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1403 | 1403.7646_arXiv.txt | The hunt for Earth analogue planets orbiting Sun-like stars has forced the introduction of novel methods to detect signals at, or below, the level of the intrinsic noise of the observations. We present a new global periodogram method that returns more information than the classic Lomb-Scargle periodogram method for radial velocity signal detection. Our method uses the Minimum Mean Squared Error as a framework to determine the optimal number of genuine signals present in a radial velocity timeseries using a global search algorithm, meaning we can discard noise spikes from the data before follow-up analysis. This method also allows us to determine the phase and amplitude of the signals we detect, meaning we can track these quantities as a function of time to test if the signals are stationary or non-stationary. We apply our method to the radial velocity data for GJ876 as a test system to highlight how the phase information can be used to select against non-stationary sources of detected signals in radial velocity data, such as rotational modulation of star spots. Analysis of this system yields two new statistically significant signals in the combined Keck and HARPS velocities with periods of 10 and 15~days. Although a planet with a period of 15~days would relate to a Laplace resonant chain configuration with three of the other planets (8:4:2:1), we stress that follow-up dynamical analyses are needed to test the reliability of such a six planet system. | 14 | 3 | 1403.7646 |
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1403 | 1403.7193_arXiv.txt | The Standard Model has been completed by the discovery of an apparently elementary Higgs boson at the LHC. On one hand, the absence of evidence for additional degrees of freedom at the LHC challenges many proposals for new weak-scale physics beyond the Standard Model. On the other hand, the recent discovery of primordial tensor modes in the CMB by BICEP2 \cite{bicep2} points to the existence of new physics at a scale that suggestively coincides with apparent gauge coupling unification in supersymmetric extensions of the Standard Model \cite{Dimopoulos:1981yj}. That the scale indicated by cosmological observations coincides with the scale indicated by low-energy observations is extremely suggestive. In this paper we pursue the idea that cosmology may provide even more concrete evidence for the existence of supersymmetry (SUSY) well above the weak scale.\footnote{In some string models, having $m_{\frac{3}{2}} < H$ causes problems for moduli stabilization \cite{Kallosh:2004yh}, which some authors take as evidence against low scale SUSY. We will ignore such concerns here. } Cosmological inflation \cite{inflation} offers a novel opportunity to search for SUSY in the universe. The discovery of primordial tensor modes in the CMB by BICEP2 \cite{bicep2} strongly supports the idea that inflation occurred at very high energies. For the reported central value of $r= 0.2^{+0.07}_{-0.05}$, the inflationary Hubble scale is given by $H \sim 1.1 \times10^{14}$ GeV. Since any field with mass less than the inflationary Hubble scale can be produced during inflation, cosmological observables are sensitive to particles produced at these incredible energies. Although the potential reach in energy of inflation is well-known, it has been less appreciated in the particle physics community that cosmological observables can directly test the presence of additional particles and interactions at these scales (see \cite{Abazajian:2013vfg} for a recent review). One crucial observation is the single-field consistency condition \cite{Maldacena:2002vr,Creminelli:2004yq}, which states that if inflation is described by a single degree of freedom then the bispectrum of the scalar curvature perturbation, $\zeta$, satisfies \beq\label{equ:singlecc} \lim_{\k_3 \to 0} \langle \zeta_{\k_1} \zeta_{\k_2} \zeta_{\k_3} \rangle' \to P_\zeta(k_1) P_\zeta (k_3) \left[ (n_s -1) + {\cal O}(k_3^2) \right] \ . \eeq Deviations from the consistency condition offer a relatively clean method for detecting additional fields present during inflation. The most commonly studied deviation is the case of local non-Gaussanity, where $(n_s - 1) \to \fNL^{\rm local}$, which is most easily produced by additional {\it massless} scalars. On the other hand, {\it massive} scalars with $0 < m \leq \tfrac{3}{2} H$ give rise to a bispectrum with soft limit \cite{Chen:2009zp} \beq\label{equ:multi} \lim_{\k_3 \to 0} \langle \zeta_{\k_1} \zeta_{\k_2} \zeta_{\k_3} \rangle' \to P_\zeta(k_1) P_\zeta (k_3) \left[ (n_s -1) + {\cal O}(k_3^{\alpha}) \right] \ , \eeq where $\alpha \equiv \tfrac{3}{2} - \sqrt{\tfrac{9}{4} - \frac{m^2}{H^2}}$. Measuring $\alpha < 2$ both tells us that there is an extra degree of freedom and indicates its mass\footnote{Strictly speaking, weakly coupled massive particles only produce $\alpha \leq \tfrac{3}{2}$. Taking $m > \tfrac{3}{2} H$ does not extend this limit, as these massive fields can be integrated out, up to exponentially suppressed contributions \cite{Chen:2012ge}. There is no obstacle to producing the full range $0 \leq \alpha< 2$ with additional fields, as was demonstrated concretely in \cite{Green:2013rd}. } during inflation. The above phenomenon provides a novel search technique for supersymmetry at high scales \cite{Baumann:2011nk}. Although the non-observation of superpartners at the LHC is beginning to challenge scenarios of weak-scale supersymmetry, there remains strong motivation for so-called {\it split supersymmetry} scenarios where most or all superpartners lie outside the reach of the LHC \cite{split, minisplit}. If this is the course Nature has chosen, verifying the existence of supersymmetry at high scales requires new experimental probes. Provided that more direct sources of SUSY breaking scale are below the inflationary Hubble parameter, then the dominant source of SUSY breaking during inflation is set by the curvature, namely $H$. As a result, we expect to find additional scalar particles with masses set by $H$ that naturally produce signatures at a detectable level \cite{Baumann:2011nk, Assassi:2013gxa}. A detection of $\alpha \sim 1$ would then provide tantalizing evidence that SUSY is relevant to our universe, even if it is never probed directly at the LHC (see also \cite{Iliesiu:2013rqa} for a different approach). In this paper, we will explore the capability of cosmological observables to shed light on SUSY at high scales. In section \ref{sec:susy} we review the low-energy evidence in support of supersymmetry at high scales, including the success of precision gauge coupling unification and the observed Higgs mass. In particular, the observed Higgs mass provides a suggestive upper bound on the present scale of SUSY breaking. In section \ref{sec:reach}, we then discuss the reach of cosmological observations in terms of the scale of SUSY breaking. We will discuss the assumptions that go into the predicted signals and how these compare with existing indirect probes. In section \ref{sec:forecast}, we will forecast our ability to detect $0 < \alpha < 2$ in an ideal 3d experiment, with an eye towards large scale structure surveys. In section \ref{sec:fakes}, we present possible alternative explanations of such a signal and how one could try to distinguish them. We conclude in section \ref{sec:outlook} with a discussion of the prospects for observation. Although the BICEP2 measurement of primordial tensor modes provides strong motivation, our study remains relevant irrespective of future changes in the central value of $r$. | 14 | 3 | 1403.7193 |
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1403 | 1403.6434.txt | {The Triplet extension of the MSSM (TMSSM) alleviates the little hierarchy problem and provides a significant enhancement of the loop-induced diphoton rate of the lightest CP-even Higgs h. In this paper we pursue the analysis of the TMSSM Higgs phenomenology by computing for the first time the h to Z and gamma decay. Interestingly we find that the rates of loop-induced decays are correlated and their signal strengths can rise up to 40\% -- 60\% depending on the channel. We furthermore study the dark matter phenomenology of the TMSSM. The lightest neutralino is a good dark matter candidate in two regions. The first one is related to the Higgs and Z resonances and the LSP is mostly Bino. The second one is achieved for a mass larger than 90 GeV and the LSP behaves as the well-tempered neutralino. An advantage of the triplet contribution is that the well-tempered neutralino can be a Bino-Triplino mixture, relieving the problem of achieving $M_2\sim M_1$ in unified scenarios. The dark matter constraints strongly affect the Higgs phenomenology, reducing the potential enhancements of the diphoton and of the Z + gamma channels by 20\% at most. In the near future, dark matter direct searches and collider experiments will probe most of the parameter space where the neutralino is the dark matter candidate. } | \label{sec:intro} The discovery of the Higgs boson~\cite{CMS:yva,ATLAS:2013sla} has closed a long era: its mass is no more a free parameter. Its value $m_h\simeq 126\,$GeV is in agreement with the mass range predicted in supersymmetric scenarios~\cite{Espinosa:1992hp}. Nevertheless, the minimal version of these models, the so-called Minimal Supersymmetric Standard Model (MSSM), turns out to be ailing by the LHC discovery. The value $m_h\simeq 126\,$GeV is indeed well above the one that the MSSM {\it naturally} predicts, and heavy third generation squarks and large stop mixing are required to reproduce the measured mass~\cite{Arbey:2011ab,Carena:2011aa,Kang:2012bv,Fan:2014txa}. The MSSM electroweak sector therefore needs an unpleasant amount of fine tuning and a little hierarchy problem plagues the model. In non-minimal supersymmetric scenarios this problem can be alleviated. They can indeed involve new contributions (absent in the MSSM) that rise the tree-level prediction of the Higgs mass. For this reason smaller radiative corrections and less tuning in the electroweak sector are required. The drawback of this important achievement is (partial) loss of predictivity since extra free parameters have been introduced. A compromise between naturalness and predictivity is thus to consider scenarios extending the MSSM as little as possible. If one does not enlarge the gauge symmetry group of the Standard Model (SM), the only extension boosting the tree-level Higgs mass is to couple new chiral superfields to the Higgs sector of the superpotential. To this aim only singlets and $SU(2)_L$ triplets with hypercharges $Y=0,\pm 1$ are allowed by gauge invariance~\cite{Espinosa:1991wt}. Whereas the former option has been deeply studied, see e.g.~\cite{Ellwanger:2009dp} and references therein, the latter is less known and has received special attention only after ATLAS and CMS initially measured sizeable deviations in the diphoton Higgs rate~\cite{Aad:2012tfa,Chatrchyan:2012ufa}. Indeed, the triplet superfield involves extra charginos that can largely enhance the diphoton channel~\cite{antonio1,Kang:2013wm,Basak:2013eba} without requiring peculiar features such as large deviations in the main Higgs decay rates, huge stop masses, ultra light charginos or very heavy Higgsinos as it occurs in other scenarios~\cite{Carena:2011aa,Batell:2013bka,Casas:2013pta,Belanger:2014roa}. In particular, such an enhancement can be achieved in both decoupling and non-decoupling regime (i.e.~with large and small CP-odd Higgs mass $m_A$) while resembling the dominant SM Higgs couplings~\cite{antonio2}. Although the observed Higgs signal strengths~\cite{CMS:yva,ATLAS:2013sla} might appear SM-like because of an accidental compensation between production and decay rates that per se differ from the SM predictions, it is still worth to analyze scenarios where each Higgs decay but the loop-induced ones, and Higgs production is SM-like. In this simplified approach, indeed, it is easier to highlight the origin of a potential deviation (in loop-induced channels) that lies on the top of the global suppression/enhancement present in all channels. Such a deviation is somehow expected since loop-induced processes are particularly sensitive to new physics that can not perturb the dominant Higgs channels. This method has been applied in ref.~\cite{antonio1} to show that charginos can provide up to 45\% diphoton enhancement in the $Y=0$ Triplet extension of the MSSM (TMSSM). The same approach is applied in the present paper. We extend the analysis of ref.~\cite{antonio1} to a broader parameter space and we find that a slightly larger enhancement of about 60\% can be achieved via chargino contributions. More interestingly, we show that this departure from the SM prediction is tightly correlated to the deviation in the $h\to Z\gamma$ channel. In any case, the $\Gamma(h\to Z\gamma)$ rate can never be larger than about 1.4 times its SM value~\footnote{For studies on the $Z\gamma$ channel in other non-minimal supersymmetric frameworks see i.e.~refs.~\cite{Cao:2013ur,Belanger:2014roa}.}. These upper bounds are obtained without imposing any Dark Matter (DM) constraint on the TMSSM field content. Nevertheless they are compatible with the DM observables if the Higgs phenomenology is somehow disentangled from the DM puzzle. This is achieved for instance by invoking gravitinos, axions and axinos as DM candidates~\cite{Covi:1999ty,Baer:2008yd,Steffen:2008qp,Feng:2010gw}, or by postulating cosmological scenarios with non-standard DM production~\cite{Gelmini:2010zh}. On the contrary, if the DM candidate is required to be the Lightest Supersymmetric Particle (LSP) of the TMSSM within the traditional cosmological assumptions, the above bounds should be revisited. To this aim we pursue the analysis of the Higgs phenomenology for the case having the lightest neutralino as DM particle. In order to capture the most stringent features related to the $h\to Z\gamma$ and $h\to \gamma\gamma$ enhancements, we require the relic density to rely only on the chargino, neutralino and SM fields. In other words, besides analyzing the DM annihilation via Higgs and $Z$ boson resonances, we study a kind of well-tempered neutralino in the TMSSM. By definition the well-tempered neutralino in the MSSM is a tuned mixture of gaugino and Higgsino that achieves the correct relic density away from resonances and coannihilations with other supersymmetric particles~\cite{ArkaniHamed:2006mb}. The successful parameter space consists of either the Bino and Higgsinos, or the Bino and Wino having almost degenerate mass terms. Other issues however jeopardize these two scenarios: the former is strongly constrained by limits on the DM Spin-Independent (SI) elastic scattering, and the latter seems unnatural since supersymmetry breaking mechanisms unlikely lead to degenerate Bino and Wino soft masses. Introducing the TMSSM fermionic triplet, hereafter dubbed Triplino, provides new features to the DM phenomenology. In fact, the Triplino can play the role of the Wino component, making the tuning between Bino, Wino and Higgsinos masses unnecessary and opening up a new viable DM parameter space for the well-tempered neutralino. Moreover, the Triplino mass parameter is a superpotential term that in principle can be produced by supersymmetry breaking sources different from those generating the gaugino masses~\footnote{For instance, one can produce the gaugino masses via gauge mediation and the mass parameters of Higgsinos and Triplinos via the Giudice-Masiero mechanism. Notice that the TMSSM does not seems to be in tension with gauge mediation due to the Higgs mass $m_h\approx 126$. Indeed, no large trilinear parameters are required to naturally achieve the observed Higgs mass~\cite{antonio1}.}. Interestingly, we find that in the TMSSM the DM constraints strongly impact the loop-induced Higgs processes. Independently on the regions where the LSP achieves the observed relic density the $h \to \gamma \gamma$ and $h\to Z\gamma$ enhancements cannot be larger than 20\%. This mostly occurs because larger enhancements need light Higgsino and Triplino mass parameters, which tend to push the SI elastic scattering off nuclei of the lightest neutralino above the LUX exclusion limit~\cite{Akerib:2013tjd}. The rest of the paper is organized as follows. In section~\ref{sec:model} we review the basic features of the TMSSM and its most natural parameter space. We also present some improvements in the determination of the lightest Higgs mass $m_h$. Section~\ref{sec:higgssign} describes the Higgs signatures in the TMSSM, with emphasis to the Higgs invisible width and loop-induced decay channels. In particular, the first calculation of the $\Gamma(h\to Z\gamma)$ width in the TMSSM is presented here. Section~\ref{sec:num} is dedicated to set up the method of our numerical analysis, as well as the parameter choice. Section~\ref{sec:resh} studies in detail the signal strengths of loop-induced Higgs decays and their correlation. We then move to discuss the DM phenomenology and its impact on the Higgs signatures in section~\ref{sec:dm}. Section~\ref{sec:Concl} is finally devoted to summarize our findings. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | \label{sec:Concl} We are entering the era of precision Higgs physics: the measurements of the Higgs mass, couplings and decay modes have already started to be highly sensitive to new physics beyond the SM. In this paper we have considered the Higgs phenomenology of the $Y=0$ triplet extension of the MSSM, dubbed TMSSM, in which the new coupling between the triplet and the MSSM Higgses can alleviate the little hierarchy problem and modify the chargino and neutralino sector. We have first accurately determined the couplings and pole masses of the stops, charginos, neutralinos and lightest CP-even Higgs $h$. Then, we have tackled the subtle effects of the Triplino in the $h \to \gamma\gamma$ and $h \to Z\gamma$ loop-induced processes. We have shown that the additional Triplino component in the chargino sector provides a maximal enhancement of 60\% in the $R_{ \gamma \gamma}$ signal strength, which is slightly larger than previously estimated (i.e.~$R_{\gamma\gamma}\lesssim 1.45$)~\cite{antonio1}. An enhancement up to 40\% can be achieved in the $R_{ Z\gamma}$ signal strength, which we find to be highly correlated with the diphoton channel, even though it is always smaller than $R_{\gamma\gamma}$. The parameter region leading to the largest $R_{ \gamma \gamma}$ and $R_{ Z\gamma}$ is characterized by $\tan\beta\lesssim 2$ and $\mu\sim\mu_\Sigma\sim M_2\sim 250$\,GeV, and in particular by light charginos close to the LEP bound. The enhancement in the TMSSM is significantly larger than the one achievable in the MSSM ($\sim 20\%$ for $R_{\gamma\gamma}$) for the same chargino lower mass bound~\cite{Casas:2013pta}. The measurements of these processes are likely to improve in the next years. LHC is indeed expected to probe the SM prediction of $\Gamma(h\to Z\gamma)$ once $\mathcal O(100~$fb$^{-1})$ data is collected~\cite{Campbell:2013hz}, and to measure the $g_{h\gamma\gamma}$ effective coupling within a 10\% accuracy after a high luminosity 3000~fb$^{-1}$ run~\cite{Peskin:2012we}. With these further data the Higgs diphoton signal strength will plausibly converge to the SM value. In such a case, sizeable deviations in $h\to Z\gamma$ would not be compatible with the TMSSM. On the contrary, if data will still exhibit a positive deviation from the SM, there would be a clear indication of physics beyond the SM. The above predictions and the tight correlation between $R_{\gamma\gamma}$ and $R_{Z\gamma}$ could be thus crucial to rule out or provide hints for the scenario considered here. Besides the Higgs decays, we have investigated the DM phenomenology in the TMSSM, focusing on the interplay of the neutralino and chargino sectors enlarged by the triplet components. Similarly to the MSSM, the LSP is a viable DM candidate in the Higgs or $Z$ pole region, and in the so-called well-tempered regime. The Higgs and Z pole regions are characterized by a Bino DM and are poorly sensitive to the Triplino, as the Higgs-Higgsino-Bino is the only relevant coupling. However, the well-tempered neutralino, where the LSP achieves the correct relic density via coannihilation with the lightest chargino, presents a new feature. Indeed the Triplino component of the LSP can substitute the Wino in the well-tempered neutralino and can solve the problem of having $M_1\sim M_2$ from grand unified model perspective. Indeed the requirement of DM comes at the expenses of satisfying the LUX exclusion limit for SI elastic cross-section on nuclei. The dominant contribution is due to Higgs exchange, which imposes a lower bound on $\mu$. Interestingly we found that this has an impact for the $R_{\gamma\gamma}$ and $R_{Z\gamma}$ enhancements: the Higgs-chargino coupling is reduced as well suppressing the signal strengths to at most $20\%$. Notice that these values are once again larger than the ones provided by the MSSM with DM constraints~\cite{Casas:2013pta}, when the Higgs production is SM-like. The scenario considered here nicely illustrates the complementarity of DM direct searches with LHC. For instance the next generation of direct detection experiments, such as XENON1T, will probe a consistent portion of the neutralino TMSSM parameter space. Moreover it will be capable of constraining the Higgs invisible decay branching ratio up to 1\%, in a time scale comparable to the LHC one. In general the TMSSM is less constrained by current LHC bounds on simplified models or supersymmetric searches. Indeed the presence of the Triplino can modify the couplings and the decay modes. This has been already observed for stops in the TMSSM~\cite{deBlas:2013epa} even though a precise estimate of their current mass bound is still missing. On the other hand no study exists for the chargino and neutralino mass bounds. Although we have checked that the present constraints~\cite{Aad:2014nua,Khachatryan:2014qwa} do not apply to our analysis ballpark, a dedicated investigation would be required in order to accurately determine the allowed parameter region. Present data should primarily affect the chargino parameter region with light lightest-neutralino and with small $h\to Z\gamma$ and $h\to \gamma\gamma$ enhancements. With more LHC data strongest bounds are expected, in particular for the DM mass close to the $Z$ or $h$ resonance. On the other hand, in order to probe the coannihilation region (where the spectrum is compressed), ILC data and analyses similar to that proposed in ref.~\cite{Porto:2014fca} would be crucial. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 14 | 3 | 1403.6434 |
1403 | 1403.3409_arXiv.txt | We report the discovery of a narrow stellar stream crossing the constellations of Sculptor and Fornax in the Southern celestial hemisphere. The portion of the stream detected in the Data Release 1 photometry of the ATLAS survey is at least 12 degrees long, while its width is $\approx$ 0.25 deg. The Color Magnitude Diagram of this halo sub-structure is consistent with a metal-poor $\feh \lesssim -1.4$ stellar population located at a heliocentric distance of 20 $\pm$ 2 kpc. There are three globular clusters that could tentatively be associated with the stream: NGC 7006, NGC 7078 (M15) and Pyxis, but NGC 7006 and 7078 seem to have proper motions incompatible with the stream orbit. | The last decade has witnessed the arrival of unprecedented amounts of high quality imaging data from the Sloan Digital Sky Survey (SDSS). Thanks to the depth and the exquisite stability of the SDSS broad-band photometry across thousands of square degrees on the sky, previously unseen low-level fluctuations in the Galactic stellar density field have been unearthed \citep[see e.g.][]{NewbergLumps, Willman2005, BelokurovFOS, BelokurovReview}. In the Milky Way halo one striking example of a small-scale over-density is the GD-1 stellar stream \citep{GD1}, only a fraction of a degree in width, but running over 60 degrees from end to end in the SDSS footprint. Such stellar trails are formed in the process of a satellite disruption in the tidal field of the Galaxy. The mechanics of the stream formation and the subsequent dynamical evolution in the host potential have been carefully studied \citep[e.g.][]{EyreBinney,SandersBinney}. The consensus that has emerged from both theoretical considerations and admittedly, very few tests on the actual data \citep[e.g.][]{KoposovGD1}, is that these structures can be used to yield powerful, unbiased constraints of the matter distribution in the Galaxy. In principle there exists a methodology to model the entire spectrum of tidal debris in the Galactic halo \citep[see e.g.][]{HelmiWhite}, from very narrow structures like GD-1 to broad luminous streams like that of Sgr \citep[exposed in e.g.][]{MajewskiSgr}, including the stellar over-densities that do not necessarily even trace out a stream \citep[e.g.][]{Rewinder}. However, the thinnest streams appear doubly interesting. First, given that these are not affected by the progenitor's gravity, methods are now in place to infer the Galactic potential without the need to resort to approximating the stellar tracks with a single orbit \citep[e.g.][]{Bovy2014,Sanders2014}. Second, along these narrow tidal tails it is easiest to observe density fluctuations due to interactions with dark matter sub-halos in the Galaxy \citep[see e.g.][]{Carlberg2009,Yoon2011}, provided that so-called epicyclic feathering is taken care of \citep[see e.g.][]{Mastro}. SDSS data has already been thoroughly mined to yield a handful of cold stellar streams. These include, for example, the tails of the Pal 5 globular cluster \citep{Od2001,Od2003,GrillmairPal5}, as well as tails around NGC 5466 \citep{Belokurov5466}, NGC 5053 \citet{Lauchner2006}, Pal 14 \citep{Sollima2011} and Pal 1 \citep{NO2010}. In addition \citet{GrillmairStyx} found a group of narrow streams Acheron, Cocytos, Lethe and more recently the discovery of the Pisces Stellar Stream was announced \citep{Bonaca2012, Martin2013}. \begin{figure*} \centering \includegraphics[width=0.49\linewidth]{figures/map.eps} \includegraphics[width=0.49\linewidth]{figures/ext.eps} \caption[]{\small \textit{Left:} Background subtracted density map of stars optimally weighted by proximity to an isochrone with $\feh =-2.1$ and an age of $12.5$ Gyr at a distance of 20 kpc. Darker shades of grey correspond to enhanced stellar densities. A narrow stellar stream is clearly visible crossing the area diagonally from $\alpha \sim 15^{\circ}$ to $\alpha=30^{\circ}$. The red line shows the great circle with the pole at $(\alpha,\delta)=(77 \fdg 16, 46 \fdg 92)$ aligned with the stream. The Fornax and Sculptor dwarf spheroidals are visible at $(\alpha,\delta)\approx (40^\circ, -35^\circ)$ and $(15^\circ, -33^\circ)$ respectively while much of the darker shading to the north and west is due to the southern Sgr stream. \textit{Right:} A map of the Galactic dust extinction $E(B-V$) around the stream area \citep{Schlegel} does not seem to reveal any features coincident with the stream.} \label{fig:map} \end{figure*} In this paper we present the discovery of a new stellar stream in the Southern celestial hemipshere based on photometry from Data Release 1 (DR1) of the VST ATLAS survey \citep{Shanks2013}. ATLAS is one of the three imaging surveys being currently undertaken within the remit of the ESO VST Public Surveys Program, the other two being KiDS \citep{Kids} and VPHAS+ (Drew et al 2014). The aim of the ATLAS survey is to obtain photometry in the SDSS $ugriz$ filters down to $r\sim 22$ for approximately 4,500 square degrees of the southern sky. The primary motiviation for the survey is to identify large numbers of $z\lesssim 2$ QSOs and Luminous Red Galaxies out to redshifts of $z\sim 0.6$ for studies of the cosmological matter power spectrum. The ability to use this data to also detect low-level Galactic Halo stellar sub-structure spanning several hundred square degree survey fields is testament to the quality and the stability of the ATLAS photometric and astrometric calibration. Section 2 briefly describes the data being used, while Section 3 gives the details of the newly discovered stream. The final section summarises the main results. | We have presented the discovery of a new narrow stream in the VST ATLAS DR1 data. The portion of the stream detected has the following properties. \smallskip \noindent (1) The stream lies on a great circle with a celestial pole at $(\alpha,\delta)=77 \fdg 16, 46 \fdg 92$. \smallskip \noindent (2) The heliocentric distance to the stream is $\sim 20$ kpc. \smallskip \noindent (3) The width of the debris distribution on the sky is 0.25 degrees which at the distance of the stream corresponds to $\sim$ 90 pc physical size. \smallskip \noindent (4) The CMD of the stream appears to be well described by an old and metal-poor isochrone. But in order to pinpoint the precise metallicity and age of the stream we will need spectroscopic measurements of the stream members. \smallskip \noindent (5) There are 3 Galactic globular clusters that could plausibly act as stream progenitors, based on their proximity to the stream orbit: NGC 7006 , NGC 7078 (M15) and Pyxis. However, given the proper motions available in the literature, NGC~7006 and NGC~7078 are unlikely to be associated with the stream. It is more difficult to rule out Pyxis at this stage. | 14 | 3 | 1403.3409 |
