compton scattering as mechanisms for electron thermalization, and the exact treatment of Compton scattering and pair production. We have shown that our results are in very good agreement with both analytic estimations and numerical simulations that have been presented in the literature. Inclusion of an accurate calculation of radiative transport, a multi-zone treatment of the emitting region, and the emission from the reverse shock will be the topics of future research. We presented the results of test simulations where external MeV photons, distributed similarly to prompt GRB emission, interact with the shock-accelerated electrons and are Compton scattered to TeV energies. Such VHE emission is likely to dominate the synchrotron self-Compton emission at times up to $t \gtrsim 10 \:\textrm{s}$ in GRBs with similar parameter sets and might be an interesting target for observations of low-redshift bursts with future instruments such as the Cherenkov Telescope Array \citep{2013APh....43..252I}. We also showed that the inverse Compton flux is not always quenched by electron-positron pair production, owing to the decrease in the pair production opacity before the deceleration time. A very prominent inverse Compton peak appears during the prompt emission, the location of which is determined by the maximum energy given to the photons by the peak electrons. Even though the test simulations show that the inverse Compton scattering of photons distributed according to a Band function is likely to produce a high-energy spectral component, the current model does not explain the typical long duration of the observed LAT emission in the case of GRB 090902B. A possible explanation for the LAT data is that the unscattered prompt photons are traveling within a smaller solid angle than the scattered photons, and the arrival time of the upscattered radiation is delayed as a result \citep{2005ApJ...618L..13B,2013arXiv1307.2663B}. This effect is not accounted for in the current version of the code. However, the large-angle effect cannot explain the LAT emission if it lasts much longer than the prompt GRB. As another example, we compared the synchrotron emission in a wind-type and a constant-density ISM medium for a set of typical parameters and studied the effect of different simulation parameters on the thermalization of low-energy electrons. As long as the synchrotron cooling time is shorter than the lifetime of the flow, the low-energy end of a power-law electron distribution is strongly thermalized due to synchrotron self-absorption heating, which is usually neglected in similar numerical codes. The shape of the emerging radiation spectrum is still nearly identical to the case with pure power-law electrons, although a small bump is seen in the synchrotron spectrum thanks to the Maxwellian shape of the low-energy electrons. However, the simulation shows that it is necessary to include the self-absorption heating term in the electron kinetic equation to prevent the electrons from cooling down to the low-energy edge of the grid as long as thermalization is important. Additionally, we have demonstrated that a flattening or rebrightening segment in the forward shock light curve is expected in the case of purely Maxwellian electron injection in a waveband where the observed flux gradually becomes dominated by the inverse Compton component, as described by \citet{2011A&A...531A..76P}. Such flattenings are observed in a large fraction of X-ray light curves, but their physical origin is still under debate. In the ISM simulations, the flux during the flat segments is much lower than in the case of a typical shallow decay phase observed by {\it Swift}/XRT. However, the shallow decay flux in the wind case could be detectable. | 14 | 3 | 1403.5506 |
1403 | 1403.4927_arXiv.txt | The BICEP-2 team has reported the detection of primordial cosmic microwave background B-mode polarization, with hints of a suppression of power at large angular scales relative to smaller scales. Provided that the B-mode polarization is due to primordial gravitational waves, this might imply a blue tilt of the primordial gravitational wave spectrum. Such a tilt would be incompatible with standard inflationary models, although it was predicted some years ago in the context of a mechanism that thermally generates the primordial perturbations through a Hagedorn phase of string cosmology. The purpose of this note is to encourage greater scrutiny of the data with priors informed by a model that is immediately falsifiable, but which \textit{predicts} features that might be favoured by the data-- namely a blue tensor tilt with an induced and complimentary red tilt to the scalar spectrum, with a naturally large tensor to scalar ratio that relates to both. | The BICEP-2 team just announced the detection of primordial cosmic microwave background (CMB) B-mode polarization, implying a tensor-to-scalar ratio of $r = 0.2 \pm 0.05$ \cite{BICEP}. The positive detection of primordial gravitational waves constitutes a major advance for early universe cosmology, giving us a new diagnostic tool with which to scrutinize models of the very early universe against observational data. Conventional adiabatic cosmological fluctuations do not predict any B-mode polarization at the linear level in cosmological perturbation theory. Hence in the context of the simplest models, primordial B-mode polarization must be due to gravitational waves \footnote{Note, however, that beyond the simplest single field scalar models, there are other sources of B-mode polarization e.g. from cosmic strings \cite{Holder}. B-mode polarization will also be produced by lensing of E-mode polarization, which in turn is directly generated from cosmological fluctuations, and that this B-mode lensing signal has in fact recently been discovered by the South Pole \cite{Hanson} and the Polarbear telescopes \cite{Dobbs}.}. Although the BICEP-2 collaboration's analysis took $n_T=0$ as a prior in its simulated data, we wish to ask whether a suppression of power in the BB angular power spectrum at large angular scales relative to smaller scales might be seen in the data, in particular in the B2 x Keck cross correlation function at long wavelengths (which is less sensitive to systematic noise\footnote{Which we note can only boost the auto-correlation function.}) and in the B2 x B2 correlation function, although the latter is more susceptible to contamination from foregrounds (see Fig. \ref{bicepfig}). Whether this suppression is statistically significant remains to be seen. If it is, it could be interpreted as indicative of a positive tilt of the primordial tensor spectrum at the largest angular scales. If this does turn out to be the case, then this result would be very hard to interpret in the context of the standard inflationary paradigm of early universe cosmology (see also \cite{contaldi} for an analysis of the additional tension between measuring a large $r$ with the small $k$ scalar power spectrum). Assuming that space-time is described by General Relativity and that matter obeys the ``weak energy condition", inflation generically predicts a red spectrum of gravitational waves, i.e. $n_T < 0$. This arises from the fact that the amplitude of the gravitational waves on a scale $k$ is set by the amplitude of the Hubble expansion rate $H$ at the time when that scale exits the Hubble radius, and that during inflation ${\dot{H}} < 0$. For single field slow roll models this relation is precisely \begin{figure}[t] \includegraphics[height=6cm]{speccomp.pdf} \caption{The BB auto and cross correlation functions as seen by the BICEP collaboration (courtesy of \cite{BICEP}).} \label{bicepfig} \end{figure} \be \label{sftilt} n_T = -2\epsilon \, , \ee with $\epsilon := -\dot H/H^2$. A related challenge for standard inflationary cosmology in light of the BICEP-2 data is that the tensor-to-scalar ratio $r = 0.2$ implies a large field excursion over the duration in which the observed modes in the CMB were produced \cite{Lyth}: \be \Delta\phi \gtrsim M_{pl} \, . \ee Constructing a model that safely accomplishes this is challenging to say the least from the perspective of effective field theory, as field excursions comparable to the cut-off of the theory typically generate large anomalous dimensions for operators that were initially suppressed (by appropriate powers of the cutoff), potentially spoiling the requisite conditions for inflation to occur as it progresses\footnote{The so called sensitivity of large field models to `Planck slop' \cite{BCQ}.}. However, this is not to say that this might not be accomplished in the context of some fundamental theory construction-- see \cite{SW} for an interesting claim (and \cite{JC} for a counter-claim)-- within the context of string theory, large field excursions certainly appear to be problematic \cite{swampland}. In this note we wish to remind cosmologists of a mechanism to generate the primordial perturbations from the thermodynamics of \textit{closed} strings in a quasi-static background, which-- \begin{itemize} \item Naturally generates a large tensor to scalar ratio; \item \textit{Predicts} a blue tilt to the tensor spectrum, \item with a complimentary red tilt to the scalar spectrum, both of which relate to $r$. \end{itemize} This construction relies upon a background that consists of a quasi-static initial state in the Einstein frame, whose specific realization can be addressed in the context of particular string constructions (see \cite{Florakis, KPT1} for some recent attempts), but whose existence we will take for granted in the following as far as the study of fluctuations is concerned, just as one typically does in the context of inflationary cosmology\footnote{Requiring that inflation exists in the context of a consistent quantum theory requires considerable tuning at the level of the low energy effective description \cite{ETA} (its so-called UV sensitivity).}. In fact, this is the very premise of the effective theory of the adiabatic mode \cite{Senatore}-- the so called effective theory of inflation. In what follows, we will first address plausible constructions that could give rise to the requisite background as motivation for the subsequent section-- the main focus of this note-- where we argue that \textit{the thermodynamics of closed strings in the early universe can naturally generate a large, blue titled tensor mode background.} Our goal is to provide observations with a novel, predictive, and falsifiable model which can inform the formulation of priors when analyzing the data in a manner that is easily contrasted against the predictions of inflationary cosmology. Whether there are hints for a blue tensor tilt in the data is to be viewed as secondary to the goal of providing a `straw-model' with which to contrast the predictions of inflation against, the scientific utility of which needing no further elaboration. | In generating the primordial perturbations from the thermodynamics of closed strings, one can isolate the background dependence to one free parameter, the ratio of string length and Planck length, and a function $T(k)$ which is close to the Hagedorn temperature and is a decreasing function of $k$. The precise $k$ dependence depends of $T$ encodes the details of the transition between the Hagedorn phase and the radiation phase, and will determine the relative tilts of the spectra (\ref{scalartilt}) and (\ref{tfinal}), which are to first approximation, equal and opposite (\ref{tfinal}). With these inputs we can also compute the amplitudes of the scalar and tensor spectra at any pivot scale, which is unique for only \textit{three spatial large compact dimensions} \cite{Ali}. As emphasized in \cite{BNPV}, the key result is that the tensor spectrum has a blue tilt, whereas the scalar fluctuations retain a red tilt. This feature distinguishes the string thermodynamic generation of the primordial perturbations from standard inflationary realizations. | 14 | 3 | 1403.4927 |
1403 | 1403.6325_arXiv.txt | }[2]{{\footnotesize\begin{center}ABSTRACT\end{center} \vspace{1mm}\par#1\par \noindent {~}{\it #2}}} \newcommand{\TabCap}[2]{\begin{center}\parbox[t]{#1}{\begin{center} \small {\spaceskip 2pt plus 1pt minus 1pt T a b l e} \refstepcounter{table}\thetable \\[2mm] \footnotesize #2 \end{center}}\end{center}} \newcommand{\TableSep}[2]{\begin{table}[p]\vspace{#1} \TabCap{#2}\end{table}} \newcommand{\FigCap}[1]{\footnotesize\par\noindent Fig.\ % \refstepcounter{figure}\thefigure. #1\par} \newcommand{\TableFont}{\footnotesize} \newcommand{\TableFontIt}{\ttit} \newcommand{\SetTableFont}[1]{\renewcommand{\TableFont}{#1}} \newcommand{\MakeTable}[4]{\begin{table}[htb]\TabCap{#2}{#3} \begin{center} \TableFont \begin{tabular}{#1} #4 \end{tabular}\end{center}\end{table}} \newcommand{\MakeTableSep}[4]{\begin{table}[p]\TabCap{#2}{#3} \begin{center} \TableFont \begin{tabular}{#1} #4 \end{tabular}\end{center}\end{table}} \newenvironment{references}% { \footnotesize \frenchspacing \renewcommand{\thesection}{} \renewcommand{\in}{{\rm in }} \renewcommand{\AA}{Astron.\ Astrophys.} \newcommand{\AAS}{Astron.~Astrophys.~Suppl.~Ser.} \newcommand{\ApJ}{Astrophys.\ J.} \newcommand{\ApJS}{Astrophys.\ J.~Suppl.~Ser.} \newcommand{\ApJL}{Astrophys.\ J.~Letters} \newcommand{\AJ}{Astron.\ J.} \newcommand{\IBVS}{IBVS} \newcommand{\PASP}{P.A.S.P.} \newcommand{\Acta}{Acta Astron.} \newcommand{\MNRAS}{MNRAS} \renewcommand{\and}{{\rm and }} {We present an analysis of the detached eclipsing binaries V44 and V54 belonging to the globular cluster M55. For V54 we obtain the following absolute parameters: $M_p=0.726\pm 0.015\,M_\odot$, $R_p=1.006\pm 0.009\,R_\odot$, $L_p=1.38\pm 0.07\,L_\odot$ for the primary, and $M_s=0.555\pm 0.008\,M_\odot$, $R_s=0.528\pm 0.005\,R_\odot$, $L_s=0.16\pm0.01\,L_\odot$ for the secondary. The age and apparent distance modulus of V54 are estimated at 13.3 -- 14.7 Gyr and $13.94\pm0.05$ mag, respectively. This derived age is substantially larger than ages we have derived from the analysis of binary systems in 47 Tuc and M4. The secondary of V44 is so weak in the optical domain that only mass function and relative parameters are obtained for the components of this system. However, there is a good chance that the velocity curve of the secondary could be derived from near-IR spectra. As the primary of V44 is more evolved than that of V54, such data would impose much tighter limits on the age and distance of M55. } {binaries: close – binaries: spectroscopic – globular clusters: individual (M55) – stars: individual (V44-M55, V54-M55) } | The general goal of the Cluster AgeS Experiment (CASE) is to determine the basic stellar parameters (masses, luminosities, and radii) of the components of globular cluster binaries to a precision better than 1\% in order to measure ages and distances of their parent clusters, and to test stellar evolution models (Kaluzny et al. 2005). Within the CASE series, this is the third paper devoted to the study of detached eclipsing binaries (DEBs) with main-sequence or subgiant components. In the previous two papers we analyzed three such systems in M4 (Kaluzny et al. 2013) and another one in 47 Tuc (Thompson et al. 2010). The eclipsing binaries M55-V44 and M55-V54 (henceforth V44 and V54) were discovered in the field of the globular cluster M55 (Kaluzny et al. 2010). Follow-up spectroscopy indicated that both systems are radial velocity members of the cluster, and their membership was confirmed by the proper motion study of Zloczewski et al. (2012). V44 and V54 are located in the turn-off region on the color-magnitude diagram (CMD) of the cluster. With orbital periods of 2.17~d and 9.27~d they are well detached, and their analysis can provide interesting constraints on the age and distance of M55. In this paper we analyze light and velocity curves of the SB2 system V54 and determine the absolute parameters of its components. We also obtain the mass function and a light-curve solution of the SB1 system V44. THe photometric and spectroscopic data are described in Section~\ref{sec:obs}. The analysis of the data is detailed in Section~\ref{sec:analysis}. The age and distance of M55 are derived in Section~\ref{sec:age}, and our results are summarized in Section~\ref{sec:sum}.\\ | \label{sec:sum} We have analysed photometric and spectroscopic observations of the eclipsing binaries V44 and V54 in M55. V54 is an SB2 system, and we obtained absolute parameters of the components of this system, leading to an estimate of age and distance modulus of the cluster. The resulting age of 13.3 -- 14.7 Gyr derived from $M-R$ and $M-L$ diagrams is compatible with the age 13 -- 14 Gyr obtained from CMD fitting with Dartmouth or Victoria-Regina isochrones. Apparent distance moduli derived with the same two methods and the same two sets of isochrones are also compatible within the errors. The best chance to tighten the age limits is offered by the mass-radius relation, which is independent of uncertainties in distance and extinction that plague ages obtained via isochrone fitting to the CMD. Since at the location of the primary on the $M-R$ plane the isochrones are almost perpendicular to the mass axis (that's why the corresponding ellipses in Fig. \ref{fig:ellipses} are very narrow), the age uncertainty arises mainly from the error in the mass of the primary. The only way to reduce it is by taking additional spectra of V54. Thompson et al. (2010) estimate that doubling of the present set of velocity measurements would improve the mass estimates by 33 per cent. The uncertainty of the location of V54 components on the $M-L$ plane is mainly due to the relatively poor accuracy of their luminosities which were found using absolute radii from Table \ref{tab:abs_parm} and effective temperatures estimated from $(B-V)-T_{eff}$ calibration of Castelli and Kurucz (2004). Unfortunately, the existing empirical calibrations are poorly defined for low mass stars with ${\rm [Fe/H]<-1.5}$ (Casagrande et al. 2010; Gonzalez-Hernandez \& Bonifacio 2009). A significant improvement is expected when results from the GAIA astrometric mission will become publicly available. Also our determination of the distance to M55 depends heavily on the adopted teperature scale. A better accuracy could be achieved based on the surface brighness in $V$, which is known to tightly correlate with the $V-K$ color (Di Benedetto 2005). Given that V54 is located in the outer uncrowded part of the cluster, accurate near-IR photometry of these systems with large telescopes is entirely feasible. Independent and potentially more accurate estimates of age and distance of M55 may follow from the analysis of the V44 system, whose primary is more evolved than the primary of V54 (see Fig.~\ref{fig:CMD}). While the secondary of V44 is too faint to measure its velocity in the visible domain, it seems to contribute enough light for successful measurements in the near IR. Accurate parameters of both systems could yield not only tight age limits, but also impose useful constraints on the primordial He content of the cluster. The age we derived for M55 is by $\sim2$ Gyr older than the ages we obtained from the analysis of eclipsing binaries in 47~Tuc (Thompson et al. 2010) and M4 (Kaluzny et al. 2013). This can be compared with recent result of Hansen et al. (2013) who, based on cooling sequences of white dwarfs, found that NGC 6397 ([Fe/H]=-2) is about 2 Gyr older than 47 Tuc. Hansen et al. stress the benefits of deriving ages from WDs, which are largely insensitive to metallicity, but are still subject to uncertainties in distance and reddening. In our case, sensitivity to metallicity is present but the uncertainties in distance and reddening are not. Since M55 and NGC 6397 have very similar metallicities, it is significant that these two more-or-less independent methods arrive at a similar age difference between metal-poor and metal-rich globular clusters. \Acknow { JK, MR and WP were partly supported by the grant DEC-2012/05/B/ST9/03931 from the Polish National Science Center. We thank Dr. Joao Alves for sending us the information about reddening in the M55 region. This series of papers is dedicated to the memory of Bohdan Paczy\'nski. } | 14 | 3 | 1403.6325 |
1403 | 1403.3393_arXiv.txt | Distance measurements to molecular clouds are important, but are often made separately for each cloud of interest, employing very different different data and techniques. We present a large, homogeneous catalog of distances to molecular clouds, most of which are of unprecedented accuracy. We determine distances using optical photometry of stars along lines of sight toward these clouds, obtained from PanSTARRS-1. We simultaneously infer the reddenings and distances to these stars, tracking the full probability distribution function using a technique presented in \citet{Green:2014}. We fit these star-by-star measurements using a simple dust screen model to find the distance to each cloud. We thus estimate the distances to almost all of the clouds in the \citet{Magnani:1985} catalog, as well as many other well-studied clouds, including Orion, Perseus, Taurus, Cepheus, Polaris, California, and Monoceros R2, avoiding only the inner Galaxy. Typical statistical uncertainties in the distances are 5\%, though the systematic uncertainty stemming from the quality of our stellar models is about 10\%. The resulting catalog is the largest catalog of accurate, directly-measured distances to molecular clouds. Our distance estimates are generally consistent with available distance estimates from the literature, though in some cases the literature estimates are off by a factor of more than two. | \label{sec:intro} Molecular clouds are the site of star formation, where all stars are born \citep{Blitz:1999}. The study of molecular clouds then informs critical elements of astrophysics, like the initial mass function of stars and the build-up of galaxies. Intense study has focused on the Milky Way's molecular clouds, the nearest and most accessible sites of star formation. The distances to these clouds are fundamental to deriving their basic physical parameters---like mass and size---from observations. But estimating the distance to molecular gas is difficult, and a number of different techniques have been explored and applied, often only to individual clouds of interest. These techniques are varied. A common method is to estimate cloud distances kinematically. In this technique a cloud's recessional velocity is measured by the Doppler shift of its spectral lines and converted to a distance by assuming that the cloud follows the Galactic rotation curve. This technique is widely applicable and has been used to estimate the distances to large numbers of molecular clouds (e.g., \citet{RomanDuval:2009}), but it is problematic in the presence of peculiar velocities and non-circular motions. A second method is to find the distance to objects associated with a cloud and to place the cloud at the same distance; for instance, many clouds have formed young OB associations of stars for which distances can be estimated. A third method is to estimate a cloud's distance from its reddening and absorption of starlight. Light passing through molecular clouds is extinguished by dust and gas; in particular, optical and infrared light is reddened by dust. This allows stars in the foreground of the cloud to be distinguished from stars in its background. By finding the distances to these stars, the distance to the cloud can be determined. Recently \citet{Lallement:2014} and \citet{Vergely:2010} have mapped the 3D distribution of the ISM in the solar neighborhood using this basic technique. A systematic study \citep{Lombardi:2001, Lombardi:2011, Lada:2009} using data from 2MASS has led to precise distance estimates for a number of clouds by counting the number of unextinguished foreground stars toward large molecular clouds and comparing with predictions for the distribution of stars from the Besan\c{c}on Galactic model \citep{Robin:2003}. We have developed a related technique: we simultaneously infer the distance and reddening to stars from their \PS\ \citep[PS1]{PS1_system} photometry and bracket clouds between foreground unreddened stars and background reddened stars. The use of only \PS\ photometry gives us access to three-quarters of the sky and hundreds of millions of stars, but has the disadvantage that the distances and reddenings we infer have strongly covariant uncertainties. We track the full probability distribution function of distance and reddening to each star, and model the results as produced by a screen of dust associated with the cloud, with an angular distribution given by the Planck dust map \citep{Planck:2011}. We then perform an MCMC sampling to determine the range of probable distances to the cloud. This paper is part of an ongoing effort to study the dust using \PS\ photometry. The basic method is presented in \citet{Green:2014}, while E. Schlafly et al. (2014, in preparation) demonstrates that the technique closely reproduces the widely-used reddening map of \citet[SFD]{Schlegel:1998}. This work serves additionally to demonstrate the 3D power of the method, recovering the distances to the Galaxy's molecular clouds. We measure the distances to many well-studied molecular clouds in the Galaxy: Orion, $\lambda$ Orionis, Taurus, Perseus, California, Ursa Major, the Polaris Flare, the Cepheus Flare, Lacerta, Pegasus, Hercules, Camelopardis, Ophiuchus, and Monoceros R2. We additionally estimate the distances to most of the clouds of the \citet{Magnani:1985} catalog of high Galactic latitude molecular clouds, though in some cases the cloud does not fall within the PS1 footprint. Our distances are often consistent with, but more precise than, other available distance estimates, though we find that occasionally the literature distance estimates are off by as much as a factor of two. In this work we avoid clouds in the inner Galaxy. In principle we could apply this technique there as well, but these clouds require more sophisticated modeling of the potentially many molecular clouds on each line of sight through the disk. We accordingly defer analysis of the inner Galaxy to later work. We describe in \textsection\ref{sec:ps1} the \PS\ survey, which provides the optical photometry on which this work is based. In \textsection\ref{sec:method} we describe our method for determining the distances to the dust clouds. In \textsection\ref{sec:results}, we apply our technique to sight lines through molecular clouds in the \PS\ footprint, and present a catalog of cloud distances. In \textsection\ref{sec:sysunc} and \textsection\ref{sec:discussion}, we discuss the systematic uncertainties in the method and the implications of the results in light of the literature. Finally, we conclude in \textsection\ref{sec:conclusion}. | \label{sec:conclusion} We present a catalog of distances to molecular clouds. We obtain secure distance estimates to most of the high-latitude MBM molecular clouds, with accuracy typically limited by systematics to 15\%. We further obtain distances to a number of well-studied molecular clouds at high latitudes or in the outer Galaxy. We obtain secure distances to Lacerta, Pegasus, the Cepheus Flare, the Polaris Flare, Ursa Major, Camelopardis, Perseus, Taurus, $\lambda$ Orionis, Orion A, and Monoceros R2, in general confirming but refining the accepted distances to these clouds. We highlight the complexity of the Cepheus Flare region, separating the molecular gas there into nearby (360~pc) and distant (900~pc) parts. We correct the literature distance to Ursa Major and clarify the distance to the Polaris Flare, placing them at the distance of the nearby component of the Cepheus Flare. We make the first distance estimates of which we are aware for a number of clouds, including clouds in Camelopardis and Pegasus, as well as a number of the MBM clouds. Our distance estimates reach statistical uncertainties better than 5\% in 0.2\degree\ radius lines of sight, though we caution that we only expect absolute accuracies of 10\% owing to systematic uncertainties in our technique. This work has only scraped the surface of what is possible with the combination of this technique and the PS1 photometry. We have focused on a sampling of sight lines through well-studied molecular clouds, but in principle the entire $\delta > -30\degree$ sky is amenable to this analysis. The 5\% relative distance accuracy suggests that we can make 3D maps of major molecular clouds. At the distance of Orion, we would expect to be able to obtain 20~pc distance resolution, compared with the $\sim 80$~pc projected angular size of the Orion A and B complex. This initial work has aimed only to get accurate overall distances to major molecular clouds, but already hints of distance gradients across clouds like Orion and Perseus are present. Moreover, the data used in this study come from the PS1 single-epoch data. The PS1 Science Consortium is rapidly improving deeper data coming from stacks of the $\sim 7$ PS1 exposures of each part of the sky in each filter. This stacking process increases the limiting magnitude of the survey by about a magnitude, allowing access to stars 50\% further away than we currently consider. Additionally, PS1 parallax and proper motion studies are beginning to bear fruit, opening the possibility of incorporating astrometric information into our distance estimates. Such an effort would dramatically improve our distance and reddening measurements to individual stars and serve as a useful pathfinder to the Gaia mission. We additionally look forward to eventually including other sources of data in our analysis. Our technique is naturally extendable to accept other sources of photometry. Of particular interest is infrared photometry, which, as demonstrated by \citet{Lombardi:2011}, can strongly constrain the distances to molecular clouds. Infrared data would allow us to probe higher column density clouds and improve our ability to discriminate variation in stellar temperature from variation in reddening. Existing data from 2MASS would already be helpful; upcoming infrared surveys like the UKIRT Hemisphere Survey and the VISTA Hemisphere Survey, which are better matched to the depth of PS1, will provide still better leverage. ES acknowledges support from the DFG grant SFB 881 (A3). GMG and DPF are partially supported by NSF grant AST-1312891. N.F.M. gratefully acknowledges the CNRS for support through PICS project PICS06183. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France. 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, the University of Maryland, and Eotvos Lorand University (ELTE). | 14 | 3 | 1403.3393 |
1403 | 1403.1675_arXiv.txt | Measurements of the mass and age of young stars from their location in the HR diagram are limited by not only the typical observational uncertainties that apply to field stars, but also by large systematic uncertainties related to circumstellar phenomena. In this paper, we analyze flux calibrated optical spectra to measure accurate spectral types and extinctions of 283 nearby T Tauri stars. The primary advances in this paper are (1) the incorportation of a simplistic accretion continuum in optical spectral type and extinction measurements calculated over the full optical wavelength range and (2) the uniform analysis of a large sample of stars, many of which are well known and can serve as benchmarks. Comparisons between the non-accreting TTS photospheric templates and stellar photosphere models are used to derive conversions from spectral type to temperature. Differences between spectral types can be subtle and difficult to discern, especially when accounting for accretion and extinction. The spectral types measured here are mostly consistent with spectral types measured over the past decade. However, our new spectral types are 1-2 subclasses later than literature spectral types for the original members of the TW Hya Association and are discrepant with literature values for some well known members of the Taurus Molecular Cloud. Our extinction measurements are consistent with other optical extinction measurements but are typically 1 mag.~lower than near-IR measurements, likely the result of methodological differences and the presence of near-IR excesses in most CTTSs. As an illustration of the impact of accretion, spectral type, and extinction uncertainties on the HR diagrams of young clusters, we find that the resulting luminosity spread of stars in the TW Hya Association is 15-30\%. The luminosity spread in the TWA and previously measured for binary stars in Taurus suggests that for a majority of stars, protostellar accretion rates are not large enough to significantly alter the subsequent evolution. | Classical T Tauri stars are the adolescents of stellar evolution. The star is near the end of its growth and almost fully formed, with a remnant disk and ongoing accretion. The accretion/disk phase typically lasts $\sim 2-5$ Myr, though some stars take as long as 10 Myr before losing their disks and emerging towards maturity. Strong magnetic activity leads to pimply spots on the stellar surfaces. Some T Tauri stars are still hidden inside their disks, not yet ready to emerge. Manic mood swings change the appearance of the star and are often explained with stochastic accretion. Depression has been seen in lightcurves on timescales of days to years. Sometimes every classical T Tauri star seems as uniquely precious as a snowflake. T Tauri star properties were systematically characterized in seminal papers by e.g. \citet{Cohen1979,HBC1988,Basri1990,Valenti1993,Hartigan1995,Kenyon1995,Gullbring1998}. In the last decade, dedicated optical and IR searches revealed thousands of young stars, typically confirmed with spectral typing \citep[e.g.][]{Hillenbrand1997,Briceno2002,Luhman2004,Rebull2010}. However, significant differences in extinction and accretion properties between different papers and methods has lead to confusion in the properties of even the closest and best studied samples of young stars. Some of this confusion is exacerbated by stochastic and rotation variability of T Tauri stars. While manic and depressive periods provide fascinating diagnostics of the stellar environment and star-disk interactions, they also pose significant problems for assessing the stellar properties and evolution of the star/disk system. How disk mass, structure and accretion rate change with age and mass requires accurate spectral typing and luminosity measurements \citep[e.g.][]{Furlan2006,Sicilia2010,Oliveira2013,Andrews2013}. While median cluster ages provide an accurate relative age scale between regions \citep[e.g.][]{Naylor2009}, age spreads within clusters may be real or could result from observational uncertainties \citep[e.g.][]{Hartmann1998,Hillenbrand2008,Preibisch2012}. \begin{table*}[!t] \caption{Observation Setup and Log} \label{tab:obslog.tab} \begin{tabular}{cccccccccccc} &&&& \multicolumn{3}{c}{Blue Setup} &&& \multicolumn{3}{c}{Red Setup}\\ Telescope & Dates& Instrument & Slit & Grating & Wavelength & Res. &&&Grating & Wavelength & Res. \\ \hline Palomar & 18-21 Jan. 2008 & DoubleSpec & 1--4$^{\prime\prime}$ & B600 & 3000--5700 & 700 &&& R316& 6200--8700 & 500 \\ Palomar & 28-30 Dec. 2008 & DoubleSpec & 4$^{\prime\prime}$ & B600 & 3000--5700 & 700 &&& R316& 6200--8700 & 500 \\ Keck I & 23 Nov.~2006 & LRIS & $0\farcs7-1^{\prime\prime}$ &B400 & 3000--5700 & 900 &&& R400 & 5700--9400 & 1000 \\ Keck I & 28 May~2008 & LRIS & $1^{\prime\prime}$ &B400 & 3000--5700 & 900 &&& R400 & 5700--9400 & 1000 \\ \hline \end{tabular} \end{table*} The uncertainties in stellar parameters affect our interpretation of stellar evolution. For example, \citet{Gullbring1998} found accretion rates an order of magnitude lower than those of \citet{Hartigan1995} and attributed much of this difference to lower values of extinction. The \citet{Gullbring1998} accretion rates of $10^{-8}$ M$_\odot$~yr$^{-1}$ means that steady accretion in the CTTS phase accounts for a negligible amount of the final mass of a star. However, subsequent near-IR analyses have revised extinctions upward \citep[e.g.][]{White2001,Fischer2011,Furlan2011}. These higher extinctions would yield accretion rates of $10^{-7}$ M$_\odot$~yr$^{-1}$, fast enough that steady accretion over the $\sim 2-3$ Myr CTTS phase would account for $\sim 20-50$\% of the final stellar mass, or more with the older ages measured by \citet{Bell2013}. The uncertainties in stellar properties introduce skepticism in our ability to use young stellar populations to test theories of star formation and pre-main sequence evolution. For classical T Tauri stars, minimizing the uncertainties in spectral type, extinction, and accretion (often referred to as veiling of the photosphere by accretion) requires fitting all three parameters simultaneously \citep[e.g.][]{Bertout1988,Basri1989,Hartigan2003}. In recent years, such fits have received increasing attention and have been applied {\it HST} photometry of the Orion Nebula Cluster \citep{daRio2010,Manara2012}, broadband optical/near-IR spectra of two Orion Nebular Cluster stars \citep{Manara2013b}, and to near-IR spectroscopy \citep{Fischer2011,McClure2013}. In this project, we analyze low resolution optical blue-red spectra to determine the stellar and accretion properties of 283 of the nearest young stars in Taurus, Lupus, Ophiucus, the TW Hya Association, and the MBM 12 Association. This first paper focuses on spectral types and extinctions of our sample. The primary advances are the inclusion of blue spectra to complement commonly used red optical spectra and accretion estimates to calculate the effective temperatures and luminosities with a single, consistent approach for a large sample of stars. Discrepancies are found between our results and near-IR based extinction measurements. We then discuss how these uncertainties affect the reliability of age measurements. This work was initially motivated to calculate accretion rates from the excess Balmer continuum emission, which will be described in a second paper. A third paper in this series will discuss spectrophotometric variability within our sample. \begin{table} \caption{Flux Calibration} \label{fluxcal.tab} \begin{tabular}{ccc} Wavelength & G191B2B & LTT 3864 \\ \hline \AA & \multicolumn{2}{c}{Absolute Scatter in Fluxes} \\ \hline 3500 & 0.067 & 0.087\\ 4300 & 0.056 & 0.046\\ 5400 & 0.041 & 0.047\\ 6300 & 0.063 & 0.091 \\ 8400 & 0.061 & 0.089 \\ \hline Flux ratio & \multicolumn{2}{c}{Scatter in Flux Ratios} \\ \hline $F_{7020}$/$F_{7140}$ & 0.007 & 0.005\\ $F_{8400}$/$F_{6300}$ & 0.016 & 0.014\\ $F_{6300}$/$F_{5400}$ & 0.057 & 0.101 \\ $F_{4300}$/$F_{5400}$ & 0.034 & 0.051\\ $F_{3500}$/$F_{5400}$ & 0.048 & 0.087\\ \hline \end{tabular} \end{table} | This paper provides a consistent and robust set of spectral types and extinctions for 283 young stars, including many of the most well studied. The primary advances in this paper are the implementation of simultaneous measurements of the extinction, accretion continuum flux, and spectral type for accreting stars and a sufficient sample size to obtain a robust set of extinction-corrected spectral templates. The effects of veiling on spectral type and extinction are reduced when analyzing spectra with coverage from 4000--9000 \AA. A similar approach was recently used by \citet{Manara2013b} to investigate two stars with previously reported ages of 30 Myr. Their accurate spectral type and luminosity yielded an age of 2--3 Myr, consistent with the age of the parent cluster. An updated grid of photospheric M star templates will eventually be needed to account for the evolution of colors with pre-main sequence contraction. Unfortunately, this problem is challenging to solve because the T Tauri stars with known $A_V=0$ mag.~are those in the 7--10 Myr old TWA and the $\eta$ Cham association. No T Tauri star in a young ($<3$ Myr) region can be assumed to have $A_V=0$ (or any other $A_V$) based only on its colors, independent of a model template. The high binary fraction of young WTTSs \citep{Kraus2012} also affects their use as photospheric templates. Although we minimize the effects of gravity dependence by using T Tauri stars as templates, the gravity dependence between 1--10 Myr may still be significant and is not accounted for. Photometric samples of non-accretors are likely reliable, but degeneracy between spectral type and accretion continuum flux can lead to spectral type uncertainties of at least two subclasses. The approach to measuring spectral types and extinction in this paper can reach a luminosity accuracy of $\sim 0.1-0.2$ dex for most classical T Tauri stars, and should serve as a particularly useful guide in the analysis of broadband spectra obtained by {\it VLT}/X-Shooter (e.g., Manara et al. 2013ab) and for analysis of GAIA observations. The grid of spectral types should be improved and based on more direct measurements of effective temperature by comparing high resolution spectra to models. The spectral type-effective temperature conversions are also uncertain at present because model atmopsheres fail to reproduce some large spectral features for spectral types later than M4. The relationship between spectral type and extinction needs particular improvement between K5--M0.5, where the accuracy of our grid relative to other publications is especially uncertain. Our results rely upon the assumption that the accretion continuum flux is flat. However, the strength of the broadband accretion continuum should be measured with simultaneous broadband spectra. Finally, extinction measurements should include the effect of gravity on photospheric emission, following the gravity-dependent colors obtained by \citet{Covey2007} and \citet{Pecaut2013}. Ideally, some optically thin line ratios could be found and used to measure extinction, independent of gravity. | 14 | 3 | 1403.1675 |
1403 | 1403.0390_arXiv.txt | {Protoplanetary disks are vital objects in star and planet formation, possessing all the material, gas and dust, which may form a planetary system orbiting the new star. Small, simple molecules have traditionally been detected in protoplanetary disks; however, in the ALMA era, we expect the molecular inventory of protoplanetary disks to significantly increase. } {We investigate the synthesis of complex organic molecules (COMs) in protoplanetary disks to put constraints on the achievable chemical complexity and to predict species and transitions which may be observable with ALMA. } {We have coupled a 2D steady-state physical model of a protoplanetary disk around a typical T Tauri star with a large gas-grain chemical network including COMs. We compare the resulting column densities with those derived from observations and perform ray-tracing calculations to predict line spectra. We compare the synthesised line intensities with current observations and determine those COMs which may be observable in nearby objects. We also compare the predicted grain-surface abundances with those derived from cometary comae observations.} {We find COMs are efficiently formed in the disk midplane via grain-surface chemical reactions, reaching peak grain-surface fractional abundances $\sim$~10$^{-6}$--10$^{-4}$ that of the H nuclei number density. COMs formed on grain surfaces are returned to the gas phase via non-thermal desorption; however, gas-phase species reach lower fractional abundances than their grain-surface equivalents, $\sim$~10$^{-12}$--10$^{-7}$. Including the irradiation of grain mantle material helps build further complexity in the ice through the replenishment of grain-surface radicals which take part in further grain-surface reactions. There is reasonable agreement with several line transitions of \ce{H2CO} observed towards T~Tauri star-disk systems. There is poor agreement with \ce{HC3N} lines observed towards LkCa~15 and GO~Tau and we discuss possible explanations for these discrepancies. The synthesised line intensities for \ce{CH3OH} are consistent with upper limits determined towards all sources. Our models suggest \ce{CH3OH} should be readily observable in nearby protoplanetary disks with ALMA; however, detection of more complex species may prove challenging, even with ALMA `Full Science' capabilities. Our grain-surface abundances are consistent with those derived from cometary comae observations providing additional evidence for the hypothesis that comets (and other planetesimals) formed via the coagulation of icy grains in the Sun's natal disk. } {} | \label{introduction} Protoplanetary disks are crucial objects in star formation. They dissipate excess angular momentum away from the protostellar system, facilitate the accretion of matter from the natal cloud onto the new star, and contain all the material, dust and gas, which will likely go on to form a surrounding planetary system \citep[for a review, see, e.g.,][]{williams11}. The study of the detailed chemistry of these objects has gained impetus in recent years driven by the impending completion of the Atacama Large Millimeter/Submillimeter Array (ALMA). ALMA, with its unprecedented sensitivity and spatial and spectral resolution, will reveal, for the first time, the composition of protoplanetary disks on $\sim$~milliarcsecond scales, probing material $\lesssim$~10~AU of the parent star in objects relatively close to Earth ($\approx$~140~pc). This spatial resolution will be achievable using the most extended configuration (with maximum baseline, $B$~$\approx$~16~km) at its highest operational frequencies ($\nu$~$>$~275~GHz). This will allow the study of the detailed composition of the cold molecular material within the `planet-forming' region of nearby disks, which will advance our understanding of the process of planetary system formation, and help answer questions regarding the morphology and composition of our own Solar System. The molecules observed in protoplanetary disks have thus far been restricted to small species and associated isotopologues due to their relatively high abundance and simple rotational spectra leading to observable line emission. The sources in which these molecules have been detected are also limited to a handful of nearby, and thus well-studied, objects. Molecules have been observed at both infrared (IR) and (sub)mm wavelengths with the IR emission originating from the inner warm/hot material (T~$\gtrsim$~300~K, R~$\lesssim$~10~AU) and the (sub)mm emission originating from the outer cold disk (T~$<$~300~K, R~$\gtrsim$~10~AU). The molecules detected at (sub)mm wavelengths include CO, \ce{HCO+}, CN, HCN, CS, \ce{N2H+}, SO and \ce{C2H} \citep[see, e.g.,][]{kastner97,dutrey97,vanzadelhoff01,thi04,fuente10,henning10}. Also detected are several isotopologues of the listed species, e.g., $^{13}$CO, C$^{18}$O, H$^{13}$CO$^{+}$, \ce{DCO+} and DCN \citep[see, e.g.,][]{vandishoeck03,thi04,qi08}. Several relatively complex molecules have also been observed: \ce{H2CO} \citep{dutrey97,aikawa03,thi04,oberg10,oberg11}, \ce{HC3N} \citep{chapillon12}, and $c$-\ce{C3H2} \citep{qi13b}. Line emission in the (sub)mm can be observed from the ground and such observations have historically been conducted using single-dish telescopes, e.g., the JCMT (James Clerk Maxwell Telescope), the CSO (Caltech Submillimeter Observatory), the IRAM (Institut de Radioastronomie Millim\'{e}trique) 30~m telescope, and APEX (Atacama Pathfinder Experiment). More recently, several interferometers have been available, e.g., the SMA (Submillimeter Array), CARMA (Combined Array for Research in Millimeter-wave Astronomy), and PdBI (Plateau-de-Bure Interferometer). These latter facilities have enabled spatially-resolved mapping of very nearby objects including the archetypical protoplanetary disk, TW Hydrae, located at a distance of $\approx$~56~pc \citep{oberg10,oberg11,hughes11,qi11}. Due to its proximity, TW Hydrae was observed during ALMA Science Verification which utilised between six and nine antennae working in conjunction to map the line emission from this source \citep{oberg12,rosenfeld12}. Early results from ALMA also include the first detection of the location of the CO snowline\footnotemark[1]~in the disk around HD~163296 using \ce{DCO+} line emission \citep{mathews13}, and in the disk around TW~Hydrae using \ce{N2H+} line emission \citep{qi13c}. \footnotetext[1]{The CO snowline marks the transition zone in the disk midplane ($T$~$\approx$~17~K) beyond which CO is depleted from the gas via freezeout onto dust grains.} The launch of the {\em Herschel} Space Observatory allowed the first detection of ground-state transitions of ortho- and para-\ce{H2O} (at 557~GHz and 1113~GHz, respectively) in the disk of TW Hydrae \citep{hogerheijde11}. \citet{bergin13} also report the first detection of HD in TW~Hydrae using {\em Herschel}, allowing, for the first time, a direct determination of the disk mass without relying on analysis of dust thermal emission or CO rotational line emission. \citet{bergin13} determine a disk mass $<$~0.05~M$_\odot$ confirming that TW~Hydrae, although considered a rather old system ($\sim$~10~Myr), contains sufficient material for the formation of a planetary system. Also detected in the far-IR using {\em Herschel}, is the molecular ion, \ce{CH+}, in the disk of the Herbig~Be star, HD~100546 \citep{thi11}, and multiple lines of OH and warm \ce{H2O} have also been detected in numerous sources \citep[][]{fedele12,meeus12,rivieremarichalar12}. Most detections of line emission in the mid-IR have been conducted with the {\em Spitzer} Space Telescope and molecules observed include OH, \ce{H2O}, \ce{C2H2}, HCN, CO, and \ce{CO2} \citep{lahuis06,carr08,salyk08,pontoppidan10,bast13}. Species detected at IR wavelengths are also limited to abundant, small, simple molecules with strong rovibrational transitions and/or vibrational modes, which are able to survive the high temperatures encountered in the inner disk. \citet{mandell12} also report the detection of near-IR emission lines of \ce{C2H2} and HCN, for the first time, using ground-based observatories (CRIRES on the Very Large Telescope and NIRSPEC on the Keck II Telescope). The greatest chemical complexity (outside of our Solar System) is seen in massive star-forming regions towards the Galactic centre \citep[e.g., Sgr~B2(N),][]{turner91} and in objects called `hot cores' and `hot corinos', considered important stages in high-mass ($M_{\ast}$~$\gtrsim$~10~$M_{\odot}$) and low-mass ($M_{\ast}$~$\lesssim$~10~$M_{\odot}$) star formation, respectively \citep[see, e.g., ][]{herbst09}. Hot cores are remnant, often clumpy, cloud material left over from the explosive process of high-mass star formation which is heated by the embedded massive star. They are warm ($T$~$\sim$~100~K), dense ($n$~$\gtrsim$~10$^{6}$~cm$^{-3}$), relatively large ($R$~$\sim$~0.1~pc) objects which are heavily shielded by dust from both the internal stellar radiation and the external interstellar radiation ($A_\mathrm{v}$~$\sim$~100~mag). Hot corinos, considered the equivalent early stage of low-mass star formation, possess similar densities and temperatures to hot cores, yet, are much less massive and smaller in spatial extent (typically, $R$~$\sim$~100~AU). The line emission from the hot corino arises from a very compact region on the order of 1" in size for a source at the distance of Taurus (140~pc). Hence, we are limited to studying a handful of nearby sources \citep[see, e.g.,][]{ceccarelli05}. Nevertheless, hot corinos are certainly as chemically complex as their more massive counterparts (if not more so), attested by the detection of glycolaldehyde, \ce{HOCH2CHO}, in IRAS 16293+2422 during ALMA Science Verification \citep{jorgensen12}. Hot cores and corinos are typified by the detection of rotational line emission from complex organic molecules (henceforth referred to as COMs), the formation of which remains one of the great puzzles in astrochemistry. The generally accepted mechanism is that simple ices formed on grain surfaces in the molecular cloud at 10~K, either via direct freezeout from the gas phase or via H-addition reactions on the grain (e.g., CO, \ce{H2O}, \ce{H2CO}, \ce{CH3OH}), undergo warming to $\approx$~30~K where they achieve sufficient mobility for grain-surface chemistry to occur via radical-radical association to create more complex ice mantle species (e.g., \ce{HCOOCH3}). The grain-surface radicals necessary for further molecular synthesis are thought to be produced by dissociation via UV photons created by the interaction of cosmic rays with \ce{H2} molecules. Dissociation and/or ionisation via energetic electrons, created along the impact track as a cosmic ray particle penetrates a dust grain, is an alternative scenario \citep[see, e.g.,][]{kaiser97}. Further warming to $T$~$\gtrsim$~100~K allows the removal of these more complex species from the ice mantle via thermal desorption thus `seeding' the gas with gas-phase COMs. Typically, the observed rotational line emission is characterised by a gas temperature of $\gtrsim$~100~K with COMs observed at abundances $\sim$~10$^{-10}$ to $\sim$~10$^{-6}$ times that of the \ce{H2} number density \citep[see, e.g., ][]{herbst09}. Comparing the physical conditions in hot cores/corinos with those expected in the midplane and molecular regions of protoplanetary disks, it appears a similar chemical synthesis route to COMs may be possible; however, to date, targeted searches for gas-phase COMs in nearby protoplanetary disks have been unsuccessful \citep[see, e.g.,][]{thi04,oberg10,oberg11}. The possible reasons for this are severalfold: (i) gas-phase COMs are relatively abundant in disks; however, due to their more complex spectra and resulting weaker emission and the small intrinsic size of disks, existing telescopes are not sufficiently sensitive to detect line emission from COMs on realistic integration time scales, (ii) gas-phase COMs are relatively abundant in disks; however, previous targeted searches have not selected the best candidate lines for detection with existing facilities, and (iii) gas-phase COMs achieve negligible abundances in disks. The latter reason may be related to the major difference between hot cores/corinos and disks: the presence of external UV and X-ray radiation. Certainly, observations using ALMA, with its superior sensitivity and spectral resolution, will elucidate which scenario is correct. The confirmation of the presence (or absence) of COMs in disks is of ultimate astrobiological importance; is it possible for prebiotic molecules to form in the disk and survive assimilation into planets and other objects such as comets and asteroids? Looking at our own Solar System, it appears possible. Many relatively complex molecules have been observed in the comae of multiple comets: \ce{H2CO}, \ce{CH3OH}, HCOOH, \ce{HC3N}, \ce{CH3CN}, \ce{C2H6} \citep[see, e.g.,][and references therein]{mumma11}. The brightest comet in modern times, Hale-Bopp, displayed immense chemical complexity with additional detections of \ce{CH3CHO}, \ce{NH2CHO}, \ce{HCOOCH3}, and ethylene glycol, \ce{(CH2OH)2} \citep{crovisier04a,crovisier04b}. In addition, the simplest amino acid, glycine (\ce{NH2CH2COOH}), was identified in samples of cometary dust from comet 81P/Wild 2 returned by the Stardust mission \citep{elsila09}. The detection of gas-phase glycine is considered one of the `holy grails' of prebiotic chemistry; however, thus far, searches for gas-phase glycine towards hot cores have been unsuccessful \citep[see, e.g.,][]{snyder05}. In \citet{walsh10} and \citet{walsh12}, henceforth referred to as WMN10 and WNMA12, we calculated the chemical composition of a protoplanetary disk using a gas-phase chemical network extracted from the UMIST Database for Astrochemistry (\citeauthor{woodall07} \citeyear{woodall07}\footnotemark[2]), termed `{\sc Rate}06', and the grain-surface chemical network from \citet{hasegawa92} and \citet{hasegawa93}. We included the accretion of gas-phase species onto dust grains and allowed the removal of grain mantle species via both thermal and non-thermal desorption. In WMN10, our aim was to study the effects of cosmic-ray-induced desorption, photodesorption, and X-ray desorption on the chemical structure of the disk, whereas, in WNMA12, we extended our investigations to cover the importance of photochemistry and X-ray ionisation on disk composition. In both works, we focussed our discussions on species detected in disks, the most complex of which, at that time, was formaldehyde, \ce{H2CO}. \footnotetext[2]{\url{http://www.udfa.net}} {\sc Rate}06 includes several gas-phase COMs, including methanol (\ce{CH3OH}), formaldehyde (\ce{H2CO}), formic acid (HCOOH), methyl formate (\ce{HCOOCH3}), dimethyl ether (\ce{CH3OCH3}) and acetone (\ce{CH3COCH3}). These represent the most simple alcohol, aldehyde, carboxylic acid, ester, ether and ketone, respectively. The network also includes several larger members of these families, e.g., ethanol (\ce{C2H5OH}) and acetaldehyde (\ce{CH3CHO}). In WMN10 and WNMA12, we adopted the grain-surface network of \citet{hasegawa92} and \citet{hasegawa93} which includes the grain-surface synthesis of several of these more complex species. However, this network concentrates on simple atom-addition reactions, more likely to occur at the lower temperatures encountered in dark clouds. Hence, to date, the grain-surface chemistry that has been included is by no means comprehensive regarding the grain-surface synthesis of COMs. In this work, we study the efficiency of the synthesis of COMs in protoplanetary disks using a chemical network typically used for hot core and hot corino chemical models. In Sect.~\ref{protoplanetarydiskmodel}, we describe our protoplanetary disk model (Sect.~\ref{physicalmodel}) and chemical network (Sect.~\ref{chemicalmodel}). In Sects.~\ref{results} and \ref{discussion}, we present and discuss our results, respectively, and in Sect.~\ref{summary} we state our conclusions. | \label{discussion} In this section, we discuss and compare our results with observations of molecular line emission from protoplanetary disks and cometary comae, and with results from other models with similar chemical complexity. We also discuss the astrobiological significance of our work. \subsection{Comparison with observations} \label{comparisonwithobservations} Our exploratory calculations suggest that complex organic molecules may be efficiently formed on grain surfaces in protoplanetary disks. However, it is difficult to observe ice species in disks and indeed, this has only been achieved for water ice in a handful of (almost) edge-on systems \citep[see, e.g.,][]{terada07}. Instead, in the cold outer regions of disks, we are limited to observing gas-phase molecules which possess a permanent dipole moment, only. This also presents difficulties if the gas-phase form of the molecule is not present in sufficient quantities and/or also possesses a complex rotational spectrum. The only relatively complex molecules detected in disks, to date, are formaldehyde, \ce{H2CO}, cyanoacetylene, \ce{HC3N}, and cyclopropenylidene, $c$-\ce{C3H2}. \citet{dutrey97} detected several rotational lines of \ce{H2CO} in the disks of DM~Tau and GG~Tau deriving a column density of $\sim$~10$^{12}$~cm$^{-2}$. \citet{aikawa03} and \citet{thi04} present detections of formaldehyde in the disk of LkCa~15 determining column densities of 7.2~--~19~$\times$~10$^{12}$~cm$^{-2}$ and 7.1~--~51~$\times$~10$^{11}$~cm$^{-2}$, respectively. The large spread in column density is due to the difficulty in using a simple model to derive column densities from observations, even when several lines of the species are detected. From Fig.~\ref{figure8} and Table~\ref{table2}, we can see the column density in the outer disk i.e., $>$~10~AU (10$^{12}$~--~10$^{13}$ cm$^{-2}$) compares well with those values constrained from observation. More recently, \citet{oberg10,oberg11} and \citet{qi13a} present detections of \ce{H2CO} using the SMA in a selection of protoplanetary disks in the well-studied Taurus region and in the Southern sky. They confirmed the previous detections of \ce{H2CO} in the disks of DM~Tau and LkCa~15, and they also present new detections: one line in the disk of AA~Tau, two lines in GM~Tau, and two lines in TW~Hya. They also detected formaldehyde in the disks of IM~Lup, V4046~Sgr, and HD~142527. Their detected lines and line intensities towards T~Tauri disks are listed in Table~\ref{table3}. Their values range from $\sim$~100~mJy~km~s$^{-1}$ to $\sim$~1~Jy~km~s$^{-1}$ depending on the source and transition. The authors do not infer any column densities using their data and explain the difficulties in doing so. Instead, they present integrated intensities with which we can compare our calculated line intensities. We also note here the first detection of \ce{HC3N} in a selection of T~Tauri disks by \citet{chapillon12}. The authors state that their observations are most sensitive to a radius of around 300~AU and derive column densities of $\lesssim$~3.5~$\times$~10$^{11}$~cm$^{-2}$, $\approx$~8~$\times$~10$^{11}$~cm$^{-2}$ and $\approx$~13~$\times$~10$^{11}$~cm$^{-2}$ for DM~Tau, LkCa~15 and GO~Tau, respectively. Comparing this with our calculations at 305~AU in Table~\ref{table2}, we determine a column density of 1~$\times$~10$^{11}$~cm$^{-2}$, which is within the upper limit derived for DM~Tau, but around one order of magnitude lower than the column densities for the remaining two sources. We also calculated the rotational line spectra for \ce{HC3N} and found the lines were much weaker than the observed line intensities ($\ll$~10~mJy~km~s$^{-1}$ versus 60~--~100~mJy~km~s$^{-1}$). We note here that LkCa~15 is a particularly peculiar object: the discovery of a large cavity in continuum emission within a radius of $\approx$~50~AU has reclassified this object as a transition disk \citep{pietu06} in which planet formation is likely at an advanced stage \citep{kraus12}. Analysis of CO line observations also identified the lack of a vertical temperature gradient in this disk \citep[see, e.g.,][]{pietu07}. In addition, GO~Tau hosts a particularly large, massive molecular disk \citep[$R_\ce{CO}$~$\sim$~900~AU,][]{schaefer09}. Hence, our disk model is likely not a good analogue for both these sources, providing further explanation for the disagreement between our model results and observations. Recently, \citet{qi13b} reported the detection of cyclopropenylidene, $c$-\ce{C3H2}, in a protoplanetary disk for the first time. The authors identified several lines of this species in ALMA Science Verification observations of the disk of HD~163296, a Herbig~Ae star. This allowed the authors to derive a column density $\sim$~10$^{12}$~--~10$^{13}$~cm$^{-2}$. Herbig Ae/Be stars are more massive and luminous than T~Tauri stars, hence, our disk model is not a suitable analogue for this source. However, it is interesting to consider whether this species may also be detectable in disks around T~Tauri stars. Our model predicts a column density of $\approx$~1~$\times$~10$^{11}$~cm$^{-2}$ at a radius of 100~AU and $\approx$~3~$\times$~10$^{11}$~cm$^{-2}$ at 305~AU, around two orders of magnitude lower than that derived for HD~163296. Gas-phase methanol, \ce{CH3OH}, has not yet been detected in a protoplanetary disk. However, there have been multiple searches in several well-studied objects giving well-constrained upper limits to the line intensities and column densities. \citet{thi04} searched for four lines of methanol (2$_{02}$-1$_{01}$ A, 4$_{22}$-3$_{12}$ E, 5$_{05}$-4$_{04}$ A, 7$_{07}$-6$_{06}$ A) in the disks of LkCa15 and TW Hya using the IRAM 30~m and JCMT single-dish telescopes. In all cases, upper limits only were determined, leading to derived upper column densities between $\approx$~1~$\times$~10$^{13}$ and $\approx$~4~$\times$~10$^{14}$~cm$^{-2}$. Again, our calculated column densities agree with these values in that we predict column densities generally lower than the upper limits derived from the observations. \citet{oberg10,oberg11} also included a line transition of methanol in their SMA line survey of protoplanetary disks. They targeted the 4$_{22}$-3$_{12}$ transition of E-type \ce{CH3OH} at 218.440~GHz in a range of T~Tauri and Herbig~Ae/Be disks and were unable to detect the line in all cases. In Table~\ref{table3}, we compare our modelled line intensities with observations towards sources in which \ce{H2CO} has been detected and in which \ce{H2CO} and \ce{CH3OH} upper limits have been determined. We restrict this list to T~Tauri stars which possess a substantial gaseous disk. We have rescaled our modelled intensities by the disk size and distance to source using the values listed in Table~\ref{table3}. We have listed the sources roughly in order of decreasing spectral type, from K3 (GM Aur) to M1 (DM Tau). We have converted the IRAM 30~m and JCMT line intensities from \citet{thi04} using the standard relation \begin{equation} \left ( \frac{T}{1\,\mathrm{K}} \right ) = \left ( \frac{S_\nu}{1\,\mathrm{Jy}\,\mathrm{beam}^{-1}} \right ) \left [ 13.6 \left ( \frac{300\,\mathrm{GHz}}{\nu} \right )^{2} \left ( \frac{1"}{\theta^2} \right ) \right ] \end{equation} where $T$ is the line intensity in K, $S_\nu$ is the line intensity in Jy~beam$^{-1}$, $\nu$ is the line frequency in GHz and $\theta$ is the beam size in arcseconds. The modelled line intensities for \ce{H2CO} agree reasonably well (within a factor of three) with most transitions towards most sources. For the hotter stars (GM~Tau, LkCa~15, V4046~Sgr, and TW Hya) there is better agreement for the higher frequency transitions than for the lower frequency transitions. For the cooler stars (DM Tau and GG~Tau), there is also reasonable agreement with the lower frequency transitions. The change in line intensity ratios moving from hotter stars to cooler stars reflects the change in disk temperature structure and thus excitation conditions. For the lines in which we see poor agreement, the calculations tend to underestimate the observed line intensities. We would not expect absolute agreement with any particular source because we have adopted `typical' T~Tauri star-disk parameters in our model. However, the level of agreement between our calculations and observations is sufficient for us to conclude that our model is providing a reasonable description of the formation and distribution of \ce{H2CO} in protoplanetary disks around T~Tauri stars and the resulting line emission expected from these objects. Comparing the methanol upper limits and calculated line intensities, we see that our calculations fall well within the upper limits for all sources. Our calculations suggest that the lines of methanol targeted in previous surveys of disks are likely too weak to have been observable. However, our calculations also suggest several potential candidate lines we expect to be strong enough for detection with ALMA (see~Sect.~\ref{linespectra} and Fig.~\ref{figure9}). \begin{table*} \footnotesize \caption{\ce{H2CO} and \ce{CH3OH} rotational transitions in protoplanetary disks.} \centering \begin{tabular}{lcccccccc} \hline\hline Object & Distance & $R(\ce{CO})$ & $i$(CO) & Transition & Frequency & Observed intensities & Modelled intensities & References \\ & (pc) & (AU) & (deg) & & (GHz) & (mJy~km~s$^{-1}$) & (mJy~km~s$^{-1}$) & \\ \hline \multicolumn{9}{c}{\ce{H2CO}}\\ \hline GM Aur & 140 & 630 & 49 & 3$_{03}$-2$_{02}$ & 218.222 & 560 & 130 & 1,2 \\ & & & & 4$_{14}$-3$_{13}$ & 281.527 & 570 & 410 & 2 \\ LkCa 15 & 140 & 905 & 52 & 2$_{12}$-1$_{11}$ & 140.839 & 844 & 110 & 3,4 \\ & & & & & & 820 & & 5 \\ & & & & 3$_{03}$-2$_{02}$ & 218.222 & 696 & 280 & 4 \\ & & & & & & 660 & & 2 \\ & & & & 3$_{22}$-2$_{21}$ & 218.475 & $<$ 498 & 30 & 4 \\ & & & & 3$_{12}$-2$_{11}$ & 225.697 & 498 & 480 & 4 \\ & & & & 4$_{14}$-3$_{13}$ & 281.527 & 1120 & 810 & 2 \\ & & & & 5$_{15}$-4$_{14}$ & 351.768 & 5264 & 1100 & 4 \\ AS 205 & 125 & 250 & 25 & 3$_{03}$-2$_{02}$ & 218.222 & $<$ 160 & 27 & 2,6 \\ & & & & 4$_{14}$-3$_{13}$ & 281.527 & $<$ 300 & 78 & 2 \\ AS 209 & 125 & 340 & 38 & 3$_{03}$-2$_{02}$ & 218.222 & $<$ 210 & 49 & 2,6 \\ & & & & 4$_{14}$-3$_{13}$ & 281.527 & $<$ 480 & 140 & 2 \\ V4046 Sgr & 73 & 370 & 33 & 3$_{03}$-2$_{02}$ & 218.222 & 1001 & 170 & 2,7 \\ & & & & 4$_{14}$-3$_{13}$ & 281.527 & 950 & 500 & 2 \\ TW Hya & 56 & 215 & 6 & 3$_{12}$-2$_{11}$ & 225.697 & $<$ 1026 & 170 & 4,8 \\ & & & & 4$_{14}$-3$_{13}$ & 281.527 & 1220 & 290 & 9 \\ & & & & 5$_{15}$-4$_{14}$ & 351.768 & $<$ 726 & 390 & 4 \\ & & & & & & 540 & & 9 \\ AA Tau & 140 & 995 & 75 & 3$_{03}$-2$_{02}$ & 218.222 & $<$ 520 & 340 & 2,10 \\ & & & & 4$_{14}$-3$_{13}$ & 281.527 & 160 & 980 & 2 \\ DM Tau & 140 & 890 & 35 & 2$_{12}$-1$_{11}$ & 140.839 & 300 & 110 & 3,11 \\ & & & & 2$_{02}$-1$_{01}$ & 145.602 & 110 & 83 & 11 \\ & & & & 3$_{13}$-2$_{12}$ & 211.211 & 480 & 400 & 11 \\ & & & & 3$_{03}$-2$_{02}$ & 218.222 & 350 & 270 & 2 \\ & & & & 4$_{14}$-3$_{13}$ & 281.527 & 290 & 790 & 2 \\ GG Tau & 140 & 800 & 37 & 2$_{12}$-1$_{11}$ & 140.839 & 340 & 85 & 11,12 \\ & & & & 2$_{02}$-1$_{01}$ & 145.602 & 420 & 67 & 11 \\ & & & & 3$_{13}$-2$_{12}$ & 211.211 & 790 & 320 & 11 \\ IM Lup & 190 & 900 & 54 & 3$_{03}$-2$_{02}$ & 218.222 & 530 & 150 & 2,13 \\ & & & & 4$_{14}$-3$_{13}$ & 281.527 & 1370 & 440 & 2 \\ \hline \multicolumn{9}{c}{\ce{CH3OH}}\\ \hline GM Aur & 140 & 630 & 49 & 4$_{22}$-3$_{12}$ E & 218.440 & $<$~90 & 15 & 1,14 \\ LkCa 15 & 140 & 905 & 52 & 2$_{02}$-1$_{01}$ A & 96.741 & $<$~247 & $<$ 1.00 & 3,4 \\ & & & & 4$_{22}$-3$_{12}$ E & 218.440 & $<$~498 & 30 & 4 \\ & & & & & & $<$~150 & & 14 \\ & & & & 5$_{05}$-4$_{04}$ A & 241.791 & $<$~497 & 71 & 4 \\ AS 205 & 125 & 250 & 25 & 4$_{22}$-3$_{12}$ E & 218.440 & $<$~180 & 2.9 & 2,6,14 \\ AS 209 & 125 & 340 & 38 & 4$_{22}$-3$_{12}$ E & 218.440 & $<$~250 & 5.3 & 2,6,14 \\ V4046 Sgr & 73 & 370 & 33 & 4$_{22}$-3$_{12}$ E & 218.440 & $<$~380 & 18 & 7,14 \\ TW Hya & 56 & 215 & 6 & 7$_{07}$-6$_{06}$ A & 338.409 & $<$~362 & 31 & 8,14 \\ AA Tau & 140 & 995 & 75 & 4$_{22}$-3$_{12}$ E & 218.440 & $<$~180 & 36 & 10,14 \\ DM Tau & 140 & 890 & 35 & 3$_{03}$-2$_{02}$ A & 145.103 & $<$~240 & 20 & 3,15 \\ & & & & 4$_{22}$-3$_{12}$ E & 218.440 & $<$~100 & 29 & 14 \\ IM Lup & 190 & 900 & 54 & 4$_{22}$-3$_{12}$ E & 218.440 & $<$~310 & 16 & 13,14 \\ \hline \end{tabular} \label{table3} \tablebib{ (1) \citet{dutrey08}; (2) \citet{oberg10}; (3) \citet{pietu07}; (4) \citet{thi04}; (5) \citet{aikawa03}; (6) \citet{andrews09}; (7) \citet{rodriguez10}; (8) \citet{andrews12}; (9) \citet{qi13a}; (10) \citet{kesslersilacci04}; (11) \citet{dutrey97}; (12) \citet{guilloteau99}; (13) \citet{panic09}; (14) K. \"{O}berg (2012, private communication); (15) \citet{dutrey00}.} \end{table*} \subsection{Comparison with other models} \label{comparisonwithothermodels} Here, we compare our results with other protoplanetary disk models, concentrating on work which has published lists of column densities or fractional abundances for relatively complex species. Historically, chemical models of disks have concentrated on simple, abundant species (and isotopologues), since these are readily observed in many systems (e.g., CO, \ce{HCO+}, CN, CS, and HCN). As we enter the era of ALMA, it is likely that the molecular inventory of protoplanetary disks will significantly increase, requiring much more sophisticated complex chemical models, such as that presented here. In Table~\ref{table4}, we compare column densities of various complex molecules at a radius of $\approx$~250~AU with other protoplanetary disk models of comparable chemical complexity and with similar chemical ingredients. \citet[][W07]{willacy07} presented a chemically complex model of a protoplanetary disk, including a comprehensive deuterium chemistry. We compare our column densities with Model~C in that work, which includes both grain-surface chemistry and non-thermal desorption. \citet[][SW11]{semenov11} present results from a disk model which uses a network with a number of chemical reactions ($\approx$~7300) approaching the number in the network presented here ($\approx$~9300). We compare our results with their `laminar' model in which they neglect turbulent mixing, since we do not consider mixing in this work. We also list the column densities from our most recent paper, WNMA12, which is most similar to the work presented here in that the disk physical model is identical as are the methods used to calculate the chemistry. In WNMA12 we used a chemical network based on `{\sc Rate}06', the most recent release of the UMIST Database for Astrochemistry (UDfA) available at that time, whereas, here, our network is derived from the Ohio State University (OSU) network and includes a vastly more comprehensive grain-surface chemical network to simulate the build up of complex molecules. The network used in W07 is also derived from UDfA, albeit an earlier version \citep[{\sc Rate}95,][]{millar97}, whereas, the network used by SW11 is also based on the OSU network. Care must be taken when comparing results from different protoplanetary disk models, as they often differ in physical ingredients as well as the chemistry. The work presented here generally predicts higher column densities for COMs than those presented in W07 and SW11 despite relatively similar (within an order of magnitude) column densities for CO, \ce{H2CO}, and \ce{HC3N}. In this work, we calculate significantly higher column densities for \ce{CH3OH}, \ce{HCOOH}, \ce{CH3CN}, \ce{CH3CHO}, \ce{NH2CHO}, \ce{HCOOCH3}, \ce{C2H5OH}, \ce{CH3OCH3}, and \ce{CH3COCH3}. The network used by SW11 is based on that presented in \citet{garrod06} which does not contain many pathways to the larger species introduced in \citet{garrod08}. Also, they adopt $E_{d}$~=~0.77~$E_{D}$ for their grain-surface diffusion rates \citep[which originates from][]{ruffle00}, where $E_{D}$ is the binding energy of the molecules to the grain surface. This is a rather conservative value and partly explains their much lower abundances of complex species. In addition, they do not consider quantum tunnelling of H atoms on grain surfaces, nor through reaction energy barriers \citep[a full description of the chemical model is provided in][]{semenov10}. The neglect of H atom tunnelling through reaction energy barriers explains the particularly low column density of methanol in SW11 ($\sim$~10$^{8}$~cm$^{-2}$). W07 include atom-addition grain-surface reactions only and thus neglect radical-radical pathways to form larger COMs. Comparing our results with those from our previous work (WNMA12), we see a significant increase in the column density of \ce{CH3OH}, \ce{HCOOCH3}, \ce{CH3OCH3}, and \ce{CH3COCH3} when using the gas-grain network presented here. The higher column density of grain-surface methanol, {\em s-}\ce{CH3OH}, can be attributed to the higher binding energy of CO adopted here. In previous work, we used the value measured for pure CO ice (855~K) as opposed to the value measured in water ice (1150~K). The binding energy regulates the abundance of {\em s-}CO on the grain and thus the amount available for conversion to {\em s-}\ce{CH3OH}, as well as the grain-surface radicals, {\em s-}\ce{HCO}, {\em s-}\ce{CH3O}, and {\em s-}\ce{CH2OH}. Regarding the formation of {\em s-}\ce{HCOOCH3}, the grain-surface association reaction, {\em s-}\ce{HCO}~+~{\em s-}\ce{CH3O}, is included in both models. The difference in column density is due, again, to the different sets of binding energies adopted. The results from our exploratory calculations presented in Sect.~\ref{verticalresults} demonstrate the importance of radiation processing for the production of {\em s-}\ce{CH3OCH3} and {\em s-}\ce{CH3COCH3} in the disk midplane. The midplane is the densest region of the disk and thus contributes significantly to the vertical column density. In previous work we did not include the processing of ice mantle material by UV photons and X-rays. We also see a decrease in the column density of gas phase \ce{HC3N}, and a corresponding increase in the grain-surface column density, compared with our previous values. This is due to the increased binding energy for \ce{HC3N} adopted here (4580~K compared with 2970~K). Our previous value shows better agreement with the column densities constrained from observations (~$\sim$~10$^{12}$~cm$^{-2}$). Protoplanetary disks are turbulent environments and the effects of vertical turbulent mixing on disk chemical structure has been investigated by multiple groups \citep[see, e.g.,][]{ilgner04,willacy06,semenov06,aikawa07,heinzeller11,semenov11}. \citet{semenov11} conducted a comprehensive investigation of disk chemical structure with and without turbulent mixing and identified a plethora of species which are sensitive to mixing. \citet{semenov11} also used a chemical network including several complex molecules (see~Table~\ref{table4}). Of the gas-phase molecules of interest here, they found that the column densities of \ce{HCOOH}, \ce{HC3N}, and \ce{CH3CN}, were significantly affected by the inclusion of turbulent mixing. However, they concentrated their discussions on species which reached column densities $\gtrsim$~10$^{11}$~cm$^{-2}$. In this work, we assume the dust grains are well mixed with the gas and, for the calculation of the chemical structure, we assume the grains are compact spherical grains with a fixed radius. In reality, the dust grains will have both a size distribution and variable dust-to-gas mass ratio caused by gravitational settling towards the midplane and dust-grain coagulation (grain growth). Several groups have also looked at the effects of dust-grain evolution on protoplanetary disk chemistry \citep[see, e.g.,][]{aikawa06,fogel11,vasyunin11,akimkin13}. A parameterised treatment of grain growth affects the geometrical height of the molecular layer but appears to have little effect on the column densities of gas-phase molecules \citep{aikawa06}. Larger grains may lead to a reduced volume grain-surface area (for a fixed dust-to-gas mass ratio) which will lower the accretion rate of molecules onto dust grains, thereby potentially lowering the formation rate of COMs. However, this effect depends on the assumed morphology and porosity of the grains. Grain coagulation models suggest that the growing dust grains retain a `fluffy' morphology (with a low filling factor, $\ll$~1) such that the volume grain-surface area may not significantly decrease before compression occurs \citep[see, e.g.,][]{ossenkopf93,ormel07,kataoka13}. Grain settling towards the midplane allows the deeper penetration of UV radiation leading to warmer grains in the disk midplane. This subsequently results in a smaller depletion (freezeout) zone and a larger molecular layer situated closer to the midplane \citep{fogel11}. \citet{akimkin13} performed a coupled calculation of the structure of a protoplanetary disk including dust evolution and radiative transfer, and subsequently calculated the chemical evolution. They find that the abundances of several species are enhanced in models including dust evolution, including the relatively complex species, \ce{NH2CN} and \ce{HCOOH}. We intend to explore the effects of grain evolution on the formation and distribution of COMs in future models. A final issue to consider is the validity of our set of initial abundances. Disk formation is thought to be a vigorous, energetic, and potentially destructive process. Accretion shocks are thought to occur as material falls from the envelope onto the disk. Heating by the shock may raise the temperature of the dust grains above the sublimation temperature of ices and energised ions may sputter ices from dust grain surfaces \citep[see, e.g.,][]{neufeld94,tielens94}. Hence, using initial abundances representative of dark cloud (or prestellar) conditions may not be realistic because dust grains may be stripped of ices as they pass through an accretion shock during the disk formation stage. \citet{visser09} studied the 2D chemical evolution during the protostellar collapse phase to determine the chemical history of simple ices contained within the disk at the end of collapse. They concluded that accretion shocks that occur as material falls from the envelope onto the disk are much weaker than commonly assumed. For the outer disk, the main contribution to heating is via stellar heating \citep[see Fig.~3 in][]{visser09}. Sputtering of dust grains by energetic ions can also occur. \citet{visser09} also considered this and concluded that the shock velocities experienced by the gas, $\approx$~8~km~s$^{-1}$, are not sufficient to energise ions, such as \ce{He+}, to energies required for the removal of water molecules from grain surfaces. As a result, much of the material contained within the outer disk ($\gtrsim$~10~AU) at the end of collapse consists primarily of ``pristine'' interstellar ice \citep[see Fig. 4 in][]{visser11}. Hence, beginning our simulations with initial molecular abundances representative of prestellar conditions is an appropriate assumption. \begin{table*} \footnotesize \caption{Calculated column densities (cm$^{-2}$) at $R$~=~250~AU.} \centering \begin{tabular}{llcccccc} \hline\hline & & \multicolumn{4}{c}{Gas phase} & \multicolumn{2}{c}{Grain surface} \\ \multicolumn{2}{c}{Species} & W07$^1$ & SW11$^2$ & WNMA12$^3$ & This work & WNMA12$^3$ & This work \\ \hline\\ Carbon monoxide & \ce{CO} & 2.9(17) & 1.2(18) & 1.3(18) & 3.1(17) & 2.1(14) & 2.4(17) \\ Formaldehyde & \ce{H2CO} & 3.9(12) & 2.4(13) & 5.9(12) & 7.2(12) & 1.7(14) & 2.9(17) \\ Methanol & \ce{CH3OH} & 1.5(11) & 6.1(08) & 1.9(12) & 1.7(13) & 8.4(16) & 7.7(17) \\ Formic acid & \ce{HCOOH} & $\cdots$ & 1.4(11) & 9.8(13) & 1.6(13) & 2.1(16) & 6.1(16) \\ Cyanoacetylene & \ce{HC3N} & 7.0(11) & 2.9(11) & 2.5(12) & 1.7(11) & 2.9(08) & 4.0(12) \\ Acetonitrile & \ce{CH3CN} & 4.9(10) & 6.0(10) & 5.6(12) & 7.3(11) & 5.8(15) & 2.5(15) \\ Acetaldehyde & \ce{CH3CHO} & $\cdots$ & 8.0(07) & 6.2(10) & 3.4(11) & 2.2(15) & 3.7(16) \\ Formamide & \ce{NH2CHO} & $\cdots$ & 2.0(10) & $\cdots$ & 9.2(11) & $\cdots$ & 1.4(16) \\ Methyl formate & \ce{HCOOCH3} & $\cdots$ & 8.8(04) & 3.5(08) & 2.7(11) & 1.8(11) & 1.1(16) \\ Ethanol & \ce{C2H5OH} & $\cdots$ & 4.4(06) & 1.4(12) & 6.1(10) & 1.8(16) & 6.9(15) \\ Dimethyl ether & \ce{CH3OCH3} & $\cdots$ & 4.2(02) & 7.0(09) & 7.6(10) & 1.4(14) & 1.2(16) \\ Acetone & \ce{CH3COCH3} & $\cdots$ & 1.3(03) & 2.8(06) & 4.2(09) & 2.6(09) & 6.2(15) \\ \hline \end{tabular} \label{table4} \tablefoot{$a(b)$ represents $a\times10^b$} \tablebib{ (1) \citet[][W07]{willacy07}; (2) \citet[][SW11]{semenov11}; (3) \citet[][WNMA12]{walsh12}} \end{table*} \subsection{Complex molecules in comets} \label{comets} It is generally accepted that minor bodies in the Solar System, such as asteroids and comets, likely formed in conjunction with the planets in the Sun's primordial disk and can be considered remnant material left over from the process of planet formation. When a comet's orbit is perturbed in such a way that it is injected into the inner Solar System, the gradual warming of the nearing Sun evaporates solid surface material and creates an expansive cometary coma of gaseous volatile material enveloping the comet nucleus. Photolysis of the sublimated material (termed `parent' species) and subsequent chemistry creates ions and radicals and new molecules (termed `daughter' species). It is now understood that comets are complex objects composed of ice (mainly \ce{H2O}, \ce{CO2}, and CO), refractory material (such as silicates), and organic matter. To date, more than 20 parent molecules have been observed in cometary comae including the relatively complex species, \ce{H2CO}, \ce{CH3OH}, HCOOH, \ce{CH3CHO}, \ce{HC3N}, \ce{CH3CN}, \ce{NH2CHO}, \ce{HCOOCH3}, and \ce{(HOCH2)2} (ethylene glycol), which are relevant to this work. Of these species, \ce{CH3CHO}, \ce{NH2CHO}, \ce{HCOOCH3}, and \ce{(HOCH2)2} have been observed in only a single object, comet Hale-Bopp, with percentage abundances of 0.02~\%, 0.015~\%, 0.08~\%, and 0.25~\% (relative to \ce{H2O}), respectively \citep[see, e.g.,][]{bockelee04,crovisier04a,crovisier04b,crovisier06,mumma11}. In Fig.~\ref{figure10} we present the range of calculated abundances for each of these grain-surface species relative to water ice (red lines) and compare these with our initial adopted dark cloud ice ratios (green asterisks) and data derived from cometary comae observations (blue asterisks and lines). The fractional abundances from the disk model are determined by the relative vertical column densities at each radius. We restrict our data to $R$~$\gtrsim$~20~AU which corresponds to the radius beyond which grain-surface COMs achieve significant column densities (see Fig.~\ref{figure8}). This also correlates with the region where comets are postulated to have originally formed and resided in modern dynamical models of the Solar System \citep[see, e.g.,][]{gomes05,walsh11}. The single points and upper limits for the comet observations refer to data derived from observations of comet Hale-Bopp. We find that our range of calculated abundances (relative to water ice) are consistent with those derived from observations, with some overlap between the two datasets for most species. Exceptions include {\em s-}\ce{CH3CHO} and {\em s-}\ce{NH2CHO} for which a single observation only is available. In both cases, our data range is larger than the observed ratio, with the lower limit of our data within a factor of a few of the measured ratio. Another exception is {\em s-}\ce{CH3CCH}, for which an upper limit towards Hale-Bopp only has been derived \citep[][]{crovisier04b}. Again, we find our calculated ratio range is larger than the upper limit. In this case, the lower limit of our data is much further away from that derived from observation, by a factor of $\approx$~30. It is also interesting to compare our range of calculated abundances in the disk model with our initial abundances adopted from a dark cloud model (see Table~\ref{table1}). The {\em s-}\ce{H2CO}/{\em s-}\ce{CH3OH} ratio indicates there is significant chemical processing of the dark cloud grain-surface material within the disk with this ratio decreasing from cloud to disk. For all other species (except {\em s-}\ce{CH3CCH}) the dark cloud abundance is lower than the lower limit reached in the disk model indicating that disk physical conditions are necessary for thermal grain-surface chemistry to efficiently form the complex molecules observed in cometary comae. It certainly appears that our grain-surface chemistry is appropriate for describing the grain-surface formation of most COMs observed in cometary comae, supporting the postulation that comets are formed via the coagulation and growth of icy dust grains within the Sun's protoplanetary disk. One outstanding issue is the high abundance of ethylene glycol (\ce{(HOCH2)2}) observed towards comet Hale-Bopp, with a percentage abundance of 0.25\% relative to water ice. This ratio is similar to that observed for \ce{H2CO} and around an order of magnitude higher than the ratio derived for \ce{CH3CHO} and \ce{NH2CHO}. Also, \ce{(HOCH2)2} is observed to be at least 5 times more abundant than the chemically-related species, \ce{HOCH2CHO} \citep{crovisier04b}. In this network, we include a single barrierless route to the formation of {\em s-}\ce{(HOCH2)2} via the grain-surface association of two {\em s-}\ce{CH2OH} radicals. Under the conditions throughout much of the disk, the mobility of this radical is significantly slower than smaller radicals, such as, {\em s-}\ce{CH3} and {\em s-}\ce{HCO}, due to its significantly larger binding energy to the grain mantle ($\approx$~5000~K). The large binding energy is due to the presence of the -OH functional group allowing hydrogen bonding of this species with the bulk water ice \citep[see, e.g.,][]{garrod08}. Hence, the reaction forming {\em s-}\ce{(HOCH2)2} cannot compete with other barrierless radical-radical association reactions which form, for example, {\em s-}\ce{C2H5OH} and {\em s-}\ce{HOCH2CHO}. We find a negligible abundance of {\em s-}\ce{(HOCH2)2} is produced throughout our disk model. In the network used here, radical-radical association pathways only have been included for the formation of many COMs, in addition to pathways involving sequential hydrogenation of precursor molecules. However, an alternative grain-surface route to the production of \ce{(HOCH2)2} (and other COMs) has been proposed by \citet{charnley97} involving sequential atom-addition reactions. For example, \ce{(HOCH2)2} is postulated to form via the hydrogenation of {\em s-}\ce{OCCHO}, which in turn is formed from {\em s-}\ce{CO} via atom-addition reactions, i.e., \begin{equation*} s\mbox{-}\ce{CO} \xrightarrow{s\mbox{-}\ce{H}} s\mbox{-}\ce{HCO} \xrightarrow{s\mbox{-}\ce{C}} s\mbox{-}\ce{HC2O} \xrightarrow{s\mbox{-}\ce{O}} s\mbox{-}\ce{OCCHO} \end{equation*} followed by \begin{equation*} s\mbox{-}\ce{OCCHO} \xrightarrow{s\mbox{-}\ce{H}} s\mbox{-}\ce{CHOCHO} \xrightarrow{2s\mbox{-}\ce{H}} s\mbox{-}\ce{HOCH2CHO} \xrightarrow{2s\mbox{-}\ce{H}} s\mbox{-}\ce{(HOCH2)2}. \end{equation*} In this sequence, 2{\em s-}\ce{H} implies a barrier penetration reaction by a hydrogen atom followed by the exothermic addition of an additional H atom. This sequence of atom-addition reactions is postulated to lead to different ratios of resultant grain-surface COMs relative to the radical-radical network used here and may provide a route to the formation of {\em s-}\ce{(HOCH2)2}. However, as discussed in \cite{herbst09}, many of these reaction rates and reaction barriers remain unmeasured. The efficacy of this type of formation route to COMs under protoplanetary disk conditions is yet to be studied and we intend to explore this in future work. Of course, it is also possible that significant processing of the cometary surface by UV photons (and potentially, cosmic rays) over the comet's lifetime may lead to a surface composition which differs from the initial grain mantle composition in the protoplanetary disk. In addition, thermally driven chemical processing of the comet's interior may occur. This may be caused by heating due to radioactive decay of radionuclides, such as $^{26}$Al \citep[see, e.g.,][]{wallis80,prialnik87}. \begin{figure*} \subfigure{\includegraphics[width=1.0\textwidth]{./comet_ratios.eps}} \caption{Range of abundances of grain-surface complex molecules relative to water ice from our model calculations (red lines) compared with those derived from observations of cometary comae (blue asterisks and lines) and our initial dark cloud ice ratios (green asterisks). The comet data is from \citet{bockelee04} and \citet{crovisier06}. The single points and upper limits for the comet ratios represent data derived from observations of comet Hale-Bopp \citep{crovisier04b}.} \label{figure10} \end{figure*} \subsection{Implications for astrobiology} \label{astrobiology} One of the most complex molecules detected to date is aminoacetonitrile, \ce{NH2CH2CN}, which was observed towards the hot core in the massive star-forming region, Sgr B2(N), with a fractional abundance $\sim$~10$^{-9}$ \citep{belloche08}. \ce{NH2CH2CN} has been postulated as a potential precursor to the simplest amino acid, glycine, \ce{NH2CH2COOH}. In turn, amino acids are the building blocks of proteins, considered a key component for the commencement of life. Multiple routes to the formation of glycine (and other simple amino acids) under interstellar conditions have been proposed including via Strecker synthesis \citep[see, e.g.,][]{peltzer84}, UV-irradiated ice mantles \citep[see, e.g.,][]{bernstein02,munoz-caro02}, and gas-phase chemistry \citep[see, e.g,][]{blagojevic03}. Recently, \citet{garrod13} investigated the formation of glycine in hot cores via grain-surface radical-radical reactions, i.e., an extension to the reaction scheme used here, incorporating the ice chemistry proposed in \citet{woon02} to describe the formation of glycine in UV-irradiated ices. \citet{garrod13} calculated a peak fractional abundance for gas-phase glycine $\sim$~10$^{-10}$~--~10$^{-8}$ with the molecule returned to the gas phase at temperatures $\gtrsim$~200~K. He also included gas-phase formation of glycine \citep{blagojevic03} and determined it to have a negligible effect on the resulting abundances. The detection of glycine is considered one of the holy grails of astrochemistry and astrobiology; however, searches for gas-phase glycine, thus far, have been unsuccessful \citep[see, e.g.,][]{snyder05}. The predictions from \citet{garrod13} are consistent with upper limits derived from these observations. He proposes that due to the high binding energy of glycine, the emission from hot cores is expected to be very compact, and thus, an ideal target for detection with ALMA. Certainly, a similar grain-surface formation route to {\em s-}\ce{NH2CH2CN} and thus, {\em s-}\ce{NH2CH2COOH}, may be possible under protoplanetary disk conditions and should be explored in future models, particularly considering the recent identification of glycine in a sample returned from comet 81P/Wild 2 \citep{elsila09} and the detection of numerous amino acids in meteorites, some of which are either very rare on Earth or, indeed, unknown in terrestrial biochemistry \citep[for an overview, see, e.g.,][]{ehrenfreund00}. Models would help ascertain whether it is possible for simple amino acids to form on dust grains in the Sun's protoplanetary disk and become incorporated into comets and asteroids. Such models could also provide further evidence for the delivery of prebiotic molecules to Earth via asteroid and/or cometary impact, rather than forming `in situ' early in the Earth's evolution. However, based on our molecular line emission calculations (see Sect.~\ref{linespectra}), even if such prebiotic molecules were present in quantities similar to that expected in hot cores, the detection of the gas-phase form of these species in protoplanetary disks would be incredibly challenging, if not impossible, even with ALMA full capabilities. | 14 | 3 | 1403.0390 |
1403 | 1403.7786_arXiv.txt | {Combining Planck CMB temperature~\cite{Ade:2013kta} and BICEP2 B-mode polarization data~\cite{Ade:2014gua,Ade:2014xna} we show qualitatively that, assuming inflationary consistency relation, the power-law form of the scalar primordial spectrum is ruled out at more than $3\sigma$ CL. This is an important finding, since the power-law form of the scalar primordial spectrum is one of the main assumptions of concordance model of cosmology and also a direct prediction of many inflationary scenarios. We show that a break or step in the form of the primordial scalar perturbation spectrum, similar to what we studied recently analyzing Planck data~\cite{Hazra:2013nca}, can address both Planck and BICEP2 results simultaneously. Our findings also indicate that the data may require more flexibilities than what running of scalar spectral index can provide. Finally we show that an inflaton potential, originally appeared in~\cite{Bousso:2013uia}, can generate both the step and the break model of scalar primordial spectrum in two different limits. The discussed potential is found to be favored by Planck data but marginally disfavored by BICEP2 results as it produces slightly lower amplitude of tensor primordial spectrum. Hence, if the tensor-to-scalar ratio ($r$) quoted by BICEP2 persists, it is of importance that we generate inflationary models with large $r$ and at the same time provide suppression in scalar primordial spectrum at large scales. } | The primary goal of physical cosmology is to find an accurate model of the Universe. The current standard model of cosmology, also known as concordance model, is a spatially flat FLRW Universe consist of weakly interacting cold dark matter, cosmological constant and baryons all in the context of power-law form of the primordial perturbations. This simple power law form is a natural result if, for example, one assumes slow-roll Inflation. While the features in the primordial power spectrum (PPS) have been subject of various studies, cosmological observations before BICEP2 B-mode polarization results~\cite{Ade:2014gua,Ade:2014xna} have been all essentially consistent to the power-law form of the primordial perturbation spectrum. The very high tensor-to-scalar ratio from the BICEP2 B -Mode polarization at large angular scales may change this picture. In fact it seems to be hard to assume a power-law form of the primordial spectrum and have a good fit to both Planck temperature~\cite{Ade:2013kta} (and WMAP~\cite{Hinshaw:2012fq} low-$\ell$ polarization) and BICEP2 B-Mode observations simultaneously. Assuming that BICEP2 has detected B-modes corresponding to a tensor-to-scalar ratio $r_{0.002}\sim0.2$ (defined at pivot scale $0.002{\rm Mpc^{-1}}$) indicates that the high tensor component will add power to large angular scale temperature anisotropy. On the other hand Planck data indicates a mild suppression in large angular scale power compared to standard power law $\Lambda$CDM model~\cite{Ade:2013kta}. In fact there were hints of large scale suppression in scalar power since the first year results of WMAP~\cite{Peiris:2003ff}. Several model independent reconstruction methods~\cite{Hazra:2013nca,reconstruction-all} using different datasets too hint towards suppression in scalar power. This mild suppression limits the tensor primordial spectra to be higher than $r_{0.002}< 0.12$ (95\% C.L.)~\cite{Ade:2013uln} assuming a primordial perturbation power law. In this paper we address the consistency of the power-law PPS with combination of Planck and BICEP2 data and show qualitatively that the power-law PPS is in fact ruled out at more than $3\sigma$. This is an important result, since it may require additional degrees of freedom to the current standard model of cosmology (assuming the observational results persist). At the same time most (slow-roll) inflationary scenarios also result to power-law form of the PPS and our findings show that all these models are now in fact in tension with observational data. In our analysis we study some well motivated phenomenological forms of the primordial spectrum such as a broken PPS model and a Tanh step model (we studied both models recently in ~\cite{Hazra:2013nca}) as well as some potential theoretical models such as~\cite{Bousso:2013uia} to see if they can perform well fitting both Planck and BICEP2 data. Assuming these models and by deviating from some basic assumptions such as inflationary consistency relation we study how well we can address the combined data and what are the affects on the constraints of the cosmological parameters. We show that step model and model with a break in scalar PPS are suitable candidates to fit both Planck and BICEP2 data and this motivates some particular inflationary model building. We perform full Monte Carlo sampling in our analysis to estimate the consistency of the the power-law form of PPS with Planck and BICEP2 data and for some other cases we limit ourselves to some particular model samples when we study inflationary scenarios. In a companion paper we focus on inflationary scenarios that may be able to fit all different data satisfactorily. This paper is organized as follows. We discuss first the assumed phenomenological models and then we test the consistency of the power-law scalar PPS with Planck and combination of Planck and BICEP2 data. We then briefly discuss about potential inflationary scenarios that can fit both Planck and BICEP2 data. We end the paper by results and conclusions. | \label{sec:conclusions} In this paper we show that power-law form of the primordial spectrum (scalar PPS) cannot fit properly Planck temperature and BICEP2 B-Mode polarization data simultaneously. In fact scalar power-law form of PPS is disfavored at more than $3\sigma$ in comparison with the kink model we have studied. There have been hints in the Planck temperature data alone (along with WMAP low $\ell$ E-mode polarization data) that a broken form of the PPS can fit the data pretty well as we have discussed it in details in~\cite{Hazra:2013nca}. However, due to cosmic variance, it was not possible to favor this phenomenological model to power-law form of the primordial spectrum with high confidence using temperature data alone. Using BICEP2 B-mode polarization data this degeneracy seems to be broken now and we have estimated that the power-law form of PPS is ruled out at more than $3\sigma$ CL, assuming inflationary consistency relation. It is indeed interesting to see that a simple broken form of the PPS can indeed fit both data pretty well unlike power-law form of the primordial spectrum. This is evident by looking at the overlap of the confidence contours in Fig.~\ref{fig:icon1d2d} fitting Planck data alone and fitting combination of Planck and BICEP2 data for two cases of power-law and broken form of the PPS. Ruling out the power-law form of the scalar primordial spectrum is an important result since it is one of the main assumptions of the concordance model of cosmology and at the same time there are many inflationary scenarios that result to power-law form of the primordial spectrum. If confirmed, the detection of CMB B-mode polarization is a major discovery and opens windows. An example is that these new results seems to be quite severe for the existing standard model of cosmology for investigation of fine structure of inflation. We have also shown that a Tanh step form of the PPS, again what we proposed and discussed in~\cite{Hazra:2013nca}, can also fit properly both Planck and BICEP2 data simultaneously. These results reflects the fact that era of the Vanilla concordance model of cosmology is near to end (if the observational data persist) and we need more flavour to explain our observable Universe. We may need about 3 additional parameters to express the initial perturbation such as the value of $r$ (tensor-to-scalar ratio), an additional spectral index for the low-$k$ wavenumbers and a $k_{\rm b}$, the wavenumber where the break/transition can occur. We should note that there is an advantage for the broken or step forms of the scalar PPS to assumption of running of the scalar spectral index since by assuming running we limit ourself to a particular shape and thereby allow less flexibility. It is interesting to note that by assuming running there is one less degree of freedom in comparison to the broken PPS, but the power-law form of the PPS is ruled out with higher confidence by assuming the broken PPS. We should note here that the running of the spectral index, $d n_{\rm S}/d\ln k$ needs to have a large negative value $\sim -0.02$~\cite{Ade:2014xna} in order to fit both Planck and BICEP2. This large negative running introduces large suppression in scalar PPS at small scales which needs to be balanced by including another parameter, for example, running of running or neutrinos. By just considering the running of the scalar spectral index one can get about 6.5 improvement in the $\chi^2$ in comparison to power law, while by considering a broken power spectrum we can get an improvement of about 12-13. It is in fact double improvement in the $\Delta \chi^2$ by having only one more extra parameter. Our result signifies the importance of the one additional degree of freedom in the broken scalar PPS (compared to running), {\it i.e.} the position of the break $k_{\rm b}$. This shows that the assumption of the running of the spectral index may not suffice to explain the data properly. Our results also indicate that relaxing the inflationary consistency relation can help the power-law form of the PPS to be less inconsistent to the data (but still considerably inconsistent) but it would not improve the fit much for the broken or step forms of the PPS. This is a good news since it seems by assuming these simple non-power-law forms of the PPS, there will not be any tension between various CMB data and we can still hold on the theoretically important inflationary consistency relation. From the theoretical perspective, once we fix an inflationary model potential, we fix the amplitude and tilt of the scalar and tensor power spectra simultaneously. There is no more freedom to change $r$ within a model. Our preliminary results presented in this work show that within an inflationary scenario, previously discussed in~\cite{Bousso:2013uia} we can get the scalar PPS that can resemble the broken and the step like scalar PPS discussed in this paper, but it can not generate large tensor component that can address BICEP2 data well. If observational constraints for $r$ changes to values close to $~0.1$ rather than the current central value of $0.2$, there seems to be more space for inflationary scenarios to explain all data simultaneously. Thus, the BICEP2 discovery of primordial gravitational waves, while confirming the general observational prediction~\cite{S79} of the of the early Universe scenario with the de Sitter (inflationary) stage preceding the hot radiation dominated stage, shows that the inflationary stage is not so simple and may not be described by a one-parametric model. We focus on inflationary model (Whipped Inflation) building in a separate paper~\cite{Hazra:2014jka} wherein we discuss that using canonical scalar fields, generation of large tensors, suppression in scalar power at large scales and at the same time low level of non-Gaussianity is achievable. We should also mention that expansion beyond linear order term in the slow roll part of the potential in~\cite{Bousso:2013uia} can also help in generating a large tensor amplitude and hence can address BICEP2 data better than the potential we have used in this paper~\cite{Bousso:2014jca}. We also wish to perform a complete parameter estimation for the inflationary models we discussed, where we expect the B-mode polarization data from POLARBEAR~\cite{Ade:2014afa} will help us in providing tighter constraints. Just before finishing this paper it is important to address about an alternative approach which has been proposed to reconcile Planck and BICEP2 data within the frame of the power-law form of the primordial spectrum by assuming extra massive sterile neutrino. Soon after release of the Planck data it was realized that by assuming an additional massive sterile neutrino, the scale invariant form of the primordial spectrum ($n_s=1$) can be consistent to the data while the Hubble parameter should have higher values (which could make it even more consistent to the local Universe estimations of $H_0$)~\cite{Ade:2013zuv}. After release of BICEP2 data one could guess that this model may work well to fit both Planck and BICEP2 as well. In~\cite{Zhang:2014dxk,Dvorkin:2014lea} it was shown that such model can indeed fit the combination of Planck and BICEP2 data reasonably well while staying within the context of the power-law form of the primordial spectrum. We should note that this model has limited flexibilities to suppress low $\ell$ scalar multipoles and also to describe the fine structure of the temperature spectrum. We would have soon E-Mode polarisation data from Planck (where the effect of broken form of the primordial spectrum and having massive sterile neutrino would be different on the data) that will help us to differentiate between these two main alternatives with high confidence. We should mention here that there have been some publications~\cite{Contaldi:2014zua,Miranda:2014wga,Abazajian:2014tqa,Ashoorioon:2014nta} in last few days after release of the BICEP2 data that we may share some of the results, however we should emphasize here that this paper is in fact a straightforward extension of the ~\cite{Hazra:2013nca} considering BICEP2 data. | 14 | 3 | 1403.7786 |
1403 | 1403.7053_arXiv.txt | We use $N$-body simulations to compare the evolution of spatial distributions of stars and brown dwarfs in young star-forming regions. We use three different diagnostics; the ratio of stars to brown dwarfs as a function of distance from the region's centre, $\mathcal{R}_{\rm SSR}$, the local surface density of stars compared to brown dwarfs, $\Sigma_{\rm LDR}$, and we compare the global spatial distributions using the $\Lambda_{\rm MSR}$ method. From a suite of twenty initially statistically identical simulations, 6/20 attain $\mathcal{R}_{\rm SSR} << 1$ \emph{and} $\Sigma_{\rm LDR} << 1$ \emph{and} $\Lambda_{\rm MSR} << 1$, indicating that dynamical interactions could be responsible for observed differences in the spatial distributions of stars and brown dwarfs in star-forming regions. However, many simulations also display apparently contradictory results -- for example, in some cases the brown dwarfs have much lower local densities than stars ($\Sigma_{\rm LDR} << 1$), but their global spatial distributions are indistinguishable ($\Lambda_{\rm MSR} = 1$) and the relative proportion of stars and brown dwarfs remains constant across the region ($\mathcal{R}_{\rm SSR} = 1$). Our results suggest that extreme caution should be exercised when interpreting any observed difference in the spatial distribution of stars and brown dwarfs, and that a much larger observational sample of regions/clusters (with complete mass functions) is necessary to investigate whether or not brown dwarfs form through similar mechanisms to stars. | One of the outstanding questions in star formation is whether the mechanism through which brown dwarfs (objects not massive enough to burn hydrogen in their cores) form is more like that of higher (e.g. Solar) mass stars, or more like that of giant planets. This can be addressed by comparing the various properties of brown dwarfs (BDs) with stars, such as multiplicity \citep{Duchene13b}, kinematics \citep{Luhman07} and spatial distribution \citep{Kumar07}. Several studies \citep[e.g.][]{Luhman06,Bayo11,Parker11b,Parker12c} have shown that BDs have a similar spatial distribution to stars in some star-forming regions; but there are other regions where the BDs appear to be more spread out \citep{Kumar07,Caballero08,Kirk12}. Furthermore, several studies \citep{Andersen11,Suenaga13} have determined the ratio of stars to BDs (the `substellar ratio' $\mathcal{R}_{\rm ss}$) as a function of distance from the centre of the Orion Nebular Cluster (ONC) and there is tentative evidence for a decrease in $\mathcal{R}_{\rm ss}$ as a function of distance from the cluster centre, though measuring the substellar mass function in this region (and others) remains challenging \citep[e.g.][]{Alves12,DaRio12,Lodieu12}. Taken at face value, these results suggest that brown dwarfs have different spatial distributions to stars in some (but not all) star forming regions and clusters. This could imply that brown dwarfs form through a different mechanism to stars in those regions, \citep[e.g.][]{Thies08}, or perhaps that dynamical interactions alter their spatial distribution in some regions \citep[e.g.][]{Adams02,Reipurth01,Goodwin05c}, but not others. In order to test this, $N$-body simulations (which can be repeated many times with different random number seeds to guage the level of stochasticity in the initial conditions) of the evolution of young star forming regions should be analysed with the same method(s)/techniques(s) used to analyse observational data. In this paper, we use three different diagnostics to compare the spatial distributions of stars and BDs in numerical simulations of the evolution of star-forming regions. We measure the ratio of stars to BDs ($\mathcal{R}_{\rm ss}$) as a function of distance from the cluster centre; we compare the `local density ratio' of stars and BDs using the $\Sigma_{\rm LDR}$ method \citep{Maschberger11,Parker14b}, and we compare the global spatial distributions using the `mass segregation ratio' $\Lambda_{\rm MSR}$ \citep{Allison09a}. We then re-examine the ONC data from \citet{Andersen11} to look for differences in the local density of BDs compared to stars using $\Sigma_{\rm LDR}$, and the relative spatial distribution using $\Lambda_{\rm MSR}$. | \label{conclude} We have used three different diagnostics to look for differences in the spatial distributions of stars compared to brown dwarfs in $N$-body simulations of star-forming regions. We find that determining the $\mathcal{R}_{\rm ss}$ ratio as a function of distance from the cluster centre cannot be used on its own to draw conclusions on the spatial distribution of BDs compared to stars. In a cluster with a radially decreasing $\mathcal{R}_{\rm ss}$ ratio, the brown dwarfs may have a spatial distribution that is indistinguishable from stars ($\Lambda_{\rm MSR} = 1$, $\Sigma_{\rm LDR} = 1$). Similarly, the inverse can also be true; the BDs have a significantly different spatial distribution compared to stars in that they are more spread out ($\Lambda_{\rm MSR} << 1$ and/or $\Sigma_{\rm LDR} << 1$), but the $\mathcal{R}_{\rm ss}$ ratio increases or remains constant towards the outskirts of the cluster. These findings lead us to strongly advocate the use of more than one diagnostic when assessing the spatial distributions of BDs compared to stars in star-forming regions. When applied to data from the ONC, the $\mathcal{R}_{\rm SSR}$ ratio \emph{and} $\Sigma_{\rm LDR}$ ratio -- and tentative evidence from $\Lambda_{\rm MSR}$ -- do suggest that the BDs are more spread out than stars. However, this dataset is spatially incomplete, and a more comprehensive survey of the ONC would be highly desirable. Randomly distributing masses drawn from an IMF can result (in 1/20, or 5\,per cent, of simulations) in a radially decreasing $\mathcal{R}_{\rm ss}$ ratio before dynamical evolution; which may or may not be mirrored in the $\Lambda_{\rm MSR}$ and $\Sigma_{\rm LDR}$ measurements. Furthermore, dynamical evolution leads to significant differences between the spatial distributions of stars and BDs in more than 25\,per cent of our simulations. This implies that a large observational sample of regions/clusters is needed to assess whether the primordial spatial distributions of stars and BDs are different (which would suggest that their formation mechanisms are different). | 14 | 3 | 1403.7053 |
1403 | 1403.0517_arXiv.txt | Strongly interacting galaxies undergo a short-lived but dramatic phase of evolution characterized by enhanced star formation, tidal tails, bridges and other morphological peculiarities. The nearest example of a pair of interacting galaxies is the Magellanic Clouds, whose dynamical interaction produced the gaseous features known as the Magellanic Stream trailing the pair's orbit about the Galaxy, the Bridge between the Clouds, and the Leading Arm, a wide and irregular feature leading the orbit. Young, newly formed stars in the Bridge are known to exist, giving witness to the recent interaction between the Clouds. However, the interaction of the Clouds with the Milky Way is less well understood. In particular, the Leading Arm must have a tidal origin, however no purely gravitational model is able to reproduce its morphology and kinematics. A hydrodynamical interaction with the gaseous hot halo and disk of the Galaxy is plausible as suggested by some models and supporting neutral hydrogen (H~I) observations. Here we show for the first time that young, recently formed stars exist in the Leading Arm, indicating that the interaction between the Clouds and our Galaxy is strong enough to trigger star formation in certain regions of the Leading Arm --- regions in the outskirts of the Milky Way disk ($R\sim 18$ kpc), far away from the Clouds and the Bridge. | The Magellanic Clouds (MCs) offer a unique opportunity to study galaxy interactions in unprecedented detail due to their proximity to the Milky Way (MW). Thus, detailed mapping of their gaseous content, the 3D kinematics of their stellar content, and the chemical-abundance makeup of these components are readily available for the Clouds. The most obvious features of their interaction are the H~I structures known as the $\sim 200^\circ$-long Magellanic Stream (MS), the Bridge, and the Leading Arm (LA) (Nidever et al. 2010). The recent work on the modeling of the Clouds' interaction by Diaz \& Bekki (2012) makes a compelling case for tidal model, where the MS, Bridge and LA are made primarily of material pulled out from the SMC during two close encounters between the two Clouds. The first encounter took place $\sim 2$ Gyr ago, and the second $\sim 200$ Myr ago. This work used the most recent absolute proper-motion determinations for the Clouds: one HST-based (Kallivayalil et al. 2006), the other ground-based (Vieira et al. 2010). Both determinations imply exactly two encounters between the Clouds to reproduce the MS, LA and Bridge. As for their motion relative to the MW, HST measurements(Kallivayalil 2006, 2013) favor the scenario where the MCs are on the first passage about the MW, with the MS and LA determined solely by the tidal interaction between the Clouds. The ground-based proper-motion measurement allows for two pericentric passages of the Clouds about the MW in the past 2.5 Gyr, and thus some tidal influence of the Galaxy in the formation of the MS, LA and Bridge is expected. The main drawback of the tidal models is that, while they produce a leading arm, all fail to reproduce the observed multi-branches morphology of the LA, and its kinematics. A model by Diaz \& Bekki (2011) that also includes a hydrodynamical interaction of the LA with the diffuse, hot gaseous halo of the MW, better reproduces the kinematics along the LA, but not its morphology. The LA has a complex structure, possibly made of as many as four substructures according to For et al. (2013) and Venzmer et al. (2012), situated above and below the Galactic plane, and encompassing $\sim 60^\circ$ in width. It has been argued that there is a strong drag exerted by the MW gaseous disk on the H~I substructures in the LA (McClure et al. 2008, Venzmer et al. 2012). This is implied by the head-tail velocity structure of the H~I clouds in the LA, as well as the velocity gradient seen in a given LA substructure/arm (Venzmer et al. 2012). Thus it would be enlightening to search for newly formed stars in the LA, an expected result of the hydrodynamical interaction between the MS gas and the MW gaseous disk and halo. We also note that the most recent ($\sim 200 $ Myr ago) encounter between the two Clouds which created the Bridge, is abundantly accompanied by recent star formation, a fact well known since the work by Irwin et al. (1990) and subsequent follow-up by, e.g., Demers et al. (1998). In a recent study Casetti-Dinescu et al. (2012) listed 567 OB-type star candidates in a $\sim$ 7900 square degree area encompassing the periphery of the Clouds, the Bridge, the LA, and part of the MS. The photometric and proper-motion selection was aimed at finding hot (earlier than B5) and distant stars. Also, the proper-motion selection was aimed at selecting stars with motions consistent with membership to the Magellanic system. In the LA region, three stellar overdensities were found comprising a total of 45 candidates. This is a lower limit of such candidates, since the study is area-wise incomplete (Casetti-Dinescu et al. 2012). Here, we have spectroscopically observed 42 of the 45 candidates. Their spatial distribution is shown in Figure 1. Also shown is the H~I distribution from the GASS survey (McClure-Griffith et al. 2009, Kalberla et al. 2010) for which we have restricted the velocity with respect to the Local Standard of Rest, to be $ 150 \le V_{LSR} \le 400$ km/s. The three candidate overdensities we label: A at $(\Lambda_M, B_M) \sim (15^\circ,-22^\circ)$, B at $(\Lambda_M,B_M) \sim (42^\circ,-8^\circ)$, and C at $(\Lambda_M,B_M) \sim (52^\circ,28^\circ)$. In what follows, we dscribe the spectroscopic observations and the results. | 14 | 3 | 1403.0517 |
|