1403 | 1403.3123_arXiv.txt | We use analytic estimates and numerical simulations of test particles interacting with magnetohydrodynamic (MHD) turbulence to show that subsonic MHD turbulence produces efficient second-order Fermi acceleration of relativistic particles. This acceleration is not well-described by standard quasi-linear theory but is a consequence of resonance broadening of wave-particle interactions in MHD turbulence. We provide momentum diffusion coefficients that can be used for astrophysical and heliospheric applications and discuss the implications of our results for accretion flows onto black holes. In particular, we show that particle acceleration by subsonic turbulence in radiatively inefficient accretion flows can produce a non-thermal tail in the electron distribution function that is likely important for modeling and interpreting the emission from low luminosity systems such as Sgr A* and M87. \\ | \label{sec:introduction} In the limit of low-frequency magnetohydrodynamic (MHD) fluctuations, charged relativistic particles are accelerated by mirror forces resulting from magnetic compressions \citep{Achterberg1981}, \begin{equation} \label{eq:mirrorForce} \frac{d p_\parallel}{dt} = \frac{p_\perp v_\perp}{2B} \nabla_\parallel \mid\boldsymbol{B} \mid, \end{equation} where $\parallel$ and $\perp$ denote directions relative to the local magnetic field. In MHD, magnetic compressions are caused by slow modes and fast modes, with slow modes containing most of the compressive energy in subsonic turbulence. Because slow modes propagate approximately along the magnetic field in most regimes, a pure linear resonance with relativistic particles requires $\omega \simeq k_\parallel v_p = k_\parallel v_\parallel$, or equivalently $v_p \simeq c$, where $v_p$ is the parallel phase velocity of slow modes. Thus linear theory predicts no acceleration of high-energy particles by MHD-scale slow modes, because the resonance condition cannot be satisfied. As a result, fast modes have traditionally been believed to be the dominant source of relativistic particle acceleration by MHD-scale turbulent fluctuations \citep{Achterberg1981,Miller1996}. However, subsonic turbulence does not contain significant fast mode energy (see, e.g., \citealt{Yao2011a, Howes2012} for empirical constraints on the fast and slow mode energy in the solar wind). This appears to significantly limit the efficiency of relativistic particle acceleration by MHD turbulence in many astrophysical environments. In strong MHD turbulence, the waves comprising MHD turbulence are not long-lived but instead have a decay time comparable to their linear period. In this case, the linear resonance is not the appropriate condition for wave-particle interaction. Instead, the resonance is nonlinearly broadened \citep{Bieber1994a,Gruzinov1999b,Shalchi2004,Shalchi2004a,Qin2006,Yan2008a,Lynn2012}. Resonance broadening allows waves to interact with relativistic particles when $\omega_{\rm nl} \gtrsim k_\parallel c$, where $\omega_{\rm nl}^{-1}$ is the non-linear correlation time of the turbulence. In this paper, we estimate the resulting particle acceleration analytically (\S \ref{sec:transportProperties}) and numerically using simulations of relativistic test particles interacting with MHD turbulence (\S \ref{sec:numericalMethods} \& \ref{sec:numericalResults}). Our results are potentially relevant to a wide range of astrophysical plasmas; in \S \ref{sec:conclusions} we briefly assess the implications of our results for non-thermal emission from accretion disks around black holes. | \label{sec:conclusions} Our results demonstrate that subsonic MHD turbulence efficiently accelerates relativistic particles with a Fermi-like momentum diffusion coefficient $D_p \propto p^2$. This is true for both $\beta \lesssim 1$ and $\beta \gtrsim 1$ and is thus a robust property of charged particles interacting with low frequency MHD turbulence. We have restricted our analysis to particles whose (relativistic) cyclotron frequencies are larger than the frequencies of the turbulent fluctuations. In practice this limits our analysis to particles that are not too relativistic. Our key analytic result is that nonlinear broadening of quasi-linear resonances implies that slow modes in strong MHD turbulence can interact efficiently with relativistic particles, despite being unable to satisfy the linear resonance condition (see \citealt{Chandran2000} for a similar result). {\bf In particular, resonance broadening allows long wavelength turbulent magnetic field compressions satisfying $k_\parallel c \lesssim \omega_{\rm nl}$ to accelerate particles, where $\omega_{\rm nl}$ is the non-linear decay rate of the turbulence at a given scale.} Because slow modes tend to be energetically more important than fast modes in subsonic turbulence, this suggests that interactions with slow modes may dominate the overall particle acceleration by low-frequency, weakly compressible MHD turbulence. This is contrary to the standard quasi-linear theory results in the literature (e.g., \citealt{Achterberg1981}). However, the particle acceleration efficiency by fast modes depends sensitively on their turbulent power spectrum, which is not fully understood. In particular, if the fast mode spectral index is $\alpha \sim 3/2$ (which is not the case in our simulations, though it is suggested by some studies), fast modes may be more efficient than slow modes at accelerating particles even if their total energy density is smaller (see eq. \ref{eq:fastModeDiffusion}). For relativistic particles, momentum diffusion of the form $D_p \propto p^2$ produces a power-law spectrum $dN/dp \propto p^{-1}$ so long as the acceleration time of particles (which is independent of particle energy) is shorter than the radiative loss timescale and the escape time from the acceleration region \citep{Blandford1987}. The total energy in the accelerated particle population depends on the efficiency with which `seed' relativistic particles are created. Because suprathermal particle acceleration is inefficient for non-relativistic particles interacting with MHD turbulence \citep{Lynn2012}, it is not clear if the net acceleration efficiency (by turbulent mechanisms alone) will be substantial for plasmas with non-relativistic temperatures, because the turbulence itself does not self-consistently seed relativistic particles. By contrast, for relativistically hot plasmas, the formation of a non-thermal tail of relativistic particles by the mechanism studied here is likely to be quite efficient. One particularly important application of our results is thus to accretion flows onto black holes, where the electrons can in some cases have $k T \gtrsim m_e c^2$ even though the disk turbulence itself is non-relativistic. \subsection{Implications for Black Hole Accretion Flows} Weakly compressible MHD turbulence is generic in black hole accretion flows as a consequence of the nonlinear evolution and saturation of the magnetorotational instability \citep{Balbus1998}. Non-thermal particle acceleration by such turbulence is of particular astrophysical interest in at least two circumstances. First, at low accretion rates onto a black hole or neutron star, the accretion flow can adopt a low-collisionality state in which much of the emission can be dominated by a non-thermal population of electrons, if such a population is present (e.g., \citealt{Yuan2003}). Secondly, in luminous radiatively efficient accretion flows, non-thermal emission from the disk surface layers (a ``corona") can contribute significantly to the synchrotron and high energy inverse Compton emission. We briefly discuss the implications of our results for these applications. The momentum diffusion coefficient calculated in \S \ref{sec:transportProperties} and Figure \ref{fig:Dp_vs_c} corresponds to a rate of energy gain given by \begin{equation} \label{eq:Edotacc} \dot E_{\rm acc} \sim \frac{c \, D_p}{p} \equiv A \, p \, \frac{v_A^2}{L}\left(\frac{\delta B_\parallel}{ B_0}\right)^2 \end{equation} where $A$ is a dimensionless coefficient that encapsulates the efficiency of the particle acceleration and can be calibrated using our test particle simulations. In particular, Figure \ref{fig:Dp_vs_c} corresponds to $A \sim 1/3$ for $c/v_A \sim 10-100$, the values expected in the inner regions of accretion disks around black holes. The exact value of $\delta B_\parallel/B_0$ in accretion disk turbulence is somewhat uncertain. For the $\beta \sim 10-100$ conditions expected, $\delta B_\parallel \sim 0.3 \, B_0$ is plausible. However, the exact value depends in part on the effect of collisionless damping on the compressibility of accretion disk turbulence, which is not well understood. Moreover, small-scale fluctuations generated by the mirror instability may contribute significantly to the magnetic field compressions in collisionless disks \citep{Kunz2014,Riquelme2014}. The acceleration of particles by disk turbulence requires that the acceleration time is shorter than the viscous time. Given the acceleration rate in equation \ref{eq:Edotacc} this is likely achieved in the inner regions close to the black hole. In addition, the acceleration of particles by disk turbulence is limited by radiative losses, in particular synchrotron and inverse Compton emission. Focusing on the former, we find that the maximum Lorentz factor of accelerated electrons is given by \begin{equation} \label{eq:gmax} \gamma_{\rm max} \sim A \, \left(\frac{m_e}{m_p}\right) \tau_T^{-1} \left(\frac{\delta B_\parallel}{B_0}\right)^{2} \end{equation} where $\tau_{T} \equiv \sigma_T n_e L$ is the Thompson optical depth across the outer scale of the turbulent fluctuations $L$. Equation \ref{eq:gmax} implies that non-thermal emission from accelerated electrons is likely to be particularly important in low-luminosity systems where $\tau_T \ll 1$. As a concrete example, in models of the emission from Sgr A*, $\tau_T \sim 10^{-5}-10^{-6}$ (e.g., \citealt{Yuan2003,Gammie2009}) so that $\gamma_{\rm max} \sim 100$. This implies that the particle acceleration found here may substantially modify the electron distribution function for electrons that emit in the mm-infrared. This is particularly important to understand in the context of interpreting the variable infrared emission and resolved mm images of Sgr A* (e.g., \citealt{Doeleman2008, Do2009}). In the near future, more detailed calculations of test particle electron acceleration in shearing box simulations can be used to quantify the uncertain dimensionless coefficient A in the above acceleration efficiency. A second potential application of our results is to high energy emission from luminous accreting black holes, which can be produced by a combination of thermal and non-thermal processes. However, phenomenological models of this emission suggest that $\tau_T \sim 0.1-1$ in the emission region \citep{Haardt1991,Esin1997}. As a result, it is unlikely that the particle acceleration found here is sufficiently rapid to compete with radiative losses by synchrotron and inverse Compton emission. | 14 | 3 | 1403.3123 |
1403 | 1403.1561.txt | %We find allowed parameter region in Next-to-Minimal Supersymmetric Standard Model with a long-lived slepton, where dark matter relic density, the Higgs mass, and light elements abundances (including $^7$Li and $^6$Li) are obtained correctly. {\bf ((Shimomura-san))} We show that the Li problems can be solved in the next-to-minimal supersymmetric standard model where the slepton as the next-to-lightest SUSY particle is very long-lived. Such a long-lived slepton induces exotic nuclear reactions in big-bang nucleosynthesis, and destroys and produces the $^7$Li and $^6$Li nuclei via bound state formation. We study cases where the lightest SUSY particle is singlino-like neutralino and bino-like neutralino to present allowed regions in the parameter space which is consistent with the observations on the dark matter and the Higgs mass. | \label{sec:introduction}%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The Standard Model has had enormous successes in describing the interactions of the elementary particles, predicting almost every experimental results with accuracy. % The recent discovery of the Higgs particle finally crowned the accumulation of successes~\cite{Aad:2012tfa,Chatrchyan:2012ufa}. % On the other hand, it left us a number of questions that suggest the presence of a more fundamental theory behind. % Among such questions is the nature of dark matter; it became a compelling question during this decade after the precise observations of the universe reported their results~\cite{Bennett:2012zja,Ade:2013zuv}. % The theory behind the Standard Model ought to account for this question. % Supersymmetric models have been attractive candidates for such theories. % The Minimal Supersymmetric Standard Model (MSSM) is the simplest extension and most analyzed. % The lightest supersymmetric particle should be neutral and stable, and thus can be the dark matter. % Other extensions are also of interest, one of which is the Next-to-Minimal Supersymmetric Standard Model (NMSSM). % An extra singlet chiral supermultiplet is introduced in the NMSSM, and thereby account for the $\mu$-problem~\cite{Kim:1983dt} that complicates the MSSM. % In addition, the NMSSM better reconciles with the observed Higgs mass of $125~\mathrm{GeV}$. % The MSSM predicts a Higgs mass lighter than that of the $Z$ boson at the tree level, and employs loop effects to raise it up to the observed value. % On the other hand, the Higgs mass in the NMSSM has additional terms contributed from the singlet, and it potentially offers a straightforward interpretation of the observations. % A series of works by the present authors explored impacts of the supersymmetry on the nucleosynthesis in the early universe ~\cite{Jittoh:2005pq,Jittoh:2007fr,Jittoh:2008eq,Jittoh:2010wh,Jittoh:2011ni,Kohri:2012gc}. % Focus in these works has been on the case where the next-to-lightest supersymmetric particle (NLSP) is charged and long-lived so that it survives until the time of nucleosynthesis after the big-bang. % It takes part in the nuclear reactions and alters the present-day abundance of the light elements. %produce the lightest %supersymmetric particle (LSP), which is stable, neutral, and is %observed as dark matter today. % %Such reactions, if occurred, should be engraved also on the %present-day abundance of the light elements, which are successfully %predicted by the standard model of the big-bang nucleosynthesis (BBN) %without new physics like supersymmetry. % Possible disagreement indeed persists on the abundance of lithium compared with the calculation based on the standard big-bang nucleosynthesis (BBN) scenario. % The standard calculation predicts the ratio of abundance ${\rm Log}_{10}(\mathrm{^{7}Li/H})$ to be $-9.35\pm 0.06$~\cite{Jittoh:2011ni}, while the observation indicates $-9.63 \pm 0.06$~\cite{Melendez:2004ni}. % Lithium 6 provides another possible disagreement; its observed ratio of abundance $\mathrm{^{6}Li/^{7}Li} = 0.046 \pm 0.022$ is about $10^{2}$--$10^{3}$ larger than the theoretical prediction~\cite{Asplund:2005yt}. % These discrepancies can be the trace of the interaction between nuclei and the NLSP which is absent in the standard BBN scenario. % Our scenario can thereby account for the abundance of the dark matter and of the lithium in a single framework. % We analyzed if this scenario works within the MSSM with staus as the NLSP and neutralinos as the lightest supersymmetric particle (LSP), and found the parameter region that can account for these observational handles to the new physics~\cite{Jittoh:2005pq,Jittoh:2007fr,Jittoh:2008eq,Jittoh:2010wh,Jittoh:2011ni}. % % Now that the mass of Higgs particle is determined, we are to examine whether it is compatible with our scenario. % Our previous work analyzed the constrained minimal supersymmetric standard model~\cite{Konishi:2013gda}. % There we found allowed regions in the parameter space, and presented phenomenological predictions such as mass spectra and branching ratios. % In the present paper, we further extend our scenario to the NMSSM with the flavor violation and search for its further applications. % We demonstrate that the NMSSM under our scenario can simultaneously account for the three phenomenological clues: the abundance of dark matter, that of lithium, and Higgs mass. % We explore parameter space and discover parameter points that qualify the requirements. % Special interest is in the case where the NMSSM singlet is the major component of the neutralino LSP in expectation of the difference from the MSSM. % The case of bino-like neutralino is analyzed as well. % This paper is organized as follows. % In Sec.~\ref{sec:NMSSM}, we introduce the NMSSM and define the model of our interest. % The BBN in the presence of a long-lived slepton is also explained in this section. % The exotic reactions that are absent from the standard BBN are introduced. % Section~\ref{sec:strategy} describes our strategy to find out the parameter point that are consistent with the three phenomenological clues. % The results are presented in Sec.~\ref{sec:results}. % We show the three benchmark cases characterized by the type of neutralino LSP. % The first case considers singlino-like neutralinos. % The couplings of the singlet is small and $\tan \beta$ is large in this case. % The second case also handles singlino-like neutralinos, but the singlet couplings are large and $\tan \beta$ is small compared with the first case. % The third case deals with bino-like neutralinos. % Singlet couplings are set large and $\tan \beta$ small as in the second case. % We show that each of these three cases allows the phenomenological constraints. % We spot benchmark points in the parameter space and exhibit that they accord with the observations. % The NMSSM parameters are calculated to confirm the adequate Higgs mass, and the BBN network calculation is carried out to check the relic abundance of light elements and of the LSP dark matter. % %(Section~\ref{sec:DarkMatter} ...) % The present work is summarized in Sec.~\ref{sec:summary}. %LSPÌྪQñ³êÄ¢éÌÅÓ %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | \label{sec:summary} %%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The precise observations of the universe confirmed the presence of dark matter and raised a crucial question on its nature. % Another possible problem is the discrepancy of the abundance of $^7$Li and $^6$Li. % Models are attractive when they can account for these problems and at the same time are consistent with the observed mass of Higgs particle. % We demonstrated that the NMSSM can be a candidate for such models. % We specifically considered the case where the neutralino is the stable LSP and the lightest slepton is the NLSP, and where the mass difference of the two is so tiny that the slepton becomes long-lived enough to survive until the time of the nucleosynthesis in the early universe. % The sleptons interact with the synthesized nuclei and turn into the LSPs that stay until today as dark matter particles, altering the relic abundance of the light elements. % First, we searched for the benchmark sets of parameters that can successfully drive this scenario and simultaneously can reproduce the mass of Higgs particle within $(125.6 \pm 3.0) \, \mathrm{GeV}$. % Three cases of benchmark parameters are presented: % (a) Singlino-like neutralino, small $\lambda$-$\kappa$ with large $\tan \beta$; (b) Singlino-like neutralino, large $\lambda$-$\kappa$ with small $\tan \beta$; and % (c) Bino-like neutralino, large $\lambda$-$\kappa$ with small $\tan \beta$. % We found the successful benchmark values of $(c_{e}, \lambda, \kappa)$ in all the three cases. %(Tables~\ref{tab:points1}, \ref{tab:points2}, and \ref{tab:points3}). % We confirmed that they lead to the permissible abundance of dark matter and are consistent with other experimental bounds as presented in Tables~\ref{tab:points-s-small}, \ref{tab:points-s-large}, and \ref{tab:points-b-large}. % We then traced the BBN reaction network including the exotic nuclear reactions. % We employed $Y_{\tilde{l}^-}$ (slepton yield value)-$\delta m$ (LSP-NLSP mass difference) parameter plane to present the regions of parameters that can account for the observed abundance of light elements. The results are illustrated in terms of timescales for the relevant exotic BBN reactions as follows. % The slepton needs to be long-lived enough to form the bound state with $\mathrm{^{7}Be}$, while % Is the term of 'slepton lifetime' right? It's OK. the couplings ${G_{L,R}}_\tau$ should be large enough so that the internal conversion processes occur sufficiently. %Which is internal conversion or internal conversion processes? The slepton lifetime is thus subject to the lower bound. Meanwhile, the bound state of $(\mathrm{^{4}He}~\tilde{l}^-$) accompanies other two relevant processes: one is the $\mathrm{^{4}He}$ spallation, and the other is the catalyzed fusion. % These processes can readily overproduce the light elements and should be avoided, but suitable amount of $\mathrm{^{6}Li}$ production is favorable in order to account for its observed abundance. % Of the two, the $\mathrm{^{4}He}$ spallation generally proceeds more efficiently and is prone to overproduction, but it can be avoided when the slepton has not too long lifetime and less of it forms a bound state with $\mathrm{^{4}He}$. % In addition, such lifetime can induce appropriate abundance of $\mathrm{^{6}Li}$ via catalyzed fusion as its % observed abundance is tiny as explained around Eq.~\eqref{eq:deltaLi6}. % Upper bound on the slepton lifetime is thus brought in. % This tuning its production puts the upper limit on the slepton lifetime. % Combining the above arguments, we obtain an allowed window of the slepton lifetime. % %(a) In the case of (a) (Fig.~\ref{fig:small-lk-s_bbn1}), the specific relation between $\lambda$ and $\kappa$ is necessary to make the couplings ${G_{L,R}}_\tau$ large and thus to reduce % $\mathrm{^{7}\mathrm{Be}}$ by the internal conversion processes. % Indeed, we tune the parameters $\lambda$ and $\kappa$ to make $\mu_{\kappa}^2-\mu_{\rm eff}^2$ small and hence ${G_{L,R}}_\tau$ large at SS-1. % We still need a sizable flavor mixing, $c_e$, in order to render the slepton lifetime short and thereby reduce the number of $(\mathrm{^{4}He}~\tilde{l}^-)$. % Thus we obtain the allowed region at SS-1. % On the contrary, the flavor mixing $c_e$ is small at SS-2 and hence the slepton lifetime becomes longer. % Allowed region is thus shifted toward larger $\delta m$. % the allowed region is reduced in the tuned %parameter region with small flavor mixing $c_e$ at SS-2. % The values of $G_{L,R}$ become rapidly small outside the tuned parameter region of $\lambda$ and $\kappa$ % so that the allowed region becomes small (SS-4). %(b) In the case of (b) (Fig.~\ref{fig:large-lk-s_bbn1}), the values of $G_{L,R}$ are large due to large $\lambda$ and $\kappa$, and thus the internal conversion processes occur efficiently. % Suitable flavor mixings are also necessary as in the case of (a) to avoid the excessive amount of $({}^4 \mathrm{He}~\tilde{l}^-)$. Otherwise, we may miss the allowed region (SL-2). %(c) In the case of (c) (Fig.~\ref{fig:large-lk-b_bbn1}), the couplings ${G_{L,R}}_\tau$ hardly depend on $\lambda$ and $\kappa$, %Since the coupling ${G_{L,R}}_\tau$ for the case of (c) %does not depend mainly on $\lambda$ and $\kappa$, %the lifetime of slepton is short %because the ${G_{L,R}}_\tau$ are determined from the sizable gauge %couplings. % %Therefore, we can explain the observed abundance of ${}^6 \mathrm{Li}$ %even with no %flavor mixings although the observed abundance of ${}^7 \mathrm{Li}$ is not %realized due to the spallation process in BL-1-3. %Taking the small flavor mixings into account, %we can explain both $\mathrm{Li}$ problems in BL-4 %because these mixings make slepton lifetime shorter %and suppress the spallation process. %These results of the case (c) are same as in the case of MSSM and thus the results are same as in the case of the MSSM. %(We found that same results are shown in the case of MSSM) % We conclude that our scenario successfully works in the NMSSM and can simultaneously account for the abundance of dark matter, that of light elements, and the mass of Higgs particle. % Since all the three cases we considered here are consistent to the present phenomenological bounds, they should be distinguished through the characteristic signals of accelerator experiments. % Search for such signals are left for future works. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 14 | 3 | 1403.1561 |