1403 | 1403.0721_arXiv.txt | {Rapid advancements in light-curve and radial-velocity curve modelling, as well as improvements in the accuracy of observations, allow more stringent tests of the theory of stellar evolution. Binaries with rapid apsidal advance are particularly useful in this respect since the internal structure of the stars can also be tested.} {Thanks to its long and rich observational history and rapid apsidal motion, the massive eclipsing binary \yc represents one of the cornerstones of critical tests of stellar evolutionary theory for massive stars. Nevertheless, the determination of the basic physical properties is less accurate than it could be given the existing number of spectral and photometric observations. Our goal is to analyse all these data simultaneously with the new dedicated series of our own spectral and photometric observations from observatories widely separated in longitude.} {We obtained new series of \ubv\ observations at three observatories separated in local time to obtain complete light curves of \yc for its orbital period close to 3 days. This new photometry was reduced and carefully transformed to the standard \ubv\ system using the HEC22 program. We also obtained new series of red spectra secured at two observatories and re-analysed earlier obtained blue electronic spectra. Reduction of the new spectra was carried out in the \iraf and \spefo programs. Orbital elements were derived independently with the \fotel and \phoebe programs and via disentangling with the program \korele. The final combined solution was obtained with the program \phoebe.} {Our analyses provide the most accurate value of the apsidal period of ($47.805\pm0.030$)~yrs published so far and the following physical elements: $M_1=17.72\pm0.35$~\ms, $M_2=17.73\pm0.30$~\ms, $R_1=5.785\pm0.091$~\rs, and $R_2=5.816\pm0.063$~\rs. The disentangling thus resulted in the masses, which are somewhat higher than all previous determinations and virtually the same for both stars, while the light curve implies a slighly higher radius and luminosity for star~2. The above empirical values imply the logarithm of the internal structure constant $\log~k_2 = -1.937.$ A comparison with Claret's stellar interior models implies an age close to $2\times10^6$~yrs for both stars. } {The claimed accuracy of modern element determination of 1--2 per cent still seems a bit too optimistic and obtaining new high-dispersion and high-resolution spectra is desirable.} | The results of a very detailed study of a huge body of spectral and photometric observations of an archetypal early-type eclipsing binary with a fast apsidal motion can be summarized as follows: \begin{figure} \resizebox{\hsize}{!}{\includegraphics{logk2.eps}} \caption{The comparison of the observed range of the logarithm of the internal structure constant with its theoretical value based on the evolutionary models of \citet{claret2004}, calculated for log~M 1.2 and 1.3 (masses 15.85~\ms\ and 19.95~\ms, respectively). The observed and theoretical values agree for an evolutionary age slightly over 2\,000\,000 years.} \label{logk} \end{figure} \begin{enumerate} \item A comparison of separate analyses of radial-velocity and photometric observations and of recorded times of minima covering about 130 years shows that mutually consistent determinations of the sidereal and anomalistic period and the period of apsidal-line rotation are obtained. The most accurate results are based on the \phoebe analysis of all complete light curves. \item The new values of masses, radii, effective temperatures, and luminosities of the binary components are now quite robust and indicate a close similarity of both bodies. However, it is a bit disappointing that these new solutions do not represent a substantial improvement over the previously published determinations. The problem arises partly from the fact that all electronic spectra at our disposal have a moderate spectral resolution (from 13000 to 20000), and partly in inherent problems on the side of theory. For instance, we were unable to achieve a perfect match between the disentangled and synthetic spectra. In addition, the true uncertainties in the estimated effective temperatures are surely much higher than the formal errors in the final solution. We note that \citet{hill95} estimated effective temperatures $31000\pm2000$~K and $31600\pm2000$~K and \lgg = 4.00 [cgs] and 4.00 [cgs] for stars~1 and 2, respectively. \citet{simon94b} carried out a detailed comparison of the observed and synthetic NLTE spectra derived by \citet{kunze90, kunze94} and arrived at $T_{\rm eff1} = 34200\pm600$~K, $T_{\rm eff2} = 34500\pm600$~K, log~$g_1 = 4.18$ [cgs], and log~$g_2 = 4.16$ [cgs]. As pointed out to us by Hubeny (2013, priv. comm.), the theoretical spectra computed by Kunze do not include the metal line blanketing. This indicates that they probably {\sl overestimate} the effective temperatures by as much as $\sim$2000~K. Even when one compares the synthetic spectra from the O-star grid, and the later published spectra from the B-star grid \citep{lanz2007} in the overlapping temperature range, it turns out that the two sets of model spectra differ from each other since the B-star values are based on somewhat improved input physics. Clearly, without further progress in the model spectra, one cannot do any better. \item We also note that our new masses, radii, and effective temperatures, together with the observed rate of the apsidal motion, imply the relativistic contribution to the apsidal motion $\dot\omega_{\rm rel.} = 0.0009645\,\,^\circ$d$^{-1}$, i.e. about 4.6~\% of the total observed rate of the apsidal rotation. This implies the observed value of the logarithm of the internal structure constant (average of the two nearly identical stars) $\log~k_2 = -1.937\pm0.011.$ In Fig.~\ref{logk} we compare the range of this observed value (within the quoted error) with the evolutionary models of \citet{claret2004}. The observed and theoretical values agree with each other for an evolutionary age slightly over $2\times10^6$~yrs. \end{enumerate} | 14 | 3 | 1403.0721 |
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1403 | 1403.7047_arXiv.txt | {Gum~31 is a prominent, but still rather poorly studied \ion{H}{ii} region around the stellar cluster NGC~3324 at the northwestern periphery of the Carina nebula complex. } {Our aim was to reveal and characterize the young stellar population in Gum~31. An X-ray survey is the only efficient way to identify young stars in this region with extremely high galactic field-star contamination that can avoid the strong biases of infrared excess selected samples of disk-bearing young stars. } {We used the \textit{Chandra} observatory to perform a deep (70 ksec) X-ray observation of the Gum~31 region and detected 679 X-ray point sources. This extends and complements the X-ray survey of the central Carina nebula regions performed in the \textit{Chandra Carina Complex Project} (CCCP). Using deep near-infrared images from our recent VISTA survey of the Carina nebula complex, our comprehensive \textit{Spitzer} point-source catalog, and optical archive data, we identify counterparts for 75\% of these X-ray sources. } {The spatial distribution of the X-ray selected young stars shows two major concentrations, the central cluster NGC~3324 and a partly embedded cluster in the southern rim of the \ion{H}{ii} region. However, these two prominent clusters contain only about 30\% of the X-ray selected population, whereas the majority ($\sim 70\%$) of X-ray sources constitute a rather homogeneously distributed population of young stars. Our color-magnitude diagram analysis suggests ages of $\sim 1-2$~Myr for the two clusters, whereas the distributed population shows a wider age range up to $\sim 10$~Myr. We also identify previously unknown companions to two of the three O-type members of NGC~3324 and detect diffuse X-ray emission in two parts of the region. } {An extrapolation based on the observed X-ray luminosity function suggests that the observed region contains about 4000 young stars in total (down to $0.1\,M_\odot$). This shows that the Gum~31 area contains a substantial fraction of the total stellar population in the CNC. The distributed population of young stars in the Gum~31 region is probably a part or extension of the widely distributed population of $\sim 1 - 10$~Myr old stars, that was identified in the CCCP area. This implies that the global stellar configuration of the Carina nebula complex is a very extended stellar association, in which the (optically prominent) clusters contain only a minority of the stellar population. } | The Carina nebula complex (CNC) is one of the largest, most massive, and most active star-forming complexes in our Galaxy. Located at a moderate and well known distance of 2.3~kpc \citep{Smith06}, the nebulosity extends over about 100~pc, corresponding to several degrees on the sky. The clouds contain a total gas and dust mass of about $10^6\,M_\odot$ \citep{Preibisch12} and harbor presumably about 100\,000 young stars \citep{CCCP-Clusters,Povich11,HAWKI-survey}. The star formation rate of $\sim 0.01 - 0.02\,M_\odot/{\rm yr}$ \citep[see][]{Povich11,Gaczkowski13} constitutes as much as about 1\% of the total galactic star formation. The large population of massive stars \citep[$\ge 70$ O-type and WR stars][]{Smith06}, including the object $\eta$~Car (which is the most luminous known star in our Galaxy) creates very high levels of ionizing radiation and stellar wind power, which profoundly influence the surrounding clouds. This stellar feedback has already dispersed a large fraction of the original dense molecular clouds, out of which the stars formed, but the compression of the remaining clouds by the ionization fronts and expanding wind bubbles is also currently triggering the formation of new generations of stars in the complex \citep{Smith10b,Gaczkowski13}. A general review of the CNC is provided in the book chapter by \citet{SB08}. During the last five years, several new and sensitive surveys of different parts of the CNC have been performed at wavelengths from the X-ray to submm regime. One major milestone was the deep X-ray imaging survey of the {\it Chandra} Carina Complex Project \citep[CCCP; see][for an overview]{CCCP-intro}, which mapped an area of about 1.4 square-degrees with a mosaic of 22 individual ACIS-I pointings, using a total observing time of 1.34 Mega\-seconds (15.5 days). With a detection limit corresponding to X-ray luminosities of about $10^{30}\,\rm erg/s$, these X-ray data can detect the coronal X-ray emission of young ($\la 10^7$ yrs) stars down to $\sim 0.5\,M_\odot$. The \textit{Chandra} images revealed 14\,368 individual X-ray sources, and a sophisticated classification scheme showed that 10\,714 of these are most likely young stars in the Carina nebula \citep{CCCP-classification}. The analysis of the spatial distribution of the X-ray detected young stars showed that half of the young stellar population resides in one of about 30 clusters or stellar groups, while the other half constitutes a widely dispersed population \citep{CCCP-Clusters} that is spread throughout the entire observed area. The combination of these CCCP X-ray data with a deep near-infrared survey \citep{HAWKI-survey} obtained with HAWK-I at the ESO VLT provided information about the properties of the stellar populations in the central parts of the complex \citep{CCCP-HAWKI}, including the individual stellar clusters Tr~16 and Tr~15 \citep{CCCP-Tr16,CCCP-Tr15}. Because of the wide spatial extent of the CNC, most recent studies (and all those mentioned above) have focused on the central $\la 1.5$~square-degree area. At the northwestern periphery of the CNC, about $80'$ (or $\approx 50$~pc) from $\eta$~Car, the \ion{H}{ii} region Gum~31 \citep[first described by][]{Gum55} constitutes the most prominent object in this area of the sky (see Fig.~1). Gum~31 is a roughly circular nebula with a diameter of about $12'$ ($\approx 8$~pc), and is excited by the young stellar cluster NGC~3324 (containing three known O-type stars). Despite its interesting morphology and the publicity of \textit{Hubble Space Telescope}\footnote{{\tt hubblesite.org/newscenter/archive/releases/2008/34/} } and ground-based optical images\footnote{e.g., ESO Photo Release \,\,{\tt \footnotesize www.eso.org/public/news/eso1207/}}, the Gum~31 nebula (sometimes designated as the Gabriela Mistral nebula) remained quite poorly studied until very recently, probably because the closeness to the extremely eye-catching Carina nebula has always overshaded Gum~31. In the cluster NGC~3324, stellar spectral types are only known for the three optically brightest stars, HD~92\,206 A, B, and C, which have been classified as O6.5V (HD~92\,206 A and HD~92\,206 B) and O9.5V (HD~92\,206 C) \citep{Mathys88,Walborn82}. \citet{Carraro01} identified 25 candidate members by means of optical photometry and suggested that the cluster age is presumably a few Myr. In addition to the optically visible cluster NGC~3324 in the \ion{H}{ii} region, a few dozen of embedded young star candidates have been found in the clouds surrounding the \ion{H}{ii} region, either in infrared images \citep{Cappa08} or traced by their protostellar jets \citep{Smith10a,Ohlendorf12}. However, these objects known until recently must represent only the tip of the iceberg: given the presence of three O-type stars ($M \ge 18\,M_\odot$), the extrapolation of the canonical stellar IMF \citep{Kroupa02} would suggest a total population of $\approx 1500$ low-mass ($0.1\,M_\odot \le M \le 2\,M_\odot$) stars. It is thus quite obvious that the vast majority of the stellar population in the Gum~31 region is still completely unknown. While \citet{Walborn82} already concluded that the distance modulus values of the brightest stars in NGC~3324 are consistent with those of the clusters Tr~16 / Col 228 in the central parts of the Carina nebula, the apparent segregation of the Gum~31 nebula from the bright H$\alpha$ emission of the Carina nebula in optical images left the physical relation unclear. New information on this question came from our recent \textit{Herschel} far-infrared observations \citep{Preibisch12,Roccatagliata13} of the CNC. The field-of-view of this survey was wide enough (more than 5 square-degrees) to include the area around Gum~31. The \textit{Herschel} maps showed that the dense dusty shell surrounding the Gum~31 \ion{H}{ii} region is connected by numerous filamentary cloud structures to the dense molecular clouds in the inner parts of the Carina nebula, suggesting Gum~31 to be part of the CNC. The shell around Gum~31 and the surrounding clouds host several dozens of deeply embedded protostars and pre-stellar cores, which are detected as point-like sources in our \textit{Herschel} data \citep{Gaczkowski13}. In the recent study of \citet{Ohlendorf13}, these \textit{Herschel} data were combined with \textit{Spitzer} archive data and the Wide-field Infrared Survey Explorer (WISE) point-source catalog and lead to the identification of some 600 young stellar object candidates in a one square-degree area centered on Gum~31 by means of their infrared excesses. A clear concentration of partly embedded young stellar objects is located in a dense cloud at the southwestern edge of the Gum~31 bubble. This very young cluster, which we will designate as G286.38--0.26 in following text, is located at a position where the Gum~31 shell seems to interact with the expanding super-bubbles driven by the numerous massive stars in the central parts of the CNC. This may represent an interesting example of triggered star formation in a cloud that formed and/or is compressed by colliding large-scale shocks. The current, infrared-excess selected sample of protostars and disk-bearing young stars in the Gum~31 region, resulting from our analysis of the \textit{Herschel} and \textit{Spitzer} data \citep{Ohlendorf13}, contains objects down to stellar masses of about $1\,M_\odot$. However, since the stars in this region are presumably already several Myr old, this infrared-excess selected sample must be highly incomplete, because it is well known that the typical lifetime of circumstellar disks around young stars are just a few Myr \citep[e.g.,][]{Fedele10}. The young stars which have already dispersed their disks will no longer exhibit infrared excesses and can thus not be identified by infrared excess selection. This represents a major obstacle in the identification of the young stellar population in the Gum~31 region, which is a fundamental prerequisite for an determination of its star formation history. Furthermore, because Gum~31 lies very close to the galactic plane ($b \approx -0.2^\circ$) a great deal of confusion results from galactic field star contamination. All optical and infrared images of this region are thus completely dominated by unrelated field stars, and it is therefore impossible to identify and distinguish a population of several Myr old low-mass stars from unrelated field stars with optical or infrared photometry alone. An X-ray survey, however, can solve this problem: the strongly enhanced X-ray emission of young ($\la 10^7$~yrs) stars \citep[see, e.g.,][]{Feigelson07,Preibisch_coup_orig} provides an extremely useful discriminant between young pre-main sequence stars and the much older field stars. The median X-ray luminosity of a few Myr old solar-mass stars is nearly 1000 times higher than for solar-mass field stars \citep[see][]{PF05}, and makes these young stars relatively easily detectable targets for current observatories, even at the relative large distance of 2.3~kpc. This was the motivation to perform the deep {\em Chandra} X-ray observation of the Gum~31 region presented in this paper. In Section 2 we describe the X-ray observation and data analysis and discuss basic properties of the X-ray sources. We then describe in Sect.~3 the deep near-infrared data from our recent VISTA survey of the CNC and our \textit{Spitzer} point-source catalog that are then used in Sect.~4 to identify and characterize the counterparts of the X-ray sources. In Sect.~5 we discuss the X-ray and infrared/optical properties of particularly interesting objects. In Sect.~6 we use near- and mid-infrared color-color and color-magnitude diagrams to infer the extinctions, infrared excesses, masses, and ages of the X-ray detected objects. Section 7 contains a discussion of the global properties of the X-ray selected population in Gum~31. In Sect.~8 we present and briefly discuss the diffuse X-ray emission that is detected in our {\em Chandra} image. Finally, section 9 contains a brief summary and the conclusions from this study. \begin{figure*} \parbox{9cm}{\includegraphics[width=9cm]{23133f1a.eps}} \parbox{9.2cm}{\includegraphics[width=9.2cm]{23133f1b.eps}} \caption{Left: Optical image of the Carina nebula complex (from: {\tt www.eso.org/public/images/eso0905b/}; image credit: ESO/Digitized Sky Survey 2, Davide De Martin). The region observed in the context of the \textit{Chandra} Carina Complex Project (CCCP) and the \textit{Chandra} pointing of the Gum~31 region are marked by the cyan outlines.\newline Right: Optical image of the Gum~31 region obtained with the Wide Field Imager on the MPG/ESO 2.2~m telescope at La Silla Observatory (from: {\tt www.eso.org/public/images/eso1207a/}; image credit: ESO) with outline of the \textit{Chandra} pointing in cyan. The individual X-ray point sources are marked by green crosses. The yellow boxes mark the two regions shown in Fig.~\ref{n3324-g286.fig} (box {\sf A} corresponds to the region of the cluster NGC~3324, while box {\sf B} corresponds to the region of the cluster G286.38--0.26) } \label{outlines.fig}% \end{figure*} | 14 | 3 | 1403.7047 |
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1403 | 1403.5797_arXiv.txt | Accretion disks arising from neutron star- neutron star mergers or black hole- neutron star mergers produce large numbers of neutrinos and antineutrinos. In contrast to other astrophysical scenarios, like supernovae, in mergers the antineutrinos outnumber the neutrinos. This antineutrino dominance gives neutrinos from merger disks the opportunity to exhibit new oscillation physics, specifically a matter-neutrino resonance. % We explore this resonance, finding that consequences can be a large transition of $\nu_e$ to other flavors, while the $\bar{\nu}_e$s return to their initial state. We present numerical calculations of neutrinos from merger disks and compare with a single energy model. We explain both the basic features and the conditions for a transition. | 14 | 3 | 1403.5797 |
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1403 | 1403.5274_arXiv.txt | We study the evolution of angular momenta of ($M_*=10^{10}-10^{12}\msun$) galaxies utilizing large-scale ultra-high resolution cosmological hydrodynamic simulations and find that spin of the stellar component changes direction frequently, caused by major mergers, minor mergers, significant gas inflows and torques by nearby systems. The rate and nature of change of spin direction can not be accounted for by large-scale tidal torques, because the latter fall short in rates by orders of magnitude and because the apparent random swings of the spin direction are inconsistent with alignment by linear density field. % The implications for galaxy formation as well as intrinsic alignment of galaxies are profound. Assuming the large-scale tidal field is the sole alignment agent, a new picture emerging is that intrinsic alignment of galaxies would be a balance between slow large-scale coherent torquing and fast spin reorientation by local interactions. What is still open is whether other processes, such as feeding galaxies with gas and stars along filaments or sheets, introduce coherence for spin directions of galaxies along the respective structures. | The angular momentum or spin of galaxies is a physical quantity that is far from being fully understood but is of fundamental importance to galaxy formation and cosmological applications. While N-body simulations have shed useful light on spin properties of dark matter halos \citep[e.g.,][]{2002Vitvitska}, it is expected that, given the vastly different scales between the stellar component and dark matter halo component and different physical processes governing stellar, gas and dark matter components, the angular momentum dynamics of galaxies may be quite different and not necessarily inferable from N-body simulations with any reasonable accuracy. We herewith perform a detailed analysis of the dynamics of spin of galaxies in a full cosmological context, utilizing {\it ab initio} LAOZI cosmological hydrodynamic simulations of the standard cold dark matter model \citep[][]{2014Cen} with an unprecedented galaxy sample size and ultra-high numerical resolution. This paper is the second in the series ``On the Origin of the Hubble Sequence". | Utilizing {\it ab initio} {\color{red}\bf L}arge-scale {\color{red}\bf A}daptive-mesh-refinement {\color{red}\bf O}mniscient {\color{red}\bf Z}oom-{\color{red}\bf I}n cosmological hydrodynamic simulations ({\color{red}\bf LAOZI Simulation}) of the standard cold dark matter model, we study the evolution of angular momenta of massive ($M_*=10^{10}-10^{12}\msun$) galaxies. The simulation has an ultra-high resolution of $\le 114$pc/h and contains more than 300 galaxies with stellar mass greater than $10^{10}\msun$. We find that spin of the stellar component changes direction frequently, caused by major mergers, minor mergers, significant gas inflows and torques by nearby systems, with a median time in the range $1-10$Myr for directional change of spin vector by $1$ degree of arc. The spin of the gas component changes at a factor of $5-10$ higher rates than the stellar component. Because the processes that are responsible are mostly in the nonlinear regime. we do not expect that the findings significantly depend on precise cosmological parameters, The rate of change of spin direction can not be accounted for by large-scale tidal torques, because the latter fall short in rates by $2-3$ orders of magnitude. In addition, the nature of change of spin direction - apparent random swings - is inconsistent with alignment by linear density field. % A new paradigm emerging with respect to intrinsic alignment of galaxies is that it is determined, primarily, by a balance between slow large-scale coherent torquing (if it were the sole alignment process) and fast spin reorientation by local interactions. This suggests that a significant revision to the large-scale tidal torque based alignment theory is perhaps in order. A simple analysis presented here indicates that intrinsic alignment of galaxies is dependent on redshift, luminosity, environment and galaxy type. Specifically, it is found that the alignment (1) decreases with increasing redshift, (2) decreases with decreasing stellar mass, and (3) is larger for elliptical galaxies than for spiral galaxies. While no detailed comparisons are made, the found trends appear to be broadly consistent with and thus provide the physical basis for the observed trends. What remains open is whether other processes, such as feeding galaxies with gas and stars along filaments or sheets, introduce some coherence of their own kind for spin direction of galaxies along the respective structures. This will require a separate study in greater detail. \vskip 1cm I would like to thank Claire Lackner for providing the SQL based merger tree construction software, and \citep[][]{2011Turk} for providing the very useful analysis and visualization program yt. Computing resources were in part provided by the NASA High-End Computing (HEC) Program through the NASA Advanced Supercomputing (NAS) Division at Ames Research Center. This work is supported in part by grant NASA NNX11AI23G. The simulation data are available from the author upon request. | 14 | 3 | 1403.5274 |
1403 | 1403.4272_arXiv.txt | {Dust and stars play a complex game of interactions in the interstellar medium and around young stars. The imprints of these processes are visible in scaling relations between stellar characteristics, star formation parameters, and dust properties.}{In the present work, we aim to examine dust scaling relations on a sub-kpc resolution in the Andromeda galaxy (M31). The goal is to investigate the properties of M31 on both a global and local scale and compare them to other galaxies of the local universe.}{New \textit{Herschel} observations are combined with available data from GALEX, SDSS, WISE, and \textit{Spitzer} to construct a dataset covering UV to submm wavelengths. All images were brought to the beam size and pixel grid of the SPIRE $500~\mu\mathrm{m}$ frame. This divides M31 in $22437$ pixels of $36$ arcseconds in size on the sky, corresponding to physical regions of $137 \times 608$ pc in the galaxy's disk. A panchromatic spectral energy distribution was modelled for each pixel and maps of the physical quantities were constructed. Several scaling relations were investigated, focussing on the interactions of dust with starlight.}{We find, on a sub-kpc scale, strong correlations between $M_\mathrm{dust}/M_\star$ and NUV--r, and between $M_\mathrm{dust}/M_\star$ and $\mu_\star$ (the stellar mass surface density). Striking similarities with corresponding relations based on integrated galaxies are found. We decompose M31 in four macro-regions based on their FIR morphology; the bulge, inner disk, star forming ring, and the outer disk region. In the scaling relations, all regions closely follow the galaxy-scale average trends and behave like galaxies of different morphological types. The specific star formation characteristics we derive for these macro-regions give strong hints of an inside-out formation of the bulge-disk geometry, as well as an internal downsizing process. Within each macro-region, however, a great diversity in individual micro-regions is found, regardless of the properties of the macro-regions. Furthermore, we confirm that dust in the bulge of M31 is heated only by the old stellar populations.}{In general, the local dust scaling relations indicate that the dust content in M31 is maintained by a subtle interplay of past and present star formation. The similarity with galaxy-based relations strongly suggests that they are in situ correlations, with underlying processes that must be local in nature.} | The interstellar medium (ISM) harbours a rich variety of materials that all interact with one another through multiple chemodynamical processes. Hydrogen is by far the most abundant and occurs primarily in its neutral form, varying from warm ($\sim 8000$~K) to cold gas ($\sim 80$~K). In the very cold and dense environments, it is converted into molecular hydrogen (H$_2$). The presence of ionized hydrogen (H{\sc{ii}}) increases as new stars irradiate the neutral gas. The formation of stars and their associated nucleosynthesis are the principal processes driving the chemical enrichment of the ISM. A good fraction of the produced metals are locked up in dust grains of various sizes. The micron-sized grains are thought to be silicic and carbonaceous in nature, while there is a population of nanometre-sized polycyclic aromatic hydrocarbons (PAHs). Although it only accounts for a relatively small fraction of the ISM mass, interstellar dust plays a vital role as a catalyst in the formation of H$_2$, which is crucial in the star formation process. At the same time, dust tends to absorb up to $50\%$ of the optical and ultraviolet (UV) light of stars \citep{PopescuTuffs2002}, heavily affecting our view of the universe. It re-emits the absorbed energy at longer wavelengths in the mid-infrared (MIR), far-infrared (FIR), and submillimetre (submm) bands. To study the ISM in greater detail, spatially resolved FIR observations are crucial as they constrain the local dust distribution and properties. Recent space missions were able to uncover the MIR to submm window in great detail with the Spitzer Space Telescope \citep{Spitzer}, the \textit{\textit{Herschel}} Space Observatory \citep{Herschel}, and the Widefield Infrared Survey Explorer \citep[WISE,][]{WISE}. Correlations between the main properties of dust, gas, and stars, also known as \textit{scaling relations}, define a tight link between these constituents. In the past, relations between the dust-to-gas ratio on the one hand and metallicity or stellar mass on the other hand were the only notable relations that were being investigated \citep{Issa1990,Lisenfeld1998,Popescu2002,Draine2007,Galametz2011}. Only recently were other scaling laws more systematically investigated. Namely, the ratio of dust to stellar mass, the specific dust mass, which was found to correlate with the specific star formation rate (i.e. star formation rate divided by the stellar mass, \citealt{Brinchmann2004, daCunha2010,Rowlands2012,SmithD2012}), NUV--r colour (i.e. the difference in absolute magnitude between the GALEX NUV and SDSS $r$ band), and stellar mass surface density $\mu_\star$ \citep{Cortese2012,Agius2013}. Each of the above scaling laws was derived on a galaxy-galaxy basis, considering galaxies as independent systems in equilibrium. In order to fully understand the coupling of dust with stars and the ISM, we must zoom in on individual galaxies. This is however troublesome at FIR/submm wavelengths because of the limited angular resolution. Only in the last few years, we have been able, thanks to \textit{Herschel}, to observe nearby galaxies in the FIR and submm spectral domain whilst achieving sub-kpc resolutions for the closest galaxies ($< 5.7$~Mpc) in the $500~\mu\mathrm{m}$ band. Dust scaling relations on subgalactic scales were thus so far limited to gas-to-dust ratios for a handful of local galaxies \citep[see e.g.][]{MunozMateos2009b, Bendo2010a, Magrini2011, Sandstrom2012, Parkin2012}. Today, exploiting PACS \citep{PACS} and SPIRE \citep{SPIRE}, plus the 3.5-metre mirror onboard \textit{Herschel}, IR astronomy has gone a leap forward. Spatial resolutions and sensitivities have been reached that allow us, for the first time, to accurately characterise the dust emission in distinct regions of nearby galaxies. The observed spectral energy distribution (SED) at these frequencies can be reproduced by means of theoretical models. In particular, a modified black body can be fitted to the observed FIR ($\lambda >100~\mu\mathrm{m}$) data extracted from sub-kpc regions in galaxies \citep[see e.g.][]{Smith2010,Hughes2014}. While still a step forward with respect to previous works, this approach inevitably suffers from some simplifications. For example, if dust is heated by different sources, it cannot be truthfully represented by a single temperature component. Other studies make use of the physical dust models from \citet {DraineLi2007}, which cover the $1 - 1000~\mu\mathrm{m}$ wavelength range (e.g. \citealt{Foyle2013} for M83; \citealt{Aniano2012} for NGC 626 and NGC 6946; and, more recently, \citealt{Draine2014} for M31). These studies use data at shorter wavelengths as well, being thus able to fully sample the spectral range of dust emission, including emission from warmer dust, small transiently heated grains and PAHs. They mainly focus on the distribution and properties of the dust, not including any constraints on the radiation field which heats the dust from for example, UV/optical observations. To properly investigate scaling relations on a local scale, one would ideally need a self--consistent model to derive the desired physical quantities. The complexity of a full stellar and chemical evolution model for galaxies, however, requires some simplifications. The models should treat both stellar and dust components, taking into account their influence on each other (the so--called dust energy balance). Panchromatic emission modelling of subgalactic regions has been carried out by \citet{MentuchCooper2012} for the Whirlpool galaxy (M51). However, the stellar and dust components were treated separately. \citet{Boquien2012,Boquien2013}, have performed panchromatic pixel-by-pixel fits of nearby star forming galaxies using CIGALE \citep{Noll2009}, which does include a dust energy balance. They showed that most of the free parameters could accurately be constrained, given a sufficiently large range of priors. We will perform panchromatic SED fitting, using the MAGPHYS code \citep{Dacunha2008}. The code treats both stellar light and dust emission at the same time, forcing an energy balance. It has an extended library of theoretical SEDs based on the latest version of the stellar population models from \citet{Bruzual2003} and physically motivated, multi-component dust models. Furthermore, it applies a Bayesian fitting method including a thorough error analysis. The proximity of ISM regions is crucial in order to obtain the desired, sub-kpc spatial resolution. The closest giant molecular cloud systems are of course in our own Milky Way, but it is not possible to probe the entire Galaxy. The Magellanic clouds are the nearest galaxies as they are close satellites of the Milky Way. These objects are, however, quite irregular and lower in metallicity and in total mass, hence they do not represent the well--evolved ISM of virialised large galaxies. Andromeda (M31) is the closest large galaxy at a distance $D_\mathrm{M31}=785$~kpc \citep{McConnachie2005}, which means every arcsecond on the sky corresponds to $3.8$ pc along the major axis of M31. Classified as a SAb-type LINER galaxy, M31 is a slow-star forming spiral \citep[SFR~$=0.20\;M_\odot\mathrm{yr}^{-1}$,][]{Ford2013} with an inclination of $77 \degree$ and a position angle of its major axis of $38 \degree$ \citep{McConnachie2005}. The gas and dust components of Andromeda have been extensively studied in the past \citep[e.g.][]{WalterbosSchwering1987,Montaldo2009,Tabatabaei2010} using low--resolution data at FIR wavelengths and simplified models. Although mapped in all wavelengths from UV to the FIR in the past, high--quality submm observations are thus far not available, yet these wavelengths are crucial to constraining the properties of the cold dust. The \textit{Herschel} Exploitation of Local Galaxy Andromeda \citep[][hereafter paper I]{HELGAI} is the first programme that mapped M31 from $100~\mu\mathrm{m}$ to $500~\mu\mathrm{m}$ with \textit{Herschel}, covering a large $5.5 \degree \times 2.5 \degree$ field centred around the galaxy. Even at the sparsest \textit{Herschel} resolution ($36^{\prime\prime}$ at $500~\mu\mathrm{m}$), physical scales of only $140$~pc are resolved. Andromeda is consequently the best suited object for studying the ISM in great detail while allowing at the same time, the comparison with global properties. In \citet{HelgaII} (hereafter paper II), we performed a pixel-by-pixel SED fit to the \textit{Herschel} data and map the main dust properties of Andromeda. \citet{Ford2013} (hereafter paper III) investigate the star formation law in M31 on both global and local scales. A catalogue of giant molecular clouds was recently constructed by \citet{Kirk2014} (hereafter paper IV). We aim to expand on this work by carrying out an in-depth investigation of the dust scaling relations in Andromeda. We do this by fitting panchromatic spectral energy distribution models to each statistically independent $36$--arcsecond region in the galaxy. In this way we have produced the largest and most complete view of the stars and ISM dust in a large spiral galaxy. The arrangement of the paper is as follows. In Sect.~\ref{sec:data} we give an overview of the data used and in Sect.~\ref{sec:results} we briefly discuss the processing of these data and our SED fitting method. Appendix~\ref{app:dataprocessing} goes into more detail on the processing of multi-wavelength data. The results are given in Sect.~\ref{sec:results}, along with the parameter maps of Andromeda. We analyse the dust scaling relations of Andromeda in Sect.~\ref{sec:scaling_rel}. In Sect.~\ref{sec:discussion} we present our discussion and main conclusions. | \label{sec:discussion} In the present work, we have performed SED fitting of a panchromatic dataset, collected for our neighbour galaxy M31. New \textit{Herschel} observations were combined with GALEX, SDSS, WISE, and \textit{Spitzer} data, covering UV to submm wavelengths and allowing us to derive, by exploiting a physically self--consistent model, some physical parameters both on a global and on a local scale. To create statistically independent regions, all the data were convolved to a resolution matching that of the SPIRE $500~\mu\mathrm{m}$ waveband, the lowest in our dataset, which allowed us to probe physical scales of $\sim 137\times 608$ pc in the plane of M31. In this paper we concentrate on the analysis of the scaling relations linking the dust and stellar properties. We have fitted a multi-component theoretical SED to each pixel which allowed us to estimate several physical properties of that region. Every physical parameter for every pixel was given an uncertainty estimation based on the broadness of its corresponding PDF. Physical quantities that could not be sufficiently constrained were removed from the sample. Furthermore, 2-D parameter maps are constructed for each physical quantity. Additionally, we have decomposed Andromeda in four macro-regions: the bulge, the star forming ring (the so--called $10$~kpc ring), and the inner and outer disk regions. The same fitting routine was applied to the integrated fluxes of each of these main components, as well as to the global observed fluxes. From the point of view of the dust scaling relations, M31 is an average galaxy when compared to the local galaxies in the HRS sample. On the other hand, it lies above the average $M_\mathrm{dust}$ vs SFR relation of \citet{daCunha2010}; despite a dust mass close to the sample average, Andromeda is forming stars significantly less efficiently than the other galaxies. By investigating the properties of the distinct morphological components, we find strong hints for an inside--out star formation scenario. In this evolutionary model, the bulk of stars are being formed at early epochs in the bulge, and the more recent star formation happens at larger galactocentric distances (i.e. $\gtrsim 3$~kpc). In particular, the bulk of star formation is currently taking place in the 10~kpc ring, the morphological structure that contains most of the dust in the galaxy. While the analysis presented in this paper can give no strong clues regarding the past star formation history, we do find an enhancement of dust and star formation in the ring with respect to the galactic disk. This would be consistent with a triggering due to a close encounter with the dwarf satellite M32. \cite{Block2006} used numerical N-body simulations to model the effect of the passage of M32 through Andromeda's disk. Beginning with a model with two spiral arms, they end up with a morphological structure closely matching the 10~kpc ring and the ``hole'' which is easily visible in IR images towards the south. Similarly, \citet{Gordon2006} argued that a head--on encounter with M32, might have resulted in star forming waves propagating through the 10~kpc ring. The bulge and inner disk region have red colours and high stellar mass surface density ($\mu_\star$). The star forming ring and outer disk region are bluer and have lower $\mu_\star$. Each of these regions lies on the average trend of the HRS scaling relations. In terms of NUV--r colour, the bulge of M31 is a remarkable exception, being redder than any of the submm detected HRS galaxies. The macro-regions thus have characteristics closely resembling those of global galaxies, where the bulge and inner disk may be seen as early-type while the ring and outer disk resemble late-type galaxies. The results for M31 not only support an inside-out formation pattern for the bulge-disk morphology, but they also tell us something about the differences in the local environments. When looking at the dust--to--stellar mass ratio as a function of the stellar mass surface density (see Fig.~\ref{fig:regions}), we observe a smooth transition from high stellar mass/low dust regions in the bulge, to the low stellar mass/high dust content of the outermost regions. Only the starforming ring, containing a significant fraction of the dust of the whole galaxy, is slightly displaced from this relation. On the other hand, the four regions behave quite differently when the sSFR is plotted as a function of $\mu_\star$ (see Fig.~\ref{fig:regions}). A tendency is found for regions of higher stellar mass surface density to host less star formation, in a way mimicking an internal downsizing process in star formation, in which the regions of highest stellar density have already stopped forming stars. Within each region, the variation in $\mu_\star$ spans only about 1 order of magnitude, whereas the sSFR always varies by more than 3 (with the exception of the bulge, where there is barely any star formation). These relations seem to suggest that downsizing relations break down when considering the smallest scales. Star formation at these scales has a weak dependence on the stellar mass: small-scale environments characterised by different stellar masses can be associated with very different levels of star formation, at least as far as a quiescent galaxy like M31 is concerned. What drives the general star formation mode, which determines where the galaxy as a whole places itself on the scaling relations, must instead be the total mass of all its constituents. When considering the modelling of the observed SED on the smallest scales, our main conclusions are as follows. \begin{enumerate} \item The SED of sub-kpc regions can be successfully fitted using galaxy-based models, provided that the parameter space is adequately sampled. \item When investigating the dust heating in the bulge, we recover the theoretical $(T_\mathrm{C}^\mathrm{ISM})^6 \sim$ $\mu_\star$ relation. This indicates that old stars are the dominant heating source in this region. The dust heating is more ambiguous in the disk, where both star formation and the diffuse ISRF irradiate the dust. \item We find strong correlations, on a pixel--by--pixel scale, between $M_\mathrm{dust}/M_\star$ and NUV--r (or, equivalently, sSFR), and between $M_\mathrm{dust}/M_\star$ and $\mu_\star$. These scaling relations, involving the dusty component of the ISM, are remarkably similar to those found for entire local galaxies. This suggests that the dust scaling relations are built {\it in situ}, with underlying physical processes that must be local in nature. \item As already found for other galaxies, M31 seems to have undergone an inside--out evolution in its star formation process, possibly influenced by interactions with its satellites. \item When considering the smallest scales, a wide range in dust content, sSFR, and $M_\mathrm{dust}/M_\star$ is found within Andromeda illustrating the great diversity of sub-kpc regions. Even within the late-type ring and outer disk of M31, early-type micro-regions can be found. Vice versa, the inner part of M31 still holds a small number of late-type regions. \end{enumerate} The fact that we are able to reproduce the dust scaling relations on a sub-kpc scale states that these relations are not only partially a manifestation of a galaxy-wide equilibrium, but they also arise from local scales. Local evolutionary processes involving dust creation and destruction lie at the base of these relations. The balance between dust depletion and production reflects the relative presence of old and new stars, the latter being responsible for dust generation. This raises the question: at what scales does the balanced interplay between dust and stars break down? Answers to this may be found in similar studies of the Galactic ISM or by future, high--resolution FIR space missions. Zooming into the ISM of our own galaxy can unveil two very different results. If these scaling relations break down at the size of individual molecular clouds, it would indicate that non--local scattered light plays an important role in the dust energy balance. Alternatively, if the scaling relations stay intact, non--local light is negligible at each scale, which would call for a revision of the physical properties of interstellar dust. | 14 | 3 | 1403.4272 |
1403 | 1403.2104_arXiv.txt | Previous claims of significant evidence for mirror-parity in the large-scale cosmic microwave background (CMB) data from the \textit{Wilkinson Microwave Anisotropy Probe} (\textit{WMAP}) experiment have been recently echoed in the first study of isotropy and statistics of CMB data from \textit{Planck}. We revisit these claims with a careful analysis of the latest data available. We construct statistical estimators in both harmonic and pixel space, test them on simulated data with and without mirror-parity symmetry, apply different Galactic masks, and study the dependence of the results on arbitrary choices of free parameters. We confirm that the data exhibit evidence for odd mirror-parity at a significance which reaches as high as $\sim99$ per cent C.L., under some circumstances. However, given the bias exhibited by the pixel-based statistic on a masked sky, its sensitivity to the total power and the dependence of both pixel and harmonic space statistics on the particular form of Galactic masking and other a-posteriori choices, we conclude that these results are not in significant tension with the predictions of the concordance cosmological model. | \label{sec:Introduction} The importance of large scale symmetries in the cosmic microwave background (CMB), if those are shown to exist, cannot be overstated. From a theoretical standpoint, these scales provide us with a possible observational window to the very early Universe, and to the physics at very high energies. Deviations from the expectations of the concordance $\Lambda$ cold dark matter (\LCDM) cosmological model on these scales, such as the breaking of statistical isotropy, could indicate a special location, axis or direction, which in turn can point to the existence of exotic pre-inflationary relics or other effects in the pre-inflationary Universe. Mirror-parity is an intriguing example of a symmetry which breaks statistical isotropy. As it involves a symmetry plane, it can be naturally attributed to various early Universe models predicting large-scale modulation of the CMB (\citealt{Gordon:2005jk}; \citealt*{Ackerman:2007sf}; \citealt{Hou:2010xy,Schmidt:2013ty}). Another interesting class of models that could induce this symmetry involves a finite topology for the Universe \citep*{de-Oliveira-Costa:1996qy}. In such models, statistical isotropy on large scales may be broken by the finite fundamental domain \citep{Levin:2002nr,Riazuelo:2004qv}. When searching for a mirror-parity symmetry plane on real data, however, care must be taken to avoid spurious effects that may induce even mirror-parity due to experimental systematics (some of which involve the ecliptic plane) or foregrounds (which dominate the Galactic plane). Also, as always, inherent biases and a-posteriori choices in the definition of the statistical estimators may lead to unfounded conclusions regarding anisotropies in the data. In this paper, we revisit recent claims of borderline significant detection of mirror-parity in the CMB (\citealt{Land:2005ys}; \citealt*{Ben-David:2012eu}; \citealt{Finelli:2012ij,Planck-Collaboration:2013oq}), in data from both the \textit{Wilkinson Microwave Anisotropy Probe} \citep[\textit{WMAP}; see e.g.][]{Bennett:2011lr} and \textit{Planck} \citep{Planck-Collaboration:2013kl} experiments. The different reports have been consistent so far in identifying two prominent directions in the sky, one which maximizes even mirror-parity and another corresponding to odd mirror-parity. The even mirror-parity direction found \citep{Land:2005ys,Ben-David:2012eu,Finelli:2012ij} coincides well with the direction of the CMB dipole \citep{Kogut:1993hl,Bennett:2003qd,Planck-Collaboration:2013rc}, and its reported significance was rather mild. Meanwhile, the direction in which odd mirror-parity is maximized \citep{Ben-David:2012eu,Finelli:2012ij,Planck-Collaboration:2013oq} has been assigned much higher levels of statistical significance. The various works, however, differ greatly in the methods used to analyse the data, and utilize different statistical estimators, Galactic masks and significance estimation methods. This work aims to shed a clear light on these findings. We use the recent \textit{Planck} data release to perform a robust search for mirror-parity using both a pixel-based statistic \citep{de-Oliveira-Costa:1996qy,de-Oliveira-Costa:2004fy} and a statistic in harmonic space \citep{Ben-David:2012eu, Finelli:2012ij}. We address several issues involved, from the statistical methods used in the analyses and their possibly inherent biases, to the choice of the CMB maps and foreground masks that are investigated. We also analyse the scale dependence of the results and compare them to results on random simulations which are deliberately manipulated so that they contain different types and amplitudes of mirror-parity symmetry. The structure of this paper is as follows. In Section~\ref{sec:DataAndSimulations}, we describe the data maps used for the analysis, the foreground masks we apply and the set of random simulations we use for evaluating the significance of the results. In Section~\ref{sec:Method}, we present the two statistics we use to search for mirror-parity in the CMB data as well as our methods for significance estimation. In Section~\ref{sec:TestsOnRandoms}, we describe three tests we perform on our statistics in order to compare their effectiveness in detecting mirror-parity and their sensitivity to the power spectrum amplitude and the shape of the Galactic masks. In Section~\ref{sec:Results}, we present the results of our analysis on the real data and estimate their statistical significance. We conclude in Section~\ref{sec:Conclusion}. | \label{sec:Conclusion} Over the past decade, the pursuit of large-scale anomalies in CMB data has generated numerous claims of deviations from the expected behaviour according to the concordance \LCDM\ model \citep[see][and references within]{Bennett:2011lr,Planck-Collaboration:2013oq}, spurring a tumultuous discussion of the statistical methods used in the analyses \citep{Bennett:2011lr} and especially the approaches towards significance estimation. Some of the disagreements are grounded in principles, such as the age-long Bayesian vs.\ frequentist debate, and are likely to persist. Certain claims might be considered as questions of taste. For instance, some advocate that the statistical significance of each new reported result be normalized according to the total number of tests conducted on CMB data hitherto. This may be viewed as an attempt to compensate for the `look elsewhere' effect, in the sense of looking for many anomalies on a given data set (as opposed to the effect in the sense of the estimation of the significance of a particular one), although it is not clear how this normalization is to be tracked and whether it should carry over to each new data set. Regardless of personal convictions, it is important that results of great potential importance such as those cited above undergo careful scrutiny and robust examination if they are to advance our theoretical understanding of cosmology in a meaningful way. In this work we have scrutinized previous claims of significant levels of mirror-parity in CMB maps from both the \textit{WMAP} and \textit{Planck} experiments. Indeed, using a pixel-based statistic and the large U73 Galactic mask, one can estimate the significance of the odd-mirror-parity direction to be as high as $3\sigma$ (0.13, 0.08 and 0.05 per cent for the SMICA, LGMCA and NILC maps, respectively). However, we have shown that these results are biased due to the large mask. With a smaller mask the significance level is lower, ranging from 0.25 to 1.04 per cent for the SMICA map. In addition, we have shown that these result are sensitive to assumptions regarding the total power on large scales, which is weakly constrained by the data. We have also tested the data for mirror-parity using a harmonic statistic. We have shown this statistic to be far more robust than its pixel-based counterpart -- stable against the use of a Galactic mask and with regard to the power spectrum amplitude. The statistical significance of the harmonic results reaches a level of $\sim2.5\sigma$ (0.51 per cent) at its highest. This level, however, is sensitive to the choice of component separation method applied to the CMB data. In addition, we have shown that the significance level is highly sensitive to scale. Furthermore, with the high quality of the \textit{Planck} maps, and especially the LGMCA map, we have also allowed ourselves to test the sky maps completely unmasked. This way, as long as the maps are clean enough of Galactic foregrounds, both statistics are expected to perform the best, free of anisotropy bias in the case of the pixel-based statistic (though still suffering from sensitivity to the total power) and of the \LCDM\ assumptions used for harmonic reconstruction in the case of the harmonic statistic. The results for the unmasked maps, however, are not statistically significant, with significance levels reaching only as high as 0.19 per cent for the pixel-based statistic and 0.51 per cent for the harmonic statistic. The pixel-based results are in general agreement with those of \citet{Planck-Collaboration:2013oq}. In light of these findings, we conclude that while there is some tendency for odd parity in the CMB data, which peaks at the scale of $\ell=7$, when embracing a broader perspective and examining the complete set of data maps and Galactic masks and the properties of the statistical estimators, it appears that the evidence for anomalous mirror-parity is rather weak. Our conclusion is that it poses no real challenge to the concordance model, and should therefore not be considered a \LCDM\ anomaly. | 14 | 3 | 1403.2104 |
1403 | 1403.2797_arXiv.txt | Massive objects orbiting a near-extreme Kerr black hole quickly plunge into the horizon after passing the innermost stable circular orbit. The plunge trajectory is shown to be related by a conformal map to a circular orbit. Conformal symmetry of the near-horizon region is then used to compute the gravitational radiation produced during the plunge phase. | General relativity implies that the high-redshift region very near the horizon of a near maximally-spinning Kerr black hole is governed by an infinite-dimensional conformal symmetry \cite{hep-th/9905099, 0809.4266}. X-rays \cite{McClintock:2006xd} and iron lines \cite{Brenneman:2006hw} from such regions have already been observed, and the future may hold yet higher precision observations. It is of interest to explore any potential observational consequences of the conformal symmetry. In a companion paper \cite{Porfyriadis:2014fja}, the conformal symmetry was exploited to compute gravity wave emission for an extreme-mass-ratio-inspiral within this near-horizon region. This approximates the signal from a stellar mass object orbiting near an extreme supermassive Kerr black hole and is potentially observable at eLISA \cite{Finn:2000sy, Gair:2004iv, elisa}. Once such an object passes the innermost stable circular orbit (ISCO), it plunges into the black hole. The plunge trajectory turns out to be related by a conformal map to the circular orbit. In this paper we use the conformal map to compute the gravitational radiation produced during this post-ISCO plunge into a near-extreme Kerr black hole. This complements the computation in \cite{Hadar:2009ip} of radiation produced during the post-ISCO plunge into a nonrotating Schwarzschild black hole. In section 2 we set notation and briefly review the geometry of Kerr, near-horizon extreme Kerr (NHEK) and near-horizon near-extreme Kerr (near-NHEK). Section 3 gives the conformal map from a circular orbit of a pointlike `star' in NHEK to a plunge trajectory in near-NHEK. In section 4, as a warmup to the gravity case, we couple a scalar field to the star. 4.1 computes the radiation production within near-NHEK using a bulk gravity analysis. 4.2 computes the same process using the techniques of two-dimensional conformal field theory (CFT). The results are shown to agree in subsection 4.3. In subsection 4.4 we turn to asymptotically flat near-extreme Kerr by reattaching the asymptotically flat region to near-NHEK. The resulting outgoing scalar radiation at flat future null infinity is computed. In 4.5 we derive the late-time quasinormal mode (QNM) decomposition. In section 5 all of these steps are repeated for the spin two case of gravitational radiation. | 14 | 3 | 1403.2797 |