1403 | 1403.6148_arXiv.txt | {Galactic encounters are usually marked by a substantial increase in synchrotron emission of the interacting galaxies when compared with the typical emission from similar non-interacting galaxies. This increase is believed to be associated with an increase in the star formation rate and the turbulent magnetic fields resulting from the encounter, while the regular magnetic field is usually believed to decrease as a result of the encounter.} {We attempt to verify these expectations. } {We consider a simple, however rather realistic, mean-field galactic dynamo model where the effects of small-scale generation are represented by random injections of magnetic field resulting from star forming regions. We represent an encounter by the introduction of large-scale streaming velocities and by an increase in small-scale magnetic field injections. The latter describes the effect of an increase in the star formation rate caused by the encounter. } {We demonstrate that large-scale streaming, with associated deviations in the rotation curve, can result in an enhancement of the anisotropic turbulent (ordered) magnetic field strength, mainly along the azimuthal direction. This leads to a significant temporary increase of the total magnetic energy during the encounter; the representation of an increase in star formation rate has an additional strong effect. In contrast to expectations, the large-scale (regular) magnetic field structure is not significantly destroyed by the encounter. It may be somewhat weakened for a relatively short period, and its direction after the encounter may be reversed.} {The encounter causes enhanced total and polarized emission without increase in the regular magnetic field strength. The increase in synchrotron emission caused by the large-scale streaming can be comparable to the effect of the increase in the star formation rate, depending on the choice of parameters. The effects of the encounter on the total magnetic field energy last only slightly longer than the duration of the encounter (ca. 1~Gyr). However, a long-lasting field reversal of the regular magnetic field may result.} | \label{int} Galactic encounters are spectacular phenomena that are usually marked by a substantial increase in synchrotron emission of the interacting galaxy when compared with the typical emission from similar non-interacting galaxies. The conventional interpretation associates this increase with an increase in the star formation rate (SFR) and in the strength of turbulent magnetic fields resulting from the encounter (Schleicher \& Beck 2013). The radio--far infrared relation also holds for interacting galaxies (e.g. Ivison et al. 2010). The polarized radio emission, signature of compressed turbulent fields, is also known to be enhanced during the interaction (Vollmer et al. 2013). In contrast, the regular magnetic field, which is generated by the mean-field dynamo and contributes to the synchrotron radiation, is usually believed to decrease as a result of the encounter because the interaction disturbs one of the drivers of galactic dynamos, namely the rotation curve of the interacting galaxy. Of course, this naive expectation requires verification. In its full extent this would be a far from straightforward undertaking. Indeed, a detailed modelling of the process through an encounter would require modelling the magnetic fields and hydrodynamics of both galaxies as well as the interaction itself. This does not appear feasible in the near future. There is however a simplification which isolates the features of the interactions that appear salient for the evolution of the regular galactic magnetic field, and can allow its evolution to be followed through an encounter event, at least at an exploratory level. Models for mean-field galactic dynamos have become increasingly more detailed (and so arguably more realistic). However, even with the very significant increase in computer resources that have become available over the last 20 years or so, severe approximations remain necessary. This is likely to remain true for the foreseeable future. Broadly speaking, models split into two groups. In one group, attempts are made to model in some detail processes in the ISM, including cosmic ray transport and large-scale dynamics. Direct Numerical Simulation (DNS) in ''boxes" can be used to provide estimates of transport coefficients (see e.g. Gressel et al. 2008, Brandenburg et al. 2008, Siejkowski et al. 2010). Models in the other group are in some ways less ambitious, using a simpler and more direct mean field formulation, often with rather {\it ad hoc} expressions for transport coefficients. Brandenburg (2014) gives an up-to-date review. Of course there is considerable overlap of these approaches, e.g. Hanasz, Woltanski \& Kowalik (2009). The latter type of model is much less demanding of computing resources, and readily allows extensive exploration of parameter space, and also the study of phenomena such as large-scale gas streaming in barred galaxies (e.g. Moss et al. 1998, 2001, 2007; Kulpa-Dybel et al. 2011; Kim \& Stone 2012), and the effects of spiral arms on dynamo action (Shukurov 1998; Moss 1998; Chamandy et al. 2013a,b; Moss et al. 2013). There are also many studies of the effects of externally driven gas flows on the gas content of galaxies (e.g. Vollmer, Braine \& Soida 2012 and references therein). Some of these also solve the passive induction equation, but do not include dynamo action, which is the main issue addressed in the current paper. Our intention here is to revisit the effects of a galaxy-galaxy encounter, which generates large-scale non-circular velocities, on large-scale dynamo action in the larger galaxy. This problem seems to have been first studied by Moss et al. (1993) and Moss (1996). In the first of these papers, non-circular velocities taken from a dynamical model of the encounter between M~81 and its satellite NGC~3077 by Thomasson \& Donner (1993) were introduced into a (necessarily) rather low-resolution and crude 3D dynamo model, and in the second a thin disc approximation was used. Moss et al. (1992) used a static velocity field corresponding to the instant of closest approach, whereas Moss (1997) used the fully time dependent velocity field. Although the dynamical model of the interaction is necessarily of low resolution by contemporary standards, we feel it is adequate to investigate some generic effects, and so we introduce the fully time dependent velocity field from Thomasson \& Donner (1993) into the dynamo model recently proposed by Moss et al. (2012). In this ''hybrid" dynamo model, a crude representation of the effects of supernovae driven star formation regions in injecting smaller scale magnetic field into the ISM are included explicitly in a thin disc dynamo model. The resulting global magnetic fields display a number of novel features for mean field dynamos, including disorder on the scale of a kiloparsec or less, and small- and large-scale field reversals. We were particularly interested in seeing what additional effects might result, arising from the interaction of the non-circular velocities with the effects of the magnetic field injections. This sort of study has connections with cosmic ray driven dynamos of e.g. Siejkowski et al. (2014). We first investigate the effects of the large-scale velocity field that is generated by the encounter, on the galactic magnetic field, and then tentatively explore the outcome of a parametrization of an associated increase in the star formation rate. A problem here is that the relation between the star formation rate and small-scale dynamo action is not completely resolved. We have not attempted any calibration of the novel effects with any standard model which quantifies the effect of an interaction on dynamo action via star formation, and simply represent this by a plausible parametrization. We stress that this point deserves further attention; however, such a study is obviously outside of the scope of this paper. We describe our model in Section~\ref{model} including a brief recapitulation of the galactic dynamo model of Moss et al. (2012). Our main result is that the interaction can significantly enhance synchrotron radiation emission, by its effects both on the small-scale and regular fields. Our detailed results are presented in Sect.~\ref{results} followed by discussion and conclusions in Sect.~\ref{disc}. | \label{concl} Overall we can conclude that an interaction affects the magnetic field configuration. The effects on the field can take the form of additional reversals of the large-scale magnetic field, can lead to a concentration of magnetic fields in rings, and so on. The limiting point, however, is that we only have snapshots of magnetic field configurations for a few examples, so it is problematic to isolate effects of interactions in the evolution of galactic magnetic field configurations; we can play with the governing galactic dynamo parameters to obtain configurations which more-or-less resemble fields observed in any particular case. We note again that we have used an old, low-resolution, dynamical model, so our results can only be regarded as generic. A reasonable way to isolate the effects of interactions observationally would be by a statistical study of interacting and non-interacting galaxies. The immediate feature of interest in our results is the peak in global energy in the epoch of encounter. The total magnetic energy increases during an encounter by 15\%, 80\%, 250\% for Models 17, 101 and 202 respectively, which should lead to a corresponding increase in the total flux of synchrotron radiation. \begin{figure} (a) \includegraphics[width=0.80\columnwidth]{fou_r16x.eps}\\ (b) \includegraphics[width=0.80\columnwidth]{fou_r17x.eps}\\ (c)\includegraphics[width=0.80\columnwidth]{fou_r101x.eps} (d)\includegraphics[width=0.80\columnwidth]{fou_r201x.eps}\\ \caption{Time averaged global Fourier integrals $p_i$ of $B_\phi^2$ through and beyond the interaction (which begins at $t=13.3$ Gyr, and ends at $t\approx 14.25$ Gyr in panels (a), (c), and (d) (Models 16, 101, 201 respectively), and lasts from approximately from 6.6 Gyr to 7.5 Gyr in panel (b) (Model 17). In each panel the upper continuous curves are for $m=0$, and the lower broken, dot-dashed and dotted curves are for modes $m=1, 2, 3$ respectively. } \label{r16x_fou} \end{figure} \begin{figure} (a)\includegraphics[width=0.80\columnwidth]{fou_r202x.eps} (b)\includegraphics[width=0.80\columnwidth]{fou_r203x.eps} \caption{Running time averages of global Fourier integrals $p_i$ of $B_\phi^2$ through and beyond the interaction (which begins at $t\approx13.3$ Gyr and ends at $t\approx 14.2 $ Gyr): (a) Model 202 ($q_{\rm I}=1$), (b) Model 203 ($q_{\rm I}=2$). \protect Notation as in Fig.~\ref{r16x_fou}. } \label{enhanced} \end{figure} \begin{figure} (a)\includegraphics[width=0.80\columnwidth]{r16x_fou_Bphi.eps}\\ (b)\includegraphics[width=0.80\columnwidth]{r203x_fou_Bphi.eps} (c)\includegraphics[width=0.80\columnwidth]{r9x_fou_Bphi.eps} \caption{Running time averages of global Fourier integrals $P_i$ of $B_\phi$ through and beyond the interaction, (which begins at $t\approx 13.3$ Gyr and ends at $t\approx 14.2 $ Gyr). Upper panels $P_0$, lower panels $P_1$ (continuous), $P_2$ (broken), $P_3$ (dotted): (a) Model 16, (b) Model 203, (c) Model 9 (no field injections, data not averaged). } \label{fou_sgn_r16x+r203x} \end{figure} | 14 | 3 | 1403.6148 |
1403 | 1403.6462_arXiv.txt | Combining the latest Planck, Wilkinson Microwave Anisotropy Probe (WMAP), and baryon acoustic oscillation (BAO) data, we exploit the recent cosmic microwave background (CMB) B-mode power spectra data released by the BICEP2 collaboration to constrain the cosmological parameters of the $\Lambda$CDM model, esp. the primordial power spectra parameters of the scalar and the tensor modes, $n_s,\alpha_s, r, n_t$. We obtained constraints on the parameters for a lensed $\Lambda$CDM model using the Markov Chain Monte Carlo (MCMC) technique, the marginalized $68\%$ bounds are $ r = 0.1043\pm_{ 0.0914}^{ 0.0307}, n_s = 0.9617\pm_{ 0.0061}^{ 0.0061}, \alpha_s = -0.0175\pm_{ 0.0097}^{ 0.0105}, n_t = 0.5198\pm_{ 0.4579}^{ 0.4515}.$ We found that a blue tilt for $n_t$ is favored slightly, but it is still well consistent with flat or even red tilt. Our $r$ value is slightly smaller than the one obtained by the BICEP team, as we allow $n_t$ as a free parameter without imposing the single-field slow roll inflation consistency relation. If we impose this relation, $r=0.2130\pm_{0.0609}^{0.0446}$. For most other parameters, the best fit values and measurement errors are not altered much by the introduction of the BICEP2 data. | The BICEP2 experiment\cite{Ade:2014xna,Ade:2014gua}, dedicated to the observation of the cosmic microwave background (CMB) polarization, has announced the detection of the primordial B-mode polarization, from observations of about 380 square degrees low-foreground area of sky during 2010 to 2012 in the South Pole. The detected B-mode power is in the multipole range $30<\ell<150$, a clear excess over the base lensed-$\Lambda$CDM model at these small $\ell$s, this excess can not be explained by the lensing contribution, for the CMB lensing contribution to B-mode polarization peaks at $\ell\sim1000$, while the contributed power at $\ell\sim100$, is small. The BICEP team has also examined possible systematic error and potential foreground contaminations and excluded these as possible source of the observed B-mode power. The cross-correlation between frequency bands shows little change in the observed amplitude, implying that frequency-dependent foreground are not the dominant contributor. The presence of the B-modes induced by the primordial gravitational wave in the early universe provides a direct evidence for the inflation theory. The tensor mode contribution to the CMB anisotropy may affect the global fitting of the cosmological parameters. The BICEP group reported their measured value of tensor-to-scalar ratio as $r=0.20^{+0.07}_{-0.05}$, based on the lensed-$\Lambda$CDM+tensor model, and derived from importance sampling of the Planck MCMC chain using the direct likelihood method, but they did not give constraints on other parameters. The unexpectedly large tensor-to-scalar ratio inspires a lot of interests on re-examining the inflation models \cite{ Hertzberg:2014aha,Choudhury:2014kma,Ma:2014vua,Gong:2014cqa,Xia:2014tda,Cai:2014bda} and observation datasets\cite{Zhao:2014rna,Zhao:2010ic,Zhang:2014dxk}. In this paper, we use the newly published BICEP2 CMB B-mode data, combined with the Planck CMB temperature data\cite{Collaboration:2013uv}, the WMAP 9 year CMB polarization data\cite{Hinshaw:2013dd, Bennett:2013ew}, and the BAO data from the SDSS DR9\cite{Anderson:2013jb}, SDSS DR7\cite{Padmanabhan:2012ft}, 6dF\cite{Beutler:2011ea}, to constrain the cosmological parameters in the lensed $\Lambda$CDM model. We derive constraints on the lensed $\Lambda$CDM model using the publicly available code COSMOMC\cite{Lewis:2002ah}, which implements a Metropolis-Hastings algorithm to perform a MCMC simulation in order to fit the cosmological parameters. This method also provides reliable error estimates on the measured variables. Previous CMB observations from the Planck satellite, the WMAP satellite and other CMB experiments yielded a limit of much smaller tensor-to-scalar ratio $r<0.11$ (at $95\%$ C.L.)\cite{Collaboration:2013uv}, so there is some tension between these and the BICEP result at least in the simplest lensed $\Lambda$CDM+tensors model. As pointed out by the BICEP team\cite{Ade:2014xna}, a simple way to relax this tension is to take the running of spectral index into account, we will explore this possibility in our fit. There are also wide spread interests in the tensor power spectral index, as it is an important additional source of information for distinguishing inflation models \cite{2014arXiv1403.5922A,2014arXiv1403.5163G,2014arXiv1403.4927B}, and a blue tensor power spectrum tilt $n_t\sim 2$ have been reported using the B-mode measurement\cite{2014arXiv1403.5732G}. Here we shall also investigate this problem and obtain an estimate of $n_t$ and its measurement error. | In this paper, we use the newly published BICEP2 CMB B-mode data, Planck CMB temperature data\cite{Collaboration:2013uv}, the WMAP 9 year CMB polarization data\cite{Hinshaw:2013dd, Bennett:2013ew} to constrain the base lensed $\Lambda$CDM model. In addition to the CMB data, we also use the BAO data from SDSS DR9\cite{Anderson:2013jb}, SDSS DR7\cite{Padmanabhan:2012ft}, 6dF\cite{Beutler:2011ea}, which help to break parameter degeneracy. For most parameters, we find that the best fit values and measurement errors are not altered much by the introduction of the BICEP2 data. The most affected parameters are $r$, $\alpha_s$ and $n_t$. Combining Planck + WP + BICEP+ BAO dataset, we obtain marginalized $68\%$ bounds on some interested parameters are: \begin{eqnarray} r &=& 0.1043\pm_{ 0.0914}^{ 0.0307} ~ , \\ n_s &=& 0.9617\pm_{ 0.0061}^{ 0.0061} ~ , \\ \alpha_s &=& -0.0175\pm_{ 0.0097}^{ 0.0105} ~ , \\ n_t &=& 0.5198\pm_{ 0.4579}^{ 0.4515} ~ . \label{eq:final} \end{eqnarray} We find that a blue tensor tilt ($n_t>0$) is slightly favored, but a flat or red tilt is consistent with the data. The best fitting value of $r$ we obtain is slightly smaller than BICEP2 group obtained,and the constraint on $r$ is also looser than BICEP2 group obtained. This result is reasonable, as we have not imposed the single-field-slow-roll inflation consistency relation on $n_t$, and treated it as an independent parameter. If we impose this relation, we will obtain $r=0.2130\pm_{0.0609}^{0.0446}$($1\sigma$ error) instead. In the near future, Planck and other experiments will provide more data on CMB polarization, and help improve the constraint on these parameters. | 14 | 3 | 1403.6462 |
1403 | 1403.2594_arXiv.txt | {Possible main formation sites of fluorine in the Universe include AGB stars, the $\nu$-process in Type II supernova, and/or Wolf-Rayet stars. The importance of the Wolf-Rayet stars has theoretically been questioned and they are probably not needed in the modelling of the chemical evolution of fluorine in the solar neighborhood. It has, however, been suggested that Wolf-Rayet stars are indeed needed to explain the chemical evolution of fluorine in the Bulge. The molecular spectral data, needed to determine the fluorine abundance, of the often used HF-molecule has not been presented in a complete and consistent way and has recently been debated in the literature.} {We intend to determine the trend of the fluorine-oxygen abundance ratio as a function of a metallicity indicator in the Bulge to investigate the possible contribution from Wolf-Rayet stars. Additionally, we present here a consistent HF line list for the K- and L-bands including the often used 23358.33~\AA\,line.} {High-resolution near-infrared spectra of eight K giants were recorded using the spectrograph CRIRES mounted at VLT. A standard setting was used covering the HF molecular line at 23358.33~\AA. The fluorine abundances were determined using spectral fitting. We have also re-analyzed five previously published Bulge giants observed with the Phoenix spectrograph on Gemini using our new HF molecular data.} {We find that the fluorine-oxygen abundance in the Bulge probably cannot be explained with chemical evolution models including only AGB-stars and he $\nu$-process in supernovae Type II, i.e. a significant amount of fluorine production in Wolf-Rayet stars is likely needed to explain the fluorine abundance in the Bulge. Concerning the HF line list, we find that a possible reason for the inconsistencies in the literature, with two different excitation energies being used, is two different definitions of the zero-point energy for the HF molecule and therefore also two accompanying different dissociation energies. Both line lists are correct, as long as the corresponding consistent partition function is used in the spectral synthesis. However, we suspect this has not been the case in several earlier works leading to fluorine-abundances $\sim$0.3 dex too high. We present a line list for the K- and L-bands and an accompanying partition function.} {} | \label{sec:introduction} From a nucleosynthetic perspective fluorine is a very interesting element and its cosmic origin is truly intriguing. Its creation and destruction in stellar interiors is very sensitive to the physical conditions \citep[see for example][]{2011ApJ...729...40L}, meaning that observations of fluorine abundances can provide strong constraints to stellar models. It will also be possible to observationally constrain the main stellar nuclear production sites of fluorine in the Universe at different epochs and in different stellar populations. To do this, observations of the chemical evolution of fluorine as a function of metallicity for different stellar populations have to be confronted with model predictions. Theoretical considerations have offered three main production mechanisms which all should work under prevailing conditions during different phases of stellar evolution. Their relative importance at different stages of evolution and in different stellar populations is only starting to be investigated. The different production sites of \element[][19]{F}, the only stable isotope of fluorine, that have been proposed are: \begin{itemize} \item \textit{$\nu$ nucleosynthesis in supernovae Type II (SNe II)}\\ The core collapse of a massive star, following a SN II explosion, leads to a prodigious neutrino flux. In spite of the small cross sections, the large amount of neutrinos gives rise to a significant spallation of \element[][20]{Ne} to \element[][19]{F} \citep{1988Natur.334...45W} in the overlying (neon-rich) shells of the core. \citet{1991NuPhA.527..663H} estimates the total (mu- and tau-) neutrino energy to $E_\nu=3\times 10^{53}$\,erg. \citet{2011ApJ...739L..57K} investigate the importance of this total neutrino energy for the $\nu$ process reactions for the evolution of fluorine in the solar neighborhood. They conclude that the $\nu$ nucleosynthesis should be a major fluorine production mechanism and that its relative contribution is largest for low metallicities. \item \textit{Thermal-pulsing Asymptotic Giant Branch (TP-AGB) stars}\\ Low-mass ($2 \la M/M_{\odot} \la 4$) TP-AGB stars have been suggested to produce fluorine in different burning phases during the thermal pulse stage, by nuclear reaction chains starting from $\element[][14]{N}$ \citep{1992A&A...261..157F,1992A&A...261..164J,2011ApJ...737L...8A,2011ApJ...739L..57K,2010MmSAI..81..998G}. Fluorine is then transported up to the surface by the ${3}^\mathrm{rd}$ dredge-up. Fluorine production in AGB-stars is expected to be accompanied by the slow-neutron capture nucleosynthesis (the s-process), producing elements like Sr, Y, Zr, Nb, Ba, and La \citep[e.g.][]{1998A&A...334..153M,2000A&A...362..599G,2009ApJ...694..971A}. It has been demonstrated observationally that AGB stars do produce fluorine, see for example \citet{1992A&A...261..164J} and \citet{2011ApJ...737L...8A}. \item \textit{Wolf-Rayet (W-R) stars}\\ \citet{1993nuco.conf..503M,1996LIACo..33...89M,2000A&A...355..176M} suggested that W-R stars could contribute significantly to the galactic fluorine budget. $\element[][19]{F}$ is produced in the convective cores of W-R stars, during the core He-burning phase. Due to a large mass loss caused by a metallicity-dependent, radiatively-driven wind, the destruction of $\element[][19]{F}$ by the ($\alpha$, $p$) reaction is prevented since the convective core shrinks. The fluorine left behind is eventually exposed at the surface as the heavy mass loss strips the star of the outer layers. This mechanism depends on key parameters, such as initial mass, metallicity, and rotational velocity. Fluorine is produced from $\element[][14]{N}$, which means that the more $\element[][14]{N}$ is available the more fluorine is expected. A second metallicity-dependent effect is the metallicity-dependent winds. Both circumstances favor the fluorine production at higher metallicities. \citet{2005A&A...443..243P} show that when incorporating newer yields and including models of rotating W-R stars, the yields from this mechanism are significantly reduced, implying that W-R stars might not be a major contributor of fluorine. However, they conclude that due to large uncertainties in key nuclear-reaction rates and mass-loss rates, the question of the contribution to galactic $\element[][19]{F}$ from W-R stars is still open. \end{itemize} Using a semi-analytic multizone chemical-evolution model, \citet{2004MNRAS.354..575R} show for the first time the impact of the AGB and W-R star contributions to the Galactic chemical evolution of fluorine. They show that $\nu$ nucleosynthesis was dominant in the early universe and that AGB stars' significance successively grows. Based on the old yields and non-rotating models, they further show that the contribution of W-R stars is significant for solar and super-solar metallicities, increasing the [F/O] ratio by a factor of two at solar metallicities. Their conclusion is that all three production sites are needed in order to explain the Galactic chemical evolution of fluorine for a range of metallicities. \citet{2011ApJ...739L..57K} modeled the evolution of fluorine in the solar neighborhood including AGB stars and $\nu$ nucleosynthesis with two different neutrino energies ($E_\nu=3\times 10^{53}$\,erg and $E_\nu=9\times 10^{53}$\,erg). Note that the contributions from W-R stars are underestimated in these models, because the elements such as C, N, and possibly F that are newly produced and have been lost via stellar winds before supernova explosions are not included. The models show a good agreement with field stars of higher metallicities. At lower metallicities the models cannot reproduce the observations of \citet{2013ApJ...765...51L}, but still the model that fits best include the $\nu$ process with $E_\nu=3\times 10^{53}$\,erg. The abundance of fluorine in stars is difficult to measure due to a paucity of suitable spectral lines. Highly ionized \ion{F}{v} and \ion{F}{vi} lines in the UV have been used by \citet{2005A&A...433..641W} in extremely hot post-AGB stars and a handful of \ion{F}{i} lines between 6800-7800~\AA\,have been used in extreme helium stars and R Coronae Borealis stars \citep{2006ApJ...648L.143P,2008ApJ...674.1068P}. All other studies we are aware of have been made using HF molecular lines in the K-band and mostly the HF($1-0$) R9 line at 23358.329~\AA. Relevant for the observations we present in this paper, is the study by \citet{2008ApJ...679L..17C} who present the first study of the chemical evolution of fluorine in the Galactic Bulge, by investigating six red giants in Baade's Window (five of these spectra are re-analyzed in this paper). They find that the fluorine to oxygen abundance ratio in the Bulge follows and extends the solar neighborhood trend. The trend at higher metallicities needs other sources of fluorine in addition to the $\nu$ process contribution, which is sufficient at lower metallicities. These are the AGB star and W-R star contributions. By investigating the correlation with abundances of s-process elements, the authors conclude that, for the Bulge, the W-R wind contribution to the fluorine budget should be important and larger than for the Disk. They therefore suggest that W-R stars might have played a vital role in the chemical evolution of the Galactic Bulge. In this paper, we observationally investigate the chemical evolution of fluorine in the Bulge, by analyzing red giants from three fields. We discuss the relative contributions of the different main nucleosynthetic sites suggested, by comparing with the latest and most updated models for the evolution of fluorine in the Bulge. Our main conclusion is that a significant fluorine production in W-R stars is likely needed to explain the fluorine abundance in the Bulge, meaning that the production in AGB-stars and SNe II is probably not enough. | At low metallicity, our observed fluorine-oxygen abundance trend in the Bulge is lower than predicted in our Bulge model including the $\nu$-process, showing a steeper slope than the model. This might suggest a metal-dependent production source of fluorine. This source cannot be the $\nu$-process in SNe II because it is not metal-dependent over our metallicity range, and it cannot be AGB-stars because these produce s-elements at the same time as fluorine and would probably not give rise to the observed decline in [Zr/F] for increasing [F/H] (as shown in Figure \ref{fig:cunha_af}). Therefore our data corroborate the findings of \citet{2008ApJ...679L..17C} that W-R stars might be an important source of fluorine in the Bulge. To fully evaluate this we need galactic chemical evolution models that include full sets of yields of AGB stars, W-R stars, and supernova explosions. We believe that some of the earlier reports of high fluorine abundances might be due to the use of mis-matching molecular data for the HF-molecule, but this has to be investigated. To help with this we have presented a HF line-list with a consistent partition function for lines in the K- and L-bands. | 14 | 3 | 1403.2594 |
1403 | 1403.5461_arXiv.txt | We investigate the oxygen and nitrogen abundance distributions across the optical disks of 130 nearby late-type galaxies using around 3740 published spectra of H\,{\sc ii} regions. We use these data in order to provide homogeneous abundance determinations for all objects in the sample, including H\,{\sc ii} regions in which not all of the usual diagnostic lines were measured. Examining the relation between N and O abundances in these galaxies we find that the abundances in their centres and at their isophotal $R_{25}$ disk radii follow the same relation. The variation in N/H at a given O/H is around 0.3 dex. We suggest that the observed spread in N/H may be partly caused by the time delay between N and O enrichment and the different star formation histories in galaxies of different morphological types and dimensions. We study the correlations between the abundance properties (central O and N abundances, radial O and N gradients) of a galaxy and its morphological type and dimension. | The chemical composition of a galaxy is one of its most fundamental characteristics. Here we focus on disk galaxies. Because the chemical enrichment depends on various physical processes, such as the star formation history and the mass exchange between the galaxy and its environment, progress in our understanding of galaxy formation and evolution processes depends to a large extent on improving our knowledge of the detailed chemical properties of galaxies, such as the radial distribution of element abundances across galactic disks. Establishing the macroscopic properties of spiral and irregular galaxies that likely govern the distribution of heavy elements across their disks is very important in understanding the (chemical) evolution of galaxies. Accurate abundance determinations in a sample of galaxies are mandatory for such investigations. The classical $T_{e}$ method, often referred to as the direct method, is generally considered to provide the most reliable oxygen and nitrogen abundances in H\,{\sc ii} regions. When this method cannot be used (due to the lack of measurements of the required weak auroral lines in the spectra of H\,{\sc ii} regions) then combinations of the strong nebular line intensities in spectra of H\,{\sc ii} regions can be used as indicators of their oxygen abundances, as was first suggested by \citet{Pageletal1979MNRAS189} and \citet{Alloinetal1979AA78}. This approach is usually referred to as the ``strong-line method'' and has been widely adopted. The establishment of calibrations (i.e., of relations between metallicity-sensitive emission-line combinations and metallicity) was the subject of numerous studies \citep[][among many others]{DopitaEvans1986ApJ307,McGaugh1991ApJ380,Pilyugin2000AA362, Pilyugin2001AA369,KewleyDopita2002ApJS142, PettiniPagel2004MNRAS348,Tremonti2004ApJ613, Liangetal2006ApJ652,Stasinska2006AA454,Thuanetal2010ApJ712}. A calibration is defined not only by the adopted indicators but also by the objects that serve as calibrating data points. A sample of H\,{\sc ii} regions with abundances derived through the $T_e$ method is used to construct an empirical calibration. A set of photoionization models is used to construct a theoretical (model) calibration. Even if the same indicator is used in the empirical calibration and in the theoretical calibration and if the same spectral measurements in a given H\,{\sc ii} region are used, those calibrations can produce significantly different abundance estimations. Metallicities derived using theoretical calibrations tend to be systematically higher (up to $\sim$0.7 dex) than those derived using the empirical calibrations \citep[see reviews by][]{KewleyEllison2008ApJ681,LopezSanchezEsteban2010AA517,LopezSanchezetal2012MNRAS426}. Therefore, oxygen abundances in extragalactic H\,{\sc ii} regions obtained in different studies using different calibrations can be significantly different. Spectroscopic measurements of H\,{\sc ii} regions within and beyond optical radii of galaxies were carried out in many works (see list of references below). In these studies, usually the H\,{\sc ii} regions in one or several galaxies are measured and the radial distributions of the element abundances across the disks of those galaxies are estimated. Since often different methods for abundance determinations are used in different works, the resulting abundances from these studies are not homogeneous and cannot be directly compared to each other. Therefore, the abundances in a sample of galaxies can be analyzed only after those abundances are homogenized, i.e., all the abundances are redetermined in a uniform way. This is the first step in our present investigation. It should be noted that there have been several attempts to use uniform abundances for the determination of radial abundance gradients in a sample of galaxies; e.g., \citet{VilaCostas1992MNRAS259} for a sample of 30 galaxies, \citet{Zaritskyetal1994ApJ420} for 39 galaxies, \citet{vanZeeetal1998AJ116} for 11 galaxies, \citet{Pilyuginetal2004AA425} for 54 galaxies, and \citet{Moustakasetal2010ApJS190} for 21 galaxies. However, those samples contain a relatively small number of galaxies (whereas our present sample includes 130 galaxies). So far, little attention has been paid to the radial distributions of nitrogen abundances in the disks of galaxies, despite the fact that this provides several advantages for the study of the chemical evolution of galaxies. Indeed, since at 12+log(O/H) $\ga$ 8.3, secondary nitrogen becomes dominant and the nitrogen abundance increases at a faster rate than the oxygen abundance \citep[e.g.,][]{Henryetal2000ApJ541}, the change in nitrogen abundances with galactocentric distance should then show a larger amplitude in comparison to oxygen abundances and, as a consequence, should be easier to measure. Furthermore, there is a time delay in nitrogen production as compared to oxygen production \citep{maeder1992,vandenhoek1997,pagel1997}. Thus the comparison between the radial distributions of oxygen and nitrogen abundances in the disks of galaxies can shed additional light on the chemical evolution of galaxies. Therefore we consider here not only the radial distributions of oxygen abundances but also those of nitrogen abundances. Our paper is organized in the following way. We describe the method used for the oxygen and nitrogen abundance determinations in the H\,{\sc ii} regions of our galaxy sample in Section 2. We describe the observational data that were used to determine the abundances in the H\,{\sc ii} regions in Section 3. We discuss the abundance properties in the disks of nearby galaxies (within the optical isophotal radii) in Section 4. We summarize our results in Section 5. Throughout this paper, we will use the following standard notations for the line intensities: \\ $R_2$ = $I_{\rm [O\,II] \lambda 3727+ \lambda 3729} /I_{{\rm H}\beta }$, \\ $N_2$ = $I_{\rm [N\,II] \lambda 6548+ \lambda 6584} /I_{{\rm H}\beta }$, \\ $S_2$ = $I_{\rm [S\,II] \lambda 6717+ \lambda 6731} /I_{{\rm H}\beta }$, \\ $R_3$ = $I_{{\rm [O\,III]} \lambda 4959+ \lambda 5007} /I_{{\rm H}\beta }$. \\ With these definitions, the excitation parameter $P$ can be expressed as: $P$ = $R_3$/($R_2$+$R_3$). | We compiled published spectra of H\,{\sc ii} regions in 130 nearby galaxies. Our list contains 3904 spectra including 162 spectra of H\,{\sc ii} regions beyond the isophotal radius $R_{25}$. The oxygen and nitrogen abundances in H\,{\sc ii} regions were determined on the metallicity scale defined by H\,{\sc ii} regions with $T_e$-based abundances. The radial gradients of oxygen and nitrogen abundances across the disks of the galaxies were estimated. At the centers of metal-rich galaxies (i.e., (12 +log(O/H) $\ga$ 8.2), we found the relation between N and O abundances to be (N/H)$_{R_0}$ $\propto$ (O/H)$_{R_0}^{2.5}$. The (N/H)$_{R_{25}}$ = $f$(O/H)$_{R_{25}}$ relation between N and O abundances at the $R_{25}$ isophotal radii of high metallicity galaxies is similar to that for the abundances at their centers. The variation in (N/H) at a given (O/H) is around 0.3 dex. To test whether the scatter in N/H at a given O/H can be explained by the time delay between nitrogen and oxygen enrichment and the different star formation histories in galaxies of different morphological types and dimensions (masses), we derived a more complex relation between N and O abundances (N/H) = $f$((O/H),$T$,$R_{25}$). We found that the morphological type, $T$, is a more important ``second parameter'' in the relation for central abundances, while the log$R_{25}$ is a more important second parameter in the relation for abundances at the $R_{25}$ radii of our galaxies. Since there is a correlation between $T$ and log$R_{25}$ it is as yet unclear whether this difference is meaningful. The fact that the relation between N/H and O/H depends on additional parameter(s), namely $T$ and/or log$R_{25}$ suggests that the scatter in N/H at a given O/H can be caused, at least partly, by the time delay between nitrogen and oxygen enrichment and the different star formation histories in different galaxies. The best fit to N/H as a function of O/H is close to a linear relation at low metallicity (12 + log(O/H) $\la$ 8.0). The central oxygen abundances (O/H)$_{R_0}$ show a trend along the Hubble sequence of galaxies of late morphological types ($T$ $\ga$ 5) such that the oxygen abundances are lower in galaxies of later types. This trend disappears for early morphological types. The central oxygen abundance also correlates with optical galaxy radius for small galaxies, $R_{25}$ $\la$ 10 kpc, being lower in galaxies of smaller radii. The trend disappears for galaxies with large radii. The maximum gas-phase oxygen abundance in large (10 kpc $\la$ R$_{25}$ $\la$ 30 kpc) galaxies (or in galaxies of early (1 $\la$ $T$ $\la$ 5) morphological types) is constant, 12 + log(O/H) $\sim$ 8.85. This implies that the observed central oxygen abundance of the most oxygen-rich galaxies in our sample is a factor of $\sim$2 higher than the gas-phase oxygen abundance in the solar neighbourhood. The central nitrogen abundances (N/H)$_{R_{0}}$ show a similar behaviour. The observed central nitrogen abundance in the most nitrogen-rich galaxies of our sample is 12 + log(O/H) $\sim$ 8.42. The radial O and N abundance gradients (in units of dex~kpc$^{-1}$) within the optical radius do not show any significant correlation with the morphological type and optical radius. However, the spread in the gradients increases with decreasing galaxy radius in the sense that shallow gradients are seen both in small and large galaxies while steep gradients occur only in a small galaxies. The smaller a galaxy the steeper is the gradient that it may show. The abundance data set presented in this paper serves as the foundation for other investigations we are carrying out. In a forthcoming paper \citep{Pilyugin2014AJ}, we examine relations between the radial abundance distribution across the disk and the disk surface brightness profiles in the optical $B$ and infrared $K$ bands for a sample of nearby galaxies. | 14 | 3 | 1403.5461 |
1403 | 1403.0905_arXiv.txt | In the spirit of minimal modeling of complex systems, we develop an idealized two-column model to investigate the climate of tidally locked terrestrial planets with Earth-like atmospheres in the habitable zone of M-dwarf stars. The model is able to approximate the fundamental features of the climate obtained from three-dimensional (3D) atmospheric general circulation model (GCM) simulations. One important reason for the two-column model's success is that it reproduces the high cloud albedo of the GCM simulations, which reduces the planet's temperature and delays the onset of a runaway greenhouse state. The two-column model also clearly illustrates a secondary mechanism for determining the climate: the nightside acts as a ``radiator fin'' through which infrared energy can be lost to space easily. This radiator fin is maintained by a temperature inversion and dry air on the nightside, and plays a similar role to the subtropics on modern Earth. Since 1D radiative-convective models cannot capture the effects of the cloud albedo and radiator fin, they are systematically biased towards a narrower habitable zone. We also show that cloud parameters are most important for determining the day--night thermal emission contrast in the two-column model, which decreases and eventually reverses as the stellar flux increases. This reversal is important because it could be detected by future extrasolar planet characterization missions, which would suggest that the planet has Earth-like water clouds and is potentially habitable. | Extrasolar planets in tidally locked orbital configurations around low-mass and relatively cool M-dwarf stars are the prime targets of several ongoing and proposed planet search programs \citep{Tarteretal2007}. In advance of direct observations of these planets, numerical models, including both single column radiative--convective models \citep[e.g.,][]{Kastingetal1993,Wordsworthetal2010, HuandDing2011, Kopparapuetal2013} and three-dimensional atmospheric general circulation models \citep[GCMs,][hereafter, YCA13]{Joshietal1997, Joshi2003, MerlisandSchneider2010, HengandVogt2011, Pierrehumbert2011, Edsonetal2011, Wordsworthetal2011, Leconteetal2013, Shieldsetal2013, Yangetal2013} have been employed to investigate their potential climates. Radiative--convective models employ sophisticated radiative schemes, but neglect atmospheric dynamics. GCMs calculate atmospheric dynamics in detail, and use radiative schemes of varying levels of complexity. In addition to numerically intensive modeling, idealized models can help make clear essential physical mechanisms in complex systems such as the climates of tidally locked planets. Low-order models are relatively easy to analyze and determine the dominant mechanisms, since mechanisms can be easily added, removed, or changed. They can be compared with more complex models to aid the interpretation of complex model results, and they help to address the critical question of what is the minimum physics necessary to understand a particular problem. In this study, we develop a two-column model of the atmospheres of Earth-like tidally locked planets in order to better understand the following two questions: (1) What determines the surface temperature? and (2) What determines the thermal infrared emission contrast between dayside and nightside? These questions are critical because the surface temperature determines whether liquid water can be maintained on the surface, which determines habitability in the traditional sense, and the thermal emission contrast will be one of the first observational signals that will be used to characterize tidally locked terrestrial planets \citep[e.g.,][]{Knutsonetal2007}. The model divides the atmosphere into two columns with one representing the dayside and the other representing the nightside (Fig.~1). This is the most basic possible simplification that allows horizontal heterogeneity, and it is justified by the following facts: (1) Only the dayside receives stellar flux from the parent star, whereas the nightside is heated by atmospheric and ocean heat import from the dayside. The temperatures of the two sides are therefore determined by distinct physical processes. (2) GCM simulations have found that on the nightside horizontal surface and air temperature gradients are extremely small and atmospheric descent occurs throughout the nightside \citep{Joshietal1997, MerlisandSchneider2010}; therefore, the nightside atmosphere can be approximated as a single column. (3) On the dayside, the surface temperature is homogeneous in the vicinity of the substellar point in GCM simulations, and more broadly much of the dayside is characterized by robust moist convection and mean ascent \citep[YCA13;][]{MerlisandSchneider2010}, so that it can be roughly treated as a single column. (4) The dayside has a moist atmosphere with a robust hydrological cycle whereas the nightside atmosphere is extremely dry \citep{MerlisandSchneider2010, Edsonetal2011}. (5) Future observations of horizontal variations in the climate on tidally locked terrestrial planets are likely to be coarse, so it is reasonable to develop a model with minimal horizontal resolution. In our model, the dayside column is thought of as a moist region with deep convection, mean ascent, and a warm surface. The air temperature follows a saturation moist adiabatic temperature profile. In this way the dayside is similar to the warm pool of the tropical Pacific Ocean on Earth, where the temperature profile is close to a moist adiabat \citep{XuandEmanuel1989}. The nightside column has mean descent with dry air and a cold surface, which is analogous in some ways the subtropics of Earth. The implicit atmospheric circulation connecting the dayside and the nightside in the model approximates a global-scale Walker circulation driven by the stellar energy contrast, as obtained in GCM simulations (see the review of \cite{Showmanetal2013}). To simulate the effect of clouds, we employ a simple convective cloud scheme and tune one free parameter of the scheme to produce a similar cloud albedo to that obtained in the GCM simulations of YCA13. At a given stellar flux and a specified ocean heat transport, the model computes dayside and nightside surface temperature, dayside and nightside free-tropospheric temperature, dayside convective heat flux, and atmospheric heat transport from the dayside to the nightside. \begin{figure}[h!] \vspace{-15mm} \begin{center} \includegraphics[angle=0, width=27pc]{./af0y_lines_color_TCM.eps} \end{center} \vspace{-8mm} \caption{Schematic representation of the two-column model for the climate of tidally locked terrestrial planets. The dayside column consists of a cloudy part and a clear-sky part. S$_0$: stellar flux at substellar point; $\alpha_p$: planetary albedo; $OLR$: outgoing longwave radiation (infrared emission to space); $f_c$: effective cloud fraction; $T_c$: cloud emission temperature; $\varepsilon_2,_3$: atmospheric emissivity; PBL: planetary boundary layer; $F_c$: convective heat flux; $F_a$: atmospheric heat transport from dayside to nightside; $F_o$: ocean heat transport; and $F_d$: the fraction of atmospheric heat transport from the dayside to the nightside deposited in the nightside boundary layer due to adiabatic heating.} \label{fig1} \end{figure} Similar types of models have been employed to investigate the climate of the tropics on modern Earth. Such models have been found to be useful for identifying the mechanisms that regulate the surface temperature of the tropical Pacific Ocean, including surface evaporation \citep{HartmannandMichelsen1993}, convective clouds \citep{RamanathanandCollins1991}, stratus low clouds \citep{Larsonetal1999}, water vapor greenhouse effect \citep{Pierrehumbert1995}, horizontal atmospheric heat transport \citep{HartmannandMichelsen1993, Pierrehumbert1995}, and ocean dynamics \citep[e.g.,][]{ClementandSeager1999}. Because of the extreme simplifications of these column models, their applicability to direct comparison with detailed and spatially dense observations is somewhat limited; however, they have significantly improved our understanding of tropical climate by helping to identify the essential processes that govern the system. \begin{figure*}[] \begin{center} \vspace{-40mm} \begin{center} \includegraphics[angle=0, width=34pc]{./af1_1200_T-profiles_Rotations_2panels.eps} \end{center} \vspace{-35mm} \caption{GCM output that supports the weak-temperature-gradient approximation for tidally locked planets. (a) Vertical temperature profiles for a tidally locked configuration with a rotation period of 60 Earth-days. Red line: for the tropics (30$^{\circ}$S--30$^{\circ}$N) of the dayside; green line: for the extra-tropics of the dayside; blue line: for the entire nightside; and black line: global mean. (b) Free-tropospheric temperature contrast between dayside and nightside (mean value between the pressure levels of 100 and 600 $hPa$) as a function of rotation period. The stellar flux is 1200\,W\,m$^{-2}$ in these simulations.} \label{fig2} \end{center} \end{figure*} We tune our two-column model to a GCM, which fully resolves atmospheric circulation, radiative transfer, the hydrological cycle, and clouds. We use an Earth GCM, the Community Atmosphere Model version 3 (CAM3; \cite{Collinsetal2004}), developed by the National Center for Atmospheric Research. The GCM is coupled to a mixed layer ocean with a uniform depth of 50 m. We have modified the GCM to simulate the climate of tidally locked planets around M-stars (for details, see YCA13). The default planetary radius, gravity, and rotation period are set to two times Earth's, 13.7\,m\,s$^{-2}$, and 60 Earth-days, respectively. The atmosphere is composed of N$_2$ and H$_2$O with a surface pressure of 1 bar, and other greenhouse gases such as CO$_2$, CH$_4$, N$_2$O, and O$_3$ are set to zero. Three groups of experiments with different stellar fluxes (from 1000 to 2400~W\,m$^{-2}$), ocean heat transport (from 0 to 55~W\,m$^{-2}$), and rotation periods (from 5 to 100~Earth-days) were performed, and the results are used here for comparisons with the two-column model. The GCM and the two-column model share the same planetary and atmospheric parameters, such as planetary gravity and specific heat of air. The goal of this study is to increase understanding of the key controls on surface temperature and thermal emission flux of tidally locked Earth-like planets. The basic physical processes that we build into the model are outlined in Section 2. The model is developed in Section 3. Section 4 presents the behavior of the two-column model, comparisons of model results with GCM simulations, and sensitivity analyses of the model. In Section 5, the critical stellar flux at which the day--night thermal emission contrast becomes negative is addressed. This contrast is essential for the interpretation of phase curves of terrestrial planets that will be measured in the near future. Section 6 summarizes the main findings of this study. | We constructed a low-order two-column model to simulate the climate of tidally locked terrestrial planets near the inner edge of the habitable zone of M-dwarf stars. This model incorporates the weak-temperature-gradient approximation to calculate horizontal atmospheric heat transport, the fixed anvil temperature hypothesis to calculate the cloud longwave effect, and a simple convective cloud scheme to calculate the cloud albedo. Our main findings are as follows: (1) The relatively few parameters of the two-column model can be tuned to reproduce the basic behavior of a very complicated global climate model and respond faithfully to variations in both the stellar flux and ocean heat transport. This suggests that the physical processes built into the two-column model are the most important ones for determining the climate of tidally-locked terrestrial planets. (2) The two-column model clearly illustrates the importance of dayside clouds and a nightside ``radiator fin'' for determining the climate of tidally locked planets. Dayside clouds increase the planetary albedo, which cools the planet and delays the onset of a runaway greenhouse. The radiator fin, primarily maintained by the dryness of the atmosphere and the temperature inversion on the nightside, allows the planet to easily lose energy to space through infrared thermal emission, which also cools the planet. Atmospheric dynamics are crucial for both of these effects, so they cannot be calculated by 1D radiative-convective models. This means that 1D radiative-convective models will tend to produce overly conservative estimates of the inner edge of the habitable zone. (3) Observations of the day--night thermal emission contrast will be critical for deciphering the climate of tidally locked terrestrial planets. As the stellar flux increases, nightside emission increases faster than dayside emission, so that at a critical stellar flux there will be a reversal in the day--night thermal emission contrast, with higher emission from the nightside. Sensitivity experiments in the two-column model show that cloud variables are most likely to cause differences in the critical stellar flux between different models. This means that future observations of the day--night thermal emission contrast will be useful for distinguishing between cloud models. | 14 | 3 | 1403.0905 |
1403 | 1403.6903_arXiv.txt | In this paper we present the results of observations of seventeen \HII regions in thirteen galaxies from the SIGRID sample of isolated gas rich irregular dwarf galaxies. The spectra of all but one of the galaxies exhibit the auroral [O {\sc iii}] 4363\AA~ line, from which we calculate the electron temperature, T$_e$, and gas-phase oxygen abundance. Five of the objects are blue compact dwarf (BCD) galaxies, of which four have not previously been analysed spectroscopically. We include one unusual galaxy which exhibits no evidence of the [N {\sc ii}] $\lambda\lambda$ 6548,6584\AA~ lines, suggesting a particularly low metallicity ($<$ Z$_\odot$/30). We compare the electron temperature based abundances with those derived using eight of the new strong line diagnostics presented by \cite{2013ApJS..208...10D}. Using a method derived from first principles for calculating total oxygen abundance, we show that the discrepancy between the T$_e$-based and strong line gas-phase abundances have now been reduced to within $\sim$0.07 dex. The chemical abundances are consistent with what is expected from the luminosity--metallicity relation. We derive estimates of the electron densities and find them to be between $\sim$5 and $\sim$100 cm$^{-3}$. We find no evidence for a nitrogen plateau for objects in this sample with metallicities 0.5~$>$ ~Z$_\odot$~$>$~0.15. | The metallicity of \HII regions in small isolated dwarf galaxies is key to investigating the physical processes that govern star formation in undisturbed stellar systems.\footnote{In this paper we attempt to be explicit in our terminology, using the term ``oxygen abundance'', and referring to ``metallicity'' only in widely used terms such as ``mass--metallicity'' and to refer to total chemical abundances. In addition, the abundance of oxygen measured from spectra is the gas-phase abundance, and does not take into account the oxygen in dust grains.} The Small Isolated Gas Rich Irregular Dwarf galaxy (SIGRID) sample of small isolated gas rich irregular dwarf galaxies \citep{2011AJ....142...83N} was selected with the aim of exploring the behavior of the mass-- or luminosity--metallicity relation at the low end of the mass scale. This is based on the observation that nebular metallicity decreases with galaxy stellar mass/luminosity \citep[see, for example,][]{2004ApJ...613..898T, 2006ApJ...647..970L}. However, the low end of the mass scale shows significantly more scatter in metallicity than the high end in the Tremonti SDSS data. By selecting isolated dwarf galaxies, it was our intention to see if this scatter persisted, and whether it was an intrinsic property of small galaxies. The SIGRID study is complementary to the ``Choirs'' study which looks for tidal dwarf emission line galaxies in group environments \citep{2014ApJ...782...35S}. It is distinct from the Spitzer Local Volume Legacy survey used by \cite{2012ApJ...754...98B} and the SDSS data of \cite{2004ApJ...613..898T} in using targets specifically chosen for their isolation. It is most similar in concept to the study by \cite{2011AstBu..66..255P} of galaxies in the Lynx-Cancer void, but is limited to small very isolated dwarf objects. Other questions that the SIGRID observations are intended to address are the existence or otherwise of a primary nitrogen ``plateau'' at metallicities below Z=8.45 \citep{1998ApJ...497L...1V}, and the relationship between oxygen abundances determined using ``direct method'', based on the measurement of the electron temperature and the estimation of the ionization correction factors to account for unseen ionization stages, and ``strong line'' technique, based on a calibration of the bright emission lines and emission line ratios. There has not been good agreement to date between the two methods, attributed to the empirical nature of the strong line methods. They have been calibrated in terms of the direct method, and have not until recently had an analytical basis. The direct method has been used as a standard for temperature and metallicity measurement, against which the strong line methods have been calibrated. \cite{2013ApJS..208...10D} subsequently presented a set of strong line diagnostic grids derived from the Mappings photoionization modelling code, based on the latest atomic data \cite[see][]{2013ApJS..207...21N}. We use both the new atomic data and the new diagnostic grids in our analysis. One might expect there to be greater scatter in the mass--metallicity relation at low masses, due to (1) measurement noise in nebular spectra in fainter galaxies, and (2) different star formation histories in the galaxies. \cite{2006ApJ...647..970L} suggest that the apparent scatter diminishes in observations at longer wavelengths (4.5$\mu$m), and we present additional optical spectral evidence on this question. The behavior of the ratio of nitrogen to oxygen abundances at low metallicities also shows increased scatter at lower metallicity. The current consensus appears to be that there is a low metallicity plateau in log(N/O), indicating the existence of primary nitrogen \citep[see, for example,][]{1993MNRAS.265..199V,1998ApJ...497L...1V,2002MNRAS.330...75C,2006ApJ...636..214V,2010ApJ...720.1738P,2009MNRAS.398..949P,2012ApJ...754...98B,2013ApJ...765..140A}. However, these previous works were not confined to small isolated dwarf galaxies. Results in our earlier paper on two isolated Local Void dwarf galaxies \citep{2014ApJ...780...88N}, indicated that log(N/O) did not plateau at low metallicity, suggesting no evidence for primary nitrogen. In this paper we present additional evidence for this. The paper is structured as follows: in Section 2 we detail the sample selection, the spectroscopic observations, and the data reduction details. H$\alpha$ images of each observed target, spectra, and de-reddened nebular emission line fluxes are presented in Section 3. In Section 4 we present the principal results of the analyses: electron temperatures, gas-phase nebular metallicities with the diagnostic grids, the nitrogen to oxygen flux ratios, the {[}S {\sc ii}{]} line ratios and electron densities, and the luminosity--metallicity results. In Section 5 we discuss these results, including the anomalies, and in Section 6 we present our conclusions. A discussion of methods used to estimate errors in the emission line fluxes is given in the Appendix. | In this paper we have presented the results of observations of seventeen \HII regions in thirteen small isolated dwarf irregular galaxies, most from the SIGRID sample, all but one exhibiting the [O~{\sc iii}] auroral line. All have measured oxygen abundances $<$8.2 ($<$0.3 Z$\odot$), one has an apparent abundance of 7.44 and another very low metallicity object with Z$\sim$7.2. We have derived a method for calculating total gas-phase oxygen abundances using only the optical spectra between 3700 and 7000 \AA. This method gives very similar results to previous empirical fit methods. From an analysis of abundances and ionization coefficients using the diagnostic grids developed by \cite{2013ApJS..208...10D}, we find the direct method oxygen abundances are consistently within 0.07 dex of the strong line diagnostic results, making allowance for the oxygen locked up in dust grains. From the line ratio of the two red [S~{\sc ii}] lines we find that the electron densities occurring in the objects observed are between $\sim$5 and 100 cm$^{-3}$. The nitrogen abundance, as expressed in log(N/O), continues the trend evident in \cite{1998AJ....116.2805V}, but from this sample we find no clear evidence for a nitrogen floor. There is increased scatter at lower oxygen abundances, and some evidence for a bifurcation in the trend, possibly due to the presence of WN stars in some of the \HII regions. The slope of the luminosity--metallicity relation for these observations is very close to that for void galaxies in \cite{2011AstBu..66..255P}. The spectra from an apparently very low metallicity galaxy, J1118-17s2, show no nitrogen lines: we intend to undertake follow up observations on this galaxy to estimate the metallicity more accurately. | 14 | 3 | 1403.6903 |
1403 | 1403.6768_arXiv.txt | If the B-mode signal in the CMB polarization seen by the \textsc{Bicep2} experiment is confirmed, it has dramatic implications for models of inflation. The result is also in tension with Planck limits on standard inflationary models. It is therefore important to investigate whether this signal can arise from alternative sources. If so, this could lessen the pressure on inflationary models and the tension with Planck data. We investigate whether vector and tensor modes from primordial magnetic fields can explain the signal. We find that in principle, magnetic fields generated during inflation can indeed produce the required B-mode, for a suitable range of energy scales of inflation. In this case, the primordial gravitational wave amplitude is negligible, so that there is no tension with Planck and no problems posed for current inflationary models. However, the simplest magnetic model is in tension with Planck limits on non-Gaussianity in the trispectrum. It may be possible to fine-tune the magnetogenesis model so that this non-Gaussianity is suppressed. Alternatively, a weaker magnetic field can pass the non-Gaussianity constraints and allow the primordial tensor mode to be reduced to $r\simeq0.09$, thus removing the tension with Planck data and alleviating the problems with simple inflationary models. | Gravitational waves which can be observed in the polarization pattern of the cosmic microwave background (CMB) are often called the `holy grail' of inflation. Recently their experimental detection has been announced by the \textsc{Bicep2} collaboration~\cite{Ade:2014xna}. This result, if confirmed by subsequent experiments, will be among the most important in cosmology since the discovery of the CMB. The reported tensor to scalar ratio of $r\simeq 0.2$ is very high. Such a high $r$ would allow for a detailed study of the primordial tensor spectrum. It would also imply an inflationary energy scale of $\simeq 1.4\times10^{16}\,$GeV, about 12 orders of magnitude above the highest energies reached by the LHC (Large Hadron Collider) at CERN. Furthermore, within the scenario of inflation, these gravitational waves are produced by the amplification of quantum fluctuations of the gravitational field itself. The result would therefore be evidence that the metric is a quantum field, i.e. our first observational indication of quantum gravity, even if only at the linear level. The significance of this result demands rigorous scrutiny at the experimental level, extending the excellent work of the \textsc{Bicep2} collaboration. For simple inflationary models, the result is in tension with Planck data~\cite{Ade:2013zuv}, which require $r\lesssim0.11$. While experimental scrutiny proceeds, we require also a rigorous scrutiny at the theoretical level. One line of investigation is to revisit the simple inflationary models (see e.g.~\cite{Byrnes:2014xua}). Here we tackle another question -- are there alternative explanations of the signal? One of the first ideas that comes to mind is: could this signal arise from vector modes in the gravitational field which are not generated during inflation, but later in the evolution of the Universe by some inhomogeneous source? Vector modes are a potentially ideal source for B-polarization in the CMB, since their transfer function to B-modes is nearly 10 times larger than the one from tensor modes (see, e.g.~\cite{book}, Fig.~5.7). Topological defects provide a potential origin of vector modes. The possible contribution from defects to the \textsc{Bicep2} signal has been investigated by~\cite{Lizarraga:2014eaa} and it is found that defects cannot generate the observed signal~\footnote{{except if the effective inter-string distance is extremely large}}, but a small contribution from defects can alleviate the tension with the Planck results. Here we investigate another possibility, primordial magnetic fields. Magnetic fields which are generated causally after inflation have blue spectra, $n_B=2$, and cannot leave an observable imprint in the CMB~\cite{spec}. However, magnetic fields can also be generated during inflation by couplings of the electromagnetic field to the inflaton or to the metric~\cite{TW}. In this case, they can have an arbitrary spectrum with $n_B>-3$, where $n_B\simeq -3$ is scale invariant. The imprint of such magnetic fields in the CMB has been studied extensively; see e.g.~\cite{pedro} and references therein. It has been found that all modes (scalar, vector and tensor) contribute with similar amplitudes to the CMB temperature anisotropies and to the E-polarization, but the B-polarization of the compensated mode is dominated by the vector mode~\cite{Shaw:2009nf}. We show below that the combination of compensated and passive modes from a magnetic field alone can reproduce very well the \textsc{Bicep2} result without invoking primordial tensor modes (i.e. taking $r$ to be negligible). However, if we require that the initial magnetic field be Gaussian, then limits on non-Gaussianity are in tension with a pure magnetic field solution. On the other hand, the inclusion of a weaker magnetic field, consistent with non-Gaussianity bounds, can reduce the required inflationary tensor contribution to $r\simeq 0.09$, thus removing the present tension with temperature data from Planck. | } B-modes from magnetic fields can reproduce the \textsc{Bicep2} results with no contribution from inflationary gravitational waves, i.e. with $r\simeq0$. This requires, however, that the fields are generated during inflation with non-Gaussian statistics, in such a way that their energy-momentum tensor is nearly Gaussian. As far as we are aware, no specific mechanism to produce such fields has been proposed in the literature so far. If Gaussian magnetic fields are generated during inflation, then the non-Gaussianity induced by the fields that are required to replace the $r\simeq 0.2$ tensor mode, are in tension with the trispectrum limits from Planck~\cite{Trivedi:2013wqa}. Nevertheless, a reduced magnetic contribution together with a small amount of dust can bring the required tensor amplitude down to $r\simeq 0.09$. This mitigates the tension with Planck bounds~\cite{Ade:2013zuv}, which require a relatively strong negative running of the scalar spectrum if $r\simeq 0.2$. Finally, we note that there is a strong-coupling problem of magnetogenesis. If a magnetic field is generated during inflation by a term ${\cal L} \supset f^2(\phi)F_{\mu\nu}F^{\mu\nu}$ in the Lagrangian, the `running' of $f$ during inflation, which is needed for a nearly scale-invariant spectrum, requires that the coupling of the electromagnetic field to the electron $\propto e/f(\phi)$ has been strong during most of the inflationary epoch and perturbation theory cannot be trusted~\cite{Demozzi}. This can be alleviated by postulating a very low inflationary scale, or a very blue magnetic field spectrum~\cite{Raj1}. It has been shown~\cite{Raj2} that for an inflationary scale $\sim10^{16}\,$GeV, $f$ is so severely constrained that only fields with $B_1\lesssim 10^{-30}\,$G can be generated. However, there are also other ways to evade the strong-coupling problem, for example breaking gauge-invariance during inflation~\cite{CCD} or choosing different couplings to the electromagnetic field. Taking these caveats into account, we believe it is fair to say that in principle, the observed B-mode signal could be due to inflationary magnetic fields. However, even though contrary to defects~\cite{Lizarraga:2014eaa}, B-modes from magnetic fields have the right shape, the simplest models are problematic since they are in tension with the upper limit from Planck on the trispectrum. On the other hand, a weaker magnetic field compatible with non-Gaussianity constraints can reproduce the observed B-mode, while reducing the primordial signal to $r\simeq 0.09$, thus removing the tension with Planck data.\vspace*{-0.6cm} \[\]{\bf Acknowledgements:} It is a pleasure to thank Chiara Caprini, Anthony Challinor, Dani Figueroa, Rajeev Jain, Andrii Neronov, Christophe Ringeval, Lorenzo Sorbo and Kandaswamy Subramanian for discussions and Richard Shaw for his help with the modified CAMB code. We also thank Brian G. Keating for calling to our attention the \textsc{Polarbear} data. CB is supported by King's College Cambridge. RD acknowledges support from the Swiss National Science Foundation. RM is supported by the South Africa Square Kilometre Array Project, the South African National Research Foundation and the UK Science \& Technology Facilities Council (grant ST/K0090X/1). \vspace*{-0.21cm} | 14 | 3 | 1403.6768 |
1403 | 1403.4247_arXiv.txt | {The distribution of various interstellar gas components and the pressure in the interstellar medium (ISM) is a result of the interplay of different dynamical mechanisms and energy sources on the gas in the Milky Way. The scale heights of the different gas tracers, such as H\,{\sc i} and CO, are a measure of these processes. The scale height of [C\,{\sc ii}] emission in the Galactic plane is important for understanding those ISM components not traced by CO or H\,{\sc i}.} {We determine the average distribution of [C\,{\sc ii}] perpendicular to the plane in the inner Galactic disk and compare it to the distributions of other key gas tracers, such as CO and H\,{\sc i}.} {We calculated the vertical, $z$, distribution of [C\,{\sc ii}] in the inner Galactic disk by adopting a model for the emission that combines the latitudinal, $b$, spectrally unresolved BICE survey, with the spectrally resolved $Herschel$ Galactic plane survey of [C\,{\sc ii}] at $b = 0\degr$. Our model assumed a Gaussian emissivity distribution vertical to the plane, and related the distribution in $z$ to that of the latitude $b$ using the spectrally resolved [C\,{\sc ii}] $Herschel$ survey as the boundary solution for the emissivity at $b=0\degr$.} {We find that the distribution of [C\,{\sc ii}] perpendicular to the plane has a full-width half-maximum of 172 pc, larger than that of CO, which averages $\sim$110 pc in the inner Galaxy, but smaller than that of H\,{\sc i}, $\sim$230 pc, and is offset by -28 pc.} {We explain the difference in distributions of [C\,{\sc ii}], CO, and H\,{\sc i} as due to [C\,{\sc ii}] tracing a mix of ISM components. Models of hydrostatic equilibrium of clouds in the disk predict different scale heights, for the same interstellar pressure. The diffuse molecular clouds with [C\,{\sc ii}] but no CO emission likely have a scale height intermediate between the low density atomic hydrogen H\,{\sc i} clouds and the dense CO molecular clouds. } {} | \label{sec:introduction} The star formation rate in the Galaxy may be related to the pressure of the interstellar medium (ISM), which itself is a function of the interplay of dynamical processes and energy sources on the interstellar gas. It has also been suggested that ISM pressure plays a role in the formation of giant molecular clouds \cite{Blitz2004,Blitz2006}. Thus, the vertical ($z$) distribution of the various ISM gas components is an important parameter for understanding these dynamical processes throughout the Milky Way. In the Galaxy, interstellar clouds are distributed in a thin disk about the mid-plane at $b=0\degr$ and their scale height depends, in hydrostatic equilibrium, on a number of factors, including thermal pressure, the random motion of the clouds, magnetic pressure, ionization pressure, and the gravitational force of the stars and gas in the disk. Thus, determining the $z$--distribution, which may be different for various ISM components, provides information about these parameters. The scale height of the diffuse atomic hydrogen clouds is known from extensive maps of the H\,{\sc i} 21-cm line \cite[c.f.][]{Boulares1990,Dickey1990} and that for the dense molecular hydrogen clouds from large scale maps of the $^{12}$CO $J=1\rightarrow 0$ rotational line \cite[e.g.][]{Sanders1984,Dame1987,Bronfman1988,Clemens1988,Malhotra1994,Dame2001,Jackson2006}. The scale height for H\,{\sc i} and H$_2$ (as traced by CO) clouds are different and each varies by a factor of $\sim$3 across the inner Galaxy, and increases in the outer Galaxy \citep{Narayan2002}. In addition to tracing gas that can be observed in H\,{\sc i} and CO, the 1.9 THz emission from ionized carbon, [C\,{\sc ii}], traces the H$_2$ gas where carbon is ionized but little, or no, CO or neutral carbon is found (the CO-dark H$_2$ gas), and also traces the warm ionized medium (WIM). The scale height for clouds traced by [C\,{\sc ii}] is not well established because the necessary spectral line surveys of its 158-\micron line have not, until recently been available. The COBE FIRAS instrument made the only large-scale survey of spectrally unresolved [C\,{\sc ii}] \citep{Wright1991,Bennett1994}, however, COBE with its 7$\degr$ beam and $\sim$1000 \kms velocity resolution, is unable to resolve the latitudinal distribution. There are two moderate-scale Galactic surveys of spectrally unresolved [C\,{\sc ii}], the Far-Infrared Line Mapper (FILM) onboard the Infrared Telescope in Space (IRTS) \citep[][]{Shibai1994,Makiuti2002} and the Balloon-borne Infrared Carbon Explorer (BICE) \citep{Nakagawa1998}, and an earlier small-scale survey with the Balloon-borne Infrared Telescope (BIRT) \citep{Shibai1991}. However, there is only one spectrally resolved survey, the {\it Herschel} open time key program, Galactic Observations of Terahertz C+, hereafter GOT C+ \citep[see][]{Langer2010,Pineda2013,Langer2014}. FILM and BICE had an angular resolution of order ten to fifteen arcminutes, sufficient to determine the latitudinal $b$ distribution of [C\,{\sc ii}], but the velocity resolution was, at best, $\sim$175 \kms for BICE (for FILM it was $\sim$750 \kms and for BIRT 143 \kmsnospace). Only the GOT C+ survey had the spectral resolution ($<$1 \kmsnospace) sufficient to resolve the velocity structure of individual gas clouds and thus locate their Galactic radius using a position--velocity rotation curve. However, GOT C+ surveyed [C\,{\sc ii}] sparsely in longitude, $l$, and latitude, $b$. Here we determine the average scale height distribution of [C\,{\sc ii}] in $z$ by combining the high spectral resolution [C\,{\sc ii}] radial distribution from GOT C+ at $b = 0\degr$ with the BICE latitudinal angular distribution in $b$ by adopting a hydrostatic model result for the distribution in $z$. We use the BICE survey to determine the scale height of [C\,{\sc ii}] emission because it had better coverage in both longitude and latitude in the Galactic disk than the BIRT and FILM surveys. BICE had an angular resolution of 15$^\prime$, and spectral resolution of $\sim$175 \kmsnospace. The GOT C+ survey contains several hundred lines-of-sight of spectrally resolved [C\,{\sc ii}] emission throughout the Galactic disk from $l=0\degr$ to 360\deg and $b= 0\degr$, $\pm0.5\degr$, and $\pm1.0\degr$. However, because GOT C+ is a sparse survey it does not have sufficient coverage in latitude $b$ to derive a smooth continuous distribution in the vertical distribution $z$. In contrast, the BICE survey had insufficient spectral resolution, but had much better coverage in $b$, but only observed up to latitudes $b = \pm 3\degr$ and only within longitude $350\degr \le l \le 25\degr$. GOT C+ has a 3-$\sigma$ sensitivity $\sim$ 0.24 K \kms \citep{Langer2014} which, over the bandwidth of the velocity resolution of BICE, corresponds to 7.4$\times$10$^{-6}$ ergs s$^{-1}$ cm$^{-2}$ sr$^{-1}$. BICE has a 3-$\sigma$ detection limit $\sim$2$\times$10$^{-5}$ ergs s$^{-1}$ cm$^{-2}$ sr$^{-1}$ \citep{Nakagawa1998}. Thus GOT C+ is almost three times more sensitive than BICE and capable of detecting the [C\,{\sc ii}] emission seen by BICE. We also use results from FILM \citep{Shibai1996,Makiuti2002}, which observed [C\,{\sc ii}] at higher latitudes than BICE, to recalibrate the results of \cite{Nakagawa1998} for $|b| = 3\degr$ to $4\degr$. FILM observed [C\, {\sc ii}] along a great circle crossing the plane at $l = 50\degr$ (inner Galaxy) and 230\deg (outer Galaxy), but measured the [C\,{\sc ii}] intensity for larger latitudes than BICE. However, \cite{Shibai1996} and \cite{Makiuti2002} smoothed their data to 1\deg to improve their sensitivity, insufficient to use at low latitudes to determine the scale height in the disk. \cite{Nakagawa1998} compared the latitudinal distribution of [C\,{\sc ii}] from BICE with those of H\,{\sc i}, $^{12}$CO, and far-infrared dust emission, finding that H\,{\sc i} had the largest distribution in $b$ followed by [C\,{\sc ii}] and then far-infrared dust emission, and that CO had the smallest distribution in $b$. However, without knowing where the [C\,{\sc ii}] emission came from they could not assign a spatial scale height to the gas traced by the [C\,{\sc ii}] 158-\micron line. We begin with a summary of the BICE and GOT C+ [C\,{\sc ii}] distributions and then derive an approximate relationship between the radial and $b$ distributions that allow us to deconvolve the [C\,{\sc ii}] distribution in $z$ in the disk. Finally, we compare the [C\,{\sc ii}] scale height with those of CO and H\,{\sc i} and discuss its implications for the understanding the sources of [C\,{\sc ii}] emission. | \label{sec:discussion} There is a very simple explanation why the scale height for [C\,{\sc ii}] is more than that of CO but less than that of H\,{\sc i}. While H\,{\sc i} traces mainly the atomic hydrogen clouds in the Galaxy and CO traces the dense molecular clouds, [C\,{\sc ii}] traces both of these regions, as well as the diffuse molecular clouds that have [C\,{\sc ii}] but no CO emission (CO-dark H$_2$ clouds), and the WIM. The solution for the density distribution of clouds in hydrostatic equilibrium is a Gaussian function, $\propto