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1403 | 1403.1041_arXiv.txt | We present \RF, a tool for finding galaxy-scale strong gravitational lenses in multi-band imaging data. By construction, the method is sensitive to configurations involving a massive foreground early-type galaxy and a faint, background, blue source. \RF~detects the presence of blue residuals embedded in an otherwise smooth red light distribution by difference imaging in two bands. The method is automated for efficient application to current and future surveys, having originally been designed for the 150-deg2 Canada France Hawaii Telescope Legacy Survey (CFHTLS). We describe each of the steps of \RF. We then carry out extensive simulations to assess completeness and purity. For sources with magnification $\mu>$4, \RF~reaches 42\% (resp. 25\%) completeness and 29\% (resp. 86\%) purity before (resp. after) visual inspection. The completeness of \RF~is substantially improved in the particular range of Einstein radii $0\farcs 8 \le \REin \le 2\farcs0$ and lensed images brighter than $g=22.5$, where it can be as high as $\sim$70\%. \RF~does not introduce any significant bias in the source or deflector population. We conclude by presenting the final catalog of \RF~CFHTLS galaxy-scale strong lens candidates. Additional information obtained with Hubble Space Telescope and Keck Adaptive Optics high resolution imaging, and with Keck and Very Large Telescope spectroscopy, is used to assess the validity of our classification, and measure the redshift of the foreground and the background objects. From an initial sample of 640,000 early type galaxies, \RF~returns 2500 candidates, which we further reduce by visual inspection to 330 candidates. We confirm 33 new gravitational lenses from the main sample of candidates, plus an additional 16 systems taken from earlier versions of \RF. First applications are presented in the SL2S galaxy-scale Lens Sample paper series. | \label{sec:intro} Since the discovery of the first multiple quasar produced by strong gravitational lensing by a foreground massive galaxy \citep{walsh79}, and the discovery of the first giant arcs found at the centers of galaxy clusters \citep{soucail87,lynds86}, much progress has been made in exploiting the unique capabilities of strong gravitational lensing as a probe of the mass content of distant massive objects, independent of the nature of their constituents or their dynamical state. With the advent of deep, wide-field optical imaging surveys we have now entered a new era that enables the use of sizable samples of strong lensing events as precision diagnostics of the physical properties of the distant Universe. Gravitational lensing, by itself and in combination with other probes, can be used to great effect to measure the mass profiles of early-type galaxies, both in the nearby universe and at cosmological distances \citep[e.g.\ ][]{T+K02a,T+K02b,RKK03,T+K04,R+K05,Koo++06,J+K07,Gav++07,Tre10,Aug++10,Lag++10,Son++12,Bol++12,Dye++13,ORF13}. Until recently, however, this approach was severely limited by the small size of the samples of known strong gravitational lenses. This has motivated a number of dedicated searches which have increased the sample of known strong gravitational lens systems by more than an order of magnitude in the past decade. Different search strategies have been adopted depending on the properties of the parent survey. A fundamental distinction is that between source-oriented and deflector-oriented searches \citep[e.g.][]{SEF92}. The choice depends on the relative abundance of the population of foreground deflectors and that of background sources, and has strong implications for the potential applications of the resulting lens catalog. Historically, source-oriented surveys were considered first, as they generally require the analysis of a relatively small input catalog of magnified sources that turn out to be much brighter than the foreground deflector at a carefully chosen wavelength. The sources of choice were typically bright quasars or radiosources, that would outshine the light from the deflector. This approach is well-illustrated in the optical by the Sloan Quasar Lens Survey \citep[SQLS][]{Ina++03,Ina++12}, the near IR with MUSCLES \citep{Jac++12}, and in the radio by the Cosmic Lens All Sky Survey (CLASS) \citep{Mye++03,Bro++03}. This last survey led to the early discovery of 22 gravitationally lensed quasars, many of which have been followed up with \hst. The latest release of the SQLS lens catalog has reported the discovery of 49 new quasar lenses. The scarcity of the population of background sources \citep[$\sim0.1$ per deg$^2$,][]{O+M10} requires extremely wide field surveys in order to gather sizable lens samples. Recently, wide field imaging surveys at millimeter wavelengths with the South Pole Telescope \citep[SPT,][]{Hez++13}, and sub-millimeter wavelengths with the {\it Herschel} satellite like H-ATLAS \citep{Neg++10,Gon++12,Bus++13} and HerMES \citep{Con++11,Gav++11,War++13}, have made it possible to target the population of distant sub-mm galaxies in the redshift range $1-4$ and find large number of lenses, a result predicted by \citet{Bla96}. These surveys lead to a spatial density of strong lenses ranging between 0.1 and 0.2 deg$^{-2}$ \citep{Neg++10,War++13,Vie++13}. The recent availability of high-resolution spectro-imaging with the Atacama Large Millimeter Array (ALMA) of gravitational lenses found at those wavelengths makes this technique a very promising avenue for the coming decade \citep{Hez++13}. We expect that a similar density of lensed quasars will be reached by optical surveys as well, including the Dark Energy Survey, the HSC Survey and in the next decade LSST and Euclid \citep{O+M10}. Indeed, at optical wavelengths, where unobscured star-forming galaxies are clearly visible, the ever-increasing deep wide field imaging and spectroscopic surveys are providing a population of distant background sources that has reached several hundreds of thousands per square degree. Since these are generally much fainter than the foreground massive early type galaxy deflectors, an effective strategy is to focus on these less numerous foreground galaxies and look for signatures of a gravitationally-lensed background object. As most of these background sources are spatially resolved, they take the typical shape of a complete, or partial, arc-like, Einstein ring. The challenge of this approach generally resides in the limited spatial resolution of wide field surveys, and the somewhat similar wavelengths at which both the lens and the source shine, which makes it difficult to disentangle the source and deflector light. A particularly successful way to mitigate this problem has been to take advantage of large spectroscopic surveys, such as the Sloan Digital Sky Survey (SDSS), which took spectra of several hundreds of thousands of bright galaxies. By looking for composite spectra consisting of two objects at different redshifts within the solid angle covered by the spectroscopic fiber, it has been possible to build a large sample of galaxy-galaxy lens systems. These consist typically of a low redshift foreground massive early-type galaxy and a background star-forming galaxy at higher redshift. The Sloan Lens ACS Survey (SLACS) conducted \hst follow-up observations of such spectroscopic candidates, and discovered about 100 gravitational lenses in the redshift range $0.1\le \zd\le 0.4$ \citep{Bol++04,Bol++06,Aug++10}. Rarer configurations, involving a foreground late-type galaxy, or a background early-type galaxy, were also searched for \citep{Tre++11,Bre++12,Aug++11}. Recently, this technique has been extended to the SDSS-III survey and has been used to find more lenses at $z\lesssim 0.6$ \citep{Bro++12,Bol++12}. The main advantage of this spectroscopic approach is that important quantities, such as the deflector and source redshifts along with the velocity dispersion of the stars in the deflector, are obtained from the parent survey itself. The availability of this data allows for many scientific applications including, combined lensing+dynamical studies of these systems \citep{T+K04,Koo++06,Koo++09,Aug++10,Son++12}, without the need for targeted spectroscopic follow-up. Even though the approaches described above have been very successful, there is strong motivation to develop techniques to identify galaxy-galaxy lenses purely in imaging data. This is challenging, but it can potentially yield a larger number of objects than any other technique: \citet{MBS05} forecast more than 10 such systems per square degree at \hst-like depth and resolution. Similar numbers are expected for Euclid/LSST. We therefore expect an all-sky survey should find more than $10^5$ such systems. Finding a similar sample of systems from an all-sky spectroscopic catalog would require of order $10^8$ spectra, two orders of magnitude more than have been taken to date. Because of the difficulty of identifying galaxy-galaxy lenses in optical images, much effort has been devoted to the analysis of \hst data in order to exploit its resolution. Searches based on both visual and automated inspection have been conducted \citep[e.g.][]{RGO99,Mou++07,Fau++08,Jac08,Mar++09,Paw++12,NMT09}, yielding several tens of candidates over the few square degrees of available data. Still, ground-based imaging is a potentially promising avenue, with its relatively low angular resolution compensated by the ready availability of hundreds or thousands of square degrees in multiple bands. In the SDSS, with typically $1\farcs5$ seeing and limiting magnitude $r\sim 21.5$, only wide separation systems produced by very massive galaxies, groups or clusters of galaxies have been able to be uncovered \citep{Bel++09}. The success rate increases significantly with angular resolution, and therefore sub-arcsecond image quality is desirable. Thus, good image quality and wide area coverage, such as that provided by the Canada France Hawaii Legacy Survey (CFHTLS), are potentially more promising sources of lenses -- especially those with deflector redshifts $\sim0.5$ and above, where current samples are scant \citep{Tre++10}. This argument motivated the Strong Lensing Legacy Survey (SL2S), which comprised a search for group and cluster scale (Einstein Radius $\REin\gtrsim 3\arcsec$) lenses \citep{Cab++07,Mor++12}, and the present work, the SL2S galaxy-scale lens search. The ultimate goal of our work was to use the newly found lenses to study the formation and evolution of massive galaxies; our results can be found in \citet{Ruf++11,Gav++12,PaperIII,PaperIV}. In this paper we present \RF, the semi-automated procedure for finding galaxy-scale strong lenses in the multi-band imaging data that we have applied to the CFHTLS. By focusing on the most frequent lens-source configuration, that of a foreground red massive early-type galaxy and a faint blue background source at higher redshift, we implement a technique that subtracts off the foreground light and analyses the blue residuals by requiring that they are broadly consistent with a strong lensing event. The CFHTLS data are described in \Sref{sect:data} while \Sref{sect:method} describes the \RF~algorithm. A list of \Nf~lens candidates, plus \NfX~additional candidates from preliminary versions of \RF~and additional datasets are also presented in this section. In Section~\ref{sect:simus} we describe the extensive and realistic simulations of plausible galaxy-scale lens systems which we use to assess the performance of the method, both in terms of completeness and purity. We explore the dependence of those quantities on important parameters like the Einstein radius, source magnitude, and lens and source redshifts in order to characterize the selection function of the \RF sample. Then, in Section~\ref{sect:followup} we summarize the results of our multi-year follow-up campaign with high resolution imagers and spectrographs, present the sample of confirmed lenses, and discuss false positives and contaminants. We summarize our main results and present our conclusions in Section~\ref{sect:concl}. Throughout this paper, all magnitudes referred to are calculated in the AB system, and we assume the concordance $\Lambda$CDM cosmological background with $\Omega_{\rm m}=0.3$, $\Omega_\Lambda=0.7$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. | \label{sect:concl} We have developed and implemented a novel method for the automated detection and classification of strong galaxy-scale gravitational lenses in ground-based imaging surveys. The method is based on difference imaging between blue and red filters, automated morphological cuts, and final visual inspection. The alogorithm has been extensively tested in fields containing known lenses and via simulated strong lensing configuration to characterize its selection function, purity and completeness. The algorithm has been applied to search for strong lens candidates in the $\sim \SArea$ deg$^2$ of CFHTLS imaging. The results of the search as well as of extensive follow-up for confirmation and rejection of candidates have been presented. Our main results can be summarized as follows. \begin{itemize} \item The \RF algorithm is highly efficient, retaining just \NRF lens candidates (\XRF) from a sample of \NETGs pre-selected early-type galaxies. This corresponds to about 18 candidates per square degree. \item Hierarchical visual inspection of these objects led to the definition of a sample of \Ncands candidates of sufficient quality $({\tt q\_flag}\ge2)$ for follow-up observation. These systems are listed in table~\ref{tab:cands}. This procedure took approximately 48 person-hours to complete, or $\lesssim30$ person-minutes per square degree. The follow-up campaigns allowed us to confirm 33 definite gravitational lenses in this (main) sample. Extrapolating to the full catalog, about 220 lenses are still to be confirmed in our sample. \item During the early stages of this work, while still tuning the parameters of the \RF~algorithm, we also gathered \NfX~lens candidates with {\tt q\_flag}$\ge2$ that did not remain in the statistical main sample of \Ncands~systems. 16 of those systems have been confirmed as lenses. The full list of additional candidates is given in table~\ref{tab:cands2}. \item These SL2S galaxy-scale lenses lie in the redshift range $0.3\lesssim z \lesssim 0.8$, and thus complement well the SLACS galaxy-scale lens sample at median redshift 0.2. \item High-resolution imaging and extensive spectroscopic follow-up candidates allowed \Nlens (+\Nlensn systems from the additional non statistical sample) new gravitational lenses to be confirmed. The analysis of the confirmed sample has been carried out in \citep{Gav++12}, \citep{Ruf++11}, \citep{PaperIII} and \citep{PaperIV} and yielded valuable constraints on the evolution of the mass profile of massive ETGs back to redshift $z\sim0.8$ thanks to a combined lensing+dynamics analysis. \item We tested the response of \RF to simulated images of ETGs exhibiting a range of detectable features along the line of sight to quantify the fraction of recovered true lenses as well as the contamination of our lens sample by non-lenses. Both statistics are in broad agreement (within a factor $<2$) given the limitations of the simulations and the imprecision of the visual classification at the end of the \RF procedure. A rough number of 1.5 definite lenses per sq. degree can be obtained with a method similar to \RF on data similar to the CFHTLS. \item Those simulations allowed us to check that the overall completeness we could achieve for such a survey after visual inspection is close to 40\%. However, restricting the search to a narrow range of Einstein radii, between 0.9 and 2 arcsec, as well as arcs (or lensed features) brighter than a magnitude of 22.5 in the $g$ band, can significanly increase the completeness to a value $\sim$70\%. Lenses in a different regime could be better seeked with other surveys (deeper and better resolution for small Einstein radii and faint sources or wider for large Einstein radii). \end{itemize} This paper represents a step forward in the detection of galaxy-scale gravitational lenses in wide field ground-based imaging surveys. Upcoming (such as DES, KiDS and HSC) and future imaging datasets (such as Euclid and LSST) will cover thousands of square degrees of sky and are expected to contain a few hundred thousand such lenses \citep{LSST,Euclid}. Our ability to detect them efficiently in large numbers will depend largely on the development of algorithms such as \RF. However, the current version of \RF would require too much human supervision ($\sim$5000 person hours scaling from CFHTLS) for an individual investigator. Therefore a successful lens search in such surveys will require either the development of additional automated stages of classification to increase the purity of lens candidates by at least an order of magnitude, or the implementation of crowd-sourcing to carry out the visual inspection. Both approaches are currently being pursued by the authors: one strategy for the former is to perform a fast lens modeling analysis as the final stage of selection \citep{Mar++09}, while the latter approach is currently being investigated by the Space Warps project (Marshall et al, in preparation) using the same CFHTLS dataset as described here. | 14 | 3 | 1403.1041 |
1403 | 1403.1331_arXiv.txt | Current r-process models under-produce $A<130$ nuclei. P-process models either underproduce $A<100$ and $150<A<165$ p-nuclei, or have other limitations. We argue that the puzzles in these two mechanisms can be solved by heavy-ion spallation. Spallation happens in the explosive phase transition from a neutron star to a quark star - the Quark Nova (QN). The QN triggers the r-process which creates abundant $A>130$ nuclei, and spallation fragments these isotopes into $A<130$ nuclei and all 35 p-nuclei. Our model is universal in relation to a star's age, metallicity, and chemistry. | Currently there are three processes that are thought to dominate the nucleosynthesis of isotopes heavier than $^{56}$Fe: the r-process, the s-process, and the p-process. Roughly, the r-process dominates the production of the vast majority of $A>90$ isotopes, while the s-process produces most $56<A<90$ isotopes. Yet, the p-process still plays an important role: the p-nuclides, which are 35 in total, cannot be reproduced by neutron-capture processes and require the p-process \citep{rauscher2002nucleosynthesis}. For the p-process, researchers postulate the photo-disintegration ($\gamma$-process) of seed nuclei as the dominant mechanism for the production of p-nuclei \citep{rauscher2002nucleosynthesis}. Unfortunately, the r-process is still poorly understood \citep{arnould2012r}, and p-process models face many challenges \citep{rauscher2013constraining}. Historically, the main puzzle of the r-process has been its astrophysical site. Traditionally, the prime candidate for the r-process site has been Type II Supernova (SNII)\citep{arnould2012r}. However current simulations can only detonate underpowered, mid mass-range supernovae \citep{RevModPhys.85.245}. Furthermore, current SNII models cannot generate enough neutrons nor entropy to trigger the r-process \citep{roberts2010integrated}. The second most popular candidate is the the Neutron Star Merger (NSM). Yet, NSMs happen at a relatively late stage of the universe's evolution, which disagrees with the r-process enrichment of old, metal-poor stars \citep{jaikumar2007nucleosynthesis}. Furthermore, NSM's mass yield for elements lighter than $A \sim 140$ is too low in comparison with solar abundances \citep{korobkin2012astrophysical}. The most glaring issue with current p-process models is the underproduction of isotopes in certain mass-ranges in comparison to solar abundances. The prime candidate for the p-process, the SNII, under-produces p-nuclei in the mass ranges $150<A <165$, and $A \le 100$ \citep{fuller1995neutrino,howard1991new,rayet1995p}. Researchers have suggested other candidates, such as Type Ia supernovae \citep{goriely2002he,howard1991new}, thermonuclear explosions on a neutron star's surface \citep{schatz2001end}, neutron-winds acting on hot-matter \citep{frohlich2006neutrino}, and accretion discs around compact objects \citep{fujimoto2003p}. However, no alternative to SNII has gathered a majority consensus. Furthermore, these models still lack consistent hydrodynamic and nucleosynthetic treatment \citep{rauscher2013constraining}. Another challenge concerns the production of the very rare elements $^{138}$La and $^{180m}$Ta . The $\gamma$-process, which is thought of as the origin of most p-nuclei, cannot produce $^{138}$La and $^{180m}$Ta. Researchers have suggested that these rare isotopes were produced by the interaction of neutrinos with stable nuclei ($\nu$-process) \citep{woosley1990nu,heger2005neutrino}. Nonetheless, this $\nu$-process for $^{180m}$Ta has some uncertainties, because it is difficult to calculate $^{180m}$Ta's final isomeric state \citep{mohr2007survival,belic2002photo}. Finally, the p-nuclei yield is sensible to details in a star's structure, including metallicity, and details in stellar evolution, which poses a problem when calculating the final, Galactic p-nuclei yields \citep{rauscher2013constraining}. Perhaps, the puzzles in both the p-process and r-process are closely related. We argue that the isotopic abundances that are attributed to the r-process and p-process are strongly affected by the presence of an alternate site - the Quark Nova (QN). The Quark-Nova is the name given to the explosive phase transition from a neutron star to a quark-star \citep{ouyed2002quark}. Although the QN model has not yet been used to explain p-nuclei, it has been applied to the aforementioned r-process puzzle. The hot, and neutron-dense matter ejected by the QN was suggested as an ideal site for the r-process\citep{ouyed2002quark}. Although, the QN was an excellent source of very heavy, $A>130$ isotopes, it strongly under-produced isotopes in the $A<130$ range, which effectively left the $90<A<130$ solar abundances still unexplained. However, in this paper we consider the QN a site for heavy ion spallation. We argue that heavy-ion spallation can successfully enhance isotopes in the $90<A<130$ mass range and thus, effectively complete the QN's r-process. Furthermore, this $A>130$ yield also includes all 35 p-nuclei. Spallation in the QN is triggered in the following way: If there is a short time delay between an SN II and the evolution of its neutron star into a QN, the relativistic, neutron-rich ejecta released by the QN will interact with the ejecta of the aforementioned SNII, triggering spallation reactions. This interaction between the QN's ejecta and the SNII's ejecta is referred as the dual-shock Quark-Nova (dsQN). The QN's ejecta is rich in $A>130$ r-process isotopes, and these isotopes spall against the SNII's ejecta. In spallation, the $A>130$ r-process isotopes fragment into smaller daughter isotopes, which eventually turn into p-nuclei and other r-process nuclei. The masses of these spallation products covers the whole $A<130$ range. Spallation in dsQN has been previously explored \citep{ouyed2011}. However, not until now do we consider the fragmentation of $A>130$ r-process isotopes produced in the QN, which produces a much wider mass spectrum. | Below we discuss some of the implications of our findings: \textit{Spallation as a source of $A<130$ r-process isotopes}: Elements in the $90 \le A < 130$ range require a spallation mechanism. This has some very important consequences. First, this result implies that perhaps, the difficulty of producing $90 < A < 130$ r-process isotopes in either the QN or other, more conventional models like the NSM, is incompleteness - namely, the lack of a spallation mechanism. Therefore, the $90 < A < 130$ solar abundances are actually produced by both spallation and r-process. Spallation therefore, completes other more traditional r-process models by filling the $90 < A < 130$ gap. \textit{Spallation as a source of p-nuclei}: $A>130$ r-process isotopes that were produced in the QN ejecta act as "seeds" that are fragmented by spallation into p-nuclei. The spallation reaction highly excites the interacting nuclei, which in turn, de-excite by evaporating mostly neutrons. Thus, spallation by design, tends to produce fragments in the proton-rich side of the valley of stability \citep{cugnon1997nucleon}. Previously, scientists have explored a spallogenic origin for p-nuclei \citep{audouze1970some,hainebach1976cosmic}. However, the spallation mechanisms previously proposed require high proton fluxes that cannot be reproduced \citep{woosley1978p}. In contrast, the heavy-ion spallation our model offers is very efficient. A typical QN would eject up to $\sim 10^{-3}$M$_\odot$ of r-process mass, which is almost all spalled in dsQNe with $t_{\text{delay}} \le 5$ days, making the QN's r-process yield an excellent source of p-nuclei seeds. This efficiency leads to a high enhancement of $A < 100$ isotope production, including $^{92,94}$Mo and $^{96,98}$Ru. dsQNe with $t_{\text{delay}}\le 5$ days lead to a high number of collisions that break the $A>130$ beam into smaller nuclei, which produce the $A \le 100$ enhancement. Our model might offer some clarification on the origin of the rare $^{138}$La and $^{180m}$Ta. $^{138}$La can be produced by our spallation model without the need of the $\nu$-process. Finally, It has been argued that high energy spallation can produce the isomer $^{180m}$Ta \citep{hainebach1976cosmic,PhysRevC.26.435,takacs2011activation}. \textit{Robustness in relation to $^{56}$Ni-synthesis}: The r-process yield is in the order of $\sim 10^{-5}-10^{-3}\text{ M}_{\odot}$, which is much smaller compared to the $\sim 0.1 \text{ M}_{\odot}$ mass of the $^{56}$Ni layer. These large difference in magnitudes implies that the r-process beam will destroy only a trace amount of $^{56}$Ni, leaving the $\sim 0.1 \text{ M}_{\odot}$ figure almost intact. This makes our model, compatible with current light curve measurements \citep{woosley1991co}, and stellar nucleosynthesis models \citep{woosley1995presupernova}. \textit{Robustness in relation to metallicity and stellar evolution}: Studies have shown that r-process abundances in stars are invariant across a star's metallicity, and therefore it's age, which contradicts current SNII models which are sensitive to a star's initial metal content \citep{sneden2003extremely}. Furthermore, NSMs do not coalesce early in the universe, and thus, cannot explain r-process abundances in very old stars \citep{jaikumar2007nucleosynthesis}. Moreover, p-nuclei are sensitive to metallicity and other details in a star's evolution, which makes calculating the final, Galactic p-nuclei yield complicated. The dsQN's spallation yield doesn't depend on metallicity nor a star's chemical make-up, which makes it robust against age, metallicity, and stellar mass. Furthermore, dsQN can appear very early in the Universe \citep{ouyed2009quark,ouyed-metal-poor-stars,ouyed2013resolution}. This robustness makes the dsQN an appropriate explanation for the invariant r-process abundances across metallicity. In addition, this robustness makes the dsQN extremely convenient as a source of p-nuclei, for the calculations of the final, Galactic p-nuclei abundances become simplified. \textit{Compatibility with r-process mass yield}: Past studies show that the QN rate of event and is mass yield is compatible with the Galactic r-process mass yield \citep{jaikumar2007nucleosynthesis}. It follows, then that p-nuclei mass should be roughly compatible as well. \textit{Candidates dsQNe}: The dual-shock QN model \citep{ouyed2009predictions} predicted a very specific super-luminous "double-hump" profile in the dsQN's light curve. Recently, these predictions became true when SN 2009ip and SN 2010mc showed that characteristic "double-hump" and were deemed as dsQNe \citep{ouyed2013sn}. We suggested that CasA was a remnant of dsQN and that it's peculiar chemical nature was a result of spallation \citep{ouyed2011}. Recently, the dsQN's predictions were hinted as a possible explanation for CasA's unexpected distribution of Fe and $^{44}$Ti, distributions that were surveyed by NuSTAR \citep{laming2014astrophysics}. CasA was postulated as a $t_{\text{delay}} \sim 5$ days dsQN, which is comparable to the ranges in this paper. This similarity might point to a degree of universality and/or consistency self-check. Furthermore, if Cas A is a $t_{\text{delay}} \sim 5$ days dsQN, it should show signatures associated with p-nuclei as well. A $t_{\text{delay}} \sim 5$ days dsQN should produce a few percent of its $A>56$ spallation yield as p-nuclei. This p-nuclei yield includes the rare $^{138}$La and $^{180m}$Ta isotopes as well. Furthermore, dsQN spallation should produce excited, high angular momentum isomers that would decay and release photons that might be detectable. \textit{Challenges and future work}: An interesting avenue involves extending our calculations beyond the $^{56}$Ni layer to include the overlaying C and O layers as done in \citep{ouyed2013resolution, ouyed2012}. Furthermore, a more realistic model of the spallation between the QN ejecta and the SN ejecta would account for the dampening of the beam caused by the gaseous SN ejecta, which would eventually stop spallation and halt the beam in the outer shells. Due to limitations on the spallation cross section routines, we couldn't include isotopes much heavier than $A\sim 208$, which requires a more sophisticated spallation algorithm. Furthermore, the spallation algorithm we used does not discriminate between isomers, so the production of $^{180m}$Ta isn't calculated quantitatively. This simulation is very useful for pointing at broad, qualitative patterns, but further research requires a more sophisticated code, with hydrodynamics and more detailed nuclear physics included. Another task needs to be done is to make a statistical comparison with solar abundances through a galactic chemical evolution process. \textit | 14 | 3 | 1403.1331 |
1403 | 1403.3334_arXiv.txt | Roughly half of the radiation from evolving galaxies in the early universe reaches us in the far-infrared and submillimeter wavelength range. Recent major advances in observing capabilities, in particular the launch of the \emph{Herschel Space Observatory} in 2009, have dramatically enhanced our ability to use this information in the context of multiwavelength studies of galaxy evolution. Near its peak, three quarters of the cosmic infrared background is now resolved into individually detected sources. The use of far-infrared diagnostics of dust-obscured star formation and of interstellar medium conditions has expanded from rare extreme high-redshift galaxies to more typical main sequence galaxies and hosts of active galactic nuclei, out to $z \gtrsim 2$. These studies shed light on the evolving role of steady equilibrium processes and of brief starbursts, at and since the peak of cosmic star formation and black hole accretion. This review presents a selection of recent far-infrared studies of galaxy evolution, with an emphasis on \herschel\ results. | \label{sect:roadtoherschel} The very first steps to use far-infrared emission as a tool to study galaxy evolution date back to the pioneering IRAS mission, when 60~$\mu$m source counts in the ecliptic pole region were found to exceed no-evolution models \citep{hacking87}. The \emph{Infrared Space Observatory} (ISO) obtained the first deep surveys at both mid- and far-infrared wavelengths, detecting strong evolution of the luminosity function out to $z\sim 1$ and supporting by plausible extrapolation that such dusty galaxies constitute the cosmic infrared background. ISO also pioneered the application of the rich mid-infrared spectra as a tool of studying energy sources and physical conditions in dusty galaxies \citep[see][for a review]{genzel00}. The \emph{Spitzer Space Telescope} revolutionized mid-infrared surveys and opened the window to direct mid-infrared spectroscopy of faint high-z galaxies \citep[review by][]{soifer08}. The \akari\ mission \citep{murakami07} provided uniquely detailed mid-infrared photometric coverage. Independently and at a similar time, ground-based (sub)millimeter surveys at wavelengths 850~\mum\ -- 1.2~mm detected luminous `SCUBA galaxies' \citep{blain02}, which are among the most actively star forming systems at redshifts $z \sim 2$. For most redshifts of interest, both spaceborne mid-infrared and groundbased submillimeter surveys miss the rest frame far-infrared peak that dominates the spectral energy distributions of most galaxies. Extrapolation via template spectral energy distributions (SEDs) is then needed to characterize the energy budget of samples that were selected at these wavelengths. Also, such samples can be significantly biased compared to a truly calorimetric selection by infrared luminosity. When aiming for the crucial deep far-infrared surveys, fully cryogenic space missions with small 60--80~cm diameter primary mirrors, such as \iso, \spitzer, and \akari\/, were strongly limited by source confusion. At 250--500~\mum\/, results from the 2~m BLAST balloon telescope \citep{pascale08} provided a first deep glimpse at the extragalactic sky, but are now superseded by \herschel\ data. The European Space Agency's \emph{Herschel Space Observatory} \citep{pilbratt10}, in operation 2009--2013, for the first time met the need for a combination of sensitivity and large aperture (3.5~m) over the full far-infrared and submillimeter wavelength range, substantially reducing confusion levels. At 350~\mum\ and longer wavelengths, low spatial resolution \emph{Planck} all-sky maps and catalogs \citep{planck13a} support studies of galaxy evolution. Much of this review is based on \herschel\ photometric surveys and pointed observations using the camera modes of the PACS \citep[70, 100, 160~\mum,][]{poglitsch10} and SPIRE \citep[250, 350, 500~\mum,][]{griffin10} instruments. We will refer to this range as the `far-infrared', leaving the `(sub)millimeter' terminology to ground based surveys at typically 850\,$\mu$m or longer wavelengths. Similarly, `Submillimeter galaxy (SMG)' will refer to the type of galaxy detected in these surveys, while equivalent \herschel\ selected sources will be called `dusty star forming galaxy (DSFG)' where appropriate. Traditionally, studies in the local universe make use of a terminology of `luminous infrared galaxies (LIRGs)' defined by their total 8--1000~\mum\ infrared (IR) luminosity $L_{\rm IR}>10^{11}$~\lsun, and their `ultra-' and `hyper-' luminous ULIRG and HYLIRG equivalents above 10$^{12}$ and 10$^{13}$~\lsun, respectively \citep[e.g.][]{sanders96}. These are handy acronyms, but for the purpose of galaxy evolution studies it is important to recall that connotations of these classifications that were carefully calibrated in the local universe may not apply at high redshift. For example, local ULIRGs are found to be major mergers with unusually dense and warm interstellar medium. The same cannot necessarily be assumed for their higher redshift equivalents at same infrared luminosity. Where used in this review, the (U)LIRG acronyms should be seen as pure infrared luminosity classifications, without further implications for the nature of the galaxy under study. \subsection{An inventory of \herschel\ surveys} \label{sect:inventory} \begin{figure} \center \includegraphics[width=13.cm]{dlutz_fig1.eps} \caption{Survey area and point source depth reached by some extragalactic surveys with \herschel\ PACS (top, shown for 100~\mum\ but 160~\mum\ data are availabe for the same fields) and SPIRE (bottom, shown for 250~\mum\ but 350 and 500~\mum\ data are available). Exposure is computed from total observing time (including overheads) and survey area. The point source depths (including confusion noise) shown on the right axis should hence be seen as indicative only, since no attempt was made to capture detailed effects of different observing layouts in the various projects. Some sets of cluster observations are included, at the total area summing all objects. Surveys shown are from the projects listed in Section~\ref{sect:inventory} as well as other projects for the Akari-NEP (PI S.~Serjeant), Akari deep field south (PI T.~Takagi), SPT deep field (PI J.~Carlstrom) and SDSS stripe 82 \citep{viero13b}.} \label{fig:inventory} \end{figure} The \herschel\ general purpose observatory served a wide range of science goals, using its capabilities that include both photometric imaging and spectroscopy over the 55--672~$\mu$m wavelength range. Guaranteed and open time for in total about 15\% of \herschel\/'s observing time was devoted to imaging surveys of galaxy evolution, not counting pointed imaging or spectroscopic studies of individual high redshift sources. The list below summarizes some of the major efforts, from the initial guaranteed time and open time \herschel\ key programmes as well as from similar scale programs that were allocated after later calls for proposals. In addition, a good fraction of the science discussed below is based on smaller scale imaging projects as well as pointed imaging and spectroscopic studies which are not listed explicitly. \herschel\ extragalactic surveys exceed 1000 square degrees in total but do not reach the full sky coverage that \planck\ provides at lower spatial resolution for $\lambda\geq 350$~$\mu$m and \akari\ at lower resolution and sensitivity for $\lambda\leq 160$~$\mu$m. Extragalactic confusion limits and instrument design imply a focus for PACS surveys on small area deep surveys, while SPIRE surveys typically emphasize large area. Below we list some major surveys, and Figure~\ref{fig:inventory} gives a graphical overview of area and depth of fields covered. \begin{itemize} \item HerMES \citep[][http://hermes.sussex.ac.uk/content/hermes-project]{oliver12} is a wedding cake type multipurpose survey of blank and cluster fields with focus on SPIRE, covering 70 square degrees in its shallowest tier. A shallow 270 square degree extension has been obtained in the HeLMS project (PI M.~Viero). \item PEP \citep[][http://www.mpe.mpg.de/ir/Research/PEP/index.php]{lutz11} is a PACS survey of popular multiwavelength fields such as GOODS, COSMOS, EGS, ECDFS, Lockman-XMM and cluster fields, coordinated with SPIRE observations of the same fields from HerMES. \item GOODS-Herschel \citep[][http://hedam.oamp.fr/GOODS-Herschel/]{elbaz11} provides deep PACS and SPIRE observations of GOODS-North and ultradeep PACS coverage of part of GOODS-South. A combined PEP/GOODS-Herschel dataset including all PACS data of the two GOODS fields is published in \citet{magnelli13}. \item H-ATLAS \citep[][http://www.h-atlas.org]{eales10a} is the largest area shallow extragalactic \herschel\ survey covering about 570 square degrees in SPIRE and PACS. \item The Herschel Lensing Survey (HLS)\\ \citep[][http://herschel.as.arizona.edu/hls/hls.html]{egami10} provides deep PACS and SPIRE data of 44 X-ray luminous clusters as well as SPIRE snapshots of another 527 clusters. HLS aims at both lensed background objects and at cluster members. \item H-CANDELs (PI M.~Dickinson) provides deep \herschel\ data of the CANDELs\\ (http://candels.ucolick.org/index.html) subregions of the COSMOS and UKIDSS-UDS fields, that were not yet covered at equivalent depth by the projects mentioned above. \end{itemize} Figure~\ref{fig:hudf_allwave} visualizes an example for the post-\herschel\ status of deep surveys over the full mid-infrared to submillimeter wavelength range. Surveys with beams of width 5\arcsec\ to 30\arcsec\ are now available over the full range. In the deepest fields, they reach the confusion limit for all wavelengths except 70\,$\mu$m, where a fully cryogenic 3\,m class telescope such as the SPICA project will be needed. \begin{figure} \center \includegraphics[width=13.cm]{dlutz_fig2.eps} \caption{Current status of deepest 24--870~$\mu$m infrared surveys, visualized by $4\arcmin\times 4\arcmin$ cutouts in the HUDF region. Data are from the GOODS project (24~$\mu$m), PEP and the combined PEP and GOODS-Herschel data (70--160~$\mu$m), HerMES (250--500~$\mu$m), and the groundbased LESS survey \citep[870~$\mu$m,][]{weiss09}.} \label{fig:hudf_allwave} \end{figure} | New observing capabilities from \herschel\ and other facilities have enhanced the power of far-infrared studies of galaxy evolution, for redshifts up to $z\sim 2$ and beyond. \begin{enumerate} \item Near its peak, three quarters of the cosmic infrared background is resolved into individually detected galaxies. \item Far-infrared calorimetric star formation rates are now available for massive normal star forming galaxies at high redshift, and help calibrate other indicators. Direct far-infrared luminosity functions and star formation rate densities have been obtained out to $z\gtrsim 3$. \item The interstellar medium conditions of high-z massive star forming galaxies, as expressed in the far-infrared SED and in the ratio of mid- and far-infrared emission, are better described in relation to the evolving main sequence of star forming galaxies rather than by absolute infrared luminosity. Most star formation happens near the main sequence. \item The far-infrared emission shows that AGN hosts out to $z\sim 2$ typically are normal massive star forming galaxies. The role of major mergers is less important. \item Steps have been made towards using dust emission as a tracer of the total interstellar medium mass of high-z galaxies. \end{enumerate} | 14 | 3 | 1403.3334 |
1403 | 1403.3933_arXiv.txt | It is usually assumed that outflows from luminous AGN are either in the energy-conserving (non-radiative) or in the momentum-conserving (radiative) regime. We show that in a non-spherical geometry the effects of both regimes may manifest at the same time, and that it is the momentum of the outflow that sets the $\mbh-\sigma$ relation. Considering an initially elliptical distribution of gas in the host galaxy, we show that a non-radiative outflow opens up a wide ``escape route'' over the paths of least resistance. Most of the outflow energy escapes in that direction. At the same time, in the directions of higher resistance, the ambient gas is affected mainly by the incident momentum from the outflow. Quenching SMBH growth requires quenching gas delivery along the paths of highest resistance, and therefore, it is the momentum of the outflow that limits the black hole growth. We present an analytical argument showing that such energy-conserving feedback bubbles driving leaky ambient shells will terminate SMBH growth once its mass reaches roughly the $M_\sigma$ mass derived earlier by \cite{King2003ApJ} for momentum-conserving AGN outflows. Our simulations also have potentially important implications for observations of AGN jet feedback and starburst galaxy feedback. The collimation of the wide angle AGN outflow away from the symmetry plane, as found in our simulations, indicates that credit for work done by such outflows may sometimes be mistakenly given to AGN jets or star formation feedback since wide angle $v \sim 0.1 c$ outflows are harder to observe and the phase when they are present may be short. | Over the past decade, there has been growing interest in and understanding of the role of feedback from active galactic nuclei (AGN) in the evolution of galaxies. Astronomers now generally agree that outflows caused by AGN feedback can drive gas out of galaxies and quench star formation in the host \citep{Silk1998A&A, Page2012Natur}, thus establishing the observed correlations between the supermassive black hole (SMBH) mass and host galaxy spheroid velocity dispersion \citep{Magorrian1998AJ, Ferrarese2000ApJ, Tremaine2002ApJ, Gultekin2009ApJ}, dynamical mass \citep{Haering2004ApJ, McConnell2011Natur} and other parameters \citep{Aller2007ApJ,Feoli2009ApJ}. A considerable effort was extended to explain these findings from a theoretical standpoint, but the problem remains far from being solved. Both semi-analytical \citep{Croton2006MNRAS, Bower2006MNRAS} and numerical simulations \citep[e.g.,][]{DiMatteo2005Natur} showed that AGN feedback must be a key ingredient in galaxy formation and evolution, especially at the high mass end. However, the precise mechanism of the feedback communication from the AGN to galaxy's gas remains elusive to this day. The situation is complicated since there is no agreement on which process -- wide angle gas outflow, jet, or radiation -- is the main mechanism of feedback, and also whether this delivers energy (heating), physical push (momentum), or both, to the ambient gas. For example, \cite{DiMatteo2005Natur, DiMatteo2008ApJ, Booth2009MNRAS} show that depositing $\sim 5 \%$ of the AGN luminosity into the ambient gas during the rapid Eddington-limited SMBH growth establishes the observed correlations. \cite{Sijacki2007MNRAS} adds to the picture jets in form of hot bubbles emitted by AGN at lower accretion rates, while jets in \cite{Dubois2012MNRASb} also transfer momentum to the ambient gas. \cite{Sazonov2005MNRAS} and \cite{Ciotti2007ApJ} propose that Compton radiative heating of ambient gas by AGN radiation field plays a significant role in limiting SMBH growth. \cite{Fabian1999MNRAS,Thompson2005ApJ,DebuhrEtal11} suggest that radiation pressure on dust is the main culprit of AGN feedback. Wide angle outflows from AGN that deliver momentum to the ambient gas are investigated by \cite{Debuhr2010MNRAS}. \cite{King2003ApJ, King2005ApJ} considers effects of a wide angle outflow on the ambient gas; both momentum and energy of the outflow are important. In this paper we focus on the effects of fast wide angle outflows from AGN on the host galaxy gas in the context of the \cite{King2003ApJ} model. Our main results are however more general, and add to a growing body of work showing that the efficiency of energy deposition into the ambient gas is actually quite low {\em if the gas is clumpy or inhomogeneously distributed}. \cite{WagnerEtal12} studied numerically (using a grid-based code) the interaction of a powerful jet with two-phase medium in the host galaxy, and found that the efficiency of energy transfer to the cold medium is only $\sim 10$\% (undoubtedly this particular number depends on the parameters of the cold phase and, perhaps, numerical resolution). \cite{WagnerEtal13} extended this work to the case of wide angle outflows, and found similar results. These authors found that ``the outflow floods through the intercloud channels, sweeps up the hot ISM, and ablates and disperses the dense clouds. The momentum of the UFO is primarily transferred to the dense clouds via the ram pressure in the channel flow, and the wind-blown bubble evolves in the energy-driven regime.'' Bourne, Nayakshin \& Hobbs (2013; submitted, BNH13 hereafter) used a completely different numerical technique -- smoothed particle hydrodynamics \citep[SPH; see, e.g.,][]{Springel2010ARA&A}, employing the ``SPHS'' algorithm of \cite{HobbsEtal13a} that is designed to reduce artifical numerical effects of the classical SPH in a clumpy medium. BNH13 obtained results very similar to that of \cite{WagnerEtal13}, and proposed that this inefficiency of AGN feedback energy deposition into the ambient medium explains how SMBH can grow to the substantial masses observed despite producing huge amounts of energy in the fast outflows that could destroy bulges of host galaxies multiple times over. \cite{Nayakshin13b} included these effects into an analytical study of AGN feedback and showed that the observed $M-\sigma$ relation can be reproduced by such energy-conserving flows {\em if star formation in clumpy medium is also taken into account.} In this paper we perform numerical simulations with a ``classical'' SPH code that has a different implementation of AGN feedback compared to either \cite{WagnerEtal13} or BNH13, and different initial conditions for the ambient gas in the galaxy. We consider initially homogeneous, rather than clumpy, medium but distribute it in a non-spherical geometry. As an example, we consider AGN feedback on an elliptically distributed ambient gas in a galaxy, so that gas density along the galactic plane is highest and drops gradually to the lowest value perpendicular to the plane. This geometry should be considered as a simplest rudimentary step closer to realistic galaxies, which are mostly non-spherical except perhaps in the case of ``red-and-dead'' elliptical galaxies. Despite these numerical and set-up differences with previous studies, we recover the main conclusions of \cite{WagnerEtal13} and BNH13. We find that the AGN feedback quickly inflates two outflow bubbles perpendicularly to the galactic plane, where the gas density is lowest. Most of the feedback energy escapes through these funnels, leaving the denser gas exposed mainly to the momentum of the AGN wind. Therefore, the dense gas behaves as if it were affected by momentum feedback only. The energy-momentum separation found in our simulations is large-scale rather than local, small scale, as in \cite{WagnerEtal13} and BNH13, but the final conclusions are similar. We also provide a simple analytical argument \citep[related to an earlier study of ``leaky feedback bubbles'' in the content of stellar feedback by][]{HMurray09} that confirms the main result of \cite{Nayakshin13b}: since the cold gas is momentum-driven, the SMBH mass required to expel the cold shell in this {\em energy-conserving} regime is similar to the momentum-driven result of \cite{King2003ApJ}, which in itself is pleasingly close to the observed $M-\sigma$ relation \citep{Ferrarese2000ApJ, Tremaine2002ApJ, Gultekin2009ApJ}. The paper is structured as follows. We begin with a brief review of the state-of-the-art of our understanding of the physics of ultra-fast outflows in spherically symmetric models in \S \ref{sec:radiation} - \ref{sec:ufos}. We then present in \S \ref{sec:toy} a toy model spherically symmetric ``leaky shell'' calculation that takes into account energy escape from the bubble via low density channels in the ambient shell. In Section \ref{sec:nummodel} we describe the setup of numerical simulations and in Section \ref{sec:results} we present their results. We follow with a discussion in Section \ref{sec:discuss} and summarize and conclude in Section \ref{sec:concl}. | \label{sec:concl} Our numerical experiments showed that energy-conserving spherically-symmetric outflows from SMBHs may create highly aspherical bubbles if the ambient gas in the host galaxy is not spherically distributed. This may lead to a whole host of theoretical and observational implications for SMBH-host galaxy connections. Firstly, SMBH driving energy-conserving outflows may self-regulate their growth to the momentum-conserving $M_\sigma$ value found by \cite{King2003ApJ}. This implies that SMBH outflows may actually not lose much energy to radiation, as assumed in the momentum-conserving picture, and be therefore pumping {\em all} of their energy into the surrounding ambient gas. As we argued here, however, most of this energy leaks out from the bulge via low density channels and is therefore deposited outside of the bulge, in the halo of the host galaxy or even beyond. This leads to dense gas only being exposed to the ram pressure of the outflow, and thus the critical SMBH mass required to push it away and halt further accretion is very similar to $M_\sigma$. Secondly, the observational appearance of SMBH feedback may be deceiving. This point is very important: the interpretation of observed feedback processes in the host galaxies may be incorrect in some cases. For example, wide angle outflow studied here forms biconical structures which may look very much like galaxy-disc outflows driven by starbursts in the discs of the host galaxies. Another potential mis-interpretation of observations could be the buoyant bubbles usually presumed to be inflated by AGN jet activity. The bubbles we obtained here are similarly energetic and are also filled with very high temperature gas where electrons could be potentially accelerated into non-thermal distributions in shocks \citep[as suggested for the Fermi Bubbles by][]{Zubovas2012MNRASa}. If a jet, however weak, was also present in addition to the wide angle outflows studied here, then it could be a simple matter to attribute the bubble's mechanical power to what is easier to discern in the observations -- the jet. | 14 | 3 | 1403.3933 |
1403 | 1403.1107_arXiv.txt | The {\it Herschel} SPIRE instrument consists of an imaging photometric camera and an imaging Fourier Transform Spectrometer (FTS), both operating over a frequency range of $\sim$450--1550\,GHz. In this paper, we briefly review the FTS design, operation, and data reduction, and describe in detail the approach taken to relative calibration (removal of instrument signatures) and absolute calibration against standard astronomical sources. The calibration scheme assumes a spatially extended source and uses the {\it Herschel} telescope as primary calibrator. Conversion from extended to point-source calibration is carried out using observations of the planet Uranus. The model of the telescope emission is shown to be accurate to within 6\% and repeatable to better than 0.06\% and, by comparison with models of Mars and Neptune, the Uranus model is shown to be accurate to within 3\%. Multiple observations of a number of point-like sources show that the repeatability of the calibration is better than 1\%, if the effects of the satellite absolute pointing error (APE) are corrected. The satellite APE leads to a decrement in the derived flux, which can be up to $\sim$10\% (1\,$\sigma$) at the high-frequency end of the SPIRE range in the first part of the mission, and $\sim$4\% after {\it Herschel} operational day 1011. The lower frequency range of the SPIRE band is unaffected by this pointing error due to the larger beam size. Overall, for well-pointed, point-like sources, the absolute flux calibration is better than 6\%, and for extended sources where mapping is required it is better than 7\%. | The Spectral and Photometric REceiver \citep[SPIRE][]{griffin2010} is one of three focal plane instruments which operated on board the ESA {\it Herschel} Space Observatory \citep[{\it Herschel};][]{pilbratt10} between May 2009 and April 2013. It contains an imaging photometric camera and an imaging Fourier Transform Spectrometer (FTS), with both sub-instruments using arrays of bolometric detectors operating at $\sim$300~mK \citep{turner2001} and feedhorn focal-plane optics giving sparse spatial sampling over an extended field of view \citep{dohlen2000}. This paper details the calibration scheme adopted for the SPIRE FTS, updating the early description by \citet{swinyard2010}. The FTS uses two bolometer arrays of 19 and 37 detectors to provide spectral imaging over a nominal $\sim$2 arcminute field of view. The design of the SPIRE FTS \citep[][]{ade1999, dohlen2000, swinyard2010} is shown in Fig.~\ref{fts_design}: the incoming radiation from the telescope is divided into two beams by a beamsplitter (BS1). These beams are retro-reflected from back-to-back roof top mirrors (RT) mounted on a linear translation stage (the Spectrometer MEChanism; SMEC). The SMEC imparts an optical path difference (OPD) between the two beams dependent on the mirror position, and a second beam splitter (BS2) recombines the light to form an interference pattern that is focused onto the detector arrays. There are no significant spectral features within the band of the beam splitters \citep{ade1999}. The response of the detector system to the source intensity is measured as a function of the SMEC position and is hereinafter referred to as the ``interferogram''. The interferogram is the Fourier transform of the incident spectrum, as modified by the instrumental response and other instrumental effects. The use of two beam splitters in a Mach-Zehnder configuration \citep{mach,zehnder}, plus the back-to-back roof top arrangement, means that the imparted optical path difference is four times larger than the physical movement of the mechanism, making for a compact optical arrangement. The Mach-Zehnder configuration also provides natural spatial separation of the two input and two output ports always present in an FTS. In SPIRE the second input port is terminated on a cold radiative source (SCAL) which can be heated to provide a known radiation load onto the detectors \citep{hargrave2006}. The two output ports are chromatically separated to allow the instrument to cover a broad frequency range whilst maintaining close to optimal optical coupling to the detectors. There are two optimised arrays, called SSW (Spectrometer Short Wavelength, covering 959.3--1544~GHz) and SLW (Spectrometer Long Wavelength, covering 446.7--989.4~GHz)\footnote{The band limits may be expanded slightly in future versions of the pipeline.}. Fig.~\ref{interferogram} shows a typical interferogram for a source with strong spectral lines, and the equivalent spectrum observed using the high resolution mode of the instrument. The spectral resolution, defined as the distance from the peak to the first zero crossing of the instrumental line shape, is constant in frequency at $\sim$1.184\,GHz, equivalent to 230--800\,kms$^{-1}$ \citep{observersmanual}. Note that in general (except for very bright sources), the fringing shown in the top plot of Fig.~\ref{interferogram} is successfully removed by the calibration scheme as it is stable in both science and calibration observations. \begin{figure} \centering \includegraphics[width=\hsize]{OpticalLayout.png} \caption{The optical layout of the SPIRE FTS.} \label{fts_design} \end{figure} The SCAL source was designed to increase the dynamic range of the detectors by nulling the modulated signal component of the interferogram. However, the total emission from the telescope and stray light were actually lower than the values used in the initial design of the SPIRE instrument, and the SCAL source was not needed and was therefore not actively heated. \begin{figure} \centering \includegraphics[width=\hsize]{interferogram_diagram4.pdf} \includegraphics[width=\hsize]{example_spectrum2.pdf} \caption{Top: Typical measured interferogram from the SPIRE FTS for an astronomical source with strong $^{12}$CO lines. The main features of the interferogram are highlighted. Bottom: Final spectrum for this interferogram showing the SSW and SLW bands and with transitions of $^{12}$CO labelled. The insert shows a zoom around the $^{12}$CO $J$=10--9 line.} \label{interferogram} \end{figure} A schematic view of the SPIRE FTS detector arrays is shown in Fig.~\ref{fts_array_footprint}, showing the relative positions of each detector as measured at the beginning of the mission. The detectors are arranged in a hexagonally close-packed pattern with the spacing between beam centres set to $\sim$33$^{\prime\prime}$ for SSW and $\sim$51$^{\prime\prime}$ for SLW, roughly equal to two beam widths. Vignetting and distortion within the optical design of the instrument increases away from the centre of each array, effectively limiting the nominal (unvignetted) field of view to $\sim$2$^{\prime}$. The nominal field of view is shown in Fig.~\ref{fts_array_footprint} as a dashed red line. The circles shaded in blue represent SSW and SLW detectors centred on the same sky positions and the gaps in the SSW array show the location of two dead detectors (SSWD5, SSWF4). The plot on the right in Fig.~\ref{fts_array_footprint} indicates the approximate full width at half maximum (FWHM) of the beam for each detector, and the overlap of the arrays on the sky, with 19$^{\prime\prime}$ circles for SSW and 35$^{\prime\prime}$ circles for SLW. \begin{figure} \centering \includegraphics[width=\hsize]{detector_locations_calPaper.png} \caption{A schematic view of the SPIRE FTS detector arrays showing the measured position of each detector. The right hand plot shows the two arrays as they appear on sky, with 19$^{\prime\prime}$ circles for SSW and 35$^{\prime\prime}$ for SLW. See main text for more details.} \label{fts_array_footprint} \end{figure} In this paper we describe the photometric, spectroscopic and spatial calibration of the SPIRE FTS in its nominal mode \citep[additional calibration needed for its bright-source mode is described by][]{lu}. In Section~\ref{sect_flux_cal} we describe the photometric calibration starting from the engineering data output from the detector electronics through to the derivation of calibrated spectra of astronomical objects; in Section~\ref{sect_fcfactors}, we describe the derivation of the flux conversion factors; in Section~\ref{sect_freq} we deal with the spectroscopic calibration and in Section~\ref{sect_beam} with the spatial response calibration. In Section~\ref{sect_accuracy} we summarise the accuracy and repeatability of the calibration, discuss caveats on the SPIRE data and consider aspects in which we expect to see improvements as our knowledge of the data improves. The calibration described has been implemented in the {\it Herschel} Interactive Processing Environment \citep[HIPE;][]{ott} Version 11 and a companion paper, Fulton et al. (in preparation), will detail how the procedures described here are put into practice in the SPIRE FTS data pipeline. Note that all errors in this paper are quoted as 1-sigma (1$\sigma$) limits. | The overall absolute calibration uncertainty for the FTS nominal mode is summarised in this section. We break the uncertainties into three categories: point sources observed in the sparse mode, extended sources observed in the sparse mode, and extended sources observed in mapping mode. The calibration described in this paper has been implemented in the pipeline corresponding to HIPE Version 11. To summarise for point sources observed on the centre detectors (SSWD4 and SLWC3), the measured repeatability is 6\%, with the following contributions: (i) absolute systematic uncertainty in the models from comparison of Uranus and Neptune - determined to be $\pm$3\%; (ii) the statistical repeatability determined from observations of Uranus and Neptune, with pointing corrected - estimated at $\pm$1\% (excluding the edges of the bands); (iii) the uncertainties in the instrument and telescope model, which lead to an additive continuum offset error of 0.4\,Jy for SLW and 0.3\,Jy for SSW and (iv) the effect of the {\it Herschel} APE. Note that the pointing uncertainty results in a reduction in flux and is, therefore, a one-sided statistical uncertainty on the calibrated spectrum. A large pointing offset also results in a significant distortion of the SSW spectrum of a point source and a mismatch between the SLW and SSW spectra \citep[e.g. see][]{valtchanov}. Providing one is convinced that the source in question has no spatial extent, the SLW portion of the calibrated spectrum can be used to correct any apparent gain difference between the SLW and SSW spectra. For sparse observations of significantly extended sources, the absolute uncertainty in intensity for a reasonably bright, fully extended object, observed in the central detectors is 7\%, with the following contributions: (i) the uncertainty in comparing the calibration on Uranus (a point source) to the telescope is estimated at 3\%; (ii) the uncertainty on the Uranus model itself of 3\%; (iii) the systematic reproducibility of the telescope model of 0.06\%; (iv) the statistical repeatability estimated at $\pm$1\% and (v) an additive continuum offset of $3.4\times10^{-20}$\,\wmhzsr~for SLW and $1.1\times10^{-19}$\,\wmhzsr~for SSW. In practice, truly extended sources tend to be faint and the uncertainty is therefore dominated by the additive offsets. When the source extent is larger than the main beam size, but not fully extended, or if there is structure inside the beam, then the uncertainties are dominated by the source-beam coupling \citep[see][]{wu} and are significantly greater. In mapping mode, the variations between detectors becomes important and the overall repeatability has been measured as $\pm$7\% \citep[see][for a full discussion of mapping mode observations]{benielli}. The off-axis detectors are less well calibrated, especially outside the unvignetted part of the field. The level of absolute flux accuracy and repeatability obtained with the SPIRE FTS compares favourably with the SPIRE photometer \citep{bendo}. The excellent level of calibration accuracy achieved is due to the linear transform properties of the FTS, which guarantee simultaneous calibration of the entire spectrum. Thus spectral features covering a wide range in frequency can be analysed together. This is not possible with monochromating devices where only a narrow frequency range is observed, as that can lead to calibration uncertainties caused by spectral features that are broader than, or outside of, the observed band (for instance standing waves or broad spectral features in the instrument response function). The penalty is in instantaneous sensitivity, due to the increased photon noise in an FTS. However, this noise is compensated for by the higher level of calibration fidelity that can be achieved, as well as the more widely appreciated advantage in observing speed for multi-line spectral observations. | 14 | 3 | 1403.1107 |
1403 | 1403.3428_arXiv.txt | {Young massive clusters are key to map the Milky Way's structure, and near-IR large area sky surveys have contributed strongly to the discovery of new obscured massive stellar clusters.} {We present the third article in a series of papers focused on young and massive clusters discovered in the VVV survey. This article is dedicated to the physical characterization of VVV\,CL086, using part of its OB-stellar population.} {We physically characterized the cluster using $JHK_S$ near-infrared photometry from ESO public survey VVV images, using the VVV-SkZ pipeline, and near-infrared $K$-band spectroscopy, following the methodology presented in the first article of the series.} {Individual distances for two observed stars indicate that the cluster is located at the far edge of the Galactic bar. These stars, which are probable cluster members from the statistically field-star decontaminated CMD, have spectral types between O9 and B0\,V. According to our analysis, this young cluster ($1.0$ Myr $<$ age $< 5.0$ Myr) is located at a distance of $11^{+5}_{-6}$ kpc, and we estimate a lower limit for the cluster total mass of $(2.8^{+1.6}_{-1.4})\cdot10^3 {M}_{\odot}$. It is likely that the cluster contains even earlier and more massive stars.} {} | Young massive clusters (cluster mass $M>10^3 M_{\odot}$, \citealt{hanson07}) are fundamental pieces for the study of Galactic structure. Because of their youth, they give information related to the recent Galactic massive stellar formation history. They are also excellent tracers of star formation regions. In the Milky Way, we find regions with intense stellar formation activity in the Galactic centre, where the Arches \citep{nagata93,figer02}, Center \citep{krabbe95,paumard06,figer08}, and the Quintuplet \citep{glass90,okuda90,nagata90,figer99} clusters are located; in the Carina-Sagittarius arm, which hosts Westerlund 2 \citep{westerlund61,rauw07,ascenso07}, Trumpler 14 \citep{ascenso07b,sanchawala07,sana10}, and NGC 3603 \citep{goss69,harayama08}; and the close edge of the Galactic bar. In the last region, several clusters with a massive population of red supergiant (RSG) stars are found: RSGC1 \citep{figer06,davies08}, RSGC2 \citep{davies07}, RSGC3 \citep{alexander09,clark09a}, Alicante\,7 \citep{negueruela10}, Alicante\,8 \citep{negueruela11}, and Alicante\,10 \citep{gonzalezfernandez12}, but also younger clusters with a mixed population of OB-type and RSG stars (Masgomas-1, \citealt{ramirezalegria12}). By symmetry, we expect massive star clusters at the far edge of the Galactic bar. Until now, only one massive cluster is known in this region: Mercer 81 \citep{davies12}. Owing to the distance and extinction expected for stars located at the far edge of the bar, clusters in that region must be observed with near-infrared filters, which are less affected by interstellar extinction, and with a high spatial resolution, to be able to resolve the distant cluster population. One of the most recent near-IR surveys, the ESO public survey VISTA Variables in the V\'ia L\'actea (VVV, \citealt{minnitiVVV10,saito10,saito12}), is a perfect tool for this exploration, covering the Galactic bulge and the adjacent disk region, including the far edge of the bar with a spatial resolution of 0.34 arcsec pix$^{-1}$ in \textit{ZYJH$K_S$} filters. We present infrared spectrophotometric observations for VVV\,CL086 \citep{borissova11}, a massive cluster found in the direction of the Perseus arm ($l=340\fdg001$, $b=-0\fdg293$), similar to that of Mercer 81 ($l=338\fdg400$, $b=+0\fdg100$). Data observation and reduction are described in Section \ref{observations}, near-IR photometry (CMD) and spectroscopy (spectral classification) are presented in Section \ref{resultados}. Cluster physical characterization is the topic of Section \ref{discusion}, and the general conclusions are given in the final Section \ref{conclusiones}. \begin{figure} \centering \includegraphics[width=7.1cm,angle=0]{vvvcl086_01.ps} \caption{VVV $K_S$ image for VVV\,CL086. The image is 2$\times$2 arcmin$^{2}$ size and the cluster radius of 35$\arcsec$ \citep{borissova11} is shown with a black circle. Small blue (OB-type stars) and black circles mark the position of the spectroscopically observed stars.} \label{vvvCL086_Ks} \end{figure} | We presented the physical characterization of VVV\,CL086, a new massive cluster discovered using data from the VVV survey, found at the far edge of the Milky Way bar at a distance of $11^{+5}_{-6}$ kpc. This cluster is the second one found in that region of the Galaxy (the first is Mercer 81), a region highly reddened by gas and dust, which presents a relatively low mean reddening of $A_K=1.5^{+0.0(3)}_{-0.1}$ mag, however. Our spectroscopic follow-up aimed at the brightest stars in the cluster area revealed that two objects are part of the disk population (two early-B dwarfs), and two stars form part of the cluster main-sequence population. From their spectral classification and the cluster CMD we were able to deduce that earlier stars than these two observed OB-stars are probably be present in VVV\,CL086. One star from our spectroscopic follow-up was not classified. The mass estimate was derived by integrating the Kroupa IMF fitted to our data and gives a lower limit for the cluster total mass of $(2.8^{+1.6}_{-1.4})\cdot10^3 {M}_{\odot}$. We also estimated the cluster age by fitting isochrones to the pre-main sequence turn-on point. We estimated a cluster age $>1.0$ and $<5.0$ Myr. The upper age limit agrees with the earliest main sequence star found in the cluster (i.e. O9\,V star \#04). Future spectroscopic observations are planned to confirm this and to investigate the cluster massive population in more detail. | 14 | 3 | 1403.3428 |
1403 | 1403.7316_arXiv.txt | Using parsec-resolution simulations of a typical galaxy merger, we study the triggering of starbursts by connecting the (inter-)galactic dynamics to the structure of the interstellar medium. The gravitational encounter between two galaxies enhances tidal compression over large volumes, which increases and modifies the turbulence, in particular its compressive mode with respect to the solenoidal one. This generates an excess of dense gas leading to intense star formation activity. Along the interaction, the compressive turbulence modifies the efficiency of gas-to-star conversion which, in the Schmidt-Kennicutt diagram, drives the galaxies from the sequence of discs to that of starbursts. | Starbursts are generally attributed to galaxy interactions and mergers \citep[e.g.][]{Sanders1996}. Mergers induce gravitational torques and global gas inflows toward the galaxy centres, enhancing the gas surface density \citep{Keel1985}. Simulations have demonstrated that this process triggers bursts of star formation, especially in the inner regions of advanced mergers \citep[e.g.][]{Barnes1991, Hopkins2006, Robertson2006, DiMatteo2007, Cox2008, Karl2010}. Yet, this process alone does not explain all properties of merger-induced star formation, which can be intense even in early interaction phases \citep{Ellison2008}, and is often spatially-extended, off-nuclear \citep{Wang2004, Barnes2004, Cullen2006, Elmegreen2006a, Smith2008, Hancock2009, Chien2010}. In fact, starburst galaxies can convert their gas into stars an order of magnitude faster than isolated discs with similar global gas surface densities, i.e. independently of the global compression by inflows \citep{Daddi2010b, Genzel2010, Saintonge2012}. Simulations have shown that, on top of the global enhancement of the gas surface density by the interaction-induced inflows, changes in the sub-structure of the interstellar medium (ISM) can control the starburst activity of mergers \citep{Teyssier2010, Powell2013}. Interactions increase the ISM turbulence, in agreement with observations \citep{Irwin1994,Elmegreen1995b}\new{, and could help compress the diffuse gas reservoirs \citep{Jog1992}}. Nevertheless, it remains unknown (i) whether numerical models can now reach convergence on the global starburst activity by sufficiently resolving the small-scale physics and structure of the ISM, (ii) through which physical processes galaxy mergers could increase or modify turbulence, and (iii) why triggered turbulence would lead to starburst activity rather than stabilizing clouds against collapse. Using parsec-scale simulations of a representative merger, we here address these three aspects by probing the role of the tidal field in modifying the properties of the ISM turbulence down to the scales of star-forming regions. | \label{sec:summary} Using hydrodynamical simulations, we propose a physical explanation for the enhancement of star formation activity in galaxy mergers. Our main findings are as follows: \begin{itemize} \item The global SFR evolution reaches numerical convergence at parsec-scale resolution. \item The rise of the gas mass fraction in compressive tides in extended volumes during the galactic collisions pumps turbulence into the ISM and unbalances the equipartition between compressive and solenoidal turbulence modes. \new{Such turbulence is not primarily driven by feedback.} \item The compressive turbulence allows to overcome the regulating, stabilizing effect of turbulence, and to generate an excess of dense gas. \item This excess translates into an enhanced star formation activity and drives the merger to the starburst regime in the Schmidt-Kennicutt diagram. \item \new{From the gravitational and tidal trigger to the ignition of starburst, the full sequence takes $\sim 10\mh 30 \Myr$.} \end{itemize} Here, we have only accounted for \emph{integrated} properties of the tides, turbulence and star formation. A study of the spatial distribution in the galaxies as well as the propagation of these phenomena will be described in a forthcoming contribution \citepip{\Bournaud}, showing in particular that the locations in space and time of the compressive tides, compressive turbulence and star forming regions coincide. We have drawn our conclusions using a simulation of the Antennae galaxies. Obviously, in other interacting systems, the quantitative results we presented here are modulated by the parameters of the galaxies (shape of the halo, mass ratio, etc.) and by the details of their interaction (spin-orbit coupling, impact parameter, orbital eccentricity, etc.). However, the ubiquity of compressive tides in mergers has been previously demonstrated by \citet{Renaud2009}, which suggests that the triggers and physical processes mentioned above exist in many mergers. An increase of the turbulent Mach number maintaining equipartition between compressive and solenoidal modes would widen the PDF without changing its functional form (from a log-normal), and would drive the evolution of the progenitor galaxies at constant efficiency in the Schmidt-Kennicutt diagram. This explains the similar star formation efficiencies between local spirals and redshift-two discs which have turbulent speeds about ten times larger \citep[see also \citealt{Kraljic2014}]{Renaud2012}. A second component at high density in the PDF is necessary to reach the regime of local starbursts (i.e. with a SFR comparable to that of redshift-two discs but with a less turbulent ISM). This is generated by the deviation from equipartition which occurs during a merger. \new{\citet{Krumholz2012} neglected the second component of the PDF in their conversion of volume to surface densities. By doing so, the starbursting mergers lie on the same regime as disc galaxies in their model, which leads to an apparent ``universality'' of star formation.} The physical context of mergers naturally modifies the turbulence forcing in ways that have so far been only arbitrarily implemented in idealised volume-limited simulations of the ISM \citep{Federrath2008, Seifried2011, Saury2014}. We have shown here that the (10)kpc-scale forcing induces variations in the structure of the ISM at parsec-scale, but it remains unclear whether the turbulence would retain a signature of these variations at the much smaller scale of dissipation. If it does, this might result in a variation of the proto-stellar core mass function, as proposed by \citet{Hopkins2012b}, and of the stellar initial mass function in mergers, with respect to isolated galaxies. | 14 | 3 | 1403.7316 |
1403 | 1403.2705_arXiv.txt | The orbital boundary value problem, also known as Lambert Problem, is revisited. Building upon Lancaster and Blanchard approach, new relations are revealed and a new variable representing all problem classes, under L-similarity, is used to express the time of flight equation. In the new variable, the time of flight curves have two oblique asymptotes and they mostly appear to be conveniently approximated by piecewise continuous lines. We use and invert such a simple approximation to provide an efficient initial guess to an Householder iterative method that is then able to converge, for the single revolution case, in only two iterations. The resulting algorithm is compared, for single and multiple revolutions, to Gooding's procedure revealing to be numerically as accurate, while having a significantly smaller computational complexity. | \label{intro} Lambert's problem, sometimes referred to as orbital boundary value problem, is a fascinating problem in astrodynamics that intrigued, over the years, most famous mathematicians. Just like Kepler's equation, its solution is at the very heart of fundamental astrodynamical and space engineering questions \cite{celmech1,celmech2,izzolambert}. Following the fundamental work laid down, among others, by Euler, Lambert, Lagrange and Gauss, the need of having one robust algoritmic procedure able to function for a wide set of conditions led to revisit the Lambert's problem during the space era. Among the many contributions made during that period, the work of Lancaster and Blanchard \cite{lancaster} is to be highlighted as it reduced the solution to Lambert's problem to performing iterations each one requiring the computation of one only inverse trigonometric or hyperbolic function. Later, Gooding \cite{gooding} built upon these results and published a procedure achieving high precision in only three iterations for all geometries. Gooding's algorithm makes use of Halley's iterations sided to well designed heuristics to set the initial guess of the iterated variable. His methodology to reconstruct the terminal velocity vectors is also remarkable as it is purely algebraic. The resulting procedure is extremely efficient having low computational cost and high accuracy. A number of studies \cite{peterson}, \cite{klumpp} and \cite{parrish} have tested Gooding approach extensively, suggesting its superiority with respect to other Lambert solvers. His procedure is most accurate and considered as the fastest existing approach to solve Lambert's problem \cite{arora}. Aside from Gooding's algorithm, many other proposal have been put forward to design Lambert solvers, they all differ in the details of at least one of three fundamental ingredients: a) the iteration variable (directly connected to the time of flight equation), b) the iteration algorithm c) the initial guess and d) the reconstruction of the terminal velocity vectors. More recently iimprovements on the original Gooding algorithm were also claimed \cite{arora} making use of the universal variable formulation \cite{bate} and an original cosine transformation. At the same time, a number of works recently addressed the possibility of deploying a large number of Lambert's algorithms on modern GPU architectures \cite{parrish}, \cite{arora2} and \cite{wie}. Interestingly, in the first of these works, a comparison is also made between Gooding procedure, a universal variable Lambert's solver and an early (slow) version of the algorithm here described (unpublished at that time) showing already its promising nature. In this paper, we build upon Lancaster and Blanchard work, first deriving some new results, and then proposing and testing a new algorithm. The new algorithm a) iterates on the Lancaster-Blanchard variable $x$ using b) a Householder iteration scheme c) feeded by a simple initial guess found exploiting new analytical results found. The resulting procedure is simple to implement, does not make use of heuristics for the initial guess generation and is able to converge, on average, in only 2 iterations for the single revolution case and 3 in the multirevolution case, introducing a significant reduction in the overall solver complexity. | We revisit Lambert's problem building upon the results of Lancaster and Blanchard and finding some new properties of the time of flight curves. We propose a new transformation of such curves able to further simplify the problem suggesting efficient approximations to the final solution. Using our results to design a new procedure to solve the Lambert problem we are able to build a low complexity algorithm that we find able to provide accurate solutions in a shorter time when compared to the state of the art Gooding's algorithm. | 14 | 3 | 1403.2705 |
1403 | 1403.2533_arXiv.txt | {} % { This is the second in a series of papers that attempt to unveil the kinematic structure of the Galactic bulge through studying radial velocities and proper motions. We report here $\sim15000$ new proper motions for three low foreground-extinction off-axis fields of the Galactic bulge. } { Proper motions were derived from a combination of \emph{Hubble Space Telescope} Wide Field Planetary Camera 2 (WFPC2) and Advanced Camera for Surveys (ACS) images taken 8 and 9 years apart, and they reach accuracies better than $0.9 \ mas \ yr^{-1}$ for more than $\sim 10000$ objects with magnitudes $F814W\leq24$. } {The proper motion distributions in these fields are similar to those of Galactic minor axis bulge fields. We observe the rotation of main sequence stars below the turn-off within the Galactic bulge, as in the minor axis fields. } { Our stellar proper motions measurements show a significant bulge rotation for fields as far from the galactic plane as $b\simeq-8^{\circ}$. } | It is still uncertain exactly what mechanisms led to the formation of the present-day Galactic bulge. It is not clear whether the evolution of the bulge was driven by mergers, as the hierarchical galaxy formation scenarios suggest, or secularly by disk instabilities. A clear observational picture of the bulge's current structure is needed to begin to understand its formation and evolution. Because of the high and variable extinction towards the Galactic center, this has been a difficult task and is perhaps the main obstacle to formulating a unified picture. In addition to the high foreground extinction by dust towards the Galactic bulge (which is not constant even on small scales), the bulge and the disk are projected on top of each other on the sky. Disentangling them is not straightforward: even in the color-magnitude diagram (CMD), the populations overlap (Holtzman et al 1998). Moreover, blue stragglers extend brighter than the turn-off and overlap with the main sequence region hosting the young population. This complicates the separation of populations based on photometry alone, and additional measurements are required to accurately study the different components of the bulge. Despite these difficulties, progress has been made in understanding the Galactic bulge, especially over the past years. Abundance studies, such as those by Rich (1988), McWilliam \& Rich (1994), Zoccali et al. (2008), and Bensby et al. (2011), have shown that the Galactic bulge has a wide range of metallicities. However, bulge metallicities do typically differ from disk and halo populations, showing a comparatively metal-rich population. The $\alpha$-elements in the bulge have also been consistently found to be overabundant with respect to halo and disk (Zoccali et al. 2006; Fulbright 2007; Hill et al. 2011). In particular, $\alpha$-elements are related to the formation timescale of the Galactic bulge since they are primarily produced during the explosion of SN Type II (due to short-lived massive stars). Iron production, on the other hand, is favored by SNe Type Ia explosion, where SN Type Ia typically have a timescale of an order of magnitude longer than SN Type II. Therefore, the overabundance of $\alpha$-elements in the bulge suggests a rapid formation scenario. Evidently, most bulge stars were formed before the interstellar medium (ISM) could be enriched by SN Type Ia explosions, hence the short inferred formation timescale for the bulge ($<$ 1 Gyr) (e.g. Ballero et al. 2007 and references therein). Simulations in the last few years have also complicated our view of the formation scenarios of the Galactic bulge. Shen et al. (2010) reproduced the stellar kinematics of the Bulge Radial Velocity Assay (BRAVA; Rich et al. 2007) without a classical bulge, where the boxy bulge previously reported in the literature is the end-on projection of the bar structure. Conversely, Saha et al. (2012) finds that a non-rotating small classical bulge can evolve secularly through angular-momentum exchange with the Galactic bar. At the same time, number counts along the CMD between populous clusters at several latitudes have been used to estimate the foreground disk contamination (\cite{feltzing}) and the age of the bulge. The foreground disk population mimics the young bulge population, especially near the turn-off, affecting the age determinations in the bulge. The results for two bulge fields that effectively identified the contamination by foreground populations have placed the age of the bulge population as old as $\sim10\ Gyr$. \begin{table*} \begin{center} \caption{Summary of observations\label{tab:summob}}. \small \begin{tabular}{ l l l l l l l} \hline \hline Field & Epoch & Exp.(s) & Instrument & Filter & Type & $\alpha,\delta \ (J2000)$\\ \hline Field 4-7 & 1995 July 14 & 1200($\times$2), 1300($\times$2) & WFPC2 & F555W, F814W & Undithered & 18 22 16, -29 19 22 \\ & 2004 July 11 & 50($\times$2), 347($\times$4) & ACS WFC & F814W & Dithered & \\ Field 3-8 & 1996 May 1 & 2900($\times$3) & WFPC2 & F555W, F814W & Undithered & 18 24 09, -30 16 12 \\ & 2004 July 12 & 50($\times$2), 348($\times$4) & ACS WFC & F814W & Dithered & \\ Field 10-8 & 1995 Nov 30 & 1200($\times$4) & WFPC2 & F555W, F814W & Undithered & 18 36 35, -23 57 01 \\ & 2004 July 14 & 50($\times$2), 347($\times$4) & ACS WFC & F814W & Dithered & \\ \hline \end{tabular} \end{center} \end{table*} \normalsize Despite its proximity, the Galactic bulge has classically suffered from a lack of proper-motion studies. Spaenhauer et al. (1992) were the first to obtain reliable proper motions from a Galactic bulge sample using photometric plates taken more than three decades apart. This study was the subject of a subsequent analysis by Zhao et al. (1994), who included radial velocities for a small subsample (64 stars). They found a significant vertex deviation (i.e., a triaxility signature) for the metal-rich population in their small proper motion-radial velocity combined sample. The same signature of triaxility was observed by Soto, Rich \& Kuijken (2007), who combined Spaenhauer et al. (1992) proper motions with \cite{sadler} (1996) and \cite{terndrup} (1995) radial velocities and metallicities to obtain a sample of $\sim 300$ K giants with 3-D kinematics and metallicities. More recently, the same triaxility signature related to two distinct populations, metal-rich and metal-poor, has been confirmed using high-resolution spectra by Babusiaux et al. (2010) and Hill et al. (2011). Microlensing surveys of the Milky Way bulge have also contributed to % improving our knowledge of the kinematics of the Galactic bulge. The OGLE-II experiment has produced $\sim \ 5\times10^6$ proper motions (\cite{sumi}) for 49 bulge fields, covering a range of $-11^{\circ} < l < 11^{\circ}$ and $-6^{\circ} < l < 3^{\circ}$, and reaching accuracies of $0.8-3.5\ mas\ yr^{-1}$. Using this survey as a basis, Rattenbury et al. (2007a; 2007b) studied the proper motions for a subsample of bulge red clump giant stars in 44 fields. Red clump stars were used as tracers of bulge density in order to fit a mass density distribution for the bulge. Along the same lines, Vieira et al. (2007) delivered proper motions for 21,000 stars in Plaut's window ($l,b= 0^{\circ}, -8^{\circ}$) from plates spanning 21 years. Space-based observations have also played a role in bulge proper-motion observations, boasting the combination of sharper images and reduced blending. Anderson \& King (2000; 2003) developed a technique of deriving precision astrometry. Their innovations included an effective point spread function (PSF) approach that obviates the need to integrate the PSF over pixels when evaluating it for a given image, and an empirical distortion correction for WFPC2, later on extended to the ACS Wide-Field Channel (WFC) and ACS High-Resolution Channel (HRC). These procedures were subsequently applied to measuring the component of rotation of 47 Tucanae globular cluster (Anderson \& King 2003). Similarly, Kuijken \& Rich (2002; henceforth KR02) used a modification of Anderson \& King (2000) approach to be the first to use HST observations to obtain reliable bulge proper motions. KR02's sample targeted two low foreground-extinction fields, Baade's window $(l,b=1.1^{\circ},-3.8^{\circ})$ and Sagittarius I $(l,b=1.3^{\circ},-2.7^{\circ})$. \begin{figure} \centering \includegraphics[width=8cm]{fig1.jpg} \caption{The \emph{rms} residuals of the pixels positions in $x$ (bottom) and $y$ (top) from the proper motion fits in one of our fields, Field 4-7 WF2. The solid lines are the best theoretical fit for an error estimation based on photon noise alone (\emph{bottom line}) and for photon noise including additional systematic source of errors for each epoch (\emph{two top lines}). This plot shows that the dominant source of error for the fainter undithered first epoch (\emph{red}) is photon noise, while for brighter sources a residual systematic error appears that we relate to saturation due to the long exposures. Second-epoch dithered exposures (\emph{green}) are best represented by a systematic error that has been included in the estimation of the proper motion error. \label{fig:syserror}} \end{figure} Both fields were used to successfully obtain 15,862 and 20,234 proper motions respectively, proving the feasibility of space-based proper motions with considerably shorter time baselines than those previously employed on ground-based bulge proper motions (e.g., Spaenhauer et al. 1992; $\sim30\ yr$). The samples in both fields were separated into bulge and disk components based on the mean proper motions. From this it was found that the bulge component clearly resembles an old population, such as those observed in globular clusters, and shows a significant rotation with no covariance in \emph{{l,b}}. More recently, both fields in KR02 have been the subject of new proper motion studies. Koz\l owski et al. (2006) obtained proper motions for 35 small fields in the vicinity of Baade's window. Their calculations, involving the combination of ACS/HRC images and archival observations with WFPC2, yielded 15,863 stellar proper motions. The proper motions calculated by Koz\l owski et al. (2006) show consistency with the velocity dispersions found by KR02, and in addition show a significant negative covariance term in the transverse velocity $C_{lb}=\sigma_{lb}/(\sigma_l \sigma_b)\simeq -0.10$. A negative covariance such as this may imply a tilt in the velocity ellipsoid with respect to the Galactic plane. Similarly, new results for the \emph{Sagittarius I} field (\cite{clarkson}) yielded more than $180,000$ proper motions with ACS epochs separated by just 2 years. From their initial numbers they finally selected 15,323 bulge stars using a similar procedure to the one applied by KR02. The covariance they report was very similar to the one found by Koz\l owski et al. (2006). In addition, Clarkson et al. (2008) produced velocity ellipsoids in (\emph{l,b}) as a function of distance bins. These velocity ellipsoids demonstrate a slight dependence on the distance of the objects analyzed. Similar to the finding in KR02, the median stellar sequence in Clarkson et al. (2008) for this bulge sample was best represented by an 11 Gyr old isochrone. We report here new proper motions results for stars in three low foreground-extinction windows of the Galactic bulge located in the first Galactic quadrant. These new fields differ from other HST bulge proper motion observations, such as the minor-axis fields presented in Clarkson et al. (2008) and KR02, since they sample the near end of the bar in the Galactic bulge. In addition, the derived proper motions have been obtained from cross-detector observations; suitable first-epoch WFPC2 exposures were obtained from the HST archive, and complemented with more recent ACS WFC exposures. These second-epoch observations are part of a larger program (GO-9816) that includes observations in both ends of the Galactic bar, as well as minor axis fields (see \cite{soto12} for more details). The resulting time baselines for the proper motions in this work are eight to nine years. This paper is organized as follows. In \S 2 we describe our observations. A detailed account of the procedures involved in the measurements of the proper motion can be found in \S 3. Section 4 presents the analysis performed on our sample of proper motions and the implications of the results. In addition, we compare our results with those of similar surveys in the Galactic bulge. Finally, our conclusions are summarized in \S 5. \begin{figure*} \centering \includegraphics[width=6.5cm, angle=0]{fig2a.jpg}\\ \includegraphics[width=4.2cm, angle=0]{fig2b.jpg} \includegraphics[width=4.2cm, angle=0]{fig2c.jpg} \includegraphics[width=4.2cm, angle=0]{fig2d.jpg} \includegraphics[width=4.2cm, angle=0]{fig2e.jpg}\\ \includegraphics[width=4.2cm, angle=0]{fig2f.jpg} \includegraphics[width=4.2cm, angle=0]{fig2g.jpg} \includegraphics[width=4.2cm, angle=0]{fig2h.jpg} \includegraphics[width=4.2cm, angle=0]{fig2i.jpg} \caption{Kinematic selection of cluster stars in Field 10-8 in one of the cameras of WFPC2, Wide Field 4 (WF4). The selection was performed for each chip separately using a hard cut-off; cluster stars have a small velocity dispersion, which therefore biases the relative zeropoint of the proper motions of bulge stars. Thus, the cluster stars must be excluded from the reference frame used to calculate bulge proper motions. \emph{Top row:} proper motions in $pixels \ yr^{-1}$ before the refinement of the stars used as the basis for the inter-epoch transformations for WF4 chip in Field 10-8; where black dots are bulge stars, and orange dots are those kinematically classified as cluster stars. \emph{Second row:} proper motions in pixels $\mu_x$ and $\mu_y$ as a function of the $x-y$ coordinates before the correction. \emph{Bottom row:} same as second row, but after applying the correction. \label{fig:pm2coor_wf4}} \end{figure*} | We have presented \textbf{$\sim 15000$} new proper motions for three off-axis fields of the Galactic bulge. The results for these three fields show remarkable agreement with the results in KR02, and thus suggests a bulge structure where the kinematics observed close to the center along the Galactic minor axis are repeated to some extent in higher longitudes. Despite the reduced number of proper motions in comparison with minor axis fields, the rotation of the bulge is still visible in our fields, which reach $l\sim 10^{\circ}$. We explored the possible changes in the velocity (proper motion) ellipsoid within Field 4-7 as a function of the distance along the line of sight; as is the case with the results of Clarkson et al. (2008), we found a change in the tilt of the \{l,b\} velocity ellipsoid. All of this suggests that a significant fraction of the population follows bulge-like orbits, even at the location of our three off-axis fields in the outer bulge. % If we consider the anisotropies produced by the bar in the minor-axis (inner bulge) fields, what we observe should therefore be part of the Galactic bar. The importance of the extent of the bar can be related to the evolutionary stage of the bulge, where slow-rotating long bars can evolve from rapid-rotating short bars secularly (e.g. Combes 2007; Athanassoula 2005), thus this information can provide important constraints on the bulge structure and formation. Dynamical models including this new proper motion data will be able to provide new insight into the actual bulge structure, until now poorly constrained. Finally, we have demonstrated the technical feasibility of proper-motion measurements using different cameras with different geometries for the first and second epochs. Field 10-8 with its globular cluster NGC 6656 has provided us with a direct assessment of the accuracy of our proper motion procedure. A cluster dispersion of $\sim 0.9 \ mas\ yr^{-1}$ or $\sim \ 14 \ km/sec$ at 3.2 $kpc$ supports our claim of $\sim30\ km\ s^{-1}$ accuracy for bulge stars. Unfortunately, our results are severely affected by the saturation of the long first-epoch exposures, and could be significantly improved by a third ACS epoch. | 14 | 3 | 1403.2533 |
1403 | 1403.2019_arXiv.txt | We present a novel method to solve the spatially homogeneous and isotropic relativistic Boltzmann equation. We employ a basis set of orthogonal polynomials dynamically adapted to allow for emergence of chemical non-equilibrium. Two time dependent parameters characterize the set of orthogonal polynomials, the effective temperature $T(t)$ and phase space occupation factor $\Upsilon(t)$. In this first paper we address (effectively) massless fermions and derive dynamical equations for $T(t)$ and $\Upsilon(t)$ such that the zeroth order term of the basis alone captures the particle number density and energy density of each particle distribution. We validate our method and illustrate the reduced computational cost and the ability to easily represent final state chemical non-equilibrium by studying a model problem that is motivated by the physics of the neutrino freeze-out process in the early Universe, where the essential physical characteristics include reheating from another disappearing particle component ($e^\pm$-annihilation). | \subsection{Different Equilibria} At sufficiently high temperatures, such as existed in the early universe and reproduced in modern day relativistic heavy ion collider studies of the hot, ultra-dense state of matter called quark gluon plasma (QGP), both particle creation and annihilation (i.e. chemical) processes and momentum exchanging (i.e. kinetic) scattering processes can occur sufficiently rapidly to establish complete thermal equilibrium. The most probable canonical distribution function $f_{ch}^\pm$ of fermions (+) and bosons (-) in both chemical and kinetic equilibrium is found by maximizing microcanonical entropy subject to energy being conserved \begin{equation}\label{ch_eq} f_{ch}^\pm=\frac{1}{\exp(E/T)\pm 1}, \hspace{2mm} T>T_{ch} \end{equation} where $E$ is the particle energy, $T$ the temperature, and $T_{ch}$ the chemical freeze-out temperature. For a physical system comprising {\em interacting} particles whose temperature is decreasing with time, there will be a period where the temperature is greater than the kinetic freeze-out temperature, $T_k$, but below chemical freeze-out. During this period, momentum exchanging processes continue to maintain an equilibrium distribution of energy among the available particles, which we call kinetic equilibrium, but particle number changing processes no longer occur rapidly enough to keep the equilibrium particle number yield. For $T<T_{ch}$ the particle number changing processes have `frozen-out'. In this condition the momentum distribution, which is in kinetic equilibrium but chemical non-equilibrium, is obtained by maximizing microcanonical entropy subject to particle number and energy constraints and thus two parameters appear \begin{equation}\label{k_eq} f_{k}^\pm=\frac{1}{\Upsilon^{-1} \exp(E/T)\pm 1},\hspace{2mm} T_k<T\leq T_{ch}. \end{equation} The need to preserve the total particle number within the distribution introduces an additional parameter $\Upsilon$ called fugacity. A fugacity different from $1$ implies an over-abundance ($\Upsilon>1$) or under-abundance ($\Upsilon<1$) of particles compared to chemical equilibrium and in either of these situations one speaks of chemical non-equilibrium. We emphasize that in our context, the appearance of a dynamical fugacity $\Upsilon\ne 1$ is not related to conserved quantum numbers. Conserved quantum numbers such as baryon number and charge introduce, through chemical potentials, fugacities that enhance e.g particles and deplete antiparticles. In the presence of conservation laws one has for particles $\Upsilon_p=\gamma e^{\mu_p/T}$ and antiparticles $\Upsilon_a=\gamma e^{-\mu_a/T}$ where the quantity $\gamma$ is the dynamical fugacity. For fermions (+) the distribution \req{k_eq} always satisfies the Fermi-Dirac constraint $f_{k}^+\le 1$, however the distribution of particles can approach much more closely the quantum degeneracy condition $f_{k}^+\to 1$ for $\Upsilon \gg 1$. Once the temperature drops below the kinetic freeze-out temperature $T_k$ we reach the free streaming period where particle scattering processes have completely frozen out and the resultant distribution is obtained by solving the collisionless Boltzmann equation with initial condition as given by the chemical non-equilibrium distribution \req{k_eq}. As already indicated, the two transitions between these three regimes constitute the freeze-out process -- first we have at $T_{ch}$ the chemical freeze-out and at lower $T_k$ the kinetic freeze-out. \subsection{Boltzmann Evolution and Chemical Non-Equilibrium} Exact chemical and kinetic equilibrium and sharp freeze-out transitions at $T_{ch}$ and $T_k$ are only approximations. The Boltzmann equation is a more precise model of the dynamics of the freeze-out process and furthermore, given the collision dynamics it is capable of capturing in a {\em quantitative manner} the non-thermal distortions from equilibrium, for example the emergence of actual distributions and the approximate values of $T_{ch}$, $T_k$, and $\Upsilon$. Indeed, in such a dynamical description no hypothesis about the presence of kinetic or chemical (non) equilibrium needs to be made, as the distribution similar to \req{k_eq} with $\Upsilon\ne 1$ emerges naturally as the outcome of collision processes, even when the particle system approaches the freeze-out temperature domain in chemical equilibrium. Considering this physical situation it is striking that the literature on Boltzmann solvers does not reflect on the accommodation of emergent chemical non-equilibrium into the method of solution. For an all-numerical solver this may not be a necessary step as long as there are no constraints that preclude development of a general non-equilibrium solution. However, when strong chemical non-equilibrium is present either in the intermediate time period or/and at the end of the evolution a brute force approach can be very costly in computer time. Motivated by this circumstance and past work with physical environments in which chemical non-equilibrium arose, we introduce here a spectral method for solving the Boltzmann equation that utilizes a dynamical basis of orthogonal polynomials which is adapted to the case of emerging chemical non-equilibrium. We validate our method via a model problem that captures the essential physical characteristics of interest and use it to highlight the type of situation where this new method exhibits its advantages. In the cosmological neutrino freeze-out context, the general relativistic Boltzmann equation has been used to study neutrino freeze-out in the early universe and has been successfully solved using both discretization in momentum space \cite{Madsen,Dolgov_Hansen,Gnedin,Mangano2005} and a spectral method based on a fixed basis of orthogonal polynomials \cite{Esposito2000,Mangano2002}. In Refs.\cite{Wilkening,Wilkening2} the non-relativistic Boltzmann equation was solved via a spectral method similar in one important mathematical idea to the approach we present here. For near equilibrium solutions, the spectral methods have the advantage of requiring a relatively small number of modes to obtain an accurate solution, as opposed to momentum space discretization which in general leads to a large highly coupled nonlinear system of odes irrespective of the near equilibrium nature of the system. The efficacy of the spectral method used in \cite{Esposito2000,Mangano2002} can largely be attributed to the fact that, under the conditions considered there, the true solution is very close to a chemical equilibrium distribution \req{ch_eq} where the temperature is controlled by the dilution of the system. However, the recent PLANCK CMB results \cite{Planck} indicate the possibility that neutrinos participated in reheating to a greater degree than previously believed. As we discussed recently~\cite{Birrell} this can also lead to a more pronounced chemical non-equilibrium. Efficiently obtaining this emergent chemical non-equilibrium within realm of kinetic theory motivates the development of a new numerical method that adapts to this new circumstance. With this novel Boltzmann solver we present here it is also possible to return to the exploration of the emergent chemical non-equilibrium in processes governing laboratory QGP physics. Indeed, the study of chemical non-equilibrium as a separate process from kinetic equilibrium has its roots in the field of laboratory QGP formation and observation, of which one of the proposed observables is the newly produced strange quark flavor~\cite{Muller:2011tu}. The chemical non-equilibrium analysis method~\cite{Rafelski:1991} is today the only successful statistical hadronization model for experimental results~\cite{Petran:2013dva} confirming chemical non-equilibrium for all strongly interacting particles produced by a QGP fireball~\cite{Petran:2013lja}. This paper establishes our new method for the case of ultra-relativistic particles does not address the pertinent physical applications in neutrino cosmology or quark-gluon plasma physics. The former is treated in our paper \cite{Birrell_nu_param} and the later will be a subject of future work. While in this work we address the case where the mass-scale of particles is entirely irrelevant, we have also developed a similar method for the case that the particle mass is non-negligible. Introduction of a mass scale presents no major conceptual modifications, but requires detailed technical modifications from the simpler scheme we present here that don't lend themselves well to a simultaneous presentation of both methods. The method including mass will be introduced and addressed in the context of specific applications of the method in future publications. In section \ref{boltzmann_basics} we give a basic overview of the relativistic Boltzmann equation in the format aiming to address the early Universe neutrino freeze-out process, but which can be easily recast into the format appropriate for other applications. In section \ref{the_method} we discuss our modified spectral method in detail. In subsection \ref{free_stream_approach} we recall the orthogonal polynomial basis used in \cite{Esposito2000,Mangano2002} and in subsection \ref{kinetic_eq_approach} we introduce our modified basis and characterize precisely the differences in the method we propose. We compare these two bases in subsection \ref{basis_comparison}. In subsection \ref{dynamics_sec} we use the Boltzmann equation to derive the dynamics of the mode coefficients and identify physically motivated evolution equations for the effective temperature and fugacity. In section \ref{validation} we validate the method using a model problem. In \ref{orthopoly_app} we give further details on the construction of the parametrized family orthogonal polynomials we use to solve the Boltzmann equation. | We have presented a spectral method for solving the Relativistic Boltzmann equation for a system of massless fermions diluting in time based on a dynamical basis of orthogonal polynomials. The method is adapted to systems evolving near kinetic equilibrium, but allows for potentially strong chemical non-equilibrium in a transient and/or final state as well as strong reheating i.e. decoupling of temperature scaling from dilution scaling. The method depends on two time dependent parameters, the effective temperature $T(t)$ and phase space occupancy or fugacity $\Upsilon(t)$, whose dynamics are isolated by the requirement that the lowest modes capture the energy and particle number densities. This gives the method a natural physical interpretation. In particular, the dynamical fugacity is capable of naturally expressing the emergence of chemical non-equilibrium during the freeze-out process while the effective temperature captures any reheating phenomenon. Any system in approximate kinetic equilibrium that undergoes reheating and/or transitions to chemical non-equilibrium is a good match for this method. In fact it is almost assured that our method will be considerably more computationally economical than the chemical equilibrium spectral method for any physical system in which the cost of computing the collision terms is high. We validated the method on a model problem that exhibits the physical characteristics of reheating and chemical non-equilibrium. We demonstrated that particle number and energy densities are captured accurately using only two degrees of freedom, the effective temperature and fugacity. In general, this will hold so long as the back reaction from non-thermal distortions is small i.e. as long as kinetic equilibrium is a good approximation. The method presented here should be compared to the spectral method used in \cite{Esposito2000,Mangano2002}, which uses a fixed basis of orthogonal polynomials and is adapted to systems that are close to chemical equilibrium (or with a non-dynamical chemical potential in \cite{Esposito2000}) with dilution temperature scaling. In addition to more closely mirroring the physics of systems that exhibit reheating and chemical non-equilibrium, the method presented here has a computational advantage over the chemical equilibrium method. Even when the system is close to chemical equilibrium with dilution temperature scaling, as is the case for the problem studied in \cite{Esposito2000,Mangano2002}, the method presented here reduces the minimum number of degrees of freedom needed to capture the particle number and energy densities from four to two. In turn, this reduces the minimum number of collision integrals that must be evaluated by more than half. Numerical evaluation of collision operators for realistic interactions is a costly operation and so the new `emergent chemical non-equilibrium' approach we have presented here constitutes a significant reduction in the numerical cost of obtaining solutions. Moreover, even if the chemical equilibrium approach were to be properly modified to gain mathematical advantages we show in our chemical non-equilibrium approach, it is not at all clear that the chemical equilibrium method can, with comparable numerical effort, achieve a precise solution under conditions where transient or final chemical non-equilibrium and reheating are strong. Looking to future applications, the gain in numerical efficiency we achieved should allow both space and time evolution to be considered in non-trivial dynamical models such as evolution of quark-gluon plasma fireball formed in relativistic heavy ion collisions. We expect to be able to explore within the realm of Boltzmann dynamics the question of `ideal' quark flow occurring at minimum viscosity~\cite{Romatschke:2007} and the shape of momentum distributions which systematically deviate from a thermal distribution~\cite{Wilk:2009}. For study of the ensuing hadron flow it is of relevance that we have been able to find suitable weights defining a spectrum of basis states capable of addressing the case of particles where the mass scale is relevant and chemical non-equilibrium is strong while retaining some of the advantages presented here for massless particles. We will discuss the extension to the massive case in a future work. | 14 | 3 | 1403.2019 |
1403 | 1403.2369_arXiv.txt | We report a measurement of the $B$-mode polarization power spectrum in the cosmic microwave background (CMB) using the \pb\ experiment in Chile. The faint $B$-mode polarization signature carries information about the Universe's entire history of gravitational structure formation, and the cosmic inflation that may have occurred in the very early Universe. Our measurement covers the angular multipole range $500 < \ell < 2100$ and is based on observations of an effective sky area of 25\,\sqdeg\ with 3.5\arcmin\ resolution at 150\,GHz. On these angular scales, gravitational lensing of the CMB by intervening structure in the Universe is expected to be the dominant source of $B$-mode polarization. Including both systematic and statistical uncertainties, the hypothesis of no $B$-mode polarization power from gravitational lensing is rejected at \rejectnull\ confidence. The band powers are consistent with the standard cosmological model. Fitting a single lensing amplitude parameter \ABB\ to the measured band powers, \ABBmeas, where $A_{BB} = 1$ is the fiducial \wmap-9 \lcdm\ value. In this expression, ``stat'' refers to the statistical uncertainty, ``sys'' to the systematic uncertainty associated with possible biases from the instrument and astrophysical foregrounds, and ``multi'' to the calibration uncertainties that have a multiplicative effect on the measured amplitude $A_{BB}$. | \setcounter{footnote}{0} The cosmic microwave background (CMB) radiation is emitted from the primordial plasma in the early universe when stable hydrogen first forms at a redshift of $z = 1091$ \citep{2013ApJS..208...20B}. The scalar density fluctuations present in the primordial plasma at that time, which are the seeds for later structure formation, create both CMB intensity and polarization anisotropies. These scalar fluctuations can only create polarization patterns of even spatial parity, referred to as $E$-modes \citep{Seljak1997, 1997PhRvD..55.1830Z, 1997PhRvD..55.7368K}. Precision measurements of the angular power spectrum of the intensity fluctuations, \cltt, are a cornerstone of our current understanding of cosmology. The power spectrum of the primordial $E$-modes, \clee, has also been well-characterized, as has their relationship with the temperature anisotropy pattern, \clte. These measurements are consistent with a single source for both the temperature and $E$-mode signals -- the adiabatic density fluctuations in the primordial universe \citep{2014ApJ...783...67B,2005ApJ...624...10L,2006ApJ...647..813M, brown2009, quiet_FirstSeason, quiet_SecondSeason, 2013ApJS..208...20B}. Unlike $E$-modes, odd-parity $B$-mode patterns are not produced by scalar fluctuations. Any primordial $B$-modes would be evidence for tensor or vector perturbations in the gravitational metric when the CMB was emitted \citep{1997PhRvL..78.2054S, 1997PhRvL..78.2058K}. The only source of such perturbations in the standard cosmological model is a period of cosmic inflation in the early universe. If cosmic inflation occurred, the induced tensor perturbations (or inflationary gravitational wave background) would imprint a Gaussian field of $B$-mode polarization on the primordial CMB that is largest at degree angular scales. The angular power spectrum of those primordial $B$-modes, \clbb, would then provide information about cosmic inflation, and could be a window into physics at grand-unified energy scales, when the electroweak and strong forces are expected to unify \citep{KamionkowskiKosowski}. The CMB radiation that we observe has been modified by secondary effects as compared with the primordial signal. One such effect is the gravitational lensing of CMB photons by cosmological large-scale structure (LSS). The gravitational structure traversed by the CMB radiation as it travels to us distorts the primordial CMB temperature and polarization fluctuations, and generates $B$-mode polarization anisotropies \citep{HuOkamoto2002, 2006PhR...429....1L}. This distortion adds secondary power to the $B$-mode angular power spectrum peaking at scales of $0.2^\circ$, and also imprints a non-Gaussian correlation between anisotropies in the CMB temperature and polarization. That correlation can be used to reconstruct maps of the integrated structure along the line of sight -- all the structure in the observable universe. High fidelity maps of this effect will be a powerful probe of fundamental physics, cosmology, and extragalactic astrophysics. These maps will also enable the removal of this secondary $B$-mode signal to precisely characterize any primordial \clbb\ due to inflationary gravitational waves \citep{seljak2004}. The scientific prospects from precise characterization of LSS using CMB lensing are significant. The neutrino mass, known to be non-zero from neutrino flavor oscillation measurements, can be measured by its effect on LSS formation. The high velocities of cosmic neutrinos inhibit gravitational clustering on small scales, suppressing LSS on scales smaller than $\sim 100$\,Mpc. Measurements of the gravitationally lensed CMB in both temperature and polarization have the potential to measure the sum of neutrino masses with an uncertainty comparable to the known mass splittings measured by flavor oscillation experiments, which set the minimum sum of the neutrino masses to be 58\,meV \citep{beringer2012}. Also, CMB lensing measurements break the degeneracy that exists in primordial CMB temperature anisotropies between curvature, dark energy parameters, and the sum of neutrino masses \citep{SmithHuKaplinghat_2006,calabrese2009,smith09}, and thus improve constraints on those parameters. In particular, CMB lensing measurements are complementary to other probes of dark energy because they are sensitive to its high-redshift behavior \citep{calabrese2011}. Cross-correlating LSS maps from CMB lensing measurements with other LSS mass-tracing probes will improve the calibration of these tracers and the statistical accuracy of the resulting cosmological constraints. Lensing of the CMB by LSS was first observed in temperature measurements, through the LSS-induced non-Gaussianities \citep{smith2007,das2011,das2013_arXiv,vanEngelen2012,planck2013_XVII_arXiv}, and modification of the power spectrum \cltt\ \citep{calabrese2008,reichardt2009,das2013_arXiv,story2012_arXiv,planck_XVI_arXiv}. Polarization measurements have the potential to more precisely reconstruct the LSS-induced signal because the lensing $B$-modes are not contaminated by large primordial CMB fluctuations \citep{HuOkamoto2002}. Measurement of non-Gaussianity in the CMB polarization was first reported recently by \sptpol\ and \pb\ using cross-correlation with Herschel observations of high-redshift galaxies \citep{hanson2013,PB_GalaxyCross_2014}. While the cross-correlation studies establish the existence of $B$-modes from gravitational lensing, they are not sensitive to LSS throughout cosmic time because they rely on tracers that exist over a limited range of redshift. Recently \pb\ reported a measurement of the non-Gaussianities induced by LSS in the polarized CMB from CMB data alone, which is therefore sensitive to all lensing distortions along the line of sight \citep{PB_CLdd_2014}. These early measurements imply an amplitude of \clbb, but a detection of \clbb\ \textit{itself} has not yet been published. The small amplitude of this signal compared to other sources of anisotropy in the CMB makes it very difficult to measure without contamination from the instrument or astrophysical foregrounds. In this sense, \clbb\ is more difficult to characterize than the non-Gaussianity induced by LSS, but its precise characterization is required to search for the signal from cosmic inflation, and extract all of the science possible from $B$-mode cosmology. In this paper we present a measurement of \clbb\ using \pb. The \pb\ experiment uses a millimeter-wave polarimeter to make deep maps of the CMB temperature and polarization anisotropies. \Secref~\ref{sec:instrument} describes the instrument, and \secref~\ref{sec:obs} details the observations that were performed to obtain the data reported here. \Secref~\ref{sec:calibration} describes the calibration of these data, and \secref~\ref{sec:analysis} describes the analyses we used to produce the \clbb\ measurement. Possible sources of systematic contamination of \clbb\ are evaluated in \secref~\ref{sec:fg} -- astrophysical foregrounds -- and \secref~\ref{sec:systematics} -- instrumental systematics, and found to be small compared to the measured signal. Finally, we present the \pb\ measurement of binned \clbb\ power over angular multipoles $500 < \ell < 2100$ in \secref~\ref{sec:results}, and conclude with a discussion of the measurement in \secref~\ref{sec:summary}. | PTEs resulting from the null test framework. No significantly low or high PTE values are found, consistent with a lack of systematic contamination or miscalibration in the \pb\ data set and analysis. Note that the PTE values in each patch are not independent from each other. } \begin{tabular}{cccccc} \tableline \tableline & average of & extreme of & extreme of & extreme of & total \\ Patch & $\chi_{\rm null}(b)$ & $\chi^2_{\rm null}(b)$ & $\chi^2_{\rm null}$ by {\it EB}/{\it BB} & $\chi^2_{\rm null}$ by test & $\chi^2_{\rm null}$\\ \tableline RA4.5 & 11.6\% & 16.6\% & 20.6\% & 21.8\% & 14.0\% \\ RA12 & 92.4\% & 84.2\% & 60.8\% & 23.8\% & 52.6\% \\ RA23 & 75.2\% & 61.6\% & \phantom{0}6.0\% & \phantom{0}7.0\% & 18.6\% \\ \tableline \end{tabular} \end{center} \end{table*} To probe for systematic contamination that is focused in a particular power spectrum or null test data split, we calculate the sum of $\chi_{\rm null}^2(b)$ over $500 < b < 2100$ for $EB$ and $BB$ separately, (``$\chi^2_{\rm null}$ by spectrum''), and the sum of both these spectra for a specific test (``$\chi^2_{\rm null}$ by test''). \figref~\ref{fig:dist_chi2b} shows the PTE distribution of the $\chi^2_{\rm null}$ by (a) bin, (b) spectrum, and (c) test for the three patches. We require that each of these sets of PTEs each be consistent with a uniform distribution, as evaluated using a KS test, requiring a p-value (probability of seeing deviation from uniformity greater than that which is observed given the hypothesis of uniformity) greater than 5\%. These distributions are consistent with a uniform distribution from zero to one. We create test statistics based on these quantities to search for different manifestations of systematic contamination. The five test statistics are (1)~the average value of $\chi_{\rm null}$; the extreme value of $\chi^2_{\rm null}$ by (2)~bin, (3)~spectrum, and (4)~test; and (5)~the total $\chi^2_{\rm null}$ by summing up the nine null tests. In each case, the result from the data is compared to the result from simulation, and PTEs are calculated. Finally, we combine each of the test statistics, and calculate the PTE of that final test statistic, requiring it to be greater than 5\%. \tabref~\ref{tab:pte_summary} shows summary of the PTE values of each test statistic for each patch. Comparing the most significant outlier from the five test statistics with that from simulations, we get PTEs of 32.8\%, 55.6\%, and 18.0\% for RA4.5, RA12, and RA23 respectively. We achieve the requirements described above, finding no evidence for systematic contamination or miscalibration in the \pb\ data set and analysis. \subsection{Cross-check using a second pipeline} \label{sec:ubpipe} Concurrently, we have been developing an alternate data processing pipeline that was used to cross-check the results presented here. Its full description will be given in a forthcoming publication; here we highlight its most salient features. In the time domain, the alternate pipeline applies the same filters as the primary pipeline, but corrects for them while estimating the sky signals as part of the map-making procedure, following \citet{Stompor2002}. The recovered maps provide unbiased renditions of the sky signal, with the filtered modes effectively marginalized over. This is numerically challenging so we use a divide-and-conquer approach, which results in unbiased but slightly sub-optimal maps. The maps are estimated in the HEALPix pixelization \citep{Gorski2005} with $N_\textrm{side} = 2048$, so no flat-sky assumption is adopted. We produce the maps of three Stokes parameters and the $Q$ and $U$ maps are used to estimate the polarized power spectra of the sky signals. This is done with power spectrum estimation software packages based either on the pure-pseudospectra \citep{Smith2006}, \textsc{xpure} and \textsc{x$^2$pure} \citep{xpure,Grain2012,ferte2013}, or the standard pseudospectra \textsc{xpol} \citep{tristram2005} approaches. The mode-coupling matrices are computed explicitly by directly summing the required Wigner-3j symbols based on the geometry of the observed patches, noise weights and apodizations. The final spectra are calculated as weighted averages of the cross-spectra of 8 maps made of disjoint subsets of all daily maps, and \mc\ simulations are employed to estimate the final uncertainties of the computed spectra. The results of this alternate pipeline are consistent with the results of the primary pipeline described in this publication. \begin{figure*}[htpb] \centering \includegraphics[width=7in]{BB_science.pdf} \caption{\label{fig:resultBB}Binned \clbb\ spectrum measured using data from all three patches~($\sim$25\,\sqdeg). A theoretical \wmap-9 \lcdm\ high-resolution \clbb\ spectrum with \ABB$ = 1$ is shown. The uncertainty shown for the band powers is the diagonal of the band power covariance matrix, including beam covariance.} \label{fig:bmode_spectrum} \end{figure*} \subsection{Blind analysis} \label{sec:blind} The possibility of data analyzers biasing their result toward their own preconceptions, known as ``observer bias'', is a form of systematic bias that can affect the result of an experiment \citep{Klein:2005di}. Examples of preconceptions include theoretical predictions, the statistical significance that the team expect to obtain, or consistency with previous measurements. Since it is difficult to estimate the effects of observer bias, we employed an analysis methodology designed to minimize its impact. We have adopted a blind-analysis framework, which is a standard technique to minimize observer bias. % In our framework, no one in the team viewed the measured \clbb\ values, the deflection power spectra based on $B$-modes \citep{PB_CLdd_2014,PB_GalaxyCross_2014}, or the corresponding maps, until we eliminated possible sources of observer bias by finalizing calibration, filtering, data selection, data validation and showed that all systematic uncertainties were small. This framework forced us to develop quantitative tools, including null tests and simulations, that convincingly argued for analysis choices and constraints without showing \clbb, thus removing the possibility that people within the team would be more convinced by an argument or method because of the \clbb\ that it produced. Other power spectra and maps were used as subsidiary information in this work, and they were unblinded in stages during the analysis procedure. In fact, after un-blinding \clbb, questions came up about how well we had constrained electrical crosstalk, and how robust our estimate of the binned power spectrum uncertainty was. Finding our previous argument constraining electrical crosstalk weak, we developed the simulation shown in \secref~\ref{sec:crosstalk}, where we estimated that electrical crosstalk is one of our smallest systematic uncertainties. Investigating our binned power spectrum uncertainties, because of comparisons with a second pipeline, we found an error in our uncertainty estimation code. This was an error that could have been found while we were blind, but it was not. The error did not affect the central values of the measurement. We corrected this error, resulting in a reduction in the significance of our measurement by about 18\% between un-blinding and the results presented here. The qualitative consistency of the measurement with theory was not changed, the change was motivated by a pipeline comparison, and it reduced the significance of our measurement; we do not believe that this was a significant opportunity for the result to be incorrectly affected by observer bias. \label{sec:summary} We have reported a measurement of the CMB's $B$-mode angular power spectrum, \clbb, over the multipole range $500 < \ell < 2100$. This measurement is enabled by the unprecedented combination of high angular resolution (3.5\arcmin) and low noise that characterizes the \pb\ CMB polarization observations. To validate the \pb\ measurement of this faint signal, we performed extensive tests for systematic errors. We evaluated nine null tests and estimated twelve sources of instrumental contamination using a detailed instrument model, and found that all the systematic uncertainties were small compared to the statistical uncertainty in the measurement. To motivate comprehensive evaluation of the data set and prevent observer bias in data selection and analysis, the analysis was performed blind to the \clbb\ signal; all data selection and analysis choices were fixed and all systematic error tests were completed before any team members looked at the $B$-mode power spectrum. \pb\ has reached an important CMB polarization milestone, with noise levels sufficiently low to allow reconstruction of the lensing signal with more precision from polarization than from CMB temperature \citep{HuOkamoto2002}. We previously presented evidence for gravitational lensing of the CMB in \pb\ data using the non-Gaussianity imprinted in the CMB by LSS \citep{PB_GalaxyCross_2014,PB_CLdd_2014}. Those analyses, arising from the same area of sky, are also consistent with the \lcdm\ expectation and give no evidence for significant systematic errors. We can calculate the combined significance with which those measurements of non-Gaussian $B$-modes and the \clbb\ measurements reported here reject the hypothesis that there are no CMB lensing $B$-modes. In this null hypothesis, the signals are uncorrelated (when using a realization-dependent lensing bias subtraction to calculate the deflection field), so a simple quadrature sum of the rejection significance is appropriate. This calculation results in a rejection of the hypothesis that there are no lensing $B$-modes with \rejectnullcombinedsigma\ confidence for a normal distribution. CMB $B$-mode polarization is emerging as a key observable in modern cosmology. Over the next few years, measurements of CMB $B$-mode polarization will allow us to probe LSS in detail to provide insight into fundamental physics, cosmology, and extragalactic astrophysics. Detailed analysis of the signal produced by LSS will enable precision characterization of the possible underlying \clbb spectrum from cosmic inflation. The measurement of LSS-induced $B$-mode power in \pb\ data, characterized by both its non-Gaussian signature and its \clbb\ power, represents an important step in the rapidly progressing field of CMB $B$-mode science. | 14 | 3 | 1403.2369 |