exp(-0.5(z/z_0)^2)$, with the scale factor, $z_0$ proportional to the velocity dispersion, $<v_i^2>^{0.5}$, where $i$ labels the ISM cloud component. The diffuse atomic clouds with their lower densities have higher velocity dispersions, $<v_{\rm H}^2>^{0.5}$ $\sim 10$ \kmsnospace, while those for the denser CO clouds have $<v_{\rm CO}^2>^{0.5}$ $\sim 5$ \kms \cite[see discussion and models in][]{Narayan2002}. Thus in hydrostatic equilibrium, for an equal ISM pressure, the H\,{\sc i} clouds have a larger scale height than those of the denser CO clouds. \cite{Pineda2013} found that the PDRs of dense molecular clouds emit $\sim$43$\%$ of the total [C\,{\sc ii}] throughout the plane at $b$ = 0\degr, diffuse atomic hydrogen clouds $\sim$23$\%$, diffuse molecular clouds (CO-dark H$_2$ clouds) $\sim$30$\%$., and for the WIM estimated $\sim$4$\%$. In the inner Galaxy, $R_{\rm gal} <$ 9 kpc, where we evaluate the [C\,{\sc ii}] distribution, these percentages are only slightly different. Therefore, it is no surprise that the distribution of [C\,{\sc ii}], which arises from all ISM components, would have a distribution intermediate between that of the dense CO clouds and less dense atomic H\,{\sc i} clouds and the WIM. Thus [C\,{\sc ii}] traces a mixture of clouds of different mass, density, and velocity dispersion, some of which are also traced by H\,{\sc i} and CO, and the WIM. The different scale heights then depend on the different physical properties and energetics of the clouds that enter into the hydrostatic mechanisms responsible for the distribution of gas in the plane. As seen in Figure~\ref{fig:fig_5_CII_CO_HI_vs_z}, the derived $z$ distribution in [C\,{\sc ii}] does not follow that of H\,{\sc i} for $z$ greater than about $\sim$200 pc. At this height there are few dense molecular clouds as traced by CO and likely very few diffuse molecular clouds (CO-dark H$_2$ clouds), so any [C\,{\sc ii}] would have to come from the H\,{\sc i} clouds and$/$or the WIM. \cite{Makiuti2002} compared the distribution of [C\,{\sc ii}] and H\,{\sc i} at high latitudes and conclude that the [C\,{\sc ii}] emission comes primarily from the WIM (see their Figure 3). However, the FILM results at high latitude are limited to a region near the solar radius and cannot be extrapolated across the Galaxy. The combined GOT C+ and BICE data suggest that this conclusion also holds for the inner Galaxy as well, because H\,{\sc i} clouds high above the plane have low densities and are not likely to emit [C\,{\sc ii}] efficiently. This low emissivity is due to the difference in the excitation conditions for H\,{\sc i} and [C\,{\sc ii}]. The intensity of H\,{\sc i} is proportional to its column density in the optically thin regime, \begin{equation} I([{\rm H\,I}]) = 5.5\times 10^{-19}N({\rm H\,I}), \end{equation} \noindent in units of (K \kmsnospace), and is relatively insensitive to density and kinetic temperature. In contrast the [C\,{\sc ii}] emission is very sensitive to kinetic temperature, $T_{\rm kin}$ because the energy of the upper level $^{\rm 2}P_{\rm 3/2}$, $E_u/k = 91.25 K$ is typically higher than the gas temperature in neutral clouds, and density where the atomic, n(H), and molecular, n(H$_2$), hydrogen densities are much lower than the critical densities for thermalizing the C$^+$, $n_{\rm cr}(H)\sim$ 3000 cm$^{-3}$ and the recently revised value $n_{\rm cr}(H_2)\sim$4500 cm$^{-3}$ \citep{Wiesenfeld2014}. The radiative transfer equation for the [C\,{\sc ii}] intensity for optically thin emission is given in \cite{Goldsmith2012} and \cite{Langer2014}, and, in the diffuse clouds, such that the intensity can be written as, \begin{equation} I_j([{\rm C\,II}]) = 1.73\times10^{-16} e^{{-\Delta}E/kT} \frac{ n(j)}{n_{\rm cr}(j)} N_j({\rm C^+}) , \end{equation} \noindent where $I$ is in units of K \kms and the index $j$ labels H\,{\sc i} or H$_2$. Therefore, while the H\,{\sc i} intensity depends only on the column density of atomic hydrogen, the [C\,{\sc ii}] intensity also depends very sensitively on the density and will be much smaller in low density atomic hydrogen clouds above the plane. The scale height for [C\,{\sc ii}] derived here depends on the radial distribution derived from the GOT C+ sampling at $b$=0$\degr$, which as noted above is a sparse sample. The premise of the GOT C+ survey was that a well designed unbiased sampling in longitude would represent statistically the distribution of [C\,{\sc ii}] in the Galactic plane. Therefore, the fact that GOT C+, along with our model of the $z$ distribution, reproduces the total flux observed by BICE (rescaled to the FILM calibration) supports this approach. Another potential uncertainty is the adoption of a Galactic rotation curve in \cite{Pineda2013} to locate the source of the [C\,{\sc ii}] emission. In \cite{Langer2014} we adopted a rotation curve based on gas-flow hydrodynamical models to assign a distance based on velocity. We found that it made a difference mainly in the inner Galaxy, $|l|\le$6$\degr$, but this region is mostly excluded from the BICE analysis (see above). There are also regions where clouds have peculiar velocities due to Galactic dynamics, where the rotation curve may assign the wrong distance. For example, \cite{Zhang2014} find non-rotational cloud motions at the end of the Galactic bar from parallax observations of masers at $l\sim$30\degr. While 30\degr\ is outside of the longitudinal range observed by BICE, this region contributes to our GOT C+ data set. We cannot quantitatively assess the error introduced into our radial distribution but note that, because we average lines of sight from all across the Galaxy, the edges of the bar contribute a small fraction of the emission in any given ring. We cannot calculate the radial dependence for the [C\,{\sc ii}] FWHM from the GOT C+ survey without a better sampling in $b$. However, to gain some insight on the effect of a variable FWHM($R_{gal}$), we assume that it varies similar to that for CO. \cite{Sanders1984} and \cite{Clemens1988} found that the CO scale height varied roughly as $R_{\rm gal}^{0.5}$ between 3 and 9 kpc. We replaced $z_o$ in Equation (3) with one that varied $\propto R_{\rm gal}^{0.5}$ for $R_{gal} >$3 kpc at the value for 3 kpc. We solved for the scale factor that best fit the BICE distribution in $b$, similar to what was done to determine an average scale factor. We find that the best fit is given by, FWHM($R_{\rm gal})=172(R_{gal}/4.7)^{0.5}$ pc. Thus the average value for FWHM for [C\,{\sc ii}] of 172 pc corresponds to the radial solution at $\sim$ 4 to 5 kpc, essentially in the molecular ring. For the assumed radial dependence, the FWHM ranges from $\sim$140 pc to $\sim$ 230 pc over Galactic radii 3 kpc to 8 kpc. We have combined the GOT C+ spectrally resolved [C\,{\sc ii}] survey in the Galactic plane at $b = 0\degr$ with the latitudinal distribution derived from the BICE survey of spectrally unresolved [C\,{\sc ii}] to derive, for the first time, the average scale height of [C\,{\sc ii}] over the inner Galactic plane. GOT C+ is slightly more sensitive than BICE and the total flux measured by GOT C+ is close to that of BICE given the uncertainties of the BICE calibration. Therefore these two surveys are likely tracing the same ISM gas components. The average distribution in the inner Galactic disk is well fit by a single Gaussian with FWHM([C\,{\sc ii}]) = 172 pc and an offset -28 pc below the plane ($b = 0\degr$). In this paper we find that the [C\,{\sc ii}] distribution is larger in $z$ than that of CO, but smaller than H\,{\sc i}. The origin of the [C\,{\sc ii}] emission has been attributed to different sources by various authors based on the spectrally unresolved surveys. However, the result here suggests to us a more complicated picture with [C\,{\sc ii}] tracing a mix of ISM cloud categories. The GOT C+ data for $b \ne 0\degr$ may be able to give some insights on the distribution of the different ISM components, but to determine completely the distribution of [C\,{\sc ii}] for the separate ISM components as a function of Galactic radius and $z$, we will need more detailed spectrally resolved latitudinal maps across the Galaxy and with finer steps in $b$. We also need to extend the spectral line observations to higher values of $b$ than observed in the GOT C+ survey to understand the contributions of the warm ionized medium and low density high latitude H\,{\sc i} clouds to the [C\,{\sc ii}] emissivity above the disk. | 14 | 3 | 1403.4247 |
1403 | 1403.1181_arXiv.txt | Astroparticles offer a new path for research in the field of particle physics, allowing investigations at energies above those accesible with accelerators. Ultra-high energy cosmic rays can be studied via the observation of the showers they generate in the atmosphere. The Pierre Auger Observatory is a hybrid detector for ultra-high energy cosmic rays, combining two complementary measurement techniques used by previous experiments, to get the best possible measurements of these air showers. Shower observations enable one to not only estimate the energy, direction and most probable mass of the primary cosmic particles but also to obtain some information about the properties of their hadronic interactions. Results that are most relevant in the context of determining hadronic interaction characteristics at ultra-high energies will be presented. | Cosmic rays from astrophysical sources provide a natural beam of ultra high energy particles that can be used to probe particle interactions at the highest energies. The interpretation of cosmic ray measurements requires modeling of hadronic interactions in an energy range beyond that which can be studied in accelerator experiments. The knowledge of the relevant properties of hadronic interactions in this energy range is therefore of central importance for the interpretation of the cosmic ray data. Nevertheless, it is in principle possible to obtain information about hadronic interactions from the cosmic ray observations, but dealing with the fact that this natural cosmic ray beam has an unknown energy spectrum and an unknown mass composition. The mass composition of the primary particles must be estimated from the same data set. Solving the ambiguity between composition and hadronic interaction modeling is a key problem for ultra-high energy cosmic ray observations. The Pierre Auger Observatory, the world\'s largest cosmic ray observatory, is located near Malarg\"{u}e, in the Province of Mendoza, Argentina. It was designed to investigate the origin and the nature of ultra-high energy cosmic rays by taking advantage of two available techniques to detect extensive air showers initiated by ultra-high energy cosmic rays: a surface detector (SD) array and a fluorescence detector (FD). The SD consists of an array of about 1600 water-Cherenkov surface detectors deployed over a triangular grid of 1.5 km spacing and covering an area of 3000 km$^2$. The SD is overlooked by 27 fluorescence telescopes, grouped in four sites, making up the fluorescence detector. The FD observes the longitudinal development of the shower in the atmosphere by detecting the fluorescence light emitted by excited nitrogen molecules and Cherenkov light induced by shower particles in air. The FD provides a calorimetric measurement of the primary particle energy. These two detection methods are complementary, so that combining them in hybrid mode will help resolve mass composition and hadronic interaction information. | current hadronic interaction models do not accurately describe the muon signal. | 14 | 3 | 1403.1181 |
1403 | 1403.6597_arXiv.txt | {Some runaway stars are known to display IR arc-like structures around them, resulting from their interaction with surrounding interstellar material. The properties of these features as well as the processes involved in their formation are still poorly understood.} {\referee{We aim at understanding the physical mechanisms that shapes the dust arc observed near the runaway O-star AE\,Aur (HD\,34078).}} {\referee{We obtained and analyzed a high spatial resolution ($4.4''$) map of the $^{12}$CO(1-0) emission that is centered on HD\,34078, and that combines data from both the IRAM interferometer and 30m single-dish antenna.}} {One third of the 30m flux mainly originates from two small (no larger than $5''\times10''$ or $0.013 \times 0.026\pc$), and bright (1 and 3\K{} peak temperatures) CO globulettes. The line of sight towards HD\,34078 intersects the outer part of one of them, which accounts for both the properties of diffuse UV light observed in the field by~\citet[]{France.2004} and the numerous molecular absorption lines detected in HD\,34078's spectra, including those from highly excited \Ht{}. Their modelled distance from the star \referee{(0.2\pc)} is compatible with the fact that they lie on the 3D paraboloid which fits the arc detected in the 24\mum{} Spitzer image. Four other compact CO globulettes are detected in the mapped area, all lying close to the rim of this paraboloid. These globulettes have a high density and linewidth, and are strongly pressure-confined or transient.} {The presence of molecular globulettes at such a close distance from an O star is unexpected, and probably related to the high proper motion of HD\,34078. Indeed, the good spatial correlation between the CO globulettes and the IR arc suggests that they result from the interaction of the radiation and wind emitted by HD\,34078 with the ambient gas. However, the details of this interaction remain unclear. A wind mass loss rate significantly larger than the value inferred from UV lines is favored by the large IR arc size, but does not easily explain the low velocity of the CO globulettes. The effect of radiation pressure on dust grains also meets several issues in explaining the observations. Further observational and theoretical work is needed to fully elucidate the processes shaping the gas and dust in bow shocks around runaway O stars.} | \TabObservations{} \FigContext{} The interstellar medium (ISM) was discovered in 1907 by studying atomic absorption lines (from NaI and CaII) seen in the visible spectrum of bright stars. The first interstellar molecules (CN, CH, and \CHp{}) were also detected in the same way in the years 1937--1940. In the seventies, the Copernicus satellite systematically studied \Ht{} ultraviolet (UV) absorption lines to investigate the properties of diffuse interstellar gas with low visual extinction ($\Av \le 1\magn$), and this powerful method is still largely used today (\cf{} FUSE, HST/STIS and HST/COS programs). The gas in diffuse clouds is mainly neutral, warm (typically 80\K{}), and (usually) of low density ($100-500\pccm $). Such regions correspond to the transition from atomic to molecular hydrogen, where carbon is still mostly ionised or neutral, with $\N{\emr{CO}} < \mbox{a few~} 10^{16}\pscm$, and $\N{\emr{C}} \sim 3.10^{17}\pscm$. In this framework, the foreground absorption against the O9.5 star HD\,34078 stands out for its peculiar properties. HD\,34078 was ejected about 2.5\,Myr ago from the Orion region \citep[]{Blaauw.1953,Bagnuolo.2001} and is now the fastest runaway star in the local ISM, with a velocity of $\Vstar \sim 150\kms$ \citep{Tetzlaff.2011}. The line of sight towards HD\,34078 should thus offer a means to detect small scale $(5-50\au)$ density or abundance variations in the diffuse interstellar medium in only a few decades. However, as absorption line studies progressed, it was realized that this line of sight exhibits peculiar properties compared to the usual diffuse ISM on other lines of sight. Very abundant CH and \CHp{} is measured, with some time variability \citep{Rollinde.2003}. The direct starlight suffers larger reddening than scattered light in the surrounding nebula \citep{France.2004}. In addition, UV absorption studies with FUSE reveal an unusually large amount of highly excited \Ht{}, indicating the unexpected presence of dense ($n_H \simeq 10^4\pccm$) and strongly irradiated molecular gas at about 0.2\pc{} from the star~\citep{Boisse.2005}. The latter property could be related to the recent interaction of the star with the IC\,405 reflection nebula \citep{Herbig.1958}. Indeed, the emission of hot dust at 24\mum{} imaged with Spitzer/MIPS clearly delineates a parabolic curve (see Fig.~\ref{fig.context}), interpreted as the tip of a bow shock resulting from the interaction of the fast stellar wind with the preexisting diffuse gas of IC\,405 \citep{France.2007}. In order to gain 2D information on the molecular gas structure and kinematics in this bow shock, \citet{Boisse.2009} conducted sensitive \twCO{}\Jtwo{} mapping in a narrow field of view around HD\,34078 using the IRAM-30m telescope at an angular resolution of $12''$. On top of a widespread CO component, brighter emission was detected in a clumpy ``filament'' peaking slightly below the 24\mum{} arc, confirming ongoing interaction between the star and the surrounding cloud. However, the CO velocity field showed a gradient mainly perpendicular to the star-apex axis, instead of mainly along this axis as expected for a steady-state wind bow shock~\citep[see e.g.][]{Wilkin.1996}. In addition, \citet{Boisse.2009} noted that the apparent distance between the star and the IR arc of $15''$, corresponding to 0.04\pc{} at the distance of HD\,34078~\citep[$\simeq$ 530\pc{} { assuming $M_V = -4.2$}][]{Herbig.1999}, is largely incompatible with the prediction for a stationary bow shock for the wind mass-flux derived from UV line analysis \citep{Martins.2005}. They then proposed that we might be witnessing the recent birth of the bow shock, or that radiation pressure on grains might play a role in increasing its size. Alternatively, the wind mass-flux might have been largely underestimated (as proposed by \citet{Gvaramadze.2012} to explain the IR bow shock size in $\zeta$\,Oph) or some other under-appreciated physical process could be at play. Distinguishing between these hypotheses is crucial in order to better understand the processes that govern the interaction between HD\,34078 and the surrounding ISM, how they impact the absorbing gas properties on the line of sight, and how they may affect the wind mass-flux determinations from IR bow shock sizes in other runaway O stars. As a step towards this goal, we mapped the bowshaped IR arc around HD\,34078 in \twCO{}\Jone{} with the Plateau de Bure Interferometer (PdBI, complemented with IRAM-30m single-dish data to provide the short-spacings) at $\simeq 4''$ resolution, comparable to the 24\mum{} Spitzer image and 3 times better than the \twCO{}\Jtwo{} map of~\citet{Boisse.2009}. The observations and data reduction are presented in Sect.~\ref{sec:observations}. The resulting properties of the detected CO structures are described in Sect.~\ref{sec:results}. We discuss their implications in Sect.~\ref{sec:discussion}. We summarize and conclude in Sect.~ \ref{sec:conclusions}. | \label{sec:conclusions} We described the calibration and construction of the \twCO{} \Jone{} imaging at $\sim10^{-3}\pc$ ($\sim4.4''$) of the $\sim 0.31 \times 0.26\pc$ $(\sim 120''\times100'')$ toward the runaway O star HD\,34078, using observations from both the PdBI and IRAM-30m telescopes. The IRAM-30m data mainly features two unresolved globules: One at the south eastern edge of the observed field of view, not discussed here, and a second one around the star sightline. At the PdBI resolution, the latter appears to be composed mainly of 2 bright (1 and 3\K{} peak temperature) and compact (size $\le 10''$ or $0.026\pc$) globulettes linked together by an extended faint emission that amounts to two third of the total flux. The star sightline clearly intercepts the edge of globulette \#2. However, the spectrum in the direction of the star is double-peaked indicating that the star sightline also intercepts globulette \#1. These globulettes are responsible for the absorption lines from CH, \CHp{} and cold \Ht{} at 77\K{}, and also explain the large amounts of dense excited \Ht{} at 350\K{} on this sightline, the latter probably tracing a PDR on the illuminated face of the globulettes. The measured CO column density in this direction is compatible with the stellar reddening measured in far-UV and visible. The globulettes are small enough for the surrounding reflexion nebula to be less reddened than the star \citet{France.2007}. The imaged CO globulettes appear to approximately lie along the parabola walls. They are clearly not gravitationally bound. They may be pressure confined and they probably result from the interaction between the star and the preexisting diffuse gas (\eg{}, thermal instabilities in the bow shock). We quantified the actions of two competing processes in the interaction between the star and the preexisting diffuse nebula, IC\,405. The first one is ram-pressure due to the high velocity star wind. The second one is the radiative pressure (optically thin case) on the dust grains that entrain the gas through friction. Neither a non-stationary bow shock nor the effect of radiation pressure (in the optically thin limit) can explain at the same time (1) the observed large size of the IR arc in HD\,34078, (2) its distortion from a perfect parabolic shape, (3) its size ratio of 4 compared to that around $\zeta$\,Oph, and (4) its spatial correlation with dense CO globulettes. The most straightforward explanation for these 4 properties appears to be that the wind mass-flux is { 300-1000} times larger than indicated by UV lines, and close to the theoretical prescription of~\citet{Vink.2001}. The effect of radiation pressure on optically thick globulettes would be important to investigate but may also prove insufficient. Indeed, the low radial velocities observed in the CO globulettes appear puzzling for a steady-state wind bow shock. A study over a wider field of view in both CO and \Halpha{}, including regions both upstream and downstream from the IR arc, would be important to obtain additional constraints on the formation process and dynamical state of the CO globulettes. | 14 | 3 | 1403.6597 |
1403 | 1403.4592_arXiv.txt | The BICEP2 collaboration has for the first time observed the B-mode polarization associated with inflationary gravitational waves. Their result has some discomfiting implications for the validity of an effective theory approach to single-field inflation since it would require an inflaton field excursion larger than the Planck scale. We argue that if the quantum state of the gravitons is a {\em mixed} state based on the Bunch-Davies vacuum, then the tensor to scalar ratio $r$ measured by BICEP is different than the quantity that enters the Lyth bound. When this is taken into account, the tension between effective field theory and the BICEP result is alleviated. | 14 | 3 | 1403.4592 |
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1403 | 1403.4071_arXiv.txt | A new family of inflation models is introduced and studied. The models are characterised by a scalar potential which, far from the origin, approximates an inflationary plateau, while near the origin becomes monomial, as in chaotic inflation. The models are obtained in the context of global supersymmetry starting with a superpotential, which interpolates from a generalised monomial to an O'Raifearteagh form for small to large values of the inflaton field respectively. It is demonstrated that the observables obtained, such as the scalar spectral index and the tensor to scalar ratio, are in excellent agreement with the latest observations. Some discussion of initial conditions and eternal inflation is included. | 14 | 3 | 1403.4071 |
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1403 | 1403.0652_arXiv.txt | The discovery of giant planets orbiting close to their host stars was one of the most unexpected results of early exoplanetary science. Astronomers have since found that a significant fraction of these `Hot Jupiters' move on orbits substantially misaligned with the rotation axis of their host star. We recently reported the measurement of the spin-orbit misalignment for WASP-79b by using data from the 3.9~m Anglo-Australian Telescope. Contemporary models of planetary formation produce planets on nearly coplanar orbits with respect to their host star's equator. We discuss the mechanisms which could drive planets into spin-orbit misalignment. The most commonly proposed being the Kozai mechanism, which requires the presence of a distant, massive companion to the star-planet system. We therefore describe a volume-limited direct-imaging survey of Hot Jupiter systems with measured spin-orbit angles, to search for the presence of stellar companions and test the Kozai hypothesis. | 14 | 3 | 1403.0652 |