1403 | 1403.5580_arXiv.txt | If $R$-parity is only mildly violated then the lightest supersymmetric particle (LSP) can be stable over cosmologically time-scales and still account for the dark matter relic density. We examine the possibility of generating detectable X-ray lines from $R$-parity violating decays of keV-scale LSP dark matter to neutrino-photon pairs. Specifically, we consider scenarios in which the LSP is a light gravitino, bino, or hidden sector photino. Potential signals are discussed in the context of recent claims of an unidentified 3.5 keV X-ray line in studies of stacked galaxy clusters. We comment on the difficulties in obtaining the observed relic density for keV scale bino or hidden photino dark matter and some possible resolutions. | Despite negative searches at the LHC thus far, TeV scale supersymmetry (SUSY) remains the leading resolution to the hierarchy problem, particularly in light of the discovery of a Higgs boson not much above the $Z$-mass. Whilst experimental searches and theoretical considerations suggest that most of the SUSY spectrum should be above the weak scale, it is quite conceivable that certain neutral states could be substantially lighter. Such light sparticles could have interesting cosmological and astrophysical implications. Here we shall focus on the prospect of keV-scale SUSY states which can decay in such a manner as to produce observable X-ray signals. We are inspired to think about decays of light particles by the tentative 3.5 keV line observed in the combined spectra of multiple galaxy clusters as studied by the XMM-Newton X-ray observatory \cite{Bulbul:2014sua,Boyarsky:2014jta}. Thus we shall typically phrase our discussion in terms of this benchmark point and examine models of dark matter which might accommodate this phenomenon. This signal can be interpreted in terms dark matter which decays to a photon and an (effectively) massless degree of freedom with a mass and lifetime around \beq m_{\rm DM} & \simeq7~{\rm keV}~,\\ \tau_{\rm DM} & \simeq2\times10^{27}~\text{-}~2\times10^{28}~{\rm s}~. \label{1} \eeq Some potential models have been proposed which might account for this observational anomaly \cite{Bulbul:2014sua,Ishida:2014dlp,Finkbeiner:2014sja,axion,Jaeckel:2014qea, Abazajian:2014gza,Krall:2014dba,Lee:2014xua, Baek:2014qwa,Choi:2014tva,Nakayama:2014ova, Frandsen:2014lfa,Aisati:2014nda,Kong:2014gea,Cicoli}, the possibility of decaying sterile neutrinos or axions receiving particular attention. Given the wide expectation that SUSY should play a leading role in physics beyond the Standard Model, it is interesting to explore possible SUSY explanations for this signal. As is well known, the canonical SUSY extension of the Standard Model, the MSSM, requires the imposition of an ({\em ad hoc}) discrete $Z_2$ symmetry: $R$-parity \beq (-1)^{3(B-L)+2s}~. \eeq Under this discrete symmetry the Standard Model (superpartner) states transform as even (odd) representations. This is necessary purely for phenomenological purposes since there are dimension four and five operators of the form \beq \mu' \boldsymbol {L H_u}~,\quad \lambda \boldsymbol{L L \overline{E}}~,\quad \lambda' \boldsymbol{L Q \overline{D}}~,\quad \lambda'' {\boldsymbol {\overline{U}\overline{D}\overline{D}}}~, \label{rpv} \eeq which, unless the couplings are small, are problematic as they lead to fast proton decay in conflict with experimental searches \cite{Abe:2011ts,Hikasa:1992je}. If $R$-parity is an exact symmetry then the lightest supersymmetric particle (LSP) is stable. Intriguingly a stable LSP can play the role of dark matter and the occurrence of a well motivated dark matter candidate is arguably one of the great triumphs of SUSY extensions of the Standard Model, see e.g.~\cite{Jungman:1995df}. An interesting variation is the scenario in which $R$-parity is mildly violated \cite{Hall:1983id,Dawson:1985vr,Barbier:2004ez,Dreiner:1997uz} such that the LSP is effectively stable on cosmologically time-scales and can still account for the dark matter. However, a small fraction of these states will decay presently and for an appropriate lifetime can potentially generate detectable signals. Good candidates for the LSP in the MSSM are the fermion superpartners to the known boson fields, such as the neutralino or gravitino. Assuming mild $R$-parity violation (RPV), a fermion LSP lighter than the electron will dominantly decay to a photon-neutrino pair, unless the spectrum is supplemented with additional light fermion states. For an LSP decaying to a photon and an effectively massless state to produce X-ray signals, the parent state must be around the keV scale. More specifically, to match the recent anomaly at 3.5 keV \cite{Bulbul:2014sua} the parent state should be roughly 7 keV. Since we suppose that the LSP constitutes the dark matter, this will be `warm' dark matter. It should be noted that sub-keV thermally produced `hot' dark matter leads to the erasure of density perturbations at scales shorter than its free streaming length and is in conflict with observations of small scale structure, see e.g.~\cite{Viel:2005qj}. Although the cluster anomaly at 3.5 keV might be regarded as tentative at this stage, it provides motivation for us to explore interesting, non-standard, scenarios of SUSY. We begin in Sect.~\ref{S2} by investigating if a keV gravitino can give rise to decay signals and show that generically this is not the case. In Sect.~\ref{S3}, we explore the prospect of generating X-ray lines from decaying bino dark matter and argue that this is quite possible. However, as we discuss, obtaining the correct relic density for bino dark matter requires significant model building. In Sect.~\ref{S4}, we examine the motivation for light hidden sector photini and argue that such a state could be the LSP. We show that a 7 keV hidden photino LSP can have a suitable abundance to match the observed dark matter relic density and give rise to the 3.5 keV cluster line via $R$-parity violating decays. | We have discussed the prospect for generating keV X-ray signals in SUSY extensions of the Standard Model, with specific reference to the 3.5 keV line recently reported in analysis of the observations by the XMM-Newton X-ray Telescope \cite{Bulbul:2014sua,Boyarsky:2014jta}. Given that SUSY is the leading candidate for a framework of physics beyond the Standard Model, we believe that it is interesting to consider how such a signal might arise in this setting. We have highlighted the possibility of X-ray signals being generated by the decays of light LSP dark matter via $R$-parity violating operators and argued that such scenarios are viable and motivated, although in some cases require additional model building or some amount of fine-tuning. The gravitino is one of the best motivated light SUSY states, but generically we have argued that it can not lead to signatures of this type. We have proposed rather that scenarios involving keV binos or hidden sector photinos could account for this signal. Interestingly, both these models are sensitive to $T_{\rm RH}$ and typically to be successfully realised it is required that the maximum temperature after reheating is lower than a TeV. We note in passing that whilst the axino is also a good candidate for a light LSP, similar to the gravitino, a keV axino is typically too long lived to account for the 3.5 keV line. Comparing with the model of e.g.~\cite{Endo:2013si}, the axino lifetime is parametrically \beq \tau_{\widetilde{a}}\sim 5\times10^{29}~{\rm s} &\left(\frac{m_{\widetilde{a}}}{7~{\rm keV}}\right)^{-3} \left(\frac{f_{a}}{10^8~{\rm GeV}}\right)^2\\ &\times \left(\frac{m_{\widetilde{B}}}{100~{\rm GeV}}\right)^2 \left(\frac{\langle\widetilde{\nu}\rangle/v}{10^{-6}}\right)^2~. \eeq However, as noted in \cite{Kong:2014gea,Choi:2014tva} for extreme parameter choices -- with $f_a$, $m_{\widetilde{B}}$ and $\langle\widetilde{\nu}\rangle$ at the edge of experimental exclusion for most minimal models -- axino interpretations of eq.~(\ref{1}) can be constructed. Aside from these tensions with experimental constraints, this scenario may also be disfavoured from a theoretical stand point \cite{Cheung:2011mg} as, in the absence of fine tuning or sequestering, the axino mass is expected to be $m_{\widetilde{a}}\gtrsim m_{3/2}$ \cite{Cheung:2011mg}. Thus ensuring that the axino is the LSP requires some model building. Given these considerations we do not discuss this scenario in greater detail. In closing, we highlight that there is a potential opportunity to distinguish between interpretations involving axions \cite{axion,Jaeckel:2014qea,Lee:2014xua,Cicoli}, and those invoking alternative light dark matter candidates (like the bino or hidden photino) using precision measurements of $\Delta N_{\rm eff}=3\rho_{\rm hidden}/ \rho_{\nu}$. As typically axions are dominantly produced non-relativistically via the misalignment mechanism \cite{Shifman:1979if,Kim:1979if}, their contribution to $N_{\rm eff}$ is negligible, in contrast to keV scale thermally produced dark matter. It is projected that upcoming experiments \cite{Abazajian:2013oma} will be sensitive to percent-level changes in $N_{\rm eff}$ and thus should provide some insight regarding the nature of any light hidden sector states. Further study of this 3.5 keV X-ray line is certainly warranted. A strong confirmation of this signal, particularly in conjunction with the observation of a factional increase in $N_{\rm eff}$ would be an exciting signal of physics beyond the Standard Model, possibly in the guise of supersymmetry. {\bf Acknowledgements:} We are grateful to Antonio Delgado, Zhaofeng Kang, Adam Martin, Jessie Shelton, and Yue Zhang for useful interactions and the PRD referee for helpful comments. JU is grateful for the hospitality of the Department of Physics at UIUC. This research was supported by the National Science Foundation under Grant No.~PHY-1215979. | 14 | 3 | 1403.5580 |
1403 | 1403.5549_arXiv.txt | { In this paper we will analyse the constraints on a sub-Planckian excursion of a single inflaton field, which would yield a large tensor to scalar ratio, while explaining the temperature anisotropy of the cosmic microwave background (CMB) radiation. In particular, our attempt will be to reconstruct the inflationary potential by constraining, $V(\phi_0),~V^{\prime}(\phi_0),~V^{\prime\prime}(\phi_0),~V^{\prime\prime\prime}(\phi_0)$ and $V^{\prime\prime\prime\prime}(\phi_0)$, in the vicinity of the field, $\phi_0\ll M_p$, and the field displacement, $\Delta \phi \ll M_p$, where $M_p$ is the reduced Planck mass. We will provide, for the first time, a set of new {\it consistency} relationships for sub-Planckian excursion of the inflaton field, which would help us to differentiate sub-versus-super Planckian models of inflation. For a generic single field inflationary potential, we will be able to put a stringent bound on the potential energy density: $2.07\times10^{16}~{\rm GeV}\leq\sqrt[4]{V_{\star}}\leq 2.40\times 10^{16}~{\rm GeV}$, where inflation can occur on the flat potential within, $0.066 \leq\frac{\left |\Delta\phi\right|}{M_p}\,\leq 0.092$, for the following observational constraints: (Planck+WMAP-9+high L+BICEP2). We then provide a prediction for the spectral tilt ($n_{T}$), running ($\alpha_{T}$) and running of running ($\kappa_{T}$) of the tensor modes within the window, $-0.019<n_{T}<-0.033$, $-2.97\times 10^{-4}<\alpha_{T}(=dn_{T}/d\ln k)<2.86\times 10^{-5}$, and $-0.11\times 10^{-4}<\kappa_{T}(=d^{2}n_{T}/d\ln k^{2})<-3.58\times 10^{-4}$, in a model independent way. We also provide a simple example of an {\it inflection-point} model of inflation and reconstruct the potential in a model independent way to match the current observations. } \begin{document} | The primordial inflation~\cite{Guth:1980zm,Linde,Albrecht} has two {\it key} predictions - creating the scalar density perturbations and the tensor perturbations during the accelerated phase of expansion~\cite{Mukhanov:1981xt}, for a review, see~\cite{Mukhanov:1990me}. One of the predictions, namely the temperature anisotropy due to the scalar density fluctuations has now been tested very accurately by the observations from the temperature anisotropy in the cosmic microwave background (CMB) radiation~\cite{Hinshaw:2012aka,Planck-1,Planck-infl}. The detection of tensor modes has now recently been confirmed by the ground based BICEP experiment~\cite{Ade:2014xna}, which has detected for the first time a non-zero value of the tensor-to-scalar ratio at $7\sigma$ C.L. The value obtained by the BICEP team in conjunction with (Planck+WAMP-9+high L+BICEP2) put a bound on the primordial gravitational waves, via tensor-to-scalar ratio, within a window, $0.15 \leq r(k_\star)\equiv P_T(k_\star)/P_S(k_\star)\leq 0.27$, at the pivot scale, $k_\star = 0.002 {\rm Mpc^{-1}}$~\cite{Ade:2014xna}, where $P_T$ and $P_S$ denote the power spectrum for the tensor and scalar modes, respectively. Note that large $r(k_\star)$ is possible only if the initial conditions for gravitational waves is quantum Bunch-Davis vacuum~\cite{BD}, for a classical initial condition the amplitude of the gravitational waves would be very tiny and undetectable~\cite{Ashoorioon:2012kh}, therefore the first observable proof of quantum gravity. In this paper, our aim will be to illustrate that it is possible to explain the current data sets (Planck+WAMP-9+high L+BICEP2) within a sub-Planckian VEV model of inflation, where: \be \phi_0 \ll M_p\,,~~~~~~|\Delta\phi |\approx |\phi_\star-\phi_e| \ll M_p, \ee where $\phi_\star \geq \phi_0\geq \phi_e$ represents the field VEV, and $\Delta\phi$ denotes the range of the field values around which all the relevant inflation occurs, $\phi_\star$ corresponds to the pivot scale and $\phi_e$ denotes the end of inflation, and $M_p=2.4\times 10^{18}$~GeV. Naturally, the potential has to be flat enough within $\Delta \phi$ to support slow roll inflation. The above requirements are important if the origin of the inflaton has to be embedded within a particle theory, where inflaton is part of a {\it visible sector} gauge group, i.e. Standard Model gauge group, instead of an arbitrary gauge singlet, for a review, see~\cite{Mazumdar:2010sa}. If the inflaton is {\it gauged} under some gauge group, as in the case of a minimal supersymmetric Standard Model (MSSM), Ref.~\cite{Allahverdi:2006iq}, then the inflaton VEV must be bounded by $M_p$, in order to keep the sanctity of an effective field theory description~\footnote{ An arbitrary moduli or a gauge singlet inflaton can take large VEVs ( super-Planckian ) as in the case of a chaotic inflation~\cite{Linde}. Although, in the case of {\it assisted inflation}~\cite{assisted}, see {\it chaotic assisted inflation}~\cite{kanti}, the individual VEVs of the inflatons are sub-Planckian.}. Our treatment will be very generic, with a potential given by: \begin{eqnarray}\label{rt10a} V(\phi)&=&V(\phi_0)+V^{\prime}(\phi_0)(\phi-\phi_{0})+\frac{V^{\prime\prime}(\phi_0)}{2}(\phi-\phi_{0})^{2}+\frac{V^{\prime\prime\prime}(\phi_0)}{6}(\phi-\phi_{0})^{3}\,\nonumber\\ &&~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~~+\frac{V^{\prime\prime\prime\prime}(\phi_0)}{24}(\phi-\phi_{0})^{4}+\cdots\,, \end{eqnarray} where $V(\phi_0)\ll M_p^4$ denotes the height of the potential, and the coefficients: $V^{\prime}(\phi_0) \leq M_p^3,~V^{\prime\prime}(\phi_0)\leq M_p^2,~V^{\prime\prime\prime}(\phi_0)\leq M_p,~V^{\prime\prime\prime\prime}(\phi_0)\leq {\cal O}(1)$, determine the shape of the potential in terms of the model parameters. The {\it prime} denotes the derivative w.r.t. $\phi$~\footnote{ In particular, some specific choices of the potential would be a {\it saddle point}, when $V^{\prime}(\phi_0)=0=V^{\prime\prime}(\phi_0)$, an {\it inflection point}, when $V^{\prime\prime}(\phi_0)=0$.}. Previous studies regarding obtaining large $r(k_\star)$ within sub-Planckian VEV models of inflation have been studied in Refs.~\cite{BenDayan:2009kv,Shafi,Hotchkiss:2011gz,Choudhury:2013iaa}. In Ref.~\cite{BenDayan:2009kv}, the authors could match the amplitude of the power spectrum, $P_S$, at the pivot point, but not at the entire range of $\Delta\phi$ for the observable window of $\Delta N$, where $N$ is the number of e-foldings of inflation. In Ref.~\cite{Hotchkiss:2011gz}, the authors have looked into higher order slow roll corrections by expanding the potential around $\phi_0$. They pointed out that large $r\sim 0.05$ could be obtained in an {\it inflection-point} model of inflation where the slow roll parameter, $\epsilon_V$, changes non-monotonically ( for a definition of $\epsilon_V$, see Eq.~(\ref{ra1}) ). The $\epsilon_V$ parameter first increases within the observational window of $\Delta N$ and then decreases before increasing to exit the slow roll inflation by violating the slow roll condition, i.e. $\epsilon_V\approx 1$. The value of $r$ was still small in order to accommodate the WMAP data, which had probed roughly $ \Delta N\approx 8$ as compared to the Planck, which has now probed $\Delta N\approx 17$ e-foldings of inflation. Following Ref.~\cite{Hotchkiss:2011gz}, a generic bound on tensor to scalar ratio was presented in Ref.~\cite{Choudhury:2013iaa} for an {\it inflection-point} model only, when $V^{\prime\prime}(\phi_0)=0$. In this paper we will consider the full potential of Eq.~(\ref{rt10a}), and our goal will be to determine the values of $V(\phi_\star),~V^{\prime}(\phi_\star),~V^{\prime\prime}(\phi_\star),~V^{\prime\prime\prime}(\phi_\star)$ and $V^{\prime\prime\prime\prime}(\phi_\star)$ from the current (Planck+WAMP-9+high L+BICEP2) data. In this respect we will be reconstructing the inflationary potential around $\phi_0$ and the pivot scale, $\phi_\star=\phi(k_\star)$, where $k_\star = 0.002 {\rm Mpc^{-1}}$. We will also provide for he first time the consistency relations for a sub-Planckian excursion of the inflaton field. This will provide us the with an observational discriminator which could falsify sub-Planckian models of inflation in future. Within Planck's observable region of $\Delta {N}\approx 17$ e-foldings, we will be able to constrain the power spectrum: $P_S$, spectral tilt: $n_S$, running of the spectral tilt: $\alpha_S$, and running of running of the spectral tilt: $\kappa_S$, for (Planck+WMAP-9+high~L+BICEP2) data sets: \begin{eqnarray}\label{obscons} 0.15 &\leq & r(k_\star)\leq 0.27\\ \ln(10^{10}P_{S})&=&3.089^{+0.024}_{-0.027}~~ ({\rm within}~ 2\sigma ~C.L.),\\ n_{S}&=&0.9600 \pm 0.0071~~ ({\rm within}~ 3\sigma ~C.L.),\\ \alpha_{S}&=&dn_{S}/d\ln k=-0.022\pm 0.010~~({\rm within}~1.5\sigma~C.L.),\\ \kappa_{S}&=&d^{2}n_{S}/d\ln k^{2}=0.020^{+0.016}_{-0.015}~~({\rm within}~1.5\sigma~C.L.)\,. \end{eqnarray} We will briefly recap the main equations for the tensor to scalar ratio in the most general case by taking into account of the higher order slow-roll conditions. We will then derive the {\it most general} bound on $r(k_\star)$ for a sub-Planckian VEV inflation, and the corresponding values of $H_\star$ and $V(\phi_\star)$. We will then reconstruct the shape of the potential in Section 4 by providing the constraints on $V^{\prime}(\phi_0),~V^{\prime\prime}(\phi_0),~V^{\prime\prime\prime}(\phi_0)$ and $V^{\prime\prime\prime\prime}(\phi_0)$. In Section 5, we will discuss the consistency relationships, and in section 6, we will consider a specific case of inflection point inflation for the purpose of illustration. We will provide our key equations in an Appendix. | The tensor to scalar ratio can be defined by taking into account of the higher order corrections, see~\cite{Easther:2006tv,Choudhury:2013iaa,Choudhury:2013jya}: \be\label{para 21e} r\approx16\epsilon_{V}\frac{\left[1-({\cal C}_{E}+1)\epsilon_{V}\right]^{2}}{\left[1-(3{\cal C}_{E}+1)\epsilon_{V} +{\cal C}_{E}\eta_{V}\right]^{2}}\, \ee where ${\cal C}_{E}=4(\ln 2+\gamma_{E})-5$ with $\gamma_{E}=0.5772$ is the {\it Euler-Mascheroni constant}, and slow-roll parameters $(\epsilon_{V},~\eta_{V})$ are given by in terms of the inflationary potential $V(\phi)$, which can be expressed as: \begin{eqnarray}\label{ra1} \epsilon_{V}=\frac{M^{2}_{p}}{2}\left(\frac{V^{\prime}}{V}\right)^{2}\,,~~~~~~~ \label{ra2} \eta_{V}={M^{2}_{P}}\left(\frac{V^{\prime\prime}}{V}\right)\,. \end{eqnarray} We would also require two other slow-roll parameters, $(\xi^{2}_{V},\sigma^{3}_{V})$, in our analysis, which are given by: \begin{eqnarray}\label{ja1} \xi^{2}_{V}=M^{4}_{p}\left(\frac{V^{\prime}V^{\prime\prime\prime}}{V^{2}}\right)\,,~~~~~~~~ \label{ja2} \sigma^{3}_{V}=M^{6}_{p}\left(\frac{V^{\prime 2}V^{\prime\prime\prime\prime}}{V^{3}}\right)\,. \end{eqnarray} Note that we have neglected the contributions from the higher order slow-roll terms, as they are sub-dominant at the leading order. With the help of \be\label{con1} \frac{d}{d\ln k}=-M_p\frac{\sqrt{2\epsilon_{H}}}{1-\epsilon_{H}}\frac{d}{d\phi}\,\approx -M_p\frac{\sqrt{2\epsilon_{V}}}{1-\epsilon_{V}}\frac{d}{d\phi}\,,% \ee we can derive a simple expression for the tensor-to-scalar ratio, $r$, as:~\footnote{ We have derived some of the key expressions in an Appendix, see for instance, Eq.~(\ref{para 21e}), which we would require to derive the above expression, Eq.~(\ref{con2}). } \be\label{con2} r=\frac{8}{M^{2}_{p}}\frac{(1-\epsilon_{V})^{2}\left[1-({\cal C}_{E}+1)\epsilon_{V}\right]^{2}}{\left[1-(3{\cal C}_{E}+1)\epsilon_{V} +{\cal C}_{E}\eta_{V}\right]^{2}}\left(\frac{d\phi}{d{\ln k}}\right)^{2}. \ee We can now derive a bound on $r(k)$ in terms of the momentum scale: \be\begin{array}{llll}\label{con4} \displaystyle \int^{{ k}_{\star}}_{{k}_{e}}\frac{dk}{k}\sqrt{\frac{r({k})}{8}} \\ \displaystyle = \frac{1}{M_p}\int^{{\phi}_{\star}}_{{\phi}_{e}}d {\phi}\frac{(1-\epsilon_{V})\left[1-({\cal C}_{E}+1)\epsilon_{V}\right]}{\left[1-(3{\cal C}_{E}+1)\epsilon_{V} +{\cal C}_{E}\eta_{V}\right]},\\ \displaystyle \approx \frac{1}{M_p}\int^{{\phi}_{\star}}_{{\phi}_{e}}d {\phi}(1-\epsilon_{V})\left[1+{\cal C}_{E}(2\epsilon_{V}-\eta_{V})+....\right],\\ \displaystyle \approx \frac{\Delta\phi}{M_p} \left\{ 1+\frac{1}{\Delta\phi}\left[(2{\cal C}_{E}-1)\int^{{\phi}_{\star}}_{{\phi}_{e}}d {\phi}~\epsilon_{V} -{\cal C}_{E}\int^{{\phi}_{\star}}_{{\phi}_{e}}d {\phi}~\eta_{V}\right]+....\right\}\,, \end{array}\ee where note that $\Delta\phi \approx \phi_{\star}-\phi_{e}>0$ is positive in Eq.~(\ref{con4}), and $\phi_e$ denotes the inflaton VEV at the end of inflation, and $\phi_{\star}$ denote the field VEV when the corresponding mode $k_\star$ is leaving the Hubble patch during inflation. Note that $\Delta\phi>0$ implies that the left hand side of the integration over momentum within an interval, $k_{e}<k<k_{\star}$, is also positive, where $k_e$ represents corresponding momentum scale at the end of inflation. Individual integrals involving $\epsilon_V$ and $\eta_V$ are estimated in an Appendix, see Eqs.~(\ref{hj1}) and (\ref{hj2}). In order to perform the momentum integration in the left hand side of Eq~(\ref{con4}), we have used the running of $r(k)$, which can be expressed as: \be\label{con5} r(k)=r(k_{\star})\left(\frac{k}{k_{\star}}\right)^{a+\frac{b}{2}\ln\left(\frac{k}{k_{\star}}\right) +\frac{c}{6}\ln^{2}\left(\frac{k}{k_{\star}}\right)+....}\,, \ee where \be a=n_{T}-n_{S}+1,~~~b=\left(\alpha_{T}-\alpha_{S}\right),~~~c=\left(\kappa_{T}-\kappa_{S}\right)\,. \ee These parameterisation characterises the spectral indices, $n_S,~n_T$, running of the spectral indices, $\alpha_S,~\alpha_T$, and running of the running of the spectral indices, $\kappa_S,~\kappa_T$. Here the subscripts, $(S,~T)$, represent the scalar and tensor modes. Now substituting the explicit form of the potential stated in Eq.~(\ref{rt1a}), we can evaluate the crucial integrals of the first and second slow-roll parameters ($\epsilon_{V},~\eta_{V}$) appearing in the right hand side of Eq.~(\ref{con4}). For the details of the computation, see appendix. In the present context the potential dependent slow-roll parameters: ($\epsilon_{V},~\eta_{V}$), satisfy the joint (Planck+WMAP-9) constraints, which imply that~\cite{Planck-infl}: \be \epsilon_V<10^{-2}\,,~~~5\times 10^{-3}<|\eta_{V}|<0.021\,, \ee for which the inflationary potential is concave in nature. In the next section, we will discuss model independent bounds on the coefficients ($V(\phi_\star),V^{\prime}(\phi_\star),\cdots$) for a generic sub-Planckian VEV inflationary setup, for which we will satisfy the joint constraints from: (Planck+WMAP-9+high~L+BICEP2), where $r(k_\star)=0$ is disfavoured at $7\sigma$ CL., see Eq.~(\ref{obscons}). In this paper we have obtained the most general expressions for $r(k_\star),~|\Delta \phi|/M_p$ in terms of $V(\phi_\star)$ and $H(k_\star)$ for a sub-Plackian excursion of inflation by taking into account of higher order slow roll conditions. In order to satisfy the observational data from (Planck+WMAP-9+high~L+BICEP2), see Eq.~(\ref{obscons}), one requires a non-monotonic evolution of $\epsilon_V$ parameter if one wants to build a model of inflation with a sub-Planckian VEV, as pointed out in Refs.~\cite{BenDayan:2009kv,Hotchkiss:2011gz}. In this paper we have reconstructed the potential around $\phi_0$ for $V(\phi_0),~V^{\prime}(\phi_0),~V^{\prime\prime}(\phi_0),~V^{\prime\prime\prime}(\phi_0)$ and $V^{\prime\prime\prime\prime}(\phi_0)$. In order to satisfy the current observational constraints, $0.15<r_{\star}<0.27$, we have found the scale of inflation to be within: $2.07\times10^{16}~{\rm GeV}\leq\sqrt[4]{V_{\star}}\leq 2.40\times 10^{16}~{\rm GeV}$, for the VEV of inflaton varying within: $ 0.066 \leq\frac{\left |\Delta\phi\right|}{M_p}\,\leq 0.092$. We have also estimated a new set of {\it consistency relationships} for sub-Planckian model of inflation. One particular discriminator between sub-vs-super Planckian field excursion is the consistency relationships, in particular the slope of the tensor modes, i.e. $n_T=-(r/8)(2-(r/8)-n_s)< 0$ for $n_s=0.96$, see Eq.~(\ref{wq6}). This is in contrast with the super-Planckian excursion models of inflation, where typically one would expect $n_T=-(r/8)$. Furthermore, if the data could be refined to constrain $\alpha_T$, then this would really seal the status of sub-vs-super Planckian models of inflation. We have also found that choice of $\vartheta$ is rather insensitive to constraining the model parameters. Typically, in particle physics, the nature and shape of the potential will not {\it just} be a single monomial. In principle the potential could contain quadratic, cubic and quartic interactions for a renormalizable theory, or even higher order non-renormalizable terms arising from integrating out the heavy degrees of freedom, see~\cite{Mazumdar:2010sa}. In this respect our results are important for reconstructing a particle physics motivated model of inflation in a successful way. | 14 | 3 | 1403.5549 |
1403 | 1403.7857_arXiv.txt | Active galactic nuclei (AGNs) and gamma-ray bursts (GRBs) are powerful astrophysical events with relativistic jets. In this Letter the broadband spectral properties are compared between GRBs and the well-observed blazars. The distribution of GRBs are consistent with the well-known blazar sequence including the $\nu L_\nu(5\rm GHz)-\alpha_{\rm RX}$ and $\nu L_\nu(5\rm GHz)-\nu_{\rm peak}$ correlations, where $\alpha_{\rm RX}$ is defined as the broadband spectral slope in radio-to-X-ray bands, and $\nu_{\rm peak}$ is defined as the spectral peak frequency. Moreover, GRBs occupy the low radio luminosity end of these sequences. These two correlations suggest that GRBs could have a similar radiation process with blazars both in the prompt emission and afterglow phases, i.e., synchrotron radiation. | Gamma-ray bursts (GRBs) and active galactic nuclei (AGNs) are both powered by relativistic jets from accreting black holes (Gehrels et al. 2009; Urry \& Padovani 1995). The central engines of GRBs are argued to be stellar-mass black holes (Woosley 1993) and for AGNs the central engines are supermassive black holes. GRBs are most powerful explosions with isotropic-equivalent energy $E_{\rm iso}\sim 10^{50}-10^{55}$ erg in the universe (Zhang 2011), and can be detected out to very high-redshift universe (Lamb \& Reichart 2000; Wang et al. 2012). So GRBs can probe high-redshift universe, including dark energy (Dai, Liang \& Xu 2004; Schaefer 2007; Wang \& Dai 2011). Blazars include two subtype of AGNs, i.e., flat-spectrum radio quasars (FSRQs) and BL Lac objects (BL Lacs). A subclass of AGNs, e.g., super-Eddington accreting supermassive black holes, are also proposed to be a standard candle (Wang et al. 2013). The radiation mechanism of balzars is well constrained. The spectral energy distribution (SED) of blazars is well understood, including the low-energy (infrared-soft X-ray) bump and the high-energy (MeV-GeV) bump. The synchrotron radiation can account for the low-energy peak, while the MeV-GeV peak is produced by inverse Compton radiation. But for GRBs, the radiation mechanism for the prompt emission is still highly debated. The spectrum of prompt emission can be modeled by the ``Band function'' (Band et al. 1993), whose origin is still unknown (but see Lucas Uhm \& Zhang 2013). Some studies (M\'{e}sz\'{a}ros et al. 1994; Daigne \& Mochkovitch 1998) proposed that synchrotron radiation is the leading mechanism. Other mechanisms are also proposed (Pe'er et al. 2006; Rees \& M\'{e}sz\'{a}ros 2005; Beloborodov 2010). The radiation mechanism of afterglows is well understood (Sari et al. 1998). The observed afterglow radiation is well explained by synchrotron radiation (Sari et al. 1998; Panaitescu \& Kumar 2001). Some studies (Zhang 2007; Wang \& Dai 2013; Wang et al. 2014) have proposed that the mechanisms in different scale outflow or jet systems may be the same. Some works have been done on comparison between GRBs and AGNs. Wang \& Wei (2011) compared the spectral properties of blazars and optically bright GRB afterglows, and found that GRB afterglows have the same radiation mechanism as BL Lac objects. A similar correlation of the synchrotron luminosity and Doppler factor between GRBs and AGNs has been found (Wu et al. 2011). Nemmen et al. (2012) suggested that the relativistic jets in AGNs and GRBs have a similar energy dissipation efficiency. Ma et al. (2014) extended the analysis of Nemmen et al. (2012) by adding X-ray binaries and low-luminosity AGNS. Wang \& Dai (2013) found that the GRB X-ray flares and solar X-ray flares have similar distributions, which indicate that the X-ray flares of GRBs are due to a magnetic reconnection process. These similar distributions also exist in X-ray flares from black hole systems with $10^6-10^9M_\odot$ (Wang et al. 2014). Zhang et al. (2013) found that the prompt emission of GRBs may be produced by magnetic dominated jets. In this Letter we compare the broadband spectral properties of GRBs and blazars, including the $\nu L_\nu(5{\rm GHz})-\alpha_{RX}$ and $\nu L_\nu(5{\rm GHz})-\nu_{\rm peak}$ correlations, where $\alpha_{\rm RX}$ is the radio-to-X-ray spectral slope. For a GRB, $\nu_{\rm peak}$ is the peak frequency of $\nu f_\nu$ spectrum of prompt emission, while $\nu_{\rm peak}$ is the low peak of $\nu f_\nu$ spectrum for a blazar. The aim of this Letter is to explore a possible similarity in radiation mechanism between GRBs and blazars. this Letter is organized as follows. In section 2, we present the sample of blazars and GRBs. The fitting results are given in section 3. Section 4 gives conclusions and discussions. | The physics of GRBs are poorly understood, i.e., the radiation mechanism of prompt emission, the value of Lorentz factor, jet composition, and central engine (Zhang 2011). Wang \& Dai (2013) found similar frequency distributions between X-ray flares of GRBs and solar X-ray flares, which may indicate the magnetically dominated jets in GRBs. In this Letter we compile 43 GRBs with well X-ray and radio observations. Two new correlations between GRBs and blazars may provide a new clue as to the radiation mechanism of GRB prompt emission and afterglows. For example, our clear $\nu L_\nu(5{\rm GHz})-\alpha_{\rm RX}$ and $\nu L_\nu(5{\rm GHz})-\nu_{\rm peak}$ correlations suggest that the radiation mechanism of GRBs in prompt and afterglow phases and blazars is similar, namely, synchrotron radiation. Moreover, GRBs are occupy the low-luminosity region of these correlations. In this Letter, we use the radio luminosities during the GRB afterglow phase. Although some models predict that bright radio emission may be generated within about 10\,s of the initial explosion of a GRB (Usov \& Katz 2000; Sagiv \& Waxman 2002; Shibata et al. 2011), but a detection of prompt radio emission has some impediments, such as scattering (Macquart 2007; Lyubarsky 2008). Bannister et al. (2012) have searched for prompt radio emission from nine GRBs at 1.4 GHz, and found single dispersed radio pulses with significance $>6 \sigma$ in a few minutes following two GRBs. Unfortunately, the probability of GRB origin is only 2\%. There has been no confirmed evidence for detection of GRB prompt radio emission up to now. So we use the radio emission of an afterglow in this Letter. Theoretically, internal shocks produce the GRB prompt emission, and external shocks produce the afterglow emission. The radiation mechanism of an afterglow is well constrained, i.e., synchrotron emission. But the prompt emission related to $\nu_{peak}$ is not well understood. If some similar correlations between $\nu L_\nu(5{\rm GHz})$ and $\nu_{peak}$ exist in GRBs and blazars, the radiation mechanism of prompt emission is argued as synchrotron emission. Until now, the location of blazar gamma-ray emission regions are still uncertain (Marscher et al. 2010), since some theories locate blazar gamma-ray emission regions close to the black hole/accretion disk (Blandford \& Levinson 1995) while the others place them at parsec scales in the radio jet (Jorstad et al. 2001). The X-ray emission from kiloparsec-scale blazar jets has been observed (Harris \& Krawczynski 2006). Meanwhile, the radio emission region of blazars spans kiloparsec scale. Kharb et al. (2010) found that a few blazars exhibit only radio core emission. But the X-ray emission region and radio emission region do not fully overlap. So the high-energy and radio emissions of blazars also originate from different regions. So a comparison of the correlation $\nu L_\nu(5{\rm GHz})-\nu_{\rm peak}$ between GRBs and the blazar sequence is reasonable. Liang et al. (2004) found that the peak energy $\nu_{\rm peak}$ evolves with isotropic-equivalent luminosity from the time-resolved spectra. They also found that the $L_{iso}-\nu_{\rm peak}$ correlation also holds for time-resolved spectra and time-integrated spectra. Ghirlanda et al. (2010) studied the time-resolved spectra of Fermi GRBs. The peak energy $\nu_{\rm peak}$ correlates with the luminosity within individual bursts (Ghirlanda et al. 2010). Moreover, the time-resolved $L_{iso}-\nu_{\rm peak}$ correlation is very similar for all the bursts and has a slope similar to the correlation defined by the time-integrated spectra of different bursts detected by several different satellites. The time-integrated value of $\nu_{\rm peak}$ is widely used in luminosity correlations of GRBs, for example $\nu_{\rm peak}-E_{iso}$ (Amati et al. 2002), $\nu_{\rm peak}-L_{iso}$ (Yonetoku et al. 2004) and $\nu_{\rm peak}-E_{\gamma,jet}$ (Ghirlanda et al. 2004) correlations. So if some correlation is due to a similar physical mechanism, this correlation holds no matter whether the time-resolved or time-integrated values are used. Fossati et al. (1998) found that $\nu_{\rm peak}$ is anti-correlated with the synchrotron peak luminosity for blazars. For GRBs, Liang et al. (2004) found that $\nu_{\rm peak}$ is positively correlated with the isotropic-equivalent luminosity (about total luminosity). But the luminosity in the correlation for blazars is the synchrotron peak luminosity, not the total luminosity. Because there are two peaks in a blazar spectral energy distribution and there is no correlation between $E_{iso}$ and $\nu L_\nu(5{\rm GHz})$ in GRBs (Chandra \& Frail 2012), from our simple analysis above we cannot conclude that GRBs have a different $\nu L_\nu(5{\rm GHz})-\nu_{\rm peak}$ correlation compared with that for blazars. In this paper, we find that GRBs occupy the low radio luminosity end of the blazar sequence, which is similar to that of Wang \& Wei (2011). | 14 | 3 | 1403.7857 |
1403 | 1403.2475_arXiv.txt | We performed a hydrodynamical cosmological simulation of the formation of a Milky Way-like galaxy in a warm dark matter (WDM) cosmology. Smooth and dense filaments, several co-moving mega parsec long, form generically above $z\sim 2$ in this model. Atomic line cooling allows gas in the centres of these filaments to cool to the base of the cooling function, resulting in a very striking pattern of extended Lyman-limit systems (LLSs). Observations of the correlation function of LLSs might hence provide useful limits on the nature of the dark matter. We argue that the self-shielding of filaments may lead to a thermal instability resulting in star formation. We implement a sub-grid model for this, and find that filaments rather than haloes dominate star formation until $z\sim 6$ although this depends on how stars form in WDM. Reionisation decreases the gas density in filaments, and the more usual star formation in haloes dominates below $z\sim 6$, although star formation in filaments continues until $z=2$. Fifteen per cent of the stars of the $z=0$ galaxy formed in filaments. At higher redshift, these stars give galaxies a stringy appearance, which, if observed, might be a strong indication that the dark matter is warm. | \label{sect:intro} The $\Lambda$ cold dark matter (CDM) model has had tremendous success in providing a cosmogony that links the state of the very high-redshift Universe as inferred from cosmic microwave background observations and Big Bang nucleosynthesis considerations, with the observed large-scale distribution of galaxies and its evolution at later times. In this paradigm, the observed structures grew due to gravitational amplification of initially small perturbations, perhaps seeded by inflation \citep{Guth81}, with most of the gravitating mass in the form of \lq dark matter\rq, an as of yet unknown type of matter that apart from gravitationally, interacts only very weakly if at all with baryonic matter and radiation, see e.g. \cite{FrenkWhite12} for a recent review. Observations on smaller, sub galactic scales, have proven more problematic for $\Lambda$CDM, with suggestions that the low abundance \citep[e.g][]{Klypin99, Moore99} and shallow density profiles \citep[e.g][]{Gilmore2007, de_vega2010} inferred for haloes hosting Milky Way satellites are inconsistent with the numerous substructures with cuspy density profiles that form in Milky Way-like CDM haloes. Indeed the substructures in the haloes of the Aquarius project \citep{Springel08b} and the GHALO project \citep{Stadel09} may be difficult to reconcile with those inferred to host Milky Way satellites. Even if many of the smaller DM substructures may not have a sufficiently deep potential well to form stars after reionization \citep[e.g][] {Efstathiou92, Benson02} and so may well be dark, some of the more massive ones are probably too big to be affected \citep{Okamoto08} and would hence appear to be \lq too big to fail\rq\ \citep{Boylan-Kolchin11,Boylan-Kolchin12}. However the jury on this is still out: the Milky Way's halo may simply be of lower mass than those of the Aquarius haloes and hence have fewer massive satellites \citep{Wang12}. The status of the {\em density} profiles of the satellite's haloes - cored versus cuspy - is similarly unresolved. \cite{Strigari10} claim that the stellar dynamics observations of the satellites do not imply cores at all and hence may be consistent with CDM cusps. But even if the satellites had cored profile, this might result from the action of baryonic feedback processes \citep{Pontzen12, Governato12}, and hence still be consistent with CDM. Irrespective of these astrophysical considerations, particle physics has several candidates for the dark matter. A popular kind is a super symmetric particle with negligible intrinsic velocity dispersion \citep{Bertone2005}, which would qualify as \lq cold\rq\ DM. However other viable candidates have considerable intrinsic velocities, and these would constitute \lq warm\rq\ dark matter (WDM), for example a keV scale gravitino or a sterile neutrino \citep[e.g.][]{Dodelson94}. Such intrinsic velocities - as opposed to velocities induced by gravity - have two (related) effects on the formation of structure. Firstly, the motion of WDM particles quenches the growth of structure below a \lq free-streaming scale\rq, which we loosely define here as the the maximum distance over which such a particle can travel. As a consequence, WDM haloes have far less substructures with masses below the corresponding free-streaming mass - potentially alleviating the \lq missing satellites\rq\ problem discussed before \citep{Bode01}. Secondly, the intrinsic velocities of WDM particles imply a finite phase-space density - which is conserved during halo growth \citep[e.g.][]{Tremaine79, Hogan00}. This finite phase space density might be the origin of the cores inferred to exist in Milky Way satellite haloes - if indeed they are cored. Even though these considerations stimulated much of the astronomical interest in WDM, detailed analyses yielded slightly disappointing results. \cite{Shao13} demonstrated that - although WDM haloes are indeed cored - cores as large as suggested by Milky Way satellite observations will only form if the WDM free-steaming lengths is so large that the satellites themselves would fail to form, resulting in a a \lq too small to succeed\rq\ problem (see also \cite{Maccio13}). \cite{Lovell12} demonstrated that WDM does help with the satellites' profiles: since the haloes hosting satellites form later in WDM than in CDM on average, they tend to have lower central densities - alleviating the \lq too big to fail\rq\ problem. Requiring that enough satellites form leads to a conservative lower limit to the (thermal equivalent) WDM particle mass of ~1.5~keV \citep{Lovell13}, not far from the lower limits inferred from the Ly$\alpha$ forest by \cite{Boyarsky09}. Consequently there is still a window for WDM to have some effect on galaxy formation, although it seems that by itself it will not resolve the issues with the \lq too big to fail\rq\ problem, and the alleged presence of cores in satellite haloes. The study of \cite{Gao07b} points out that the formation of the first stars could proceed very differently in a WDM Universe. The reason is that the filaments that form naturally in hierarchical models have potential wells that are so deep that gas in them forms molecular hydrogen and cools. Such filaments also form in CDM, yet there they break-up into star forming mini haloes due to CDM's small-scale power. Numerical fragmentation who's origin is discussed by \cite{Wang07} prevented \cite{Gao07b} from making firm predictions about the nature of first star formation in WDM, but they argue that star formation may be very efficient and result in a range of stellar masses, which is quite different from the inefficient formation of a \lq single\rq\ massive star expected to form in CDM mini haloes \citep[e.g.][]{Abel02}. Here we present results from zoomed cosmological hydrodynamical simulations that follow the formation of a Milky Way-like galaxy in WDM. In particular we want to test whether the filaments that form around such objects could host star formation initiated by {\em atomic} line cooling, which would be the galaxy formation analogue of first star formation in WDM discussed by \cite{Gao07b}. We introduce the simulations in Section~2 and present our results in Section~3, Section~4 summarises. | We have performed cosmological hydrodynamical simulations of the formation of a Milky Way-like galaxy in a warm dark matter (WDM) scenario. The model we investigate uses a CDM transfer function exponentially cut-off below the free-streaming scale of a WDM particle which is equivalent to that of a $1.5$~keV thermal relic. The simulation includes radiative cooling from hydrogen and helium, inverse Compton cooling off the CMB, and thermal bremsstrahlung, in the presence of an imposed uniform optically thin UV/X-ray background, but ignores cooling from metals and from molecules. The simulation uses a very simple sub-grid model for star formation and neglects feedback from star formation. We examined to what extent star formation in the dense filaments that are characteristic for this WDM model contributes to the final redshift $z=0$ stellar mass and its build-up. At very high redshifts, $z\gsim 8$ say, dense and extended filaments several co-moving mega parsecs long, form before dark matter haloes themselves appear. Gas in those filaments is cold ($T\sim 10^4$~K, the base of the cooling function in the absence of metals and molecules) and dense ($n_{\rm H}\gsim 0.1$cm$^{-3}$) enough to form stars in our sub-grid model for star formation. Filaments continue to dominate star formation, with gas accreting onto them at high speed ($\sim 150$~km~s$^{-1}$), where it shock heats and subsequently cools radiatively through atomic line cooling. The column density of gas through these filaments is very high ($\gsim 10^{18}$~cm$^{-2}$), and the presence of very long and narrow Lyman-limit systems (LLS) in WDM is very striking. It might be possible to test observationally for the presence or absence of such WDM filamentary LLS, for example by studying the LLS correlation function, and hence constrain the nature of the dark matter particle. Stars formed in filaments drain in haloes. Reionisation - in these simulations assumed to occur at $z=6$ - causes the gas density in filaments to decrease, and star formation in haloes starts to dominate; however stars continue to form in filaments until $z\sim 2$, after which their column density drop below the threshold for star formation. By $z=0$, 15 per cent of stars in the final galaxy formed in filaments. We stressed that there is no accepted theory of how stars form in filaments, nor has it been investigated if and how star formation in filaments is regulated by feedback. In addition, our simulation suffers from the well known artificial fragmentation of the dark matter. Consequently we are hesitant to draw quantitative conclusions from this run. However what the simulation demonstrates unambiguously is that in this WDM model, very long and dense filaments form around Milky Way-like proto-galaxies, which would be observable as LLS or DLAs. Gas in these filaments cools radiatively through atomic line cooling, and can shield itself from the UV-background. The thermal instability that results from this is likely to lead to star formation, even at lower redshifts $z\sim 2$. In the simulation, this results in the appearance of curiously shaped stringy \lq chain\rq\ galaxies, which, if observed, might be a strong indication that the dark matter is warm. | 14 | 3 | 1403.2475 |
1403 | 1403.6898_arXiv.txt | We present Suzaku results of the two Galactic supernova remnants (SNRs), G350.1$-$0.3 and G349.7$+$0.2. We find Al and Ni K$\alpha$ lines from both the SNRs for the first time, in addition to previously detected K-shell lines of Mg, Si, S, Ar, Ca and Fe. The spectra are well described by two optically thin thermal plasmas: a low-temperature (low-$kT$) plasma in collisional ionization equilibrium and a high-temperature (high-$kT$) plasma in non-equilibrium ionization. Since the low-$kT$ plasma has solar metal abundances, it is thought to be of interstellar medium origin. The high-$kT$ plasma has super-solar abundances, hence it is likely to be of ejecta origin. The abundance patterns of the ejecta components are similar to those of core-collapse supernovae with the progenitor mass of $\sim$ 15--25~\MO\ for G350.1$-$0.3 and $\sim$ 35--40~\MO\ for G349.7$+$0.2. We find extremely high abundances of Ni compared to Fe ($Z_{\rm Ni}/Z_{\rm Fe} \sim8$). Based on the measured column densities between the SNRs and the near sky background, we propose that G350.1$-$0.3 and G349.7$+$0.2 are located at the distance of $9\pm3$~kpc and $12\pm5$~kpc, respectively. Then the ejecta masses are estimated to be $\sim$13 \MO\ and $\sim$24~\MO\ for G350.1$-$0.3 and G349.7$+$0.2, respectively. These values are consistent with the progenitor mass of $\sim$ 15--25~\MO\ and $\sim$ 35--40~\MO\ for G350.1$-$0.3 and G349.7$+$0.2, respectively. | G350.1$-$0.3 is a radio-bright supernova remnant (SNR) in the Galaxy \citep{Clark1973}. The radio morphology is not a typical shell or crab-like but has a distorted and elongated shape \citep{Salter1986}. X-rays were detected with ROSAT \citep{Voges1999} and ASCA \citep{Sugizaki2001}, and then with XMM-Newton \citep{Gaensler2008} and Chandra \citep{Lovchinsky2011}. A point-like X-ray source, XMMU\,J172054.5$-$372652, was found $\sim\timeform{3'}$ west of the brightest region. The spectral parameters are in the range of a typical central compact object (CCO) \citep{Lovchinsky2011}. \citet{Gaensler2008} detected a $^{12}$CO emission along the eastern edge of the SNR. They suggest that the molecular gas suppressed the expansion of the remnant and formed the peculiar asymmetric morphology. The distance and age are estimated to be 4.5--10.7~kpc and $\sim900$ years old, respectively \citep{Gaensler2008}. The X-ray spectrum of the SNR is reproduced by a two-component model: a high-temperature ($kT\sim1.5$~keV) plasma in non-equilibrium ionization (NEI) and a low-temperature ($kT\sim0.4$~keV) plasma in collisional ionization equilibrium (CIE) \citep{Gaensler2008}. The abundances of the former are super-solar and those of the latter are 1~solar, which suggests ejecta and interstellar medium (ISM) origin, respectively. The metal abundances of the ejecta are as high as 10~solar although the statistical errors are quite large. \citet{Lovchinsky2011}, on the other hand, reported that the spatially resolved spectra can be reproduced by one-temperature plasma models with abundances of $\sim1$--9~solar. G349.7$+$0.2 is another radio-bright SNR in the Galaxy \citep{Shaver1985}. The distance is estimated to be $18.3\pm 4.6$~kpc \citep{Caswell1975}. OH maser emission (1720 MHz) is found toward G349.7$+$0.2 at a radial velocity of $\sim+16$~km~s$^{-1}$, suggesting that the SNR is interacting with a dense molecular cloud at the kinematic distance of 22.4~kpc \citep{Frail1996}. The X-ray image taken by Chandra shows an irregular shell with the bright eastern side \citep{Lazendic2005}. The presence of H\emissiontype{I} clouds near the SNR indicates that G349.7$+$0.2 is evolved into the intercloud medium, and is responsible for the irregular morphology. Like G350.1$-$0.3, the X-ray spectrum is described by two plasmas with different temperatures: a low-temperature ($kT\sim0.8$~keV) CIE plasma with solar abundances, and a high-temperature ($kT\sim1.4$~keV) NEI plasma. The latter has an enhanced Si abundance, suggesting ejecta origin \citep{Lazendic2005}. A point source, CXOU\,J171801.0$-$372617, is found near the center of the SNR, possibly a CCO associated with the SNR \citep{Lazendic2005}. The previous results such as the presence of CCOs and the associations with molecular clouds suggest that both of G350.1$-$0.3 and G349.7+0.2 are core-collapse (CC) SNRs. The metal abundances in the ejecta should provide crucial information for the mass of the progenitor stars. Fe and Ni are the final products of the major nuclear reaction network in the evolution of massive stars and their final supernova (SN) explosions. Therefore, these elements should be particularly important to study the mechanism of CC SNe in the close vicinity of the core region. Previous works, however, have limited statistics to study the ejecta elements. This paper presents the most accurate Fe and Ni abundances in the two CC-SN candidates, G350.1$+$0.3 and G349.7$-$0.2. For the studies, we used the Suzaku satellite \citep{Mitsuda2007} because it has the highest sensitivity for diffuse X-rays in the Fe and Ni K-shell band at 5--10~keV. In this paper, we estimate errors at 90\% confidence level while figure~2, 3 and 6 show the 1~$\sigma$ errors. | The X-ray spectra of the two SNRs are well explained by the two plasma model; a high-temperature in NEI and a low-temperature in CIE. For both the SNRs, we obtain the abundances of many heavy elements in the high-temperature NEI plasma. The most important discovery is the detection of Ni with the extreme overabundance ($12\pm7$~solar for G350.1$-$0.3 and $5.3\pm2.0$~solar for G349.7$+$0.2). We also find Al for the first time from G350.1$-$0.3 and G349.7$+$0.2, the second detection after G344.7$-$0.1 \citep{Yamaguchi2012}. \subsection{Origin of the Plasmas} \subsubsection{G350.1$-$0.3}\label{G350_origin} Since the low-temperature component for G350.1$-$0.3 is in CIE with the solar abundances, it would be ISM heated by a blast wave. The high-temperature NEI component for the SNR has high metal abundances of 1.4--12~solar. Therefore, it is likely the ejecta recently heated-up by a reverse shock. In figure~\ref{fig:pattern}, we show metal abundances in the ejecta relative to Si for G350.1$-$0.3. The determination of Mg and Al abundances in the ejecta may be significantly affected by the high inferred Ni abundance, because Ni L-shell lines become important near Mg K$\alpha$ and Al K$\alpha$ line energies. Also the quality of atomic data for Ni lines in the NEI models is not good enough. Thus, the quoted errors on the Mg and Al abundances may be larger than the pure statistical error. Taking into account of possible larger errors in Mg and Al than those given in figure~\ref{fig:pattern}, the abundance patterns roughly agree with those of the CC-SN model with a progenitor mass between 15--25~\MO \ \citep{Woosley1995}. \begin{figure}[tb] \begin{center} \FigureFile(80mm,70mm){figure4.eps} \caption{Metal abundances in the ejecta of G350.1$-$0.3 relative to Si as a function of atomic number. The dotted lines represent CC models with main sequence masses of 15~\MO, 20~\MO \ and 25~\MO\ \citep{Woosley1995}.} \label{fig:pattern} \end{center} \end{figure} \subsubsection{G349.7$+$0.2}\label{G349_origin} Like G350.1$-$0.3, the low-temperature component for G349.7$+$0.2 is in CIE with the solar abundances. Thus, this component is also likely ISM heated by a blast wave. For the high-temperature component, Mg and Ni abundances are much higher than the solar values, suggesting an ejecta component. However, as we noted in \ref{G350_origin}, the Mg abundance would have a larger error. Hence, compared to the case of G350.1$-$0.3, it would be less convincing that the high-temperature component of G349.7$+$0.2 is also an ejecta origin. Still, we show the abundance pattern of G349.7$+$0.2 in figure~\ref{fig:pattern2} comparing those of the CC-SN model with a progenitor mass between 35--40~\MO \citep{Woosley1995}. From figure~\ref{fig:pattern2}, we see the abundance pattern of the SNR roughly agrees with that of a progenitor mass of $\sim$35--40~\MO. The metal abundances other than Mg and Ni are 1~solar or slightly smaller, which may conflict with the initial assumption of the ejecta origin. The abundances are determined by fixing the abundances for lighter elements than Mg, namely He, C, N, O and Ne to be 1 solar since He--Ne do not appear as emission lines in the relevant energy band of $<1.2$~keV. If the plasma is really due to the ejecta of 35--40~\MO\ star, the abundance of He--Ne should be far larger than 1 solar and the bremsstrahlung is largely dominated by the enhanced He--Ne. We, hence, assume the abundances of these light elements following the results of \citet{Woosley1995}, and re-fit the spectra. The resultant abundances of Mg--Ni become 4.5--5.5 times of the initial values, or larger than 1~solar, which supports the ejecta origin. The abundance ratios relative to Si are not changed from that of the original data given in figure~\ref{fig:pattern2}. \begin{figure}[tb] \begin{center} \FigureFile(80mm,70mm){figure5.eps} \caption{Metal abundances in the ejecta of G349.7$+$0.2 relative to Si as a function of atomic number. The dotted lines represent CC models with main sequence masses of 35~\MO \ and 40~\MO \ \citep{Woosley1995}.} \label{fig:pattern2} \end{center} \end{figure} \subsection{Ni Over-Abundance of G350.1$-$0.3 and G349.7$+$0.2} The observed abundance of Ni are far higher than that of Fe for both the SNRs. We note that the large abundances of Ni are not due to an estimation error of the NXB, which exhibits a strong neutral Ni K$\alpha$ line at 7.47~keV. We estimate for the case of FI CCDs (XIS0 + XIS3), for simplicity. Neutral Ni K$\alpha$ line flux in the NXB is $3.1(\pm 0.1)\times10^{-6}$~photons~s$^{-1}$~cm$^{-2}$, which is almost comparable to the He-like Ni K$\alpha$ + the Fe K$\beta$ lines flux in the NXB-subtracted spectra of the SNRs($\sim3.5\times10^{-6}$~photons~s$^{-1}$~cm$^{-2}$). Since the typical ambiguity of the NXB subtraction is at most 5\% \citep{Tawa2008}, a contamination of this line to the derived flux of the He-like Ni K$\alpha$ + Fe K$\beta$ lines would be less than a few \% . Furthermore, with the good energy resolution of Suzaku, we separately detect the He-like Ni K$\alpha$ + the Fe K$\beta$ lines at 7.7~keV from the neutral Ni K$\alpha$ line at 7.47~keV. The high ratio of $Z_{\rm Ni}/Z_{\rm Fe}~\sim8$ is not found from any other SNRs. In figure~\ref{SNR_comparison}, we compare simply the flux ratio of K$\alpha$ line of Fe and Ni for G350.1$-$0.3, G349.7$+$0.2, Tycho \citep{Yamaguchi2014}, Kepler \citep{Park2013}, and Cassiopeia A \citep{Maeda2009}. For G350.1$-$0.3, G349.7$+$0.2, and Cassiopeia A, only the sum of Ni K$\alpha$ and Fe K$\beta$ are available. We therefore estimated the Fe K$\beta$ flux by referring to \citet{Yamaguchi2014}, and obtain the Ni K$\alpha$ flux separately. Since the atomic numbers of Fe and Ni are nearly the same, the flux ratio of K$\alpha$ line of Ni and Fe is approximately equal, or slightly smaller (due to a smaller ionization/excitation cross section of Ni than those of Fe) than the abundance ratio. In the solar abundance, the abundance ratio Ni/Fe is $\sim$ 4\%. The flux ratios of K$\alpha$ line of Ni and Fe for Tycho, Kepler and Cassiopeia A are slightly smaller than $\sim$ 4\%, but those of G350.1$-$0.3 and G349.7$+$0.2 are larger than $\sim$ 4\%, indicating that the abundance ratio ($Z_{\rm Ni}/Z_{\rm Fe}$) is larger than 1 (in solar unit) in these SNRs. Thus Ni-overabundance for G350.1$-$0.3 and G349.7$+$0.2 can be suggested even before the spectral fitting. \begin{figure}[tb] \begin{center} \FigureFile(85mm,85mm){figure6.eps} \caption{Comparison with the flux ratio of K$\alpha$ lines between Fe and Ni. Blue points are those between Ni K$\alpha$ plus Fe K$\beta$ and Fe K$\alpha$ while red points are only for Ni K$\alpha$ and Fe K$\alpha$. The errors are at the 1$\sigma$ level. Tycho and Kepler data are derived from \citet{Yamaguchi2014} and \citet{Park2013}, respectively. The errors are not provided for Cassiopeia A \citep{Maeda2009}.} \label{SNR_comparison} \end{center} \end{figure} The high ratio $Z_{\rm Ni}/Z_{\rm Fe}~\sim8$ is not predicted by the theoretical model by \citet{Woosley1995}. The observed high ratio can be explained if a significant fraction of Ni was ejected from the core region possibly due to an asymmetric explosion. In fact, \citet{Maeda2007} reported that a large amount of Ni is ejected from the core of SN 2006aj as a result of an asymmetric explosion. G350.1$-$0.3 and G349.7$+$0.2 has the morphology away from symmetry. Previous researches claimed that the surrounding molecular gas caused the peculiar morphologies. Instead, we propose that asymmetric explosions made such morphologies. \subsection{$N_{\rm H}$, Distance and Ejecta Mass} The absorption ($N_{\rm H}$) of compact sources in the Galactic inner plane would be affected by the dust scattering effect. The observed radius of the dust scattering halo (which includes 90\% of the total flux) is $\sim \timeform{40"}$ \citep{Xiang2007} for 4U\,1624$-$49, an X-ray binary located at or behind the Galactic ridge with large absorption of $N_{\rm H} \sim 8\times10^{22}$~cm$^{-2}$ \citep{Smale2001}. Although the source sizes of G350.1$-$0.3 and G349.7$+$0.2 are larger than this radius, we still examine the dust scattering effect. We made the spectra of the SNRs from areas larger by \timeform{180"} in radius than those in figure~\ref{img} (solid lines) (e.g. for G349+0.2, the radius of the larger area is \timeform{320"}, while the original source area is \timeform{140"} radius.). The best-fit $N_{\rm H}$ of the spectra from the larger areas are $3.2(\pm 0.1)$ and $6.3(\pm 0.2)\times 10^{22}$~cm$^{-2}$ for G350.1$-$0.3 and G349.7$+$0.2, respectively, which are consistent with those given in table~\ref{tab:para}. Therefore, the dust scattering effect is not significant in the $N_{\rm H}$ estimation for these SNRs. For the distance estimation, we assume that the interstellar gas density is proportional to the stellar density of the Galactic disk given by \citet{Kent1991}. The ratio of the X-ray absorption column density $N_{\rm H}$ between the SNRs and nearby GRXE are 1.1 and 1.6 for G350.1$-$0.3 and G349.7$+$0.2, respectively (see tables \ref{tab:bg} and \ref{tab:para}). Integrating the gas density along the line of sight, we search for the distance, where the integrated gas density becomes to $N_{\rm H}$ (at 8.5 kpc) $\times$ $N_{\rm H}$ ratio (1.1 for G350.1$-$0.3 and 1.6 for G349.7$+$0.2). Here we assume $N\rm_H$ of the GRXE is that of the midpoint of the Galactic ridge along the line of sight (8.5~kpc). Then the distances are estimated to be $8.9\pm0.3$~kpc and $11.9\pm0.4$~kpc for G350.1$-$0.3 and G349.7$+$0.2, respectively. Since the stellar density model \citep{Kent1991} does not include local enhancement of interstellar media (e.g. the 3~kpc arms; \cite{Dame2008}), we make the IR extinction curves by \citet{Chen2013}\footnote{The on-line data of \citet{Chen2013} are limited in the distance of below 10~kpc and in the longitude below \timeform{10D} from the Galactic center. Dr. Chen kindly provided us with the data near at $l$=\timeform{350D} up to distance of $\sim14$~kpc.}, and re-estimate the distance with the same method as described above. Then the re-estimated distance of G350.1$-$0.3 is $9.4\pm0.4$~kpc, consistent with that taken from the stellar density model ($8.9\pm0.3$~kpc). No IR extinction curve is available at the position of G349.7$+$0.2. We therefore use one of the nearby data at $l\sim\timeform{350D}$, and obtain a distance of $12.7\pm0.6$~kpc, which is also consistent with that derived from the stellar density model ($11.9\pm0.4$~kpc). However, the near-by data show significant spatial variations, and are different from \citet{Marshall2006}. We estimate the distance variation using these data, and found the variations to be $\sim$2--3~kpc. Thus we regard the systematic distance error for G349.7$+$0.2 using the current IR extinction data is $\sim$2--3~kpc. The variation of the $N_{\rm H}$ obtained by the X-ray observations in the $\timeform{2D}< |l| < \timeform{10D}$, $ \timeform{-0.5D} < b < \timeform{0.5D}$ region is less than 30\% (90\% error) (H. Uchiyama, private communication). This would be another source of the distance uncertainty. Taking into account of all these possible systematic errors, we adopt the distances of G350.1$-$0.3 and G349.7$+$0.2 to be 9$\pm$3 and 12$\pm$5 kpc, respectively. The distance of G350.1$-$0.3 is consistent with, while that G349.7$+$0.2 is smaller than those of the previous reports \citep{Gaensler2008,Caswell1975,Frail1996}. We will estimate the ejecta masses for both the SNRs as below. As we mentioned in \ref{G349_origin}, we should deal with lighter elements which do not appear in the relevant energy band of $>$ 1.2~keV to estimate the physical condition of the plasma such as the emission measure ($EM$). Since the ejecta abundances are similar to those of 15--25~\MO\ and 35--40~\MO\ progenitor stars for G350.1$-$0.3 and G349.7$+$0.2, respectively, we assume that the abundances of elements lighter than Mg in the NEI component (ejecta) are those of the CC-SN model of a 20~\MO\ and 40~\MO\ progenitor \citep{Woosley1995} for G350.1$-$0.3 and G349.7$+$0.2, respectively and re-fit the spectra. As a result, the emission measures become $\sim$1/4 and $\sim$1/5, 5.4($\pm$0.5)$\times$10$^{11}$~cm$^{-5}$ and 1.7$(\pm$0.2)$\times$10$^{12}$~cm$^{-5}$ for G350.1$-$0.3 and G349.7$+$0.2, respectively. We then take the ratio between the electron and atomic hydrogen densities to be $n_{\rm e}/n_{\rm H} =$1.6 and 1.7 and the number ratio of all the nucleons to hydrogen to be 2.1 and 2.5 in the 20~\MO\ and 40~\MO\ progenitor, respectively. Assuming an oblate spheroid with major and minor radii of $\timeform{2.3'}$ and $\timeform{1.5'}$ for G350.1$-$0.3 \citep{Lovchinsky2011} and a sphere with a radius of $\timeform{1.2'}$ for G349.7$+$0.2 \citep{Lazendic2005}, the ejecta masses are estimated to be $\sim$13~$f^{1/2}~d^{5/2}_9$~\MO\ for G350.1$-$0.3 and $\sim$24~$f^{1/2}~d^{5/2}_{12}$~\MO\ for G349.7$+$0.2, where $f$ is a filling factor. $d_9$ and $d_{12}$ are the distance parameters in unit of 9~kpc and 12~kpc, respectively. These are roughly consistent with those estimated by the abundance patterns of 15--25~\MO\ and 35--40~\MO, respectively. | 14 | 3 | 1403.6898 |
1403 | 1403.4253_arXiv.txt | We present long-term (months-years) X-ray spectral variability of the Seyfert 1.8 galaxy NGC 1365 as observed by {\it Swift}, which provides well sampled observations over a much longer timescale (6 years) and a much larger flux range than is afforded by other observatories. At very low luminosities the spectrum is very soft, becoming rapidly harder as the luminosity increases and then, above a particular luminosity, softening again. At a given flux level, the scatter in hardness ratio is not very large, meaning that the spectral shape is largely determined by the luminosity. The spectra were therefore summed in luminosity bins and fitted with a variety of models. The best fitting model consists of two power laws, one unabsorbed and another, more luminous, which is absorbed. In this model, we find a range of intrinsic 0.5-10.0 keV luminosities of approximately $1.1 - 3.5$ ergs s$^{-1}$, and a very large range of absorbing columns, of approximately $10^{22} - 10^{24}$ cm$^{-2}$. Interestingly, we find that the absorbing column decreases with increasing luminosity, but that this result is not due to changes in ionisation. We suggest that these observations might be interpreted in terms of a wind model in which the launch radius varies as a function of ionising flux and disc temperature and therefore moves out with increasing accretion rate, i.e. increasing X-ray luminosity. Thus, depending on the inclination angle of the disc relative to the observer, the absorbing column may decrease as the accretion rate goes up. The weaker, unabsorbed, component may be a scattered component from the wind. | \label{intro} X-ray spectral observations have shown that variability in the column of absorbing material between the X-ray source and the observer is present in a number of Seyfert galaxies \citep{risaliti02}. The detection of variable absorption on a timescales of hours has indicated that the absorbing material must be close to the nucleus, at a distance similar to that of the Broad Emission Line Region \citep[e.g.][]{lamer,elvis04,puccetti}, with claims that complete occultations by Broad Line Region clouds have been observed on timescales of days \citep{risaliti07a}. NGC 1365 is a nearby Seyfert 1.8 galaxy \citep{maiolino95} which displays a large amount of X-ray spectral variability \citep{risaliti09} on timescales of hours to years \citep{brenneman}. These variations have been interpreted as the spectrum changing from being `transmission dominated' to `reflection dominated'. When the spectrum is `transmission dominated' \citep[e.g.][]{risaliti00} the absorbing material is Compton thin and the transmitted component dominates the spectrum; when the spectrum is `reflection dominated' \citep[e.g.][]{iyomoyo} the absorbing material is Compton thick, meaning the majority of direct emission is absorbed and reflected emission dominates the spectrum \citep{risaliti07b,matt}. A number of absorption and emission lines have been seen in the spectrum. A strong Fe fluorescence emission line is present at 6.4 keV, together with a group of Fe absorption lines between 6.7 and 8.3 keV, attributed to FeXXV and FeXXVI K$\alpha$ and K$\beta$ transitions. The measured velocities of these lines has lead to speculation that they originate from a highly-ionised, high-velocity outflow from NGC 1365 \citep{risaliti05a}. Although there have been many previous X-ray spectral studies of NGC 1365, these studies have all concentrated either on detailed analysis of a single epoch spectrum or on analysis of a small number of spectra taken over a relatively short timescale (hours or days). By contrast, here we study 190 {\it Swift} spectra taken over a period of six years. Whilst {\it Swift} does not provide spectral resolution as high as that used in most previous short-time X-ray spectral studies, e.g. with {\it XMM-Newton} or {\it Suzaku}, the {\it Swift} data cover a much longer time period and a far greater flux range. The {\it Swift} data therefore allow a proper investigation of flux-related spectral variability and of long term spectral variations, over a much larger dynamic range than in previous studies. The spectrum of NGC 1365, as with most AGN, has previously been modelled using a power law component, with an intrinsic spectral index, $\Gamma$. It is not known whether $\Gamma$ varies or not during changes in X-ray luminosity. Many groups \citep[e.g.][]{miller08,turner07,fabian05,pounds04} assume that there is no change. Observations in the $2-10$ keV band generally do show some variation, although the changes with luminosity are not large \citep[e.g.][]{sobolewska,zdziarski99}. \citet{sobolewska}, for example, who simply fit a power law to the 2-10 keV spectra, find that the observed $\Gamma$ varies as 2.7$\dot \mathrm{m} ^{0.08}$ over similar time scales to that of our data. In reality, however, these measurements of $\Gamma$ are, of course, depend on other parameters which were not included in the fits, such as the reflection component and any absorption. If the variation in the observed spectral index is interpreted in terms of the sum of a variable, steep spectrum component and a relatively constant reflection component with a hard spectrum, the intrinsic spectral index can remain constant; in this case, when the flux of the variable steep spectrum component is low, the hard spectrum component dominates, causing the observed spectral index to change \citep[e.g.][]{guainazzi99,uttley99,ponti,fabian03}. Furthermore, where observations with a large spectral range have been made, allowing good definition of the primary continuum slope, the observed variation of $\Gamma$ with luminosity has not been large (e.g. 0.1 in NGC 4151, \citealt{lubinski10}), 0.2 in NGC 4507, \citealt{braito}). Theoretical Comptonisation modelling \citep[e.g.][]{beloborodov99,coppi92} shows that the photon index can depend on the ratio of L$_{diss}$ to L$_s$ (where L$_{diss}$ is the power dissipated in the corona during variations and L$_s$ is the input soft photon luminosity) to a low power (-0.1 for AGN). Unless there are very large variations in these parameters, the intrinsic spectral index should therefore not change by more than a few tenths. Thus, although it is possible that a small change in spectral index may occur over the flux range sampled by our observations, the large changes in $\Gamma$ required by the pivoting power law models are assumed to be unlikely. In many of our models, including the model we deem most accurate, we therefore assume $\Gamma$ to be constant, as this is likely to be a reasonable approximation. The variability of the spectrum of NGC 1365 has previously been modelled using a partial covering model, in which a varying fraction of the X-ray source is obscured by absorbing material \citep[e.g.][]{risaliti09}. This model has been found to fit the data for individual, short-timescale events and is therefore also tested here. | {\it Swift} X-ray observations of NGC 1365 over a period of 6 years show a large amount of spectral variability. These variations are best explained by a two-component model consisting of an unabsorbed power law and a more luminous absorbed power law; for both components, the spectral index was fixed. The normalisations of the two power laws vary together. The absorbing column of the absorbed power law varies inversely with its luminosity, an effect which is not simply due to increased ionisation. This effect can be simply explained by viewing through the edge of a wind whose launch radius varies inversely with increasing accretion rate. The unabsorbed power law could be explained, as in the standard \citet{elvis} wind model, as the scattered component from the far side of the wind. | 14 | 3 | 1403.4253 |
1403 | 1403.6256_arXiv.txt | Slow neutron captures at $A \gtrsim 85$ are mainly guaranteed by the reaction \ctanb in AGB stars, requiring proton injections from the envelope. These were so far assumed to involve a small mass ($\lesssim 10^{-3}$ \ms), but models with rotation suggest that in such tiny layers excessive $^{14}$N hampers $s$-processing. Furthermore, $s$-element abundances in Galaxies require \ct-rich layers substantially extended in mass ($\gtrsim 4 \times 10^{-3}$ \ms). We therefore present new calculations aiming at clarifying those issues and at understanding if the solar composition helps to constrain the \ctb ``pocket'' extension. We show: i) that mixing ``from bottom to top'' (like in magnetic buoyancy or other forced mechanisms) can form a \ctb reservoir substantially larger than assumed so far, covering most of the He-rich layers; ii) that stellar models at a fixed metallicity, based on this idea reproduce the main $s$-component as accurately as before; iii) that they make nuclear contributions from unknown nucleosynthesis processes ($LEPP$) unnecessary, against common assumptions. These models also avoid problems of mixing at the envelope border and fulfil requirements from C-star luminosities. They yield a large production of nuclei below $A = 100$, so that $^{86,\;87}$Sr may be fully synthesized by AGB stars, while $^{88}$Sr, $^{89}$Y and $^{94}$Zr are contributed more efficiently than before. We finally suggest tests suitable to say a final word on the extension of the \ctb pocket. | A few years ago it was believed that the ``$s$ process'' was essentially clarified. It had been understood many years before that the distribution of $s$-elements in the Solar System required multi-modal neutron exposures \citep{war78}, generated in different astrophysical sites. A {\it main} component (accounting for nuclei from Sr to Pb) had been attributed to stars climbing for the second time along the red giant branch in a phase called Asymptotic Giant Branch, or AGB \citep{ibe75,gal88}, while a {\it weak} component (explaining nuclei up to the $N$ = 50 magic neutron number) had been ascribed to massive stars, during core-He and shell-C burning \citep{rai93}. Recent work on the weak component \citep[in particular by][] {the07,pig10} showed that this process is complex and strongly dependent on reaction rate uncertainties affecting $\alpha$-captures (especially the $^{12}$C+$\alpha$ reaction), the $^{12}$C+$^{12}$C fusion and neutron captures on nuclei immediately beyond iron. Hence, final results for the weak $s$-component must necessarily wait for new measurements and might, in the mean time, profit of any improvement in the constraints from the main component. A third, {\it strong} component, originally devised to account for 50\% of the $^{208}$Pb concentration \citep{cla67} was recognized to be unnecessary \citep{gal98}, as its role could actually be played by AGB stars at low metallicity, where the scarcity of iron nuclei implies that the number of neutrons per iron seed easily becomes sufficiently large to feed Pb. This passage from the main to the strong component in AGB stars is gradual and depends on the efficiency of the neutron release, so that heavy nuclei become progressively more enriched for decreasing metal content \citep{tra01,bus01}. A series of works presented in the last twenty years also clarified that the main component produced by AGB stars is due to the activation of the \ctanb source and (more marginally) of the \neanb source \citep{kae90}. Evidence on the minor role of the \neanb reaction was provided by the abundances of neutron-rich, stable isotopes of heavy elements in the Solar System (especially $^{86}$Kr and $^{96}$Zr). Indeed, when the neutron release from the \neanb reaction was allowed to increase (a fact that requires relatively high temperatures, i.e. $T \geqslant 3.5 \times 10^{8}$ K), excessive production of these nuclei invariably resulted \citep{kae90,arl99}. Another constraint pointing to the same direction came from the Rb/Zr ratios in carbon stars \citep{abia01,abia02}. This suggested that the parent stars had to be of relatively low mass, i.e. $M \lesssim 3$ \ms, because for this mass range the typical temperature in the pulses barely reaches $T = 3 \times 10^8$ K. In these stars, thermonuclear combustion in the He-burning shell is activated in short, explosive bursts (\textit{thermal pulses}), during which an intermediate convective zone forms in the He-rich layers; these bursts are separated by long intervals (several 10$^4$ yr), called \textit{interpulse} periods, where the \ctanb source releases its neutrons and produces $s$-elements \citep[see reviews in][]{bus99,kae11}. Even now, the above scenario is not exempt from weak points. In particular, the activation of the \ctanb neutron source requires unclear partial mixing mechanisms injecting protons into the He-rich layers during the downward envelope expansion called \textit{third dredge-up} (TDU), occurring after advanced thermal pulses. The lack of knowledge on these phenomena forced researchers to parameterize the amount of \ctb available in a relatively free way. The scenario outlined above was subsequently put in question by chemical evolution models of $s$-elements in the Galaxy by \citet{tra04}. These authors could not obtain, from the integrated Galactic production, the good fit to the main $s$-process component previously envisaged on the basis of a single AGB model for a Low Mass Star (LMS) of a suitable metallicity. The Galactic production reconstructed using nuclear yields from AGB models with a rather small ($\simeq 10^{-3}$ \ms) \ct-pocket turned out to be insufficient to account for the solar abundances of n-rich elements with atomic mass numbers between 86 and 130. In a recent work \citet{bis14} confirm these results adopting a \ct-pocket extension in the range from 2.5$\times$10$^{-4}$ to 2$\times$10$^{-3}$ \ms. Very recently, observations of $s$-process abundances were obtained for a large sample of Galactic open clusters, first for Ba \citep{dor09}, then also for Y, Zr, La and Ce \citep{mai11}. As the age of an open cluster can be determined accurately, the above authors could trace, for the first time, the evolution of $s$-elements in the Galactic disk directly as a function of time. This allowed them to observe an unexpected growth of $s$-element abundances in young stars of the Galaxy. Subsequent works on open clusters by \citet{yon12,jac13,mis13}, performed with different analysis methods, confirmed that such an increase is in general real (although for elements like Zr and La further investigations are needed). Another piece of evidence in the same direction came from a complementary study by \citet{mcw13}, who found $s$-element abundances increasing with time in the Sagittarius Dwarf Galaxy. They showed that the traditional slow neutron-capture nucleosynthesis scenario fails to reproduce the observed trend, which instead requires a larger neutron inventory in AGB stars. This confirms a recent proposal by \citet{mai12}; they suggested that the \ctb reservoir formed at TDU be considerably larger than previously adopted, at least in AGB stars of very low mass ($M \lesssim 1.5$ \ms). With this assumption and a chemical evolution model for the Galaxy, they could reproduce very well the abundances observed in open clusters. Another source of doubts on the traditional way of modeling the $s$-process emerged recently, this time from the theoretical side. In a paper by \citet{pie13} it was shown that, by including rotation in stellar modeling, any partial mixing at the convective border (like those previously invoked for forming a {\it small} \ctb pocket) become affected by the Goldreich$-$Schubert$-$Fricke instability, forcing a more complete mixing. Hence, at hydrogen shell re-ignition, any layer affected by the penetration of envelope material becomes dominated not by $^{13}$C, but by $^{14}$N, which is an efficient neutron poison and would strongly reduce the $s$-process efficiency. This confirms previous suggestions by \cite{lan9,her3,sies4} and indicates that all previous attempts at modelling the \ctb pocket formation through a small ($\lesssim 10^{-3}$ \ms) exponential penetration of protons below the envelope border \citep{cri09,cri11} would be no longer acceptable. Notice that in any case, this proton penetration had already the critical property of being dependent on the TDU phenomenon, being a downward partial extension of it. This induced in any case a limit on the \ctb pocket formation, as it could occur only after pulses followed by TDU, i.e. in rather advanced stages of the AGB. In this paper we want to re-analyze the rather confused situation, which emerged from the above discoveries, by ascertaining: i) if the hypothesis of a \ctb pocket substantially larger than imagined in the past years is compatible with a reproduction of the $whole$ distribution of $s$-elements in the solar main component \citep[as the][paper only verified this point for Y, Zr, Ba, La, Ce]{mai12}. ii) if one can imagine forms of deep mixing alternative to the partial extension of the envelope, in order to suggest ways for putting the formation of the \ctb pocket on safer grounds; and iii) if an analysis of the solar-system main $s$-process component can constrain the extension of the layers partially polluted by protons at TDU, where \ctb is expected to form. In section 2 we present some simple ideas aimed at setting the stage for a physical modeling of the \ctb pocket. In section 3 we discuss recent improvements in the nuclear inputs, both for neutron-capture cross sections and for the rates of the neutron-producing reactions. In section 4 we describe our computations for $s$ processing in low-mass AGB models at suitably-chosen metallicities, supposed to represent an average of the $s$-processing efficiency over the evolution of the Galaxy. In this way we adopt, for the sake of comparison, two widely different extensions for the \ctb pocket (in the range of those so far proposed by different authors). This is aimed at understanding how these results compare with the solar distribution, an issue that is discussed in section 5. Finally, in section 6 we discuss the implications of the results found and we suggest some observational and theoretical tests that should help saying a final word on the \ctb formation in AGB stars. | In this paper we re-analyzed nucleosynthesis models for slow neutron captures in AGB stars, after new observational as well as theoretical information shed doubts on the previous scenario for the formation of the \ctb neutron source and for its actual extension. In our work we have argued that, even in presence of persisting uncertainties concerning the dynamical mechanisms promoting proton penetration into the He-rich layers at the convective border, stellar physics offers other, perhaps more secure, ways of generating transport phenomena suitable for forming a \ctb reservoir and then inducing neutron-capture nucleosynthesis. In particular, we have suggested that a fruitful line of research may be that of describing, through a quantitative MHD treatment, the development of toroidal magnetic fields, induced by stellar dynamos, in the radiative He-rich layers below the convective envelope. The above scheme for the creation of a \ct-rich layer foresees that the partially mixed zones extend down to very deep regions, essentially involving most of the He-rich layers of the AGB star, due to the formation of buoyant magnetic structures close to the outer border of the degenerate C-O core. We have also underlined that any attempt at upgrading our present understanding of $s$-processing in low-mass AGB stars must take into account the fact that the infrared LFs of these last agree with stellar model predictions only if the magnitudes remain moderate ($M_{bol} \lesssim - 5$) and hence the number of pulses undergone by the star is smaller than previously assumed \citep{gua13}. The above considerations imply that $s$-processing in AGB stars is built in a way rather different than imagined so far, namely through a smaller number of pulse-interpulse cycles, each however experiencing a more efficient nucleosynthesis episode. As these required changes are also necessary to explain the increasing abundances of $s$-elements in the Galactic disk \citep{mai12}, they seem to become mandatory. Also, they cannot be mimicked by increasing the abundance of \ctb in a small pocket: the concentrations in mass of \ctb and of $^{14}$N in the reservoir formed are fixed by H-burning rates and cannot be varied freely (as is instead often done), without violating basic physical laws. We have then performed a comparison of the results achievable for reproducing the solar main component in two cases: i) the scenario most commonly used in the last 20 years for dealing with s-processing, based on a small extension of the \ctb reservoir (that we called Case A); and ii) the newly suggested one, with \ct-rich layers reaching down to deep regions of the He-rich zone (referred to as Case B). The comparison has been performed by computing n-captures in stellar models of 1.5, 2.0 and 3.0 \msb and by averaging the $s$-process results after weighting by the IMF of Salpeter. The main result is that, if the metallicity for the two cases is chosen suitably, both provide a distribution of production factors mimicking the main $s$-process component. Due to the different neutron capture efficiency resulting from the different extension of the \ctb pocket, the number of pulses differs in the two cases, much like the metallicity does: [Fe/H] $\simeq - 0.5$ and n $\gtrapprox$ 20 pulses for the old scenario, [Fe/H] $\simeq - 0.15$ and n $\lessapprox$ 15 pulses for the new one. Moreover, the main aim of the above test was to look for an answer to the question posed in the title: can we distinguish, from comparisons with solar abundances, which scenario has to be preferred? In general, if one sticks to the results from a stellar generation at a suitably chosen metallicity then a decision is not possible, as the quality of the fits to the solar abundances of $s$-only nuclei shown in Figure 7 is essentially identical. \begin{figure}[t!!] \centering \includegraphics[scale=0.45]{Fig9-Trippella.eps} \caption{The Age-Metallicity Relation (AMR) in the Galactic disk as derived by \citet{mai12}. The two boxes roughly indicate the metallicity and age intervals where the main component is best fitted by AGB star nucleosynthesis, in the cases (A and B) discussed in the text. The symbol of the Sun is also shown in the figure. \newline (A color version of this figure is available in the online journal.)}\label{nine} \end{figure} However, a closer look reveals remarkable differences in the predictions of the two cases for the nucleosynthesis of $s$-nuclei in the Galaxy. This is already evident from Figure 8, if one considers light nuclei outside the limits of the main component; but is true also for heavier isotopes, when one derives the consequences of the calculations for the chemical evolution of the Galaxy \citep{tra04,mai11,mai12}. Both issues are actually strictly connected as outlined in the following. Let us show how. The reason for the different predictions from Case A and Case B at the lower mass end of the distribution (requiring or not a solar $LEPP$ process) is rather simple. It can be illustrated with the help of the Age-Metallicity Relation (AMR), which is reproduced in Figure 9 from the results by \citet{mai12}. The two boxes represent the metallicity intervals over which the main component is best fitted in our Case A and Case B. For Case A the AMR is sampled over a short time interval, at epochs old enough that it is still far from the conditions prevailing over most of the Galactic disk duration. The total number of stars in that short interval is therefore relatively small and the effects on Galactic abundances will not be dominant. Most AGB stars will be born later, when the abundance of Fe is higher. Due to the small pocket, the number of neutrons per iron seed in them will be so small that their yields will be almost irrelevant in the global inventory of the Galaxy. As, with low neutron exposures, they feed mainly light $s$-process nuclei, these last will be insufficiently produced, hence the requirement of a $LEPP$ integration. On the contrary for Case B the reference metallicity range, due to the large pocket, is shifted upward, to conditions typical of the main Galactic disk population, lasting for several Gyr. In this case the AGB stars shown before to mimic the main component will be the dominant ones, sufficient in number and effectiveness for taking care of the Galactic enrichment, so that no extra process is required. These are examples of a more general trend. Essentially, by choosing adequately the metallicity and the \ctb pocket extension, one can obtain production factors mimicking the solar distribution in generations of AGB stars for any choice of the \ctb reservoir. However, if we want that the chosen generation can process enough Galactic material to be really dominant in the chemical evolution of the Galaxy and in controlling solar abundances, then we must choose an effective average metallicity typical of the thin Galactic disk, where the abundance evolution is low, long time scales are involved and the number of AGBs contributing becomes huge. In that case, the abundance of iron is high and to have a sufficient number of neutrons per iron seed the \ctb pocket must be quite extended in mass. This favors Case B. Obviously, Case A cannot be excluded on these grounds, but it would need a $LEPP$ contribution. For Case A, this means that searching for $average$ Galactic conditions where the solar distribution is reproduced is not really meaningful. The above discussion gives us an opportunity to identify crucial tests that should be made, from which a conclusive judgement can be derived on the real extension of the \ctb pocket (hence also on its origin). We list below six such tests that are, in our opinion, especially suitable to provide a final answer. \begin{itemize} \item {Compute models using the Case A choice for the pocket, but with a limited number of pulses (thus reaching a luminosity compatible with present-day LFs), verifying whether a compromise can be found that fits the solar data without violating the prescriptions on C-star magnitudes. We believe this should be actually very difficult, given the shortage of neutrons; but this is in any case a crucial test to be performed quantitatively.} \item {Compute Galactic chemical evolution models including at least Sr, Ba and Pb isotopes with the two scenarios and compare them with the observations (which are unfortunately limited for Pb abundances). Very young stellar clusters (absent in previous such studies by \citep{rai99,tra01,tra04} should be included. We expect that the models of Case A will not reproduce the observations, while those of Case B will; but again this has to be demonstrated in detail.} \item {Verify whether, with an extended \ctb pocket, one can reproduce the $s$/C ratios of post-AGB stars, an achievement that proved impossible for the models of Case A \citep{per12}.} \item {Detailed, quantitative models (based on MHD calculations or on other processes capable of forcing the formation of a \ctb pocket) should be developed to see what kind of mixing can be realistically expected.} \item {The abundance pattern shown by presolar materials recovered in pristine meteorites should be compared with the predictions of the two scenarios, looking for more detailed constraints possibly coming from the isotopic admixtures measured in presolar grains.} \item {When the chemical evolution of the Galaxy is computed, models of Case A were shown to require, for explaining the solar system abundances of $s$-elements up to A $\simeq$ 120, the contribution of the unknown solar $LEPP$ process \citep{tra04}. This is not necessarily coincident with the process required at low metallicity, see e.g. \citet{mon07}. From the tests made on crucial elements by \citet{mai12} we know this is not needed by the new models of Case B. Now a critical point is: can the approach of Case A, plus a unique choice for the LEPP contribution, explain the increased abundances of $s$ elements in young Galactic stellar systems? An answer can come from fixing the LEPP contributions from solar constraints, then verifying if this is sufficient for explaining the increased abundances in young clusters. Again we predict that this procedure should fail and the results by \cite{bis14} seem to point in that direction. However they do not consider the open cluster problem directly, so that a dedicated calculation must be done before a final decision is taken.} \end{itemize} The information we can get from performing the above tests would be decisive. Should the new models, with an extended \ctb pocket and a limited number of pulses, prevail (as it may seem probable now, given the larger number of constraints they appear to match) then some general conclusions on $s$-processing should be revised. In particular: i) the main component should be considered as including $^{86,\;87}$Sr completely; ii) the expectations for $^{208}$Pb in low metallicity stars would be different and probably less extreme; iii) the $s$-process contribution to nuclei like $^{88}$Sr, $^{89}$Y, $^{94}$Zr, $^{138}$Ba and $^{140}$Ce should be revised upward and accepted to reach 90 $-$ 100\% of their abundance; iv) in view of the expected new measurements for the rates of the $^{12}$C($\alpha,\gamma$)$^{16}$O and of the $^{12}$C+$^{12}$C fusion reactions, new determinations of $s$-processing in massive stars should verify the new suggestions from AGB models for nuclei below A$\sim$ 90. {\bf Acknowledgments.} O.T. thanks $INFN$ \textit{Istituto Nazionale di Fisica Nucleare}, section of Perugia, for financial support and Marco La Cognata for constructive discussions on the ${}^{13}{\rm C}(\alpha,{\rm n}){}^{16}{\rm O}$ reaction rate. S.P. acknowledges support from the Italian MIUR through the excellent found for Nuclear Astrophysics ``LNS Astrofisica Nucleare - fondi premiali''. The authors are grateful to the referee for a very careful and constructive report, that greatly helped in clarifying many relevant issues. | 14 | 3 | 1403.6256 |
1403 | 1403.6126_arXiv.txt | In order to incorporate the effects of radiation transfer in highly dynamical flows where the light-crossing time is only marginally shorter than dynamical timescales, it is necessary to solve the time-dependent transfer equation. We describe a new algorithm for solving the coupled frequency-integrated transfer equation and the equations of magnetohydrodynamics in this regime. The transfer equation is solved in the mixed frame, including velocity dependent source terms accurate to $\mathcal{O}\left(v/c\right)$. An operator split approach is used to compute the specific intensity along discrete rays, with upwind monotonic interpolation used along each ray to update the transport terms, and implicit methods used to compute the scattering and absorption source terms. Conservative differencing is used for the transport terms, which ensures the specific intensity (as well as energy and momentum) are conserved along each ray to round-off error. The use of implicit methods for the source terms ensures the method is stable even if the source terms are very stiff. To couple the solution of the transfer equation to the MHD algorithms in the {\sc Athena} code, we perform direct quadrature of the specific intensity over angles to compute the energy and momentum source terms. We present the results of a variety of tests of the method, such as calculating the structure of a non-LTE atmosphere, an advective diffusion test, linear wave convergence tests, and the well-known shadow test. We use new semi-analytic solutions for radiation modified shocks to demonstrate the ability of our algorithm to capture the effects of an anisotropic radiation field accurately. Since the method uses explicit differencing of the spatial operators, it shows excellent weak scaling on parallel computers. The method is ideally suited for problems in which characteristic velocities are non-relativistic, but still within a few percent or more of the speed of light. The method is an intermediate step towards algorithms for fully relativistic flows. | \label{sec:introduction} The dynamical effects of radiation are important in a variety of astrophysical systems. For example, in black hole accretion disks radiation is not only an important cooling mechanism, but also is the dominant source of pressure in the radiation-dominated regime. The standard thin disk model \citep[][]{ShakuraSunyaev1973}, in which heating and radiative cooling are balanced locally, is only applicable when the accretion rate is $\sim 0.01-0.3$ of the Eddington rate. At larger accretion rates, the disk is thought to become geometrically and optically thick \citep[][]{Abramowiczetal1988,AbramowiczFragile2011}, and radial advection may also become an important cooling mechanism. At both super-Eddington and strongly sub-Eddington accretion rates, strong outflows are likely to be produced \citep[][]{BlandfordBegelman1999,OhsugaMineshige2013}. While much insight into the structure and dynamics of accretion disks over this broad range of parameters has been made using analytic steady-state models, more realistic calculations require numerical methods. In addition, since the magneto-rotational instability (MRI) \citep[][]{BalbusHawley1991,BalbusHawley1998} is generally believed to be the physical mechanism for angular momentum transport in black hole accretion disks, numerical methods for magnetohydrodynamics (MHD) that include the effect of radiation are crucial to enable study of the saturation and nonlinear regime of the MRI in such disks. Shearing box simulations of the MRI in the local patches of a radiation pressure dominated accretion disk have significantly improved our understanding of the dynamics in this regime \citep[][]{Turneretal2003,Turner2004,Hiroseetal2006,Kroliketal2007,Blaesetal2011,Jiangetal2013b}. For example, in this regime disks are found to undego thermal runaway in local shearing box simulations \citep[][]{ShakuraSunyaev1976,Jiangetal2013c}. However, understanding the saturation of thermal runaways and the implications for observations of radiation pressure dominated disks require global simulations. In order to model the thermal properties of accretion disks over a wide range of parameters accurately requires, algorithms for radiative transfer (RT) that satisfy several criteria. First, the algorithm must be accurate in both optically thick (e.g. the mid-plane) and optically thin (e.g. the photosphere) regions. Second, it must be accurate over a wide range of ratios of radiation to gas pressures, from very small values (when the accretion rate is small) to very large values (when the accretion rate is high and the disk is radiation pressure dominated). Third, it must handle the non-local and anisotropic transport of photons accurately. For example, the radiation field above the photosphere will be produced by emission from many different (including distant) parts of the disk. This is likely critical to calculate the dynamics of outflows correctly. Finally, the algorithm must be efficient and highly parallelizable in order to enable large-scale global simulations on modern computer systems. Since in general the RT equation is a function of seven independent variables (time, spatial position, angles and frequency), it is usually thought that a formal solution to compute the specific intensity is intractable, even in the case of frequency independent opacities such that the grey (frequency independent) transport equation can be used. Thus, most algorithms for radiation MHD adopt a moment formalism, in which a hierarchy of angle-integrated time-dependent moment equations are solved. This reduces the dimensionality of the problem to time and spatial coordinates, as are used for the MHD equations. However, the hierarchy of moment equations require a closure relation. A variety of different closures have been adopted, including the flux-limited diffusion (FLD) approximation \citep[][]{LevermorePomraning1981,TurnerStone2001,Krumholzetal2007,Zhangetal2011,Holstetal2011,Commerconetal2011} and more recently the M1 method \citep[][]{Gonzalezetal2007,SkinnerOstriker2013,Sadowskietal2013,McKinneyetal2013}. A variety of recent studies of the global dynamics of black hole accretion disks \citep[][]{Ohsugaetal2011,Sadowskietal2013,McKinneyetal2013} have used such approximate closures. However, it is far from clear that such closures will be accurate everywhere in the flow. The angular distribution of the radiation field can be arbitrarily complex, particularly in optically thin regions. But these schemes make the simple assumptions that the Eddington tensor can be uniquely determined by the local radiation energy density and flux. And, at present, there is no means of improving the approximation in these closures -- there is no resolution parameter that can be increased to assess convergence. Thus, it is worthwhile to develop new algorithms that are based on a direct solution of the RT equation. A key message of our work is that modern computer systems have advanced to the point that formal solution of the six dimensional grey RT equation is now feasible at every time step of a MHD simulation, and that approximate closures of unknown accuracy are no longer necessary. In this vein, we recently have described an algorithm for radiation MHD that is based on the method of short characteristics to solve the \emph{time independent} RT equation in multidimensions \citep[][]{Davisetal2012}, an approach that was also used in earlier work \citep[][]{Stoneetal1992,HayesNorman2003,HubenyBurrows2007}. At every time step, short characteristics are used to compute the specific intensity, from which direct quadrature is used to compute the components of the variable Eddington tensor (VET) which serves as the closure relation in the radiation moment equations \citep[][thereafter JSD12]{Jiangetal2012}. This method has been successfully used to study the dynamics of radiation-dominated accretion disks using the local shearing box simulations \citep[][]{Jiangetal2013b,Jiangetal2013c}. The primary assumption underlying this method is that the light-crossing time is much shorter than characteristic dynamical times in the flow, so that solutions of the time-independent RT equation can be used to compute the VET. This is an excellent approximation in most flows. However, in the inner regions of black hole accretion disks, dynamical times become comparable to the light crossing time, and solution of the time-dependent RT is required. The goal of this paper is to describe such an algorithm based on ray tracing methods. Although our algorithm is motivated by the study of black hole accretion disks, it is also applicable to any system that has characteristic flow velocity which is not much smaller than the speed of light. Of course, within a few gravitational radii of the black hole, the flow will be fully relativistic, and both general relativistic MHD and RT are required to describe the dynamics. While GRMHD flows including radiation using FLD or the M1 closure have recently been reported \citep[][]{Fragileetal2012,Sadowskietal2013,McKinneyetal2013}, developing algorithms based on the formal solution of the RT equation in GR is a formidable challenge. Instead, in this paper we describe an intermediate step. That is, we develop ray tracing methods for the formal solution of the time-dependent RT equation, but restrict ourselves to non-relativistic flows. Thus, there is a relatively narrow window of velocities where our method is applicable; since when $v \ll c$ the time step in our algorithm is too small for it to be efficient, while for $v \approx c$ the flow is relativistic. Nonetheless, there is an interesting range of radii in black hole accretion disks where our method is applicable. We report results of simulations studying this region elsewhere. Moreover, we are now extending the method described in this paper to full GR. Another popular approach to solve the time-dependent RT equation are Monte Carlo methods \citep[][and references therein]{Whitney2011}. These methods are accurate, flexible and have almost perfect parallel scaling, and therefore are a very attractive direction for the future. However, the intrinsic noise associated with Monte Carlo makes the method very expensive to model radiation pressure dominated flows accurately \citep[][]{Davisetal2012}. While it may be possible to reduce the noise and computational cost in Monte Carlo using novel approaches \citep[][]{Yusef-Zadehetal1984,Densmoreetal2007,Steinackeretal2013}, in this work we instead adopt ray tracing methods. Although ray tracing methods are more complex to implement, they do not suffer from noise. This paper is organized as follows. The equations we solve are given in Section \ref{sec:equation}. The basic angular discretization scheme is described in Section \ref{sec:angles}. A complete description of the algorithm are given in Section \ref{sec:algorithm} and we show test results in Section \ref{sec:test} to demonstrate the capability of the method. In Section \ref{sec:performance}, we test the speed and parallel performance of the code. Finally, we summarize in Section \ref{sec:summary}. | \label{sec:summary} We have developed a new algorithm for radiation MHD based on solving the \emph{time-dependent} RT equation to compute the specific intensity along discrete rays. Integration of the specific intensity over angles then yields the radiation energy and momentum source terms, which are coupled to the ideal MHD equations. We have tested the code extensively based on the suite of test problems used by JSD12. Compared with commonly used methods for solving the radiation moment equations, such as FLD and M1, the most significant advantage of our algorithm is that we calculate the angular distribution of photons self-consistently. Therefore, we do not need any ad-hoc closure relation as required in FLD and M1 methods. We expect our methods to be superior in regions where the photon mean free path becomes comparable to or larger than the characteristic scales of the fluid. Then the disparate contributions of multiple, non-local sources can give rise to complexity in the angular distribution of the photons that cannot be encapsulated using only a few of lowest order moments of the radiation field ($E_r$, $\bF_r$). We have shown results for a variety of tests that demonstrate the importance of the non-local and anisotropic properties of the radiation for the correct dynamics, for example the structure of radiation modified shocks, and the shadows cast by multiple sources of radiation. The method described here can be considered an extension of previous algorithms that use the method of short characteristics to solve the \emph{time-independent} RT equation \citep[][]{HayesNorman2003,Davisetal2012}, which is then used to calculate a variable Eddington tensor to close the radiation moment equations \citep[][JSD12]{Stoneetal1992}. The full angular distribution of the specific intensity is captured in both algorithms. However, the main advantages of the algorithm described here is that it can be used for dynamical problems where the dynamical time is not negligibly small compared to the light crossing time, and when the velocity dependent terms play an important role. Moreover, since the method developed here uses explicit differencing of spatial operators, it avoids the inversion of large matrices every time step, and therefore is much more efficient on modern parallel systems. The algorithm developed in this paper is only applicable to non-relativistic flows. As general relativity (GR) is believed to be the correct description of the gravitational field near the event horizon of black holes, extending our algorithms to GR is the next step, and indeed is already underway. There are several aspects of the algorithm developed here that we expect will be crucial for the method in full GR. This includes the implicit treatment of source terms, including scattering, and the use of upwind monotonic interpolation methods that guarantee conservation of radiation energy and momentum to roundoff error. The time step for stability with our algorithm is based on the light crossing time of each cell. This is very inefficient for systems in which the typical sound speed or flow velocity is much smaller than the speed of light. In the regime where the reduced speed of light approximation applies \citep[][]{SkinnerOstriker2013}, it can be combined with our algorithm to increase efficiency. We have used our code to study the global structure of black hole accretion disks for different accretion rates at intermediate distances from the event horizon. Results from these simulations will be reported elsewhere. | 14 | 3 | 1403.6126 |