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1403 | 1403.5307_arXiv.txt | This is the first of two papers describing the second data release (DR2) of the Australia Telescope Large Area Survey (ATLAS) at 1.4~GHz, which comprises deep wide-field observations in total intensity, linear polarization, and circular polarization over the {\it Chandra} Deep Field-South and European Large Area {\it Infrared Space Observatory} Survey-South 1 regions. DR2 improves upon the first data release by maintaining consistent data reductions across the two regions, including polarization analysis, and including differential number counts in total intensity and linear polarization. Typical DR2 sensitivities across the mosaicked multi-pointing images are $30$~$\mu$Jy~beam$^{-1}$ at approximately $12\arcsec\times6\arcsec$ resolution over a combined area of 6.4~deg$^2$. In this paper we present detailed descriptions of our data reduction and analysis procedures, including corrections for instrumental effects such as positional variations in image sensitivity, bandwidth smearing with a non-circular beam, and polarization leakage, and application of the {\tt BLOBCAT} source extractor. We present the DR2 images and catalogues of components (discrete regions of radio emission) and sources (groups of physically associated radio components). We describe new analytic methods to account for resolution bias and Eddington bias when constructing differential number counts of radio components. | \label{ch4:SecIntr} Radio surveys are a cornerstone of modern astronomy. Counts of extragalactic radio sources per steradian per unit flux density provide fundamental constraints on galaxy evolution, as they implicitly encapsulate both the underlying redshift and luminosity distributions of source populations \citep[e.g.][]{1966MNRAS.133..421L}. In total intensity, the 1.4~GHz source counts are observed to flatten below 1~mJy, though the extent of this flattening is controversial because the results from deep surveys exhibit a large degree of scatter. To illustrate, see the compilation of surveys in Fig.~3 from \citet{2013PASA...30...20N} where there is a factor of 2 variation in the counts below 1~mJy. Some studies have attributed the large scatter in the faint counts to cosmic variance, namely to intrinsic differences between survey fields caused by source clustering \citep[e.g.][]{2004MNRAS.352..131S}. However, significant differences in the counts for fields observed by separate studies, such as the Lockman Hole \citep{2009MNRAS.397..281I}, indicate that calibration and data processing errors may be largely responsible for the scatter. Issues to consider include corrections for bandwidth smearing \citep[e.g.][]{2009MNRAS.397..281I}, Eddington bias \citep[e.g.][]{2006MNRAS.372..741S}, resolution bias \citep[e.g.][]{2008ApJ...681.1129B}, and non-instrumental factors such as source clustering in the field \citep[e.g.][]{2013MNRAS.432.2625H}. The present conclusion in the literature is that the scatter in the sub-millijansky counts is likely to be significantly affected by data processing differences between surveys \citep{2006MNRAS.371..963B,2007ASPC..380..189C, 2009MNRAS.397..281I,2010A&ARv..18....1D,2012ApJ...758...23C,2013MNRAS.432.2625H}. This conclusion motivates the need for studies that describe data reduction and analysis procedures in detail, so as to facilitate robust comparisons with other works and encourage future improvements. To date, very few surveys dedicated to extragalactic polarized radio sources have been conducted, primarily because of correlator limitations that have required polarization capabilities to be sacrificed for spectral resolution. Polarization surveys at 1.4~GHz include the NRAO VLA Sky Survey \citep[NVSS;][]{1998AJ....115.1693C} observed with the Very Large Array, which encompasses 82\% of the sky at resolution full-width at half-maximum (FWHM) 45\arcsec\ to a root mean square (rms) sensitivity in polarization of 0.29~mJy~beam$^{-1}$, surveys of the European Large Area {\it Infrared Space Observatory} Survey-North 1 (ELAIS-N1) region observed using the Dominion Radio Astrophysical Observatory (DRAO) Synthesis Telescope by \citet{2007ApJ...666..201T} over 7.43~deg$^2$ with resolution FWHM $\sim50\arcsec$ to 78~$\mu$Jy~beam$^{-1}$ and in a deeper follow-up study with the same facility by \citet{2010ApJ...714.1689G} over 15.16~deg$^2$ to 45~$\mu$Jy~beam$^{-1}$, and the Australia Telescope Low-Brightness Survey \citep[ATLBS;][]{2010MNRAS.402.2792S} which encompasses two fields observed with the Australia Telescope Compact Array (ATCA) over a total of 8.42~deg$^2$ with resolution FWHM $\sim50\arcsec$ to $\sim80\mu$Jy~beam$^{-1}$. A summary of polarization surveys at other radio wavelengths is presented by \citet{2012AdAst2012E..52T}. By cross-matching polarized 1.4~GHz sources with mid-infrared counterparts, \citet{2007ApJ...666..201T} identified the population of polarized millijansky sources as being extragalactic radio sources powered by AGNs. \citet{2010ApJ...714.1689G} found that the polarized emission from these sources was likely to originate in extended radio lobes. \citet{2002A&A...396..463M} and \citet{2004MNRAS.349.1267T} found an anti-correlation between the fractional linear polarization and total intensity flux densities of NVSS sources; faint sources were more highly polarized. This finding was supported for ELAIS-N1 sources by \citet{2007ApJ...666..201T} and \citet{2010ApJ...714.1689G}, and for ATLBS sources by \citet{2010MNRAS.402.2792S}. \citet{2004MNRAS.349.1267T}, \citet{2007ApJ...666..201T}, and \citet{2010ApJ...714.1689G} found that the Euclidean-normalised differential number-counts of polarized sources flattened at linearly polarized flux densities $L$~{\footnotesize $\lesssim$}~1~mJy to levels greater than those predicted by \citet{2004NewAR..48.1289B}; the latter predicted polarized source counts to $\mu$Jy levels by convolving total intensity source counts with a fractional polarization distribution modelled on NVSS data. \citet{2008evn..confE.107O} were unable to reproduce the observed flattening in a population modelling study. The observed flattening suggests the emergence of systematic changes in polarized source properties with decreasing flux density, such as higher ordering of magnetic fields in fainter sources, or perhaps the emergence of an unexpected faint population. To examine the emerging fractional polarization anti-correlation and source count flattening trends in more detail, deeper and higher angular resolution observations of the 1.4~GHz polarized sky are required. In this work we present reprocessed and new 1.4~GHz observations of the {\it Chandra} Deep Field-South \citep[CDF-S; Galactic coordinates $l\approx224\degree$, $b\approx-55\degree$;][]{2006AJ....132.2409N} and ELAIS-South 1 \citep[ELAIS-S1; $l\approx314\degree$, $b\approx-73\degree$;][]{2008AJ....135.1276M} regions, obtained as part of the Australia Telescope Large Area Survey (ATLAS) project with the ATCA. We collectively refer to these previous ATLAS papers as Data Release 1 (DR1) and denote the present work Data Release 2 (DR2). Given that DR1 did not include polarization analysis of the ATLAS data, we have chosen to reprocess the original observations to ensure consistent and improved data reduction and analysis between both the total intensity and polarization data and the two independent ATLAS regions. In preparation for ATLAS DR2, we have developed new tools to ensure accurate calculation of the statistical significance of flux density measurements in linear polarization \citep{2012MNRAS.424.2160H} and to ensure accurate measurement of these flux densities using the {\tt BLOBCAT} source extractor \citep{2012MNRAS.425..979H}. The motivations for ATLAS DR2 are to (i) present a detailed description of our data reduction and analysis procedures to inform future deep surveys such as those being developed for SKA Pathfinder facilities around the world (see summary of facilities described by \citealt{2012A&A...543A.113B} and \citealt{2013PASA...30...20N}), (ii) compute differential number counts for total intensity and linearly polarized objects (total intensity counts were not included in DR1), and (iii) investigate the nature of faint polarized sources and consider possible explanations for the fractional polarization trend seen in previous studies. Clearly, biases introduced at an early stage of data reduction have the potential to propagate through to the final data in a non-linear fashion, affecting the ability for that data to be used for unplanned and novel experiments in the future \citep[e.g.][]{2009ApJ...692..887C}. In this paper (Paper I) we focus on point (i) from above, regarding data reduction and the development of new techniques to produce high fidelity data suitable for investigating points (ii) and (iii). Results and discussion regarding points (ii) and (iii) will be presented in Paper II \citep{halesPII}. This paper is organised as follows. In \S~\ref{ch4:SecObs} we describe our ATLAS radio data and ancillary mid-infrared and optical data. In \S~\ref{ch4:SecDat} we outline our radio data reduction and post-processing procedures to obtain mosaicked images of total intensity, linear polarization (using rotation measure synthesis), and additionally circular polarization for the two ATLAS regions. In \S~\ref{ch4:SecInst} we describe instrumental effects of time-average smearing, bandwidth smearing, and polarization leakage, our methods to account for them in our ATLAS data, and the effective survey area boundaries. In \S~\ref{ch4:SecExt} we detail how radio components were detected and extracted in total intensity, linear polarization, and circular polarization, and how their flux densities were corrected to account for subtle noise-induced systematics. In \S~\ref{ch4:SecClass} we describe our implementation of two cross-identification and classification schemes: the first to group components into sources, to associate these sources with infrared sources, and to classify them according to their multiwavelength properties; and the second to associate linearly polarized components or polarization upper limits with total intensity counterparts and to classify these associations based on their polarized morphologies. In \S~\ref{ch4:SecCNC} we describe in detail corrections required to calculate total intensity and linear polarization differential number-counts, including a new fully analytic method to account for resolution bias. In \S~\ref{ch5:SecRes} we present the ATLAS DR2 total intensity and linear polarization images, and the radio component and source catalogues. We conclude in \S~\ref{ch4:SecConcl}. For reference, a selection of important symbols used in this work is presented in Table~\ref{tbl:symbols}. \begin{table*} \centering \caption{Selection of important symbols used in this work.\label{tbl:symbols}} \begin{tabular}{@{}cll@{}} \hline Symbol & Description & Defined \\ \hline $I_{\trm{\tiny MFS}}$ & total intensity mosaic produced using multi-frequency synthesis approach & \S~\ref{ch4:SecDatSubMFS} \\ $V_{\trm{\tiny MFS}}$ & circular polarization mosaic produced using multi-frequency synthesis approach & \S~\ref{ch4:SecDatSubMFS} \\ $I_{\ms i}$, $Q_{\ms i}$, $U_{\ms i}$ & $i$'th frequency channel mosaic in total intensity, Stokes $Q$, or Stokes $U$ & \S~\ref{ch4:SecDatSubPCI} \\ $\sigma_{\ms Q,i},\sigma_{\ms U,i}$ & rms noise map of $i$'th frequency channel mosaic in Stokes $Q$ or Stokes $U$ & \S~\ref{ch4:SecDatSubRM} \\ $\sigma_{\ms Q,U,i}$ & map of combined rms noise for Stokes $Q$ and Stokes $U$ in $i$'th frequency channel & \S~\ref{ch4:SecDatSubRM} \\ $I_{\trm{\tiny CA}}$ & total intensity mosaic produced using channel average approach & \S~\ref{ch4:SecDatSubRM} \\ $L_{\trm{\tiny RM}}$ & linear polarization mosaic produced using rotation measure synthesis & \S~\ref{ch4:SecDatSubRM} \\ $\sigma_{\trm{\tiny RM}}(x,y)$ & rms noise map for $L_{\trm{\tiny RM}}$ & \S~\ref{ch4:SecDatSubRM} \\ $\varpi$ & bandwidth smearing ratio (observed divided by true surface brightness) & \S~\ref{ch4:SecInstSubBS} \\ $K_{\trm{\tiny LEAK}}$ & total intensity to linear polarization leakage mosaic & \S~\ref{ch4:SecInstSubLeak} \\ $L_{\trm{\tiny RM}}^{\trm{\tiny CORR}}$ & $L_{\trm{\tiny RM}}$ corrected for polarization leakage & \S~\ref{ch4:SecInstSubLeak} \\ $F^{\trm{\tiny AREA}}$ & survey area & \S~\ref{ch4:SecInstSubArea} \\ $S_{\trm{\tiny peak}}$,$S_{\trm{\tiny int}}$ & peak or integrated surface brightness (more generally, $S$ denotes flux density) & \S~\ref{ch4:SecExtSubFlood} \\ $A_{\trm{\tiny S}}$ & detection signal-to-noise ratio & \S~\ref{ch4:SecExtSubFlood} \\ $V^{\trm{\tiny AREA}}$ & visibility area for detection & \S~\ref{ch4:SecExtSubFlood} \\ $\theta,B$ & observed or beam full-width at half-maximum & \S~\ref{ch4:SecExtSubDeconv} \\ $\Theta$ & deconvolved angular size & \S~\ref{ch4:SecExtSubDeconv} \\ $\gamma$ & slope of differential number-counts, $dN/dS \propto S^{-\gamma}$ & \S~\ref{ch4:SecExtSubDB} \\ $S_{\trm{\tiny ML}}$ & deboosted flux density using maximum-likelihood scheme & \S~\ref{ch4:SecExtSubDB} \\ $dN_{\trm{\tiny H03}}/dS$ & differential number-count fit from \citet{2003AJ....125..465H} & \S~\ref{ch4:SecExtSubDB} \\ $dN_{\trm{\tiny H03M}}/dS$ & modified version of $dN_{\trm{\tiny H03}}/dS$ & \S~\ref{ch4:SecExtSubDB} \\ $f_{\ms \Pi}$ & distribution of fractional linear polarization ($\Pi \equiv L/I$) & \S~\ref{ch4:SecExtSubDB} \\ $L_{\trm{\tiny UL}}$ & linear polarization upper limit & \S~\ref{ch4:SecClassSubPACI} \\ $r,e$ & resolution or Eddington bias corrections & \S~\ref{ch4:SecCNCSubRB} \\ $\Theta_\trm{\scriptsize max}$ & maximum intrinsic angular size for detectable component & \S~\ref{ch4:SecCNCSubRB1} \\ $\widetilde{\sigma}$ & local rms noise divided by local bandwidth smearing ratio & \S~\ref{ch4:SecCNCSubRB1} \\ $f_{\widetilde{\sigma}}$ & probability distribution for $\widetilde{\sigma}$ & \S~\ref{ch4:SecCNCSubRB1} \\ $h$ & integral angular size distribution & \S~\ref{ch4:SecCNCSubRB1} \\ $\eta$ & angular filling factor for linearly polarized emission ($\Theta_\tnm{\scriptsize L}/\Theta_\tnm{\scriptsize I}$) & \S~\ref{ch4:SecCNCSubRB1} \\ $dN_{\tnm{\scriptsize detectable}}/dS$ & differential number-counts that are observable & \S~\ref{ch4:SecCNCSubRB1} \\ $\Theta_\tnm{\tiny med}$ & median largest angular size & \S~\ref{ch4:SecCNCSubRB1} \\ $\Theta_\trm{\scriptsize min}$ & minimum intrinsic angular size for detected component to be classified as resolved & \S~\ref{ch4:SecCNCSubRB2} \\ $dN_{\tnm{\scriptsize resolved}}/dS$ & resolved detectable number-counts & \S~\ref{ch4:SecCNCSubRB2} \\ $dN_{\tnm{\scriptsize unresolved}}/dS$ & unresolved detectable number-counts, assuming ideal case without measurement bias & \S~\ref{ch4:SecCNCSubRB2} \\ $dN_{\tnm{\scriptsize unresolved-obs}}/dS$ & unresolved detectable number-counts, accounting for measurement bias & \S~\ref{ch4:SecCNCSubRB2} \\ \hline \end{tabular} \end{table*} | \label{ch4:SecConcl} We have presented data reduction and analysis procedures for the second data release of the Australia Telescope Large Area Survey. We produced and analysed sensitive 1.4~GHz images of the CDF-S and ELAIS-S1 regions across a combined area of 6.392~deg$^2$ in total intensity ($I$), linear polarization ($L$), and circular polarization ($V$). The data for $L$ were processed using RM synthesis and RM clean. Typical sensitivities across each of the mosaicked multi-pointing images are $\sim30$~$\mu$Jy~beam$^{-1}$, falling to $<25$~$\mu$Jy~beam$^{-1}$ within smaller areas. The typical spatial resolutions are $12\arcsec\times6\arcsec$. We performed component detection and extraction independently in $I$, $L$, and $V$ using a combination of {\tt BLOBCAT} and {\tt IMFIT}, accounting for spatial variations in image sensitivity, bandwidth smearing and instrumental polarization leakage. Corrections for clean bias were not required, due to our implemented cleaning strategy. ATLAS DR2 is the first survey to have been analysed using {\tt BLOBCAT}. We catalogued a total of 2416, 172, and 0 components in $I$, $L$, and $V$, respectively, and determined flux densities for each of these components by considering their angular sizes. We catalogued 2221 sources by matching single or multiple $I$ components with SWIRE mid-infrared sources, and by matching $L$ components to their $I$ counterparts. We classified these sources as AGNs, SFGs, or stars according to four diagnostic criteria. Our source catalogue is slightly biased against the detection of multi-component sources due to our nearest-neighbour cross-identification method, and toward the classification of AGNs due to lack of optical spectroscopy for the majority of sources. We presented a comprehensive prescription for handling multi-pointing data consistently in both total intensity and linear polarization. We described our data reduction and analysis procedures in detail in order to inform future surveys and to highlight our novel extensions of processing techniques from total intensity to linear polarization. We developed new analytic techniques to account for bandwidth smearing with a non-circular beam, and resolution bias in differential number-counts. We extended the analytic framework for Eddington bias corrections from total intensity to linear polarization. In Paper II we present the ATLAS DR2 cross-identification and number-count results, and discuss statistics of the faint polarized 1.4~GHz sky. | 14 | 3 | 1403.5307 |
1403 | 1403.0714_arXiv.txt | One main goal of the New Vacuum Solar Telescope (NVST) which is located at the \emph{Fuxian Solar Observatory} is to image the Sun at high resolution. Based on the high spatial and temporal resolution NVST H$\alpha$ data and combined with the simultaneous observations from the \emph{Solar Dynamics Observatory} for the first time, we investigate a flux rope tracked by a filament activation. The filament material is initially located at one end of the flux rope and fills in a section of the rope, and then the filament is activated due to magnetic field cancellation. The activated filament rises and flows along helical threads, tracking out the twisted flux rope structure. The length of the flux rope is about 75 Mm, the average width of its individual threads is 1.11 Mm, and the estimated twist is 1$\pi$. The flux rope appears as a dark structure in H$\alpha$ images, a partial dark and partial bright structure in 304 {\AA}, while as bright structures in 171 {\AA} and 131 {\AA} images. During this process, the overlying coronal loops are quite steady since the filament is confined within the flux rope and does not erupt successfully. It seems that, for the event in this study, the filament is located and confined within the flux rope threads, instead of being suspended in the dips of twisted magnetic flux. | Coronal mass ejections (CMEs) are large-scale eruptive phenomena of the Sun and release a great deal of plasma and magnetic flux into the interplanetary space, consequently disturbing the space environment around the Earth (Gosling 1993; Webb et al. 1994). As identified in the white light observations, the structure of a CME consists of three parts, i.e., a bright leading front, a dark cavity, and a bright core (Illing \& Hundhausen 1983; Chen 2011). The dark cavity is generally deemed to be a twisted magnetic flux rope (Gibson et al. 2006; Riley et al. 2008). The inner bright core is widely believed to be filament matter suspended in flux rope dips (Guo et al. 2010; Jing et al. 2010). Filament structures are quite conspicuous in H$\alpha$ observations (Hirayama 1985; Martin 1998; Lin et al. 2005) and their dynamic interactions can be caused by magnetic reconnection between the filament-carrying magnetic fields (T{\"o}r{\"o}k et al. 2011; Jiang et al. 2013). Detailed analyses reveal that magnetic flux ropes play a critical role in the formation and acceleration of CMEs (Patsourakos \& Vourlidas 2012; Cheng et al. 2013). Magnetic flux ropes can emerge directly from below the photosphere into the upper atmosphere. Using continuous vector magnetograms from the \emph{Hinode} satellite, Okamoto et al. (2008) found that two abutting regions with opposite polarities connected by strong horizontal magnetic fields first grew and then narrowed, and the orientations of the horizontal fields along the polarity inversion line changed from the normal polarity configuration to the inverse polarity one gradually. Moreover, there were significant blueshifts at the strong horizontal magnetic field area. They suggested that they observed a magnetic flux rope that was emerging from the sub-photosphere. Helical flux ropes can also be formed through magnetic reconnection between two bundles of J-shaped loops which have been frequently observed as sigmoidal structures in the extreme ultraviolet (EUV) and X-ray lines (e.g., Canfield et al. 1999; McKenzie \& Canfield 2008; Liu et al. 2010; Green et al. 2011). In simulations, magnetic reconnection between sheared loops are performed due to the imposed boundary movements, and thus can form magnetic flux ropes (Amari et al. 2000, 2003, 2011; Fan \& Gibson 2003, 2004; Aulanier et al. 2010). Moreover, using nonlinear force-free field models, magnetic flux ropes can be reconstructed from vector magnetic field observations (Canou et al. 2009; Canou \& Amari 2010; Guo et al. 2010, 2013; Jing et al. 2010; Su et al. 2011; Jiang et al. 2013; Inoue et al. 2013). After the launch of the \emph{Solar Dynamics Observatory} ({\it SDO}; Pesnell et al. 2012), with the help of high-quality multi-wavelength data of the Atmospheric Imaging Assembly (AIA; Lemen et al. 2012), many authors have reported the existence of flux ropes in the observations (Cheng et al. 2012; Zhang et al. 2012a; Li \& Zhang 2013a, 2013b, 2013c; Patsourakos et al. 2013). According to some previous observational studies, flux ropes are hot channels in the inner corona before and during solar eruptions (e.g., Zhang et al. 2012a; Cheng et al. 2012). They can be observed in high temperature lines (e.g., 131 {\AA}), while invisible in low temperature lines (e.g., 171 {\AA}). Using differential emission measure (DEM) analysis, Cheng et al. (2012) found that the temperature of twisted and writhed flux rope is higher than 8 MK. However, in the studies of Li \& Zhang (2013a, 2013b) and Patsourakos et al. (2013), the flux ropes can be observed in all the seven EUV lines formed from 0.05 MK to 11 MK. Especially, for the two events investigated by Li \& Zhang (2013b), the flux ropes were tracked by erupting material, leading to the visibility of them, while they could not be detected in all wavelengths at the pre-eruption stage. The New Vacuum Solar Telescope (NVST; Liu \& Xu 2011) is the most important facility of the \emph{Fuxian Solar Observatory} in China. The diameter of NVST is 1 m and the pure aperture is 980 millimeter. One main goal of NVST is to image the Sun at high resolution. As one of the three channels (H$\alpha$, TiO, and G band) being used now, H$\alpha$ is used to observe magnetic structures in the chromosphere. In this Letter, we investigate in detail a flux rope tracked by a filament activation using NVST H$\alpha$ observations for the first time. Combined with the Helioseismic and Magnetic Imager (HMI; Scherrer et al. 2012) and AIA data, we also study the activation and the movement of the filament, and present the corresponding appearance of higher layers revealed in different EUV passbands. | Based on the NVST H$\alpha$ observations for the first time, we study in detail a flux rope which was tracked by a filament activation. The filament material was initially located at one end of the flux rope, and the filament was activated due to the magnetic field cancellation. The activated filament then rose and flowed along the flux rope, tracking out the twisted structure. The length of the flux rope is about 75 Mm, the average width of its individual threads is 1.11 Mm, and the estimated twist is 1$\pi$. The tracked flux rope appeared as a dark structure in H$\alpha$ images, a partial dark and partial bright structure in 304 {\AA}, while as bright structures in 171 {\AA} and 131 {\AA} images. During this process, the overlying coronal loops were quite steady since the filament was confined within the flux rope and did not erupt successfully. The bright core of a CME is thought to be cool filament matter which is suspended in the dips of magnetic flux (e.g., Xia et al. 2012; Zhang et al. 2012b). As revealed by the event in the present study, the filament is closely related with the flux rope. It seems that, the steady filament is only located at one section of the flux rope. However, instead of being suspended in the dips of twisted magnetic flux, the filament material is confined within the helical threads. Only when the filament is activated, the filament material then can flow easily along the threads and thus track out the twisted structure of the flux rope. We tend to support a picture of a pre-existing flux rope that was tracked out by filament material activated by magnetic flux cancellation. However, since the observation duration of the present event is very short and the field-of-view is not large enough, we cannot exclude the formation of new flux rope due to flux cancellation, which is a popular mechanism for flux rope formation (e.g., van Ballegooijen \& Martens 1989). As studied by Zhang et al. (2012a) and Cheng et al. (2012), flux ropes could be clearly observed before and during eruptions and were only detected in hot temperature passbands. While in this study, the flux rope was only observed during the filament activation instead of before the activation, which is similar to the study of Li \& Zhang (2013b). The flux rope was tracked out by the filament material and detected in low temperature (e.g., 304 {\AA}) and high temperature (e.g., 131 {\AA}) lines, consistent with the results of Li \& Zhang (2013a, 2013b). Moreover, our results show that there exists a striking anti-correlation between H$\alpha$ and EUV lines (see Figure 5(b)). This could imply some mild heating of cool filament material into coronal temperatures during the filament activation (e.g., Landi et al. 2010). The heating should not be flare-like, reaching temperatures of 10 MK or more. This can be demonstrated by the almost identical cuts in the AIA 171 {\AA} and 131 {\AA} channels (the 131 {\AA} channel bandpass besides containing a flare peak at around 10 MK, also contains a ``warm" peak at several 0.1 MK, similar to the main peak of the 171 {\AA} channel). In some former studies (Li \& Zhang 2013a, 2013b) and also this study, flux ropes were observed in different EUV lines and appeared as bright structures (see Figures 4(c2), and 4(d2)) or partial bright structures (see Figure 4(b2)), indicating the co-existence of hot and cool components in flux ropes. In contrary, the twisted flux rope in H$\alpha$ images consists of dark threads, which is different from the appearance in EUV images. For the flux ropes studied by Li \& Zhang (2013b), the approximate length is 570 Mm. They measured the width of individual thread of flux ropes and found that the average width is 1.16 Mm. However, the length of flux rope in the present study is only 75 Mm, much smaller than those studied by Li \& Zhang (2013b). In the H$\alpha$ images, the fine-scale threads can be generally resolved, and their mean width is determined to be 1.11 Mm, consistent with the result of Li \& Zhang (2013b). | 14 | 3 | 1403.0714 |
1403 | 1403.0929_arXiv.txt | Observations of the middle-aged supernova remnants IC 443, W28 and W51C indicate that the brightnesses at GeV and TeV energies are correlated with each other and with regions of molecular clump interaction, but not with the radio synchrotron brightness. We suggest that the radio emission is primarily associated with a radiative shell in the interclump medium of a molecular cloud, while the $\gamma$-ray emission is primarily associated with the interaction of the radiative shell with molecular clumps. The shell interaction produces a high pressure region, so that the $\gamma$-ray luminosity can be approximately reproduced even if shock acceleration of particles is not efficient, provided that energetic particles are trapped in the cooling region. In this model, the spectral shape $\ga 2$ GeV is determined by the spectrum of cosmic ray protons. Models in which diffusive shock acceleration determines the spectrum tend to underproduce TeV emission because of the limiting particle energy that is attained. | A highlight of high energy $\gamma$-ray astronomy, involving space-based observations at GeV energies and ground-based observations at TeV energies, has been the detection of middle-aged supernova remnants (SNRs) interacting with molecular clouds (MCs) \citep[see][for reviews]{uchiyama11,fernandez13}. Following pioneering observations with {\it EGRET} \citep{esposito96}, the {\it Fermi} and {\it AGILE} observatories observed $\rm GeV$ $\gamma$-ray emission from several middle-aged SNRs which are interacting with MCs, including W51C \citep{Abdo09}, W44 \citep{Abdo10a}, IC 443 \citep{tavani10,Abdo10b}, and W28 \citep{Abdo10c}. Of these, TeV emission is also detected in W51C \citep{Aleksic12}, IC 443 \citep{Albert07} and W28 \citep{Aharonian08}. In the cases of W28 \citep{Aharonian08} and W44 \citep{uchiyama12}, there is $\gamma$-ray emission external to the remnants that may be associated with the remnants. The high energy emission from these sources has generally been interpreted in terms of pion-decays from cosmic ray (CR) interactions. Two scenarios have been proposed to explain the properties of these middle-aged SNRs associated with MC interaction. In one, relativistic particles escape from a SNR and interact with a nearby MC; TeV emission is produced by interaction between escaping CR particles and the MC, while GeV emission is produced by interaction between galactic CR background particles and the MC \citep{Gabici09,torres10}. In view of the two components, this scenario naturally produces double peaked $\gamma$-ray spectra. The other scenario, discussed here, involves radiative shock waves \citep{Chevalier77,B&C82,Chevalier99,bykov00,Uchiyama10}. The compressed region downstream from a radiative shock is promising because of the high particle number density and energy density. \cite{Uchiyama10} presented a crushed cloud model for W44, IC 443, and W51C, based on remnant parameters from \cite{reach05} for W44. In this view, the SNR has a $500\kms$ nonradiative shock in most of the volume and drives $100\kms$ radiative shock waves into clouds with a density of $200$ cm$^{-3}$. Ambient CRs experience diffusive shock acceleration (DSA) as well as compression in the shell. The cooling region downstream from the shock front is presumed to be the site of $\gamma$-ray emission and radio synchrotron emission. Here, we examine the $\gamma$-ray emission properties of IC 443, W28, and W51C in order to gain insight into the emission processes (Section~2). In Section~3, we discuss the structure of the magnetically supported radiative shell and the interaction region between shell and the MC clump \citep{Chevalier99}. In Section~4, we model the evolution of relativistic protons in both regions and calculate their $\pi^0$-decay emission. We compare our results with the observations of all three remnants and discuss the results in Section 5. | We calculated the emission from IC 443 for both standard DSA and pure adiabatic cases, for our radiative shell plus MC clump interaction model and the parameters shown in Table \ref{IC443MC}. Ion neutral damping and a finite acceleration time were taken into account in the DSA simulation but not the limited cooling region argument due to its uncertainty. The ionization in the layer 1 precursor is low so we only considered ion neutral damping for the shell, finding that $p_{br}\sim \rm 10$ GeV/c gives a good fit to the GeV part of the spectrum with $w=8\%$ (Figure \ref{IC443GR}). The resulting $\gamma$-ray spectrum is narrower than the observed spectrum, falling below the observed emission at high energy due to the limited acceleration time. The observed spectra from GeV to TeV energies have slopes which are similar to the input CR spectrum. While pure adiabatic compression maintains the shape of the input CR spectrum, pion-decay emission also traces the energy distribution of the parent spectrum above a few GeV. We found that the pure adiabatic case can approximately fit the spectra of IC 443 from GeV to TeV energies, but with a higher $w\approx 21\%$ (Figure \ref{IC443GR_ad}), which implies that the remnant is in a special phase of evolution. Alternatively, a higher value for the ICM density would reduce the value of $w$. The $\gamma$-ray spectrum of W28 has a similar shape to that of IC 443, except for the low energy part \citep[Figure \ref{IC443GR};][]{Abdo10c,Aharonian08}. W51C's $\gamma$-ray spectrum also has a shape similar to IC 443 but the luminosity is larger \citep[Figure \ref{IC443GR};][]{Abdo09,Aleksic12}. W51C is a large remnant compared to the other two, and may have an unusually large energy \citep{koo97,koo05}, which could account for the high luminosity. We do not attempt detailed models for W28 and W51C, but note the similar spectral shapes for the three remnants may be a result of the similar parent CR spectrum. Other possible tests of the models are the shocked MC mass and the relative intensities of the three components: radiative shell, layer 1 and layer 2. In the pure adiabatic case the shocked MC mass required for IC 443 is $m_{MC}\approx V_2t_{MC}\rho_c 4\pi R^2_s w\approx 500 \Msun$ which is more consistent with the molecular observations \citep{dickman92,Lee08} than the $190 \Msun$ in our DSA model. In the pure adiabatic case, the $\gamma$-ray emission from IC 443 is naturally dominated by the MC interaction region, especially layer 1, while for the DSA case the emission from the MC interaction region is comparable to the shell component at low energy but becomes dominant at high energy. However, the surface brightness is coupled with the angle between the collision direction and the line of sight. Considering the uncertainty in both theoretical models and observations, these tests are not definitive. Here we have sought a model for the $\gamma$-ray emission from SNRs that is consistent with the emission properties given in Section 2. The correlation of GeV and TeV emission with molecular clump interaction implies that the $\gamma$-ray emission is related to slow radiative shock waves in dense matter. Standard DSA is not efficient at high energies. Pure adiabatic compression could reproduce the ratio of TeV to GeV emission but requires a large covering factor (Figure 3). Particle acceleration process in the middle aged SNRs are still not very clear. There are other possibilities; \cite{bykov08} have found nonthermal X-ray emission near a clump interaction region in IC 443 which they interpret as the result of ejecta knots hitting the molecular gas. More detailed observations of the remnants discussed here, as well as other remnants with molecular cloud interaction, would improve our understanding of CR acceleration in SNRs. | 14 | 3 | 1403.0929 |
1403 | 1403.0574_arXiv.txt | {It is clear from the observation of charged cosmic rays up to energies of $10^{20}$~eV that particle acceleration must occur in astrophysical sources. Acceleration of secondary particles like muons and pions, produced in cosmic ray interactions, are usually neglected, however, when calculating the flux of neutrinos from cosmic ray interactions.} { Here, we discuss the acceleration of secondary muons, pions, and kaons in gamma-ray bursts within the internal shock scenario, and their impact on the neutrino fluxes.} {We introduce a two-zone model consisting of an acceleration zone (the shocks) and a radiation zone (the plasma downstream the shocks). The acceleration in the shocks, which is an unavoidable consequence of the efficient proton acceleration, requires efficient transport from the radiation back to the acceleration zone. On the other hand, stochastic acceleration in the radiation zone can enhance the secondary spectra of muons and kaons significantly if there is a sufficiently large turbulent region.} {Overall, it is plausible that neutrino spectra can be enhanced by up to a factor of two at the peak by stochastic acceleration, that an additional spectral peaks appears from shock acceleration of the secondary muons and pions, and that the neutrino production from kaon decays is enhanced. } {Depending on the GRB parameters, the general conclusions concerning the limits to the internal shock scenario obtained by recent IceCube and ANTARES analyses may be affected by up to a factor of two by secondary acceleration. Most of the changes occur at energies above $10^7$~GeV, so the effects for next-generation radio-detection experiments will be more pronounced. In the future, however, if GRBs are detected as high-energy neutrino sources, the detection of one or several pronounced peaks around $10^{6}$~GeV or higher energies could help to derive the basic properties of the magnetic field strength in the GRB.} | Gamma-ray bursts (GRBs) are a candidate class for the origin of the ultra-high energy cosmic rays (UHECRs). A popular scenario is the internal shock model, where the prompt $\gamma$-ray emission originates from the radiation of particles accelerated by internal shocks in the ejected material~ \citep{Paczynski:1994uv,Rees:1994nw} (see \cite{Piran:2004ba,Meszaros:2006rc} for reviews). If a significant baryon flux is accelerated, the GRBs may be a plausible source for the UHECRs. In this case, substantial production of secondary pions, muons, and also kaons are expected from photohadronic interactions between the baryons and the radiation field; these will decay into neutrinos and other decay products~\citep{Waxman:1997ti,Asano:2006zzb}. Gamma-rays at TeV energies and above are co-produced in the photohadronic process, but are subject to interactions with the internal photon field from the radiation processes, including synchrotron radiation and inverse Compton scattering, in the GRB. The photons are expected to cascade down via pair production cascades so that they can be detected at $\sim$ GeV energies. Several of such GRBs have been detected be Fermi-LAT \citep{fermi_bursts}, but the associated neutrino production per burst is generally expected to be rather low, see e.g.\ \cite{bhoo2010}. In general, neutrino detection from GRBs with IceCube therefore needs to be done via the stacking of a larger number of bursts, see e.g.\ \cite{Abbasi:2011qc,Abbasi:2012zw}. Very stringent neutrino flux limits for the internal shock scenario have been recently obtained by the IceCube collaboration using the stacking approach~\citep{Abbasi:2011qc,Abbasi:2012zw}. Using timing, energy, and directional information for the individual bursts, new limits have been obtained, which are basically background-free and which are significantly below earlier predictions based on gamma-ray observations~\citep{Waxman:1997ti,Guetta:2003wi,Becker:2005ej,Abbasi:2009ig}. These predictions have been recently revised from the theoretical perspective~\citep{Hummer:2011ms,Li:2011ah,He:2012tq}, yielding about a factor of ten lower expected flux~\citep{Hummer:2011ms} depending on the analytical method compared to; see also \cite{Adrian-Martinez:2013sga} for an analysis by the ANTARES collaboration using this method. This discrepancy comes mainly from the energy dependence of the mean free path of the protons, the integration over the full photon target spectrum (instead of using the break energy for the pion production efficiency), and several other corrections adding up in the same direction; see Fig.~1 (left) in \cite{Hummer:2011ms}. It should be noted at this point, that the absolute normalization of the neutrino flux scales linearly with the ratio of the luminosity in protons to electrons (baryonic loading). In typical models \citep{Waxman:1997ti,Guetta:2003wi,Becker:2005ej,Abbasi:2009ig}, this ratio is usually assumed to be $10$, while theoretical considerations % suggest a value of $100$ \citep{schlickeiser2002} if GRBs are to be the sources of the UHECRs. In a recent study~\citep{Baerwald:2014zga}, this value is self-consistently derived from the combined UHECR source and propagation model, including the fit of the UHECR data. For an injection index of two, it is demonstrated that this value depends on the burst parameters, and that values between $10$ and $100$ are plausible.\footnote{A baryonic loading of $10$ requires, however, ``typical'' source parameters $\Gamma \sim 400$ and $L_{\gamma,\mathrm{iso}} \sim 10^{53} \, \mathrm{erg \, s^{-1}}$, whereas $\Gamma \sim 300$ and $L_{\gamma,\mathrm{iso}} \sim 10^{52} \, \mathrm{erg \, s^{-1}}$ more point towards a baryonic loading of order $100$; see Fig.~7 in \cite{Baerwald:2014zga}.} Note that these numbers depend strongly on the proton and electron/photon input spectral shapes and energy ranges, and according to basic theory of stochastic acceleration, it can easily vary between $1000$ and $0.1$~\citep{lukas_paper}. As further demonstrated in \cite{Baerwald:2014zga}, the improved modeling of the GRB spectra can be used to constrain the central parameters of the calculation, \ie, the ratio of protons to electrons and the boost factor. The effects discussed in our study could then contribute to determining another basic property of the GRB, namely the magnetic field strength. Another argument can be used when relating the neutrino and UHECR fluxes directly if the cosmic rays escape as neutrons produced in the same interactions as the neutrinos~\citep{Ahlers:2011jj}. Since this possibility is strongly disfavored~\citep{Abbasi:2012zw} it is conceivable that other escape mechanisms dominate for UHECR escape from GRBs~\citep{Baerwald:2013pu}. For instance, if the Larmor radius can reach the shell width at the highest energies, it is plausible that a fraction of the cosmic rays can directly escape. Other possible mechanisms include diffusion out of the shells. In \cite{Baerwald:2014zga}, it was demonstrated that even current IceCube data already imply that these alternative escape mechanisms must dominate if GRBs ought to be the sources of the UHECR, and that future IceCube data will exert pressure on these alternative options as well. The secondary pions, muons, and kaons produced by photohadronic interactions will typically either decay (at low energies) or lose energy by synchrotron radiation (at high energies). At the point where decay and synchrotron timescales are equal, a spectral break in the secondary, and therefore also in the neutrino spectrum is expected. This is the so-called ``cooling break'', see e.g.\ \cite{Waxman:1997ti}. Additional processes, which potentially affect the secondary spectra, are: adiabatic losses, interactions with the radiation field~\citep{Kachelriess:2007tr}, and acceleration of the secondaries in the shocks or by stochastic acceleration~\citep{Koers:2007je,Murase:2011cx,Klein:2012ug}. In this study, we focus on the quantitative impact of the secondary acceleration on the neutrino fluxes, and the conditions for a significant contribution of this effect. A substantial enhancement of the secondary spectra would increase the tension between the recent IceCube observations and the predictions, and would therefore be critical for the interpretation of the recent IceCube results. | The aim of this study has been to address the quantitative importance of the acceleration of secondary muons, pions, and kaons for the neutrino fluxes. We have therefore extended the model by \cite{Hummer:2011ms}, which predicts the neutrino fluxes from gamma-ray observations in the internal shock model and which the current state-of-the-art GRB stacking analyses in neutrino telescopes are based on, by the acceleration effects of the secondaries -- as discussed in a more general sense in \cite{Klein:2012ug}. One of the key issues has been a separate description of the acceleration zone (the shocks) and the radiation zone (the plasma downstream the shocks) in a two-zone model, since it is plausible that the shock acceleration and the photohadronic processes, leading to the secondary production, happen dominantly in different regions. Two classes of acceleration have been implemented for the secondaries: shock acceleration in the acceleration zone and stochastic acceleration in the (possibly turbulent) plasma in the radiation zone. An important component of the model has been the transport of the secondaries from the radiation zone back to the acceleration zone, which we describe by Kolmogorov diffusion (optimistic) or Bohm diffusion (conservative) -- assuming that at the highest energies, where the Larmor radius reaches the size of the region, all secondaries are efficiently transported. The shock acceleration of the secondaries is then just a consequence of the efficient proton acceleration if they can be transported back to the shocks, whereas the stochastic acceleration depends on the size of the turbulent region. In both cases, some uncertainty arises from the acceleration efficiencies, which may vary within reasonable limits. We have shown that both the muon and kaon spectra can be significantly modified by shock acceleration: the muon spectrum, because muons have a long lifetime over which they can be accelerated, and the kaon spectrum, because kaons are most efficiently transported back to the acceleration zone at their highest energies (they have the highest synchrotron cooling break). The shock acceleration leads to additional peaks determined by the critical energy, where acceleration and energy loss or escape rates are equal. These peaks translate into corresponding peaks of the neutrino spectra, smeared out by the kinematics of the weak decays. The most significant enhancement at the peak is expected from the muon spectrum if the magnetic field is low and the Lorentz boost is high, since then the critical energy may coincide with the peak energy. Too low magnetic fields, on the other hand, mean that the protons cannot be efficiently accelerated. We note that our model is fully self-consistent in the sense that it is taken into account that muons are produced by pion decays, which may be accelerated themselves. The amount of shock acceleration depends critically on the transport between radiation and acceleration zone. For Bohm diffusion or even slower transport processes, hardly any modification of the neutrino spectra is observed, since the enhancement of the muon spectrum is completely shadowed by the regular pion spectrum present at higher energies. On the other hand, the results for Kolmogorov diffusion are already close to the perfect transport case (all particles efficiently transported over the dynamical timescale). The stochastic acceleration can be very efficient for muons and kaons, since their cooling breaks occur at a smaller (decay and synchrotron loss) rates than the one for pions, which means that the stochastic acceleration can be dominant at these breaks. The consequence is an enhancement at the cooling break (if the break comes from synchrotron losses) or beyond (if it comes from adiabatic losses). In the latter case, the shock and stochastic acceleration effects can add up and lead to an additional peak in the neutrino spectrum with a significant enhancement. It is conceivable that efficient stochastic acceleration means inefficient transport, \ie, the two acceleration effects are mutually exclusive in terms of their energy ranges, and that such an effect can be only observed for a flat enough energy dependence of the diffusion coefficient (such as Kolmogorov diffusion). Depending on the specific GRB parameters, secondary particle acceleration can enhance the neutrino flux by up to an overall factor of two. The enhancement is typically largest at higher energies, around $10^{8}$ GeV or above. This enhancement is relevant for extremely high-energy searches with IceCube at energies around $10^{8}$~GeV and above. In particular, southern hemisphere searches are sensitive at these energies, as the background of atmospheric muons is sufficiently small at those high energies, see \cite{Aartsen:2013ApJ} for the latest point source sensitivity of IceCube. Even northern hemisphere searches might already be sensitive to the enhancement. Next generation instruments like KM3NeT and a high-energy extension of IceCube will be able to constrain the parameter space for secondary particle acceleration effects even further. Other future experiments which have good sensitivity above $10^{8}$~GeV concern the radio emission from neutrino-induced showers, such as ARA \citep{Allison:2011wk} and ARIANNA \citep{Gerhardt:2010js,Klein:2012bu}. These conclusions will somewhat depend on the choices of the acceleration rates, which means that we cannot exclude larger effects for individual bursts. Note that some of our choices (such as for transport and size of the turbulent region) are already on the optimistic side. | 14 | 3 | 1403.0574 |
1403 | 1403.1744_arXiv.txt | {We explore the scientific potential of next-generation high-angular resolution optical imager to study the AGN/Host connection. The availability of a significant number of X-raying AGN with natural guide stars, allowing for adaptive optics at optical wavelengths, offers an interesting perspective to complement high-resolution work currently done in the near-infrared. } \FullConference{ Nuclei of Seyfert galaxies and QSOs - Central engine \& conditions of star formation \\ November 6-8, 2012\\ Max-Planck-Insitut f\"ur Radioastronomie (MPIfR), Bonn, Germany} \begin{document} | 14 | 3 | 1403.1744 |
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1403 | 1403.3094_arXiv.txt | Nuclear star clusters (NSCs) are dense stellar clusters observed in galactic nuclei, typically hosting a central massive black hole. Here we study the possible formation and evolution of NSCs through the inspiral of multiple star clusters hosting intermediate mass black holes (IMBHs). Using an N-body code we examine the dynamics of the IMBHs and their effects on the NSC. We find that IMBHs inspiral to the core of the newly formed NSC and segregate there. Although the IMBHs scatter each other and the stars, none of them is ejected from the NSC. The IMBHs are excited to high eccentricities and their radial density profile develops a steep power-law cusp. The stars also develop a power-law cusp (instead of the central core that forms in their absence), but with a shallower slope. The relaxation rate of the NSC is accelerated due to the presence of IMBHs, which act as massive-perturbers. This in turn fills the loss-cone and boosts the tidal disruption rate of stars both by the MBH and the IMBHs to a value excluded by rate estimates based on current observations. { Rate estimates of tidal disruptions can therefore provide a cumulative constraint on the existence of IMBHs in NSCs.} | Nuclear stellar clusters (NSCs) are dense and compact stellar systems observed in galactic nuclei, many of which hosting a massive black hole (MBH) at their center \citep{Bo10}. Less massive galaxies seem to contain an NSC but not always an MBH, while the most massive ones show the presence of a central MBH \citep[see e.g.][]{Gr08, GS09, Ji11, SG13}. Our own Galactic center (GC) hosts both an NSC and a central MBH, Sgr A$^{*}$ of mass ${\rm M_{BH}\sim4\times10^{6}}$ ${\rm M_{\odot}}$. Two main hypotheses have been suggested to explain the origin of NSCs: \textit{(i)} the in-situ formation, where molecular gas coming from more external regions is channeled to the GC leading to an episodic star formation epoch \citep{Mi04}; and \textit{(ii)} the cluster-infall/merger scenario, where dense massive stellar clusters decay toward the center of their host galaxy due to dynamical friction, merge and form the NSC. These two scenarios are not mutually exclusive, and both can contribute to the formation of NSCs. Here we focus only the cluster-infall scenario. The merger scenario was proposed by \citet{Tr75}, and later explored in many works in the case of a generic galaxy without a central MBH \citep[see e.g.][]{Ca93,Ag11,Ca08a,Ca08b}. Other studies explored the case of galaxies hosting an MBH \citep{An13}, and, in particular, the origin of the NSC in the GC of the Milky Way \citep{An12,Gn13,PM14}. Relaxed stellar systems around an MBH give rise to a cusp structure (a power-law distribution with a slope of $-7/4$, for a single mass population and somewhat shallower slope for most stars in a multi-mass population, \citealt{Ba76}). However, the relaxation time of the system may be longer than the age of the stars in the NSC, in which case such a cusp may not form. The two-body relaxation time of the observed Galactic NSC, evaluated at the influence radius of the MBH ($2$-$3$~pc) is estimated to be $\sim20$-$30$~Gyr \citep{Me10}. Current observations of the GC suggest the existence of a flat core in the distribution of the old stellar population of the NSC. The observed features of the Milky Way NSC have been well reproduced by \citet{An12} in their N-body simulations where they studied 12 consecutive inspirals of massive, compact clusters, which build-up an NSC with a central core of $\sim0.5$~pc and and external $\sim r^{-1.8}$ power law outside. Recent studies suggest that this scenario may also lead to age/color and mass segregation in the NSC that can potentially be observed \citep{PM14}. In analogy with galaxies hosting a central MBH, \citet{Si75} proposed that massive clusters may host an intermediate mass black hole (IMBH) at their center. The formation of IMBHs in dense clusters has been explored in several studies \citep{Lo94,Ma01, Eb01,Br03}. Numerical simulations of the formation and evolution of young dense clusters show that runaway collisions between stars can produce a massive star ($\sim10^{2}-10^{4}$~M$_{\odot}$) that sinks to the center of the cluster due to dynamical friction \citep{GFR04,Po04, FGR06,Po06,Fu09}, and later collapses to form an IMBH with a mass of $\sim10^{2}-10^{4}$~M$_{\odot}$. Given the possible existence of such IMBHs in massive clusters, the cluster-infall scenario may be significantly affected by such objects. Here we extend the cluster-infall scenario to include this possibility, and study the effects of IMBHs on the NSC formation and evolution. We use N-body simulations similar to those described in \citet{An12}, but populate the infalling clusters (ICs) with IMBHs. The outline of the paper is as follows. We begin with a description of the initial conditions and our methods in \S \ref{sec:ICs}. We then present our results (\S \ref{sec:res}), discuss their implications (\S \ref{sec:disc}) and summarize (\S \ref{sec:sum}). | \label{sec:sum} In this work we explored the formation and evolution of NSCs through the merging of ICs containing IMBHs. The presence of IMBHs substantially modifies the structure of the NSC. Our main results are as follows. \begin{itemize} \item Following the infall and disruption of an inspiraling cluster, its IMBH continues to spiral in, accompanied by a small group of stars or compact objects, which are later stripped from the IMBH close to the central MBH. The IMBHs later scatter each other and the NSC stars, and can develop high eccentricities. \item In contrast to the IMBH-free case, the NSC develops a cusp already in its early evolutionary stages. \item The system is strongly mass segregated: IMBHs and stars evolve separately and settle on two different quasi-stationary profiles characterized by a central cusp with different slopes. \item The NSC is tangentially