1403 | 1403.4912_arXiv.txt | \corr{ The $^{13}\textrm{CO}$ molecule is often used as a column density tracer in regions where the $^{12}\textrm{CO}$ emission saturates. The $^{13}\textrm{CO}$ column density is then related to that of $^{12}\textrm{CO}$ by a uniform isotopic ratio. A similar approximation is frequently used when deriving $^{13}\textrm{CO}$ emission maps from numerical simulations of molecular clouds. To test this assumption we calculate the $^{12}\textrm{CO}/^{13}\textrm{CO}$ ratio self-consistently, taking the isotope selective photodissociation and the chemical fractionation of CO into account. We model the coupled chemical, thermal and dynamical evolution and the emergent $^{13}\textrm{CO}$ emission of isolated, starless molecular clouds in various environments. Selective photodissociation has a minimal effect on the ratio, while the chemical fractionation causes a factor of 2-3 decrease at intermediate cloud depths. The variation correlates with both the $^{12}\textrm{CO}$ and the $^{13}\textrm{CO}$ column densities. Neglecting the depth dependence results in $\leq$60 per cent error in $^{12}\textrm{CO}$ column densities derived from $^{13}\textrm{CO}$. The same assumption causes $\leq$50 per cent disparity in the $^{13}\textrm{CO}$ emission derived from simulated clouds. We show that the discrepancies can be corrected by a fitting formula. The formula is consistent with millimetre-wavelength isotopic ratio measurements of dense molecular clouds, but underestimates the ratios from the ultraviolet absorption of diffuse regions. } | The carbon monoxide molecule (CO) and its isotopes are the most widely used gas-phase tracers of total column density in the interstellar medium (ISM). In contrast to the hydrogen molecule, CO is asymmetric and hence has a permanent dipole moment. Its dipole transitions between rotational levels can be excited at temperatures (few$\times10$ K) and densities ($\approx$300 cm$^{-3}$) typical of giant molecular clouds (GMCs). The emission from the lowest transitions is relatively easily detectable at millimetre wavelengths. Due to its high fractional abundance ($\chi_{\rmn{CO}}\approx10^{-4}$, in equilibrium at high density) the emission of the most abundant CO isotope ($^{12}\textrm{CO}$) is usually optically thick, therefore only lower limits of the total column density could be derived. To achieve better total column density estimates, less abundant CO isotopes are used (usually $^{13}\textrm{CO}$ and $\textrm{C}^{18}\textrm{O}$). The simplest and most often used method is the following \citep[e.g.][]{Pineda2008,Wilson2009,Pineda2010}: The $^{12}\textrm{CO}$ emission is assumed to be fully optically thick and in local thermodynamic equilibrium (LTE), allowing the excitation temperature of $^{12}\textrm{CO}$ to be calculated. If the excitation temperature is the same for $^{13}\textrm{CO}$ and if $^{13}\textrm{CO}$ is optically thin, a simple relation between the integrated intensity along the line of sight, $W(^{13}$CO$)$, and the column density of $^{13}\textrm{CO}$, $N(^{13}\textrm{CO})$ can be derived \citep[e.g. see equation~9 in][]{Pineda2008}. The $^{13}\textrm{CO}$ column densities are then usually converted into $^{12}\textrm{CO}$ column densities using a uniform $^{12}\textrm{CO}/^{13}\textrm{CO}$ isotope ratio. Finally, the $^{12}\textrm{CO}$ column density is transformed to total column density assuming a given conversion factor between CO and H$_2$. The high uncertainties and environmental dependence of this final step have been extensively studied in the literature \citep[][and references within]{Shetty2011a,Shetty2011b,GloverMacLow2011,Feldmann2012}. However, the $^{12}\textrm{CO}/^{13}\textrm{CO}$ ratio may vary considerably. Typically it is chosen to be equal to the measured $^{12}$C/$^{13}$C ratio \citep{Pineda2010}. The carbon isotope ratio, however, shows large regional variations. Based on observations of millimetre-wavelength emission of CO isotopes, \citet{LangerPenzias1990} found a systematic gradient with galactocentric distance in the carbon isotope ratio, ranging from 24 in the Galactic Center to about 70 at ${\sim}$12~kpc. They found the average ratio to be 57 at the solar galactocentric distance, which shows $^{13}$C enhancement compared to the value of 89 measured in the Solar System \citep{Geiss1988}. Observations of CO absorption in ultraviolet electronic and near-infrared vibrational transitions \citep[e.g.][]{Scoville1983,MitchellMaillard1993,Goto2003,Sonnentrucker2007,Sheffer2007} find up to a factor of 3 higher ratios in the solar neighbourhood. Nevertheless, the most frequently adopted ratios are between 57 \citep{LangerPenzias1990} and 69 \citep{Wilson1999}, the measured average values for \corr{the interstellar medium within a few kpc of the Sun}. The gradient with galactocentric distance and the \corr{$^{13}$C enhancement in the solar neighbourhood compared to the Solar System value could} be interpreted in the framework of the carbon isotopic nucleosynthesis. $^{12}$C is the primary product of the triple-alpha process during the post Red Giant Branch (RGB) evolution of massive stars. The rarer $^{13}$C is produced from $^{12}$C as a secondary product in the CNO cycle during the RGB phase of low and intermediate mass stars. Due to the longer lifetime of low and intermediate mass stars -- which are the main contributors of $^{13}$C enrichment -- the $^{12}$C/$^{13}$C ratio is expected to decrease with time and to depend on the star formation history \citep{Audouze1975}. \begin{figure} \includegraphics[scale=0.42]{fig1.eps} \caption{Qualitative picture of carbon isotopic chemistry in molecular clouds. The outer diffuse layer (gray) is followed by the translucent (yellow and blue) and the inner, dense regions (orange). The font size of the chemical species relates to the abundance of the species in the region. The font size of chemical reactions indicates the importance of the reaction in the region, the green and red colors represent reactions producing and destroying CO, respectively. When a particular isotopic species is not indicated, then both species are affected. Based on \citet{DishoeckBlack1988}.} \label{fig:cochem} \end{figure} More important for our problem, however, is the fact that the CO isotope ratio could vary by a factor of a few even if the \corr{elemental} $^{12}$C/$^{13}$C ratio is constant in a region under investigation, due to isotope-selective chemical processes. For example, \citet{DishoeckBlack1988} describe qualitatively the CO isotopic chemistry in molecular clouds as follows (see Fig.~\ref{fig:cochem}). The preferred pathways of CO production are the ion-neutral reaction of C$^{+}$ and OH producing HCO$^{+}$, which dissociatively recombines to CO and H, and the neutral-neutral reaction of CH or CH$_{2}$ with an oxygen atom. These reactions are not isotope-selective and work (with varying efficiency) in every region of the molecular cloud. In the diffuse regions (gray, $A_{\rmn{V}} < 0.5\,\textrm{mag}$\footnote{The exact visual extinction values depend on the strength of the incident radiation field and the density distribution.}) the photodissociation of CO by interstellar far-ultraviolet (FUV) photons dominates over the production reactions and most of the carbon is in ionized form. \corr{In translucent regions (yellow, $1\,\textrm{mag} < A_{\rmn{V}} < 2\,\textrm{mag}$) the CO production rates start to compete with the photodissociation. The $^{12}\textrm{CO}$ column density becomes high enough to effectively self-shield itself from the incident interstellar FUV photons. However, $^{13}\textrm{CO}$, due to its slightly shifted absorption lines at UV wavelengths and the lower abundance, is less effectively self-shielded. This difference in self-shielding, in principle, leads to isotope-selective photodissociation. In practice, the selective photodissociation is expected to be dominant only in low density ($n < 10^{2} \textrm{cm}^{-3}$) domains or in dense regions with very strong radiation fields \citep[see][]{Roellig2013}.} Further in (blue, $2\,\textrm{mag} < A_{\rmn{V}} < 5\,\textrm{mag}$), due to the shielding by dust absorption and the increasing $^{13}\textrm{CO}$ column density, both isotopic species are effectively protected from FUV photons. In this region ionized carbon is still abundant and the \begin{equation} {^{13}\rmn{C}^{+}} + {^{12}\rmn{CO}} \rightleftharpoons {^{12}\rmn{C}^{+}} + {^{13}\rmn{CO}} + \Delta\rmn{E} \label{eq:fracreact} \end{equation} fractionation reaction \citep{Watson1976} becomes important. At temperatures typical to the corresponding cloud depths, the exothermic reaction (to the right, leads to energy release) is preferred, resulting in more $^{13}\textrm{CO}$ production, and consequently in a reduced isotope ratio. \corr{This reaction provides the main $^{13}\textrm{CO}$ production and destruction paths in this region. On the other hand, the destruction of $^{12}\textrm{CO}$ is determined by the competing effects of photodissociation, chemical fractionation and dissociative charge transfer with $\textrm{He}^{+}$, while its production is mainly due to the neutral-neutral reaction of light hydrocarbons with oxygen.} At the highest column densities (orange, $A_{\rmn{V}} > 5\,\textrm{mag}$) the \corr{CO chemistry is governed by non-isotope-selective reactions. The gas phase production of both CO isotopes happens through $\textrm{HCO}^{+}$ recombination and desorption from dust grains, while the main destruction channels are the dissociative charge transfer with He$^{+}$ and $\textrm{H}_{3}^{+}$, photodissociation by cosmic-ray induced photons, and freeze-out onto grains.} As a result, the isotope ratio approaches to the elemental ($^{12}$C/$^{13}$C) ratio. Due to these processes we expect that the CO isotope ratio varies significantly even within the same GMC. In fact, observational studies report a factor of a few region-by-region variation. In the case of the Taurus molecular cloud, the indirect measurements of \citet{Goldsmith2008} and \citet{Pineda2010} find isotope ratios between 30 and the canonical value of 69, suggesting $^{13}\textrm{CO}$ enrichment. The direct determination of the $^{12}\textrm{CO}/^{13}\textrm{CO}$ ratio is usually difficult and restricted to a certain column density range. For instance, the ultraviolet and millimetre-wavelength absorption measurements, such as presented by \citet{Sheffer2007}, \citet{Sonnentrucker2007} and \citet{LisztLucas1998}, require suitable galactic or extragalactic background sources and CO column densities, which fall into the optically thin, diffuse regime ($N(^{12}\textrm{CO}) < \textrm{few} \times 10^{16}\,\textrm{cm}^{-2}$). Observations of millimetre-wavelength emission from GMCs \citep[e.g.][]{Pineda2008,Goldsmith2008,Pineda2010}, however, usually trace the higher CO column density regions, where the isotope ratio-column density correlation is not constrained by observations, and therefore as a ``best guess'' a uniform isotope ratio is adopted. The inverse problem emerges when $^{13}\textrm{CO}$ emission is inferred from (magneto-)hydrodynamical simulations. The computational cost of the chemical modelling scales with the cube of the number of species considered \citep{GloverClark2012b}. Even when only 14 self-consistently calculated (i.e. not described by conservation laws), non-equilibrium species are included in the network, the chemistry will often be the dominating factor in terms of computational cost, taking up to 90 per cent of the total computational time \citep{Glover2010,GloverClark2012b}. For the practical reason of cost efficiency usually only the most common isotope, $^{12}\textrm{CO}$ is included in the chemical networks. When observable quantities, like emission from rarer CO isotopes are inferred from such simulations \citep[e.g.][]{Beaumont2013} the canonical isotope ratio is adopted and assumed to be constant through the whole simulation domain. In this paper we investigate the effect of \emph{selective photodissociation} and \emph{chemical fractionation} on the $^{12}\textrm{CO}/^{13}\textrm{CO}$ isotopic ratio in different environments and for different cloud properties, using turbulent hydrodynamical simulations that include a self-consistent chemical and cooling model and an approximate treatment of the attenuation of the interstellar radiation field (ISRF). One of our aims is to test and improve the frequently used assumption of uniform $^{12}\textrm{CO}/^{13}\textrm{CO}$ ratio in the context of inferring $^{13}\textrm{CO}$ emission from simulations which neglect isotopic chemistry. We also provide a prescription for deriving the isotope ratio from the $^{13}\textrm{CO}$ column density (e.g. calculated from observations). In section~\ref{sec:sim} we describe the numerical setup and the initial conditions of our simulations. Section~\ref{sec:ColDenRatio} discusses the correlation between various (total, $^{12}\textrm{CO}$ and $^{13}\textrm{CO}$) column densities and the isotope ratio for different cloud conditions. We also propose a formula for inferring $^{13}\textrm{CO}$ column/number densities from $^{12}\textrm{CO}$ in simulations neglecting fractionation chemistry, and one for calculating the isotope ratio from observations of $^{13}\textrm{CO}$. Then in section~\ref{sec:EmissionMaps}, we post-process the simulations with line radiative transfer to quantitatively compare the emergent $^{13}\textrm{CO}$ line profiles and emission maps in case of self-consistently calculated, column density dependently inferred and uniform isotope ratios. Section~\ref{sec:com} compares our results to previous theoretical works and to observations of the $^{12}\textrm{CO}/^{13}\textrm{CO}$ column density ratio. We summarize the results and draw our final conclusions in section~\ref{sec:Sum}. \begin{table} \caption{Model parameters and used snapshots\label{table1}} \label{tab1} \begin{center} \resizebox{8.4cm}{!} { \begin{tabular}{cccccc} \hline Model & $n_{0}$ [$\rmn{cm}^{-3}$] & Metallicity [$Z_{\odot}$]& ISRF [$G_{0}$]& Time [Myr] \\ \hline a & 300 & 0.3 & 1 & 2.046 \\ b & 300 & 0.6 & 1 & 1.930 \\ c & 300 & 1 & 0.1 & 2.124 \\ d & 300 & 1 & 1 & 2.150 \\ e & 300 & 1 & 10 & 2.022 \\ f & 1000 & 1 & 1 & 0.973 \\ \hline g & 300 & 1 & 1 & 2.150 \\ \hline \end{tabular} } \end{center} Summary of model parameters. Each model cloud has $10^{4}M_{\odot}$ and an SPH mass resolution of 0.5$M_{\odot}$. In each case the analysed snapshots are \corr{chosen to represent the molecular clouds preceding star formation.} $Z_{\odot}$ and $G_{0}$ refer to the solar metallicity and Draine radiation field strength (1.7 in units of the \citet{Habing1968} field) respectively. \corr{Models from a) to f) have fully molecular initial conditions while model g) is calculated with atomic initial composition.} \end{table} | \label{sec:Sum} We investigate the effects of selective photodissociation and chemical fractionation on the $^{12}\textrm{CO}/^{13}\textrm{CO}$ ratio in realistic hydrodynamical simulations of isolated molecular clouds, considering a range of cloud properties. We aim to quantitatively test the validity of the frequently assumed uniform isotope ratio, when $^{13}\textrm{CO}$ intensities are calculated from hydrodynamical simulations neglecting isotopic chemistry or when $^{12}\textrm{CO}$ column density/mass is inferred from observations of $^{13}\textrm{CO}$ emission. We find a close correlation between the $^{12}\textrm{CO}$ column density and the isotope ratio, which shows only a weak dependence on cloud conditions within the considered parameter range. \corr{The isotope-selective-photodissociation is effective and increases the isotope ratio up to 70 in some regions. However, the mass fraction of the gas with higher than elemental ratio is negligible, and has virtually no effect on the column density ratio and on the observables.} The chemical fractionation reaction -- by enhancing the $^{13}\textrm{CO}$ abundance -- reduces the isotope ratio to values as small as 20 in the $10^{15}\,\textrm{cm}^{-2} < N(^{12}\textrm{CO}) < 10^{17}\,\textrm{cm}^{-2}$ range and has a significant effect on the millimetre-wavelength emission. At high CO column densities neither of the isotope-selective reactions are effective, therefore the ratio increases to the value of the elemental isotopic abundance ratio. The isotope ratio varies similarly with the $^{13}\textrm{CO}$ column density. The correlations depend only weakly on environmental conditions, such as ISRF strength or metallicity. This is because the ability of gas to form CO correlates with its ability to shield itself form the interstellar radiation field. If the irradiation is stronger or the metallicity is lower, the regions of significant $^{12}\textrm{CO}$ and $^{13}\textrm{CO}$ formation shift towards higher total column densities by a similar amount. Furthermore, the characteristic gas temperature of regions where CO reside are very similar in every simulation, providing a consistent condition for the chemical fractionation. If $^{12}\textrm{CO}$ number densities are provided \citep[e.g. from a simulation of a molecular cloud, such as in][]{GloverClark2012a}, and if we want to calculate the corresponding $^{13}\textrm{CO}$ number densities, then using a uniform $^{12}\textrm{CO}/^{13}\textrm{CO}$ ratio results in the $^{13}\textrm{CO}$ number densities being underestimated by as much as 30-40 per cent (see Table~\ref{13comass}). This can lead to errors of up to 50 per cent in some regions of the derived $^{13}\textrm{CO}$ integrated intensity maps (see Fig.~\ref{fig:intensity_dist}). \corr{The peak brightness temperature is affected by about 30 per cent in this case. However, the adopted isotope ratio has only a few per cent influence on the line width and therefore the derived velocity dispersion. } In section~\ref{sec:fitting} we derive a fitting formula based on the simulation results to address this issue in a computational cost efficient way (i.e. to infer the correct ratio without the need of the full fractionation chemistry model). When the fitting formula is applied, then the $^{13}\textrm{CO}$ column/number density and emission can be recovered with errors smaller than 10 per cent. If we have instead $^{13}\textrm{CO}$ column densities (e.g. derived from observations), then using a uniform $^{12}\textrm{CO}/^{13}\textrm{CO}$ ratio to convert them to $^{12}\textrm{CO}$ column densities may result in the overestimation of the $^{12}\textrm{CO}$ column densities by up to 50-60 per cent (see Table~\ref{12comass}). In section~\ref{sec:13coobs} we construct a fitting formula describing the correlation of the isotope ratio and the $^{13}\textrm{CO}$ column density. By applying this formula the errors could be reduced to ${\sim}10$ per cent. We note however, that the error introduced by using a fixed $^{12}\textrm{CO}/^{13}\textrm{CO}$ ratio may be smaller then other sources of errors in the determination of the $^{13}\textrm{CO}$ column density (such as the assumption that the excitation temperatures of $^{12}\textrm{CO}$ and $^{13}\textrm{CO}$ are the same and that the $^{13}\textrm{CO}$ emission is optically thin), which can lead to uncertainties of as much as a factor of 4 \citep[see][]{Padoan2000}. Nevertheless, in the $10^{14}\,\textrm{cm}^{-2} < N(^{13}\textrm{CO}) < 10^{16}\,\textrm{cm}^{-2}$ range, where the aforementioned assumptions are valid and the chemical fractionation is effective, the fitting formula provides more accurate conversion to $^{12}\textrm{CO}$ than the assumption of uniform scaling. The proposed fitting formulae are consistent with millimetre-wavelength $^{12}\textrm{CO} / ^{13}\textrm{CO}$ column density ratio measurements, and underestimate the ratio measured in ultraviolet absorption by a factor of 2-3 (see Fig.~\ref{fig:ratio_obs}). The reason for the discrepancy with the ultraviolet data is probably that the UV measurements are tracing a qualitatively different population of clouds, more consistent with PDR models (diffuse clouds or highly irradiated massive clumps) than with giant molecular clouds. \corr{Finally, we conclude that the fitting formulae proposed above are good representations of the $^{12}\textrm{CO}/^{13}\textrm{CO}$ isotopic ratio distributions of our hydro-chemical simulations and that they can be used to infer $^{13}\textrm{CO}$ properties from (magneto-) hydrodynamical simulations in a computationally cost-efficient manner, and more precise $^{12}\textrm{CO}$ column density estimates from millimetre-wavelength $^{13}\textrm{CO}$ observations, provided that the molecular clouds under investigation are similar to those presented here.} | 14 | 3 | 1403.4912 |
1403 | 1403.6917_arXiv.txt | Compact binary mergers are the strongest candidates for the progenitors of Short Gamma Ray Bursts (SGRBs). If a gravitational wave (GW) signal from the compact binary merger is observed in association with a SGRB, such a synergy can help us understand many interesting aspects of these bursts. We examine the accuracies with which a world wide network of gravitational wave interferometers would measure the inclination angle (the angle between the angular momentum axis of the binary and the observer's line of sight) of the binary. We compare the projected accuracies of GW detectors to measure the inclination angle of double neutron star (DNS) and neutron star-black hole (NS-BH) binaries for different astrophysical scenarios. We find that a 5 detector network can measure the inclination angle to an accuracy of $\sim 5.1 (2.2)$ degrees for a DNS(NS-BH) system at 200 Mpc if the direction of the source as well as the redshift is known electromagnetically. We argue as to how an accurate estimation of the inclination angle of the binary can prove to be crucial in understanding off-axis GRBs, the dynamics and the energetics of their jets, and help the searches for (possible) orphan afterglows of the SGRBs. | \subsection{Gamma Ray Bursts} Gamma Ray Bursts (GRBs) are classified as `long' and `short' based on the duration ($T_{90}$) in which 90\% of the total observed energy is emitted~\cite{ShGRBclass93} and the spectral type. Long duration GRBs correspond to $T_{90}>2 s$ and are spectrally soft whereas the short GRBs (SGRBs henceforth) correspond to $T_{90}<2s$ and are spectrally hard. Long duration GRBs occur at relatively high redshifts and in active star forming regions in the Galaxy and many of them have been associated with core-collapse supernovae~\cite{Hjorth2003dh}. On the other hand, the progenitors of SGRBs are not fully understood. They have relatively smaller redshifts and are seen with significant (tens of kpcs) off-sets with respect to the respective Galactic centers~\cite{FongBerger2013,ChurchEtal2011}. The inferred progenitor ages are typically a few Gyr pointing to older stellar populations~\cite{LeiblerBerger2010}. These features strongly support the hypothesis that they are due to the mergers of double neutron stars (DNS) or neutron star-black hole (NS-BH) binaries \cite{EichlerShGRB89,Narayan92}. Recent kilonova observation associated with the GRB130603B~\cite{KilonovaBerger2013,TanvirKilonova2013} reinforces this conjecture (see Refs.~\cite{ShGRBrevLee07,Nakar07Review,ShGRBrevBartos2013,BergerShGRBrev2013} for reviews on SGRBs and their association with gravitational waves (GWs).). The prompt emission and afterglows of the SGRBs share a lot of features of the long GRBs. This is because for both the bursts, the central engine is a black hole that accretes from the surroundings, powering a jet which produces prompt emission due to internal shocks and the afterglow by its interaction with the circumburst medium.\footnote{In the case of SGRBs, there may also be an intermediate product in the form of a hypermassive, highly magnetized neutron star which eventually collapses to a BH~\cite{RosswogEtal03,AloyEtal05,ShibataEtal05,Giacomazzo2010,Hotokezaka2011,Hotokezaka2013}.} Hence the standard fireball model~\cite{Rhoads99,SariEtal99}, which has been very successful in interpreting the long GRBs is used to interpret the SGRB data as well. But the central black hole for the these two bursts is formed via different channels. For long GRBs, the BH is formed by core collapse of a massive star, whereas for the SGRBs they may be formed by mergers of compact binaries. \subsection{GRB-GW association} If indeed the SGRBs are produced due to the mergers of compact binaries, there will also be an associated emission of GWs. If the burst happens sufficiently close-by, the corresponding GW signals may be detectable by the upcoming advanced GW interferometers such as advanced LIGO (aLIGO)\cite{aligo}, advanced Virgo~\cite{avirgo} and KAGRA~\cite{kagra}. The expected detection rate at design sensitivity for aLIGO is approximately 40(10) per year for DNS(NS-BH) sources~\cite{LSCrates}. The {\it root mean square} distance reach~\footnote{This should be distinguished from the horizon distance which is the distance to an optimally oriented binary. $D_{\rm horizon}=\frac{5}{2}D_{\rm rms}$ (see Eq.~(4.12) of \cite{DIS00}).} for DNS binaries with a signal to noise ratio (SNR) of 8 for aLIGO detector is $\simeq200$ Mpc ($z\simeq 0.05$) whereas for the NS-BH system (total mass $11 M_{\odot}$), it is roughly $1$ Gpc ($z\simeq 0.2$)(see Fig. 1 of ~\cite{AISS05} and rescale it to a SNR of 8). This distance reach will increase with the number of detectors ($n$) in the network, roughly as $\sqrt{n}$. Hence, for a 5 detector network the distance reach can be, roughly, twice the one given above. Hence the advanced GW detector era % carries the exciting prospect of detecting GW signals from compact binary mergers and detection of associated SGRBs (and their afterglows) by electromagnetic (EM) telescopes. Typical numbers of joint SGRB-GW events were estimated to be $\simeq 0.2-1(1-3) {\rm yr^{-1}}$ for a DNS(NS-BH) progenitor~\cite{DietzEtal2013}. This opens up lot of interesting studies that are possible where the synergy of GW and EM observations may lead to significant breakthroughs in the field of astronomy, specifically related to SGRBs. Such an association would confirm, beyond question, the compact binaries as progenitors of the SGRBs. GW observations would measure the (redshifted) masses of the binary to a few percent accuracy~\cite{AISS05} which will further enable us to distinguish between the DNS and NS-BH scenarios. One of the most exciting possibilities of such joint GW-EM observation is the estimation of Hubble constant without relying on the usual astronomical distance ladders, as suggested by Schutz~\cite{Schutz86} and explored in greater detail in Refs.~\cite{DaHHJ06,Nissanke09a}(see, also Refs.~\cite{,SathyaSchutzLivRev09,ShGRBrevBartos2013,BergerShGRBrev2013} for an overview of other ideas proposed in the literature). In this paper, we investigate a different aspect related to the possible GW measurement of the inclination angle (the angle between the line of sight of the observer and the angular momentum axis of the binary), (See also an another discussion by Chen and Holz ~\cite{ChenHolz2012} on the relation between jet opening angles of the SGRBs and the detection rate of NS-NS and NS-BH binary mergers in the advanced LIGO era). We discuss the accuracies with which advanced GW detectors would be able to measure the inclination angle of the binary for various astrophysical scenarios. The inclination angle is referred to as the {\it viewing angle} of the jet (angle between the jet axis and the line of sight of the observer) in the GRB literature, assuming the jets are launched along the orbital angular momentum axis, which is the same as the total angular momentum axis for nonspinning binaries or for binaries whose spins are aligned or anti-aligned with respect to the angular momentum axis. The viewing angle can play an important role in GRB modelling if the jets are pointed away from the observer (off-axis jets). We discuss how the SGRB-GW synergy can shed light on the geometry and energetics of the SGRBs A quick summary of the results are presented in Fig.~\ref{fig:summary} where the expected accuracy with which the inclination angle of the binary may be estimated for DNS and NS-BH progenitors using GW measurements for various GW detector combinations. The paper is organized as follows. In Sec.~\ref{sec:model} we discuss the data analysis technique employed in GW astronomy and the model of gravitational waveform we employ for our study. Section~\ref{sec:iota} discusses the accuracies with which GW observation will be able to extract the inclination angle of the binary. The implications of these measurements for SGRBs are discussed in Sec.~\ref{sec:GRB} and caveats of our model and future plans are discussed in Sec.~\ref{sec:caveats}. \begin{figure}\includegraphics[scale=0.4]{shGRB_NSBH_BNS} \caption{Estimated median errors in the measurement of the inclination angle of the binary (also the viewing angle of the SGRB jet) as a function of luminosity distance (bottom axis) and redshift (top axis) for different detector configurations, different astrophysical scenarios, and for DNS and NS-BH binaries. The triangles correspond to the astrophysical scenario where there is no EM information, and the circles correspond to 3D localized SGRBs (see text for details). The filled and open data points distinguish the 3 detector GW detector configuration LHV (two advanced LIGO detectors and one advanced Virgo) and a 5 detector configuration LHVIK (adding Japanese detector KAGRA and the proposed LIGO-India to LHV). Errors for the DNSs are characterized by the lines which are cut at 500 Mpc and those of NS-BH are characterized by those truncated at 1 Gpc. We have not quoted errors greater than 50 degrees in this plot as it may not be of any astrophysical relevance. }\label{fig:summary} \end{figure} | \end{figure} | 14 | 3 | 1403.6917 |
1403 | 1403.6310_arXiv.txt | \noindent The recent BICEP2 measurement of B-modes in the polarization of the cosmic microwave background suggests that inflation was driven by a field at an energy scale of $2\times 10^{16}$ GeV. I explore the potential of upcoming CMB polarization experiments to further constrain the physics underlying inflation. If the signal is confirmed, then two sets of experiments covering larger area will shed light on inflation. Low resolution measurements can pin down the tensor to scalar ratio at the percent level, thereby distinguishing models from one another. A high angular resolution experiment will be necessary to measure the tilt of the tensor spectrum, testing the consistency relation that relates the tilt to the amplitude. | The BICEP2 experiment~\cite{Ade:2014gua,Ade:2014xna} has detected the B-modes of polarization in the cosmic microwave background (CMB), a signature~\cite{Seljak:1996gy,Kamionkowski:1996zd} of primordial gravitational waves generated during inflation~\cite{Guth,Starobinsky:1979ty,Linde,Mukhanov:1981xt,Albrecht}. If these results are confirmed and are indeed due to inflation, then the energy scale responsible for the early epoch of acceleration is $2\times 10^{16}$ GeV, some 13 orders of magnitude above energy scales probed in the largest colliders today. The result is stunning with far-ranging implications. Here I focus on what comes next: what can we learn about this new physics from the next generation or two of CMB polarization experiments? Although this topic has been addressed before~\cite{Song:2003ca,Verde:2005ff,Baumann:2008aq,Easson:2010uw,Abazajian:2013vfg,Wu:2014hta}, we are now (probably) in a new era, in which we know the amplitude of the signal so are standing on firmer ground. My conclusions are that we are about to embark on a three-step journey that can potentially uncover the laws of physics at ultra-high energies using CMB polarization: \begin{itemize} \item The first step is to confirm the BICEP2 result and determine the amplitude of the gravitational wave signal. There are several ways to do this. The most comforting would be a detection by the Planck satellite of the first B-mode peak on large scales. BICEP2 is not sensitive to this reionization-induced signal, so a detection would complement and confirm with little doubt that we have indeed observed primordial gravitational waves. Another test would be a detection with one of the other ground-based experiments, especially if the measurement were made in a different part of the sky and at a different frequency. In short, there are three axes along which we can move to confirm the BICEP2 result: angular frequency to detect the distinctive two-humped signal, photon frequency to eliminate the possibility of foreground contamination, and sky coverage again to mitigate foregrounds. For the next two steps to move forward, this first stage needs to conclude with the removal of the tensions between different data sets that predate not only BICEP2, but even Planck~\cite{Hou:2012xq,Archidiacono:2013lva}. If these tensions remain, the resolution may require ways of understanding the responsible physics beyond those outlined below; e.g. bumps in the primordial power spectrum~\cite{Contaldi:2014zua,Miranda:2014wga}. I will have nothing to say here about this stage, except the obvious: it is very important. \item The second stage will be to measure the tensor to scalar ratio, $r$, and the spectral index of the scalar perturbations, $n_s$, with increasing accuracy. Models make predictions in the $n_s,r$ plane~\cite{Linde:1983gd,Freese:1990rb,Dodelson:1997hr,Silverstein:2008sg,McAllister:2008hb,Baumann:2008aq,Kaloper:2008fb,Kaloper:2011jz,Ade:2013uln}, and increasing precision can help identify the correct model. Indeed, a host of models are already ruled out if $r$ is close to 0.2 as suggested by the BICEP2 results. There are plans to reduce the errors on $n_s$ by a factor of five using galaxy surveys~\cite{Levi:2013gra}, and future CMB polarization experiments can reduce the error on $r$ to the percent level. As we will see (e.g., Fig.~\rf{contour_nc_r}), this can be done mostly by increasing the sky coverage (BICEP2 covers less than a percent of the sky) even with a relatively large beam. This is easier said than done of course because BICEP2 looked at a low-foreground region, so this next generation of experiments will likely need to be equipped with multiple frequency channels in order to disentangle the signal from the foregrounds. A very important physical question underlying model choice is why the ``simplest model,'' with the field driving inflation subject to a quadratic potential, seems to fit the data. Everything we know about effective field theory tells us that since the field traverses a large distance in Planck units, higher order terms should be generated in the effective action, completely changing the simple dynamics of an $m^2\phi^2$ term. Different solutions to this problem make different predictions in the $n_s,r$ plane,. Therefore, this ``{\it Simple fits but Simple doesn't make sense}'' quandary may be resolved by obtaining greater precision in the ($n_s,r$) plane. \item The third stage will be to measure the running of the spectral index of the tensor perturbations, $n_t$, and test the prediction that $n_t=-r/8$~\cite{Caligiuri:2014sla}. As we will see (e.g, Fig.~\rf{contour_c5_nt}), carrying out this program will require exquisite ``cleaning'' of the B-mode signal from lensing~\cite{Hu:2001kj,Seljak:2003pn,Smith:2010gu} and therefore will require small-scale resolution, low-noise, and large sky coverage. Even if all these are achieved, the conclusion that $n_t$ will be measured to be non-zero rests on the assumption that $r$ is large. If we find in Step 1 that $r=0.1$, it will become virtually impossible (Fig.~\rf{best_five2r}) to detect non-zero $n_t$ at even the 2-sigma level, although there is still an enormous amount of physics that can still be gleaned from high resolution polarization experiments. \end{itemize} It is unlikely that these steps will occur sequentially: we should expect progress on all three fronts over the coming decade. To quantify these conclusions, I project constraints from polarization experiments in the two dimensional $r,n_t$ plane. The errors on $C_l^{BB}$, sample variance and noise, are computed using the standard formula~\cite{Knox:1995dq,Wu:2014hta}. On small scales, this simple formula captures the noise reported by BICEP2~\cite{Ade:2014gua,Ade:2014xna} with sky coverage set to 384 square degrees; noise per square degree pixel set to $0.087\times \sqrt{2}\mu$K; and beam width equal to $30'$ FWHM. The formula underestimates the noise in BICEP2 at low $l$ probably because it does not account for low $l$ removal from filtering\footnote{Thanks to Chris Sheehy for suggesting this.}, so the estimates presented here may be a bit optimistic. On the other hand, the projected error on $r$ using this Fisher formalism is only 20\% smaller than that obtained by BICEP2 in their analysis. As we will see, a twenty percent difference does not matter much for the calculations in the next section, % where the marginalized constraints on $r$ are displayed, and in following section, % where the constraints on $n_t$ are examined, most of the weight comes from small scales where the formula agrees well with the BICEP2 errors. \newcommand\fsky{f_{\rm sky}} | 14 | 3 | 1403.6310 |
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1403 | 1403.6193.txt | { Constraining the gas and dust disk structure of transition disks, particularly in the inner dust cavity, is a crucial step toward understanding the link between them and planet formation. HD~135344B is an accreting (pre-) transition disk that displays the CO 4.7~$\mu$m emission extending tens of AU inside its 30 AU dust cavity. } {We constrain HD~135344B's disk structure from multi-instrument gas and dust observations.} { We used the dust radiative transfer code MCFOST and the thermochemical code ProDiMo to derive the disk structure from the simultaneous modeling of the spectral energy distribution (SED), VLT/CRIRES CO~P(10)~4.75~$\mu$m, Herschel/PACS [\ion{O}{i}]~63~$\mu$m, Spitzer-IRS, and JCMT~$^{12}$CO~J=3-2 spectra, VLTI/PIONIER H-band visibilities, and constraints from (sub-)mm continuum interferometry and near-IR imaging. } { We found a disk model able to describe the current gas and dust observations simultaneously. This disk has the following structure. (1) To simultaneously reproduce the SED, the near-IR interferometry data, and the CO ro-vibrational emission, refractory grains (we suggest carbon) are present inside the silicate sublimation radius ($0.08<R<0.2$ AU). (2) The dust cavity (R$<$30 AU) is filled with gas, the surface density of the gas inside the cavity must increase with radius to fit the CO ro-vibrational line profile, a small gap of a few AU in the gas distribution is compatible with current data, and a large gap of tens of AU in the gas does not appear likely. (4) The gas-to-dust ratio inside the cavity is $>$ 100 to account for the 870 $\mu$m continuum upper limit and the CO P(10) line flux. (5) The gas-to-dust ratio in the outer disk ($30<R<200$~AU) is $<$ 10 to simultaneously describe the [\ion{O}{i}]~63~$\mu$m line flux and the CO~P(10) line profile. (6) In the outer disk, most of the gas and dust mass should be located in the midplane, and a significant fraction of the dust should be in large grains. } { Simultaneous modeling of the gas and dust is required to break the model degeneracies and constrain the disk structure. An increasing gas surface density with radius in the inner cavity echoes the effect of a migrating jovian planet in the disk structure. The low gas mass (a few~Jupiter~masses) throughout the HD~135344B disk supports the idea that it is an evolved disk that has already lost a large portion of its mass. } | Observations of young stars of different ages reveal that protoplanetary disks evolve from optically thick, gas-rich disks to optically thin, gas-poor debris disks \citep[e.g, see review by][]{WilliamsCieza2011}. The transition between these two classes of objects is believed to occur relatively fast (10$^5$ years) compared to the disk's life time (5 - 10 Myr) \citep[e.g.,][]{Cieza2007,Damjanov2007,CurrieAguilar2011}. Infrared space observatories have unveiled several young stars with spectral energy distributions (SED) characterized by IR-excess at $>$ 10 $\mu$m and significantly reduced excesses at shorter wavelengths \citep[e.g.,][]{Strom1989, Calvet2005, Sicilia-Aguilar2006, Hernandez2006, Espaillat2007, Fang2009, Merin2010, Rebull2010, Cieza2010}. These sources are commonly named ``transition disks", because they are believed to be the objects in the transition phase between a young star with an optically thick gas-rich disk and a star with an optically thin gas-poor debris disk. Generally, the lack of strong emission in the near or mid-IR SED is interpreted as evidence of a gap or cavity from a few up to tens of AU in the disk. Follow-up imaging of several transition disks with sub-mm interferometers have confirmed that transition disks indeed have a deficit of sub-mm continuum emission at a few mJy levels inside tens of AU \citep[e.g.,][]{Pietu2007,Brown2009,Hughes2009, Andrews2011, Isella2012}, thus providing further evidence for such inner disk cavities, at least on the large dust grains. Several scenarios have been discussed in the literature to explain the presence of a cavity and the transitional disk SED shape: grain growth \citep[e.g.,][]{DullemondDominik2005,Birnstiel2012}, migrating giant planets that opened a gap \citep[e.g.,][]{Varniere2006,Zhu2011,Dodson-RobinsonSalyk2011}, dust filtration by a giant planet \citep[e.g.,][]{Rice2006, Zhu2012,Pinilla2012}, disk dissipation due to a photoevaporative disk wind \citep[e.g.,][]{AlexanderArmitage2007,Owen2012}, and dust free inner holes due to radiation pressure \citep[][]{ChiangMurray-Clay2007}. One important step toward understanding the transitional disk phenomenon is to understand how the gas and dust structures compare. Do they follow each other in density structure? Do they thermally de-couple? Does the gas-to-dust ratio stay constant as a function of the radius? Several young stars with transition disks are emission-line stars that exhibit signs of accretion \citep[e.g.,][]{Najita2007,Muzerolle2010,Cieza2012b}. As a consequence, the cavities imaged in the sub-mm should have gas as the material from the optically thick outer disk flows through the cavity to be accreted by the central object \citep[e.g.,][]{LubowAngelo2006}. In fact, several transition disks display emission of CO at 4.7$~\mu$m \citep[e.g.,][]{Salyk2009} a common tracer of warm gas in the inner disk \cite[e.g.,][]{NajitaPPV}. Furthermore, in a few transition disks, CO ro-vibrational emission has been spatially resolved up to distances of tens of AU \citep[][]{Pontoppidan2008}. Moreover, recent studies of dust-scattered light in the near-IR have detected emission from small dust grains inside the cavities that were previously imaged with sub-mm interferometry, with the additional and remarkable property that no sharp edge is seen at the location of the sub-mm cavity's inner radius \citep[e.g.][]{Muto2012,Garufi2013}. This suggests that what is observed is likely a change in the dust size distribution inside the cavity \citep[][]{Dong2012}. Near-IR interferometry provides evidence of dust emission (and inhomogeneities) inside the dust cavity of transition disks \cite[e.g.,][]{Kraus2013}. Furthermore, recent ALMA observations have spatially resolved gas inside the cavity of transition disks \citep[][]{Casassus2013,Bruderer2013}. In summary, the observational evidence indicating that the dust cavities of transitions disks are indeed {\it not empty} has been steadily growing during past years. This opens interesting questions such as how much gas and dust are present in the cavity; what is their distribution and chemical composition; if dust is present, what is its size distribution; if there is a gap in the gas or in the dust, how large is it? { The goal of this paper is to address the problem of the gas and dust structures in transition disks by performing a detailed study of the transition disk HD~135344B (SAO 206462). Our aim is to constrain the disk structure by modeling simultaneously and in a coherent way multiwavelength \& multi-instrument gas and dust observations of the inner and the outer disk. Most specifically, we investigate gas and dust disk content and structure inside the cavity by simultaneously modeling the SED, new VLTI-PIONIER near-IR interferometry data, the CO~4.7~$\mu$m emission, and constraints from sub-mm continuum interferometry. Furthermore, we seek to constrain the gas mass and the gas-to-dust ratio in the outer disk by employing the SED, the [\ion{O}{I}] $63~\mu$m and $145~\mu$m lines, and (sub-)mm observations}. HD~135344B is an accreting \cite[$2\times10^{-8}$ M$_{\odot}$/yr,][]{Sitko2012} F4V young star \citep[8$^{+8}_{-4}$ Myr,][]{vanBoekel2005} that has a ``pre-transition disk" (i.e., a transition disk with near-IR excess). It has the remarkable characteristic of exhibiting CO~4.7~$\mu$m emission extending tens of AU inside the 45 AU sub-mm dust cavity \citep[][]{Pontoppidan2008} and scattered light emission from small dust grains down to 28 AU \citep[][]{Muto2012,Garufi2013}. HD~135344B display spiral structures \citep[][]{Muto2012,Garufi2013} and asymmetries in its outer disk dust emission \citep[][]{Brown2009,Andrews2011,Perez2014}. It is a transition disk that exhibits variability in its near-IR SED and in optical and near-IR line emission on time scales of months \citep[][]{Sitko2012}. It is a source on which [\ion{O}{I}] $63~\mu$m emission has been detected in Herschel/PACS observations \citep[][]{Meeus2012}. Finally, it is a transition disk without close-in low-mass stellar companions \citep[][]{Vicente2011}. We start the paper by a brief summary of current observational constraints on the disk of HD~135344B in Section 2. Then, in Section 3, we describe the modeling tools and the general modeling procedure. In Section 4, we present and discuss the different models that were tested. In Section 5, we examine the disk constraints derived from the final disk model. In Section 6, we discuss our results in the context of the study of transitional disks and planet formation. Finally a summary and conclusions are presented in Section 7. \vspace{-0.3cm} | We conducted a modeling project aimed at constraining the gas mass and the gas and dust disk structure of the transition disk HD~135344B from multi-instrument and multiwavelength observations of gas and dust. % We found that the previously suggested inner disk structure { \citep{Brown2007}}, namely a narrow dust inner disk of from 0.18 to 0.45 AU followed by a large 45 AU dust gap replenished with gas, fails to reproduce the CO ro-vibrational emission observed as the line profile produced from this disk model is a broad double peak. We have found a disk model that is able to reproduce current observational constraints. This disk is composed of three zones: \vspace{-1mm} \begin{enumerate} \item A first zone between 0.08 and 0.2 AU composed of small carbonaceous grains (and gas) with a total dust mass of a few 10$^{-12}$ M$_{\odot}$ { (a few solar abundances of carbon).} The presence of this inner carbonaceous grains provides \begin{itemize} \item[{\it a)}] {a fit to the near-IR H-band visibilities and closure phases,} \item[{\it b)}] {a fit of the near-IR SED, while allowing the warm CO at several AU to emit and contribute to the $4.7~\mu$m line profile,} \item[{\it c)}] an agreement with the higher temperatures ($T>$1\,500 K) expected in this zone. \end{itemize} \item { A second zone extending from 0.2 to 30 AU (i.e., the dust cavity) replenished with gas ($10^{-5}$-$10^{-4}$ M$_{\odot}$) with a surface density increasing as a function of the radius and dust mass of astronomical silicates of maximum $10^{-7}$ M$_{\odot}$}. An increasing surface density profile with radius is required to fit the shape of the CO ro-vibrational emission lines. The fit to the SED constrained the scale height between 0.09 and 0.13 at 10 AU with a flaring exponent 1.12. The gas-to-dust ratio in this zone is larger than 100, however, the exact value is not well constrained. We found models up to gas-to-dust ratios of 15\,000 that are consistent with the observations either by decreasing the silicate dust mass by a factor 100, or by increasing both the gas mass by a factor of a few and the power-law exponent of the surface density distribution. The upper bound to the gas mass at $R<30$ AU is given by the flux of the [\ion{O}{i}] 63~$\mu$m, combined with the requirement that the surface density of the inner disk should be equal or lower than the surface density of the outer disk at 30 AU. The dust surface density at R$<$30 AU is lower than the one expected from extrapolating the dust surface density from the outer disk. This zone can contain up to a 3\% in mass ($\sim10^{-9}$ M$_{\odot}$) in carbon grains and keep the fit to the SED. \item A third zone from 30 AU to 200 AU (the outer disk) with astronomical silicate grains, a dust mass of $2\times10^{-4}$ M$_{\odot}$, a gas mass $10^{-4}-10^{-3}$ M$_{\odot}$ { (gas-to-dust ratio $<10$)}, surface density exponent of -1.0, and flaring of 1.0. In this zone large ($0.1<a<1000~\mu$m) and small ($0.1<a<10~\mu$m) dust grains have different radial and vertical spatial distributions. Large grains are located at $45<R<200$ AU in a disk with scale height 0.05 and 75\% of the gas and dust mass. Small grains are located at $30<R<200$ AU in a disk with higher scale height (0.09 to 0.13, i.e., the same H/R as zone 2) and 25\% of the gas and dust mass. The vertical structure of the outer disk echoes the expected effects of dust growth and settling. A gas-to-dust ratio much lower than 100 in the outer disk is required to fit the [\ion{O}{i}] 63~$\mu$m line flux and to reproduce the CO ro-vibrational line profile simultaneously.\\ \end{enumerate} { The models suggest that the best fit to the gas observations in HD~135344B is provided by a disk in which the gas surface density and the scale height have no large discontinuities at 30 and 45 AU. In other words, there is no large gap in the gas distribution of HD~135344B. The cavity observed in the near-IR and sub-mm is replenished by gas and (some) dust. The presence of a small gap of a few AU in the gas is consistent with current data, a large gap in the gas of tens of AU does not appear likely. The global gas-to-dust ratio, i.e., integrated over the full disk, is much lower than 100. This provides further evidence that HD~135344B is an evolved protoplanetary disk that has already lost a large amount of its gas mass. % The disk structure proposed for HD~135344B could be applied to other pre-transition disks with CO ro-vibrational emission extending to several AU. The increasing surface density profile of the gas in the inner disk, the difference in the radial distribution of large and small grains in the outer disk, and a small gap in the gas between the inner and outer disk are compatible with the dynamical interaction of a single Jovian planet and the disk. However, we should be cautious because other mechanisms could be responsible for the gas and dust distribution observed in HD~135344B. The current ensemble of observations show that the transitional disk features (i.e., gap) observed in the continuum reflect only changes on the distribution of large and small dust particles in the disk. The gas has a spatial distribution that can differ from that of the dust, particularly that of large grains observed in the sub-mm. Furthermore, we find that the gas distribution and mass in transition disks can be very different from primordial disks}. Our modeling results predict that gas and (some) dust emission should be detected inside the cavity with high-sensitivity and high-spatial resolution sub-mm observations. ALMA measurements of continuum emission inside 30 AU would be a great help in determining the dust size distribution and dust mass inside the cavity, quantities that at present are free parameters in the model. ALMA gas observations at high spatial resolution, such as the CO 6-5 line, used in conjunction with other gas tracers will enable the gas mass inside the cavity, and the gas mass and surface density in the outer disk to be better constrained. {Future detailed dynamical disk models that include the dynamics of the gas, the interaction between the gas and the dust, the photochemistry of the disk, and carbon destruction would be a big help for establishing whether the carbon innermost disk suggested by our modeling is viable or whether including refractory grains different from carbon is necessary.} In this paper we showed that the simultaneous modeling of gas and dust observations is required to address the problem of the dust and gas structure of protoplanetary disks. | 14 | 3 | 1403.6193 |
1403 | 1403.3229_arXiv.txt | We present \src, a new case of radio-loud narrow line Seyfert~1 (RL NLS1) with a relatively high radio power (P$_{1.4 GHz}$=2.1$\times$10$^{25}$ W Hz$^{-1}$) and large radio-loudness parameter (R$_{1.4}$=600$\pm$100). The radio source is compact with a linear size below $\sim$1.4 kpc but, contrary to most of the RL NLS1 discovered so far with such a high R$_{1.4}$, its radio spectrum is very steep ($\alpha$=0.93, S$_{\nu}\propto\nu^{-\alpha}$) and not supporting a ``blazar-like'' nature. Both the small mass of the central super-massive black-hole and the high accretion rate relative to the Eddington limit estimated for this object (3.2$\times$10$^{7}$ M$_{\sun}$ and 0.27, respectively, with a formal error of $\sim$0.4 dex on both quantities) are typical of the class of NLS1. Through a modelling of the spectral energy distribution of the source we have found that the galaxy hosting \src\ is undergoing a quite intense star-formation (SFR=50 M$_{\sun}$ y$^{-1}$) which, however, is expected to contribute only marginally ($\sim$1 per cent) to the observed radio emission. The radio properties of \src\ are remarkably similar to those of compact steep spectrum (CSS) radio sources, a class of AGN mostly composed by young radio galaxies. This may suggest a direct link between these two classes of AGN, with the CSS sources possibly representing the misaligned version (the so-called ``parent population'') of RL NLS1 showing blazar characteristics. | Narrow-line Seyfert 1 (NLS1) galaxies have been identified as a peculiar sub-class of AGN on the basis of their relatively narrow ($\leq$2000 km s$^{-1}$) Balmer lines, relatively weak [OIII]$\lambda$5007\AA\ emission compared to the Balmer lines ([OIII]$\lambda$5007\AA/H$\beta<$3) and strong optical FeII (\citealt{Osterbrock1985, Goodrich1989, Pogge2000, Pogge2011, Veron-Cetty2001}). The presence of strong iron emission and the low [OIII]$\lambda$5007\AA/H$\beta$ flux ratio suggest that NLS1 are not obscured objects like type~2 AGN, where the observed narrow emission lines are produced in regions (the Narrow Line Region, NLR) that are located far from the nucleus. In NLS1, instead, the observed Balmer lines are likely produced in the Broad Line Region (BLR) clouds that are located close ($<$0.1 pc) the accreting super-massive black-hole (SMBH). If the BLR traces the potential well of the central SMBH, a low velocity dispersion is the result of a combination of a relatively small SMBH mass with a high accretion rate, close to the Eddington limit (L$_{Edd}$). Indeed, NLS1 are usually characterized by SMBH masses between 10$^{5}$ and 10$^{8}$ M$\sun$ and accretion rates larger than $\sim$0.1L$_{Edd}$ (e.g. \citealt{Marziani2001, Boroson2002, Greene2004, Collin2004, Grupe2004a, Botte2004, Grupe2004, Zhou2006, Whalen2006, Komossa2006, Yuan2008} but see also \citealt{Decarli2008, Marconi2008, Calderone2013}). An interesting characteristic of NLS1 is their preference of being radio-quiet (RQ) AGN. To date, only $\sim$50 radio-loud (RL) NLS1 have been discovered (\citealt{Foschini2011}). \citet{Komossa2006} has computed the fraction of radio-loud objects within the NLS1 population and compared it to the fraction of radio-loud sources within the population of broad line Seyfert~1 (BLS1) belonging to the same catalogue, and they found a significant difference ($\sim$7 per cent of radio-loud AGN, among NLS1, and 20 per cent among BLS1). On a more general basis, the fraction of radio-loud AGN seems to depend both on the mass of the central SMBH (e.g. \citealt{Chiaberge2011}) and on the accretion rate relative to the Eddington limit: large SMBH masses combined to low accretion rates are believed to give the highest probability of producing AGN with large values of radio-loudness while AGN characterized by high accretion rates and small SMBH masses, like NLS1, seem to be less effective in producing radio-loud sources (e.g. \citealt{Lacy2001, Sikora2007}). A second characteristic of the few RL NLS1 discovered so far is that most of them present a compact radio emission (linear size below a few kpc, e.g. \citealt{Doi2012} and references therein). Among the most radio-loud NLS1 there are many cases showing strict similarities with the class of blazars (BL Lac objects and flat spectrum radio quasar, FSRQ): like blazars, these RL NLS1 show a flat or inverted radio spectrum, high brightness temperatures ($T_b>$10$^7$ K, e.g. \citealt{Doi2013, Yuan2008}) and they are detected in gamma-rays by {\it Fermi}-LAT (\citealt{Abdo2009, Abdo2009a, Foschini2011}). Since blazars are usually believed to be radio galaxies whose relativistic jets are pointing towards the observer (e.g. \citealt{Urry1995}), a possible conclusion is that also some of the RL NLS1 discovered so far are ``oriented'' and relativistically beamed sources. In this case we must expect a large number of mis-oriented and unbeamed sources, the so-called parent population, that in the standard beaming model is constituted by the class of lobe-dominated radio galaxies. To date, however, only in six RL NLS1 an extended emission has been detected (\citealt{Whalen2006, Anton2008, Doi2012}) and only one RL NLS1 (SDSSJ120014.08--004638.7) is a lobe-dominated FR~II radio galaxy (\citealt{Doi2012}). It is not clear whether the lack of RL NLS1 in radio-galaxies is suggesting an intrinsic difference with respect to the other RL AGNs or if it is just a selection effect: if the BLR has a disk-like geometry, as suggested by some authors (e.g. \citealt{Pozonunez2013}, \citealt{Decarli2008}), then selecting AGN with the narrowest emission lines may preferentially lead to the selection of face-on systems and, therefore, with blazar-like radio morphology, in case of radio-loud AGNs. Alternatively, RL NLS1 are intrinsically different from more powerful RL AGN and lacking any extended emission. In this case, the mis-oriented population could appear as NLS1 with very weak (since unbeamed) radio cores i.e. sources classified as RQ NLS1 (\citealt{Foschini2013d}). Another interesting possibility is that RL NLS1 are young or ``frustrated'' radio-galaxies that have not yet formed (or they are not able to form) radio lobes on large scales. This possibility is suggested by the similarities found between the class of RL NLS1 and that of the compact steep-spectrum (CSS) sources or the gigahertz-peaked spectrum (GPS) sources (e.g. \citealt{Oshlack2001, Gallo2006, Komossa2006, Yuan2008}). Clearly, a deeper investigation of the RL NLS1 and, in particular, of the candidates for the parent population is mandatory to discriminate among the various possibilties and to unveil the true nature of this class of AGN. To date, only a few RL NLS1 with non-blazar properties (e.g. with steep radio spectra) have been discovered and studied. This number reduces to just few units if we require a value of the radio loudness parameter that unambiguously sets them in the class of radio-loud AGN (e.g. R$_{1.4 GHz}>$100, \citealt{Yuan2008, Komossa2006}). In this paper, we present the discovery of a rare example of RL NLS1 (SDSSJ143244.91+301435.3, z=0.355) with a radio-loudness parameter well in the radio-loud regime (R=160, computed at 5~GHz, and R$_{1.4}\sim$600, computed at 1.4~GHz) and showing non-blazar properties. The study of this newly discovered object may help in shedding light on the possible nature of the parent population of this peculiar class of AGN. In Section~2 we discuss the optical classification on the basis of the SDSS spectrum, while in Section~3 we study the properties of the host galaxy. In Section~4 we estimate the physical properties of the central SMBH, i.e. its mass and the accretion rate normalized to the Eddington limit. In Section~5 we analyse the radio data trying to disentangle the contribution from different sources present in the field while in Section~6 we search for high-energy (gamma-ray and X-ray) emission from \src\ by exploiting the existing public catalogues (ROSAT, XMM-{\it Newton}, {\it Fermi}-LAT). Discussion and conclusions are reported in Section~7. Along the paper we assume a flat $\Lambda$CDM cosmology with H$_0$=71 km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\Lambda}$=0.7 and $\Omega_{M}$=0.3. Errors are given at 68 per cent confidence level. \begin{figure} \centering \includegraphics[width=6.5cm, angle=-90]{opt_spectrum.eps} \caption{Optical (SDSS) spectrum in the observer's frame of SDSSJ143244.91+301435.3 (z=0.355). The strongest emission lines are labeled. } \label{opt_spectrum} \end{figure} \begin{figure} \centering \includegraphics[width=6.5cm, angle=-90]{fe_fit.eps} \caption{Result of the spectral analysis around the H$\beta$ region aimed at estimating the intensity of the iron emission at 4570\AA. The dashed line (blue in the electronic version) is the fitted continuum while the solid line (red in the electronic version) indicates the continuum+iron emission based on the iron template presented in \citet{Veron-Cetty2004}.} \label{fe_fit} \end{figure} \begin{figure} \centering \includegraphics[width=6.5cm, angle=-90]{hb_bestfit.eps} \caption{Iron and continuum subtracted spectrum around the H$\beta$ region. The results of the fitting procedure described in the text are also shown. Dotted lines are the narrow components (blue in the electronic version), dashed lines (red in the electronic version) represent the two broad components of H$\beta$ and the broad blue wings of the two [OIII] lines, the solid line (green in the electronic version) represents the total fit. } \label{iron_subtracted} \end{figure} | SDSSJ143244.91+301435.3 is one of the few (13 in total discovered so far) radio-loud NLS1 with a radio-loudness parameter (R$_{1.4}$) of the order of 500 or greater. The radio emission, however, is different from that usually observed in RL NLS1 with such a high value of R, since it shows a steep ($\alpha$=0.93) spectrum. Based on the size of the radio emission ($\leq$1.4 kpc) this object is a compact radio source. The steep radio spectrum and the lack of strong variability and polarization, instead, disfavor the hypothesis that the observed compactness is due to the orientation of a relativistic jet towards the observer, as in blazars. This makes SDSSJ143244.91+301435.3 one of the few examples of RL NLS1 with non-blazar properties and sets it as an interesting case to study the parent population of these sources. SDSSJ143244.91+301435.3 hosts a relatively low-mass SMBH (3.2$\times$10$^7$ M$_{\sun}$) accreting at a rate close to the Eddington limit (Eddington ratio of 0.27), in agreement with what is typically found in both RQ and RL NLS1 objects. We have also investigated the host galaxy properties through a modelling of the SED. We have found that the AGN dominates the emission in the visible range and at 2-3 $\,\mu$m while the host galaxy shows up in the 0.6-2 $\,\mu$m spectral range (we also detect a hint of extension in the $i$-band image) and in the MIR ($>$10 $\,\mu$m). The best modelling is obtained using a star-forming galaxy template (M82) with a SFR, estimated from the IR emission, of $\sim$50 M$_{\sun}$ y$^{-1}$. The presence and the intensity of the SF is in agreement with what is usually observed in radio-quiet NLS1 (e.g. \citealt{Sani2010}). Given the presence of the SF we have evaluated whether this could explain the observed radio emission but we have concluded that the SF is expected to contribute only marginally ($\sim$1 per cent) to the observed radio fluxes. This result, combined to the large value of the radio-loudness parameter, confirms the presence of a jet in \src. The compact morphology and the steep radio spectrum are characteristics remarkably similar to those observed in the CSS radio sources. It is nowadays clear that the large majority of CSS (and GPS) sources are young objects in which the radio emitting plasma is digging its way through the ISM of the parent galaxy. Observations of hot-spot advance speed (e.g. \citealt{Polatidis2003} and references therein) and radiative ages (e.g. \citealt{Murgia2003, Murgia1999}) for a few tens of objects show typical ages of the order of $10^3 - 10^5$ years for GPS and CSS objects respectively. The association of some RL NLS1 with CSS sources has been already proposed by other authors (\citealt{Moran2000, Oshlack2001, Gallo2006, Komossa2006}). \citet{Yuan2008} also stressed the similarities between RL NLS1 and CSS sources and suggest that RL NLS1 with a flat radio spectrum could be CSS sources whose jet is pointing towards the observer. Another interesting characteristic that makes \src\ similar to CSS sources is the presence of strong blue wings in the [OIII]$\lambda\lambda$4959\AA,5007\AA\ doublet, with FWHM of the order of 1150-1350 km s$^{-1}$ and offsets, with respect to the core of the lines, of $\sim$250-270 km s$^{-1}$. The presence of these blue-shifted and broad wings is very common in compact radio objects, including CSS sources (\citealt{Gelderman1994, Holt2008, Kim2013}) and it is usually considered as a signature of the impact of a radio jet on the interstellar medium of the host galaxy; a jet-driven outflow is often invoked as the best explanation for the observed blue-shifted, broad wings of the narrow lines (\citealt{Holt2008, Nesvadba2011, Kim2013}). Therefore, in sources like \src\ we might be witnessing the AGN-induced feedback acting in the early stages of the evolution of a radio source. The hypothesis of CSS sources as the parent population of RL NLS1 would naturally explain why RL NLS1 have been detected as compact radio systems without significant extended emission (except for very few cases). In this framework, \src\ would be oriented at an intermediate angle between face-on and edge-on systems, the former being the flat-spectrum blazar-like RL NLS1 and the latter appearing as CSS sources with an obscured optical spectrum (i.e. type~2 AGN). In the first case, it is difficult to distinguish the CSS source morphology, being the radio emission dominated by the beaming effects; in the second class of sources it is difficult, if not impossible, to establish the NLS1 nature, due to the presence of obscuration. Due to the particular orientation of the source, instead, a detailed radio follow-up of \src\ at m.a.s. resolution can provide important pieces of information about the intrinsic (i.e. not affected by beaming) radio properties of RL NLS1. A systematic search for sources like \src, where both NLS1 and CSS characteristics can be observed, and a radio follow-up of all these objects is then mandatory to confirm the possible link between CSS sources and NLS1 on a firm statistical basis. | 14 | 3 | 1403.3229 |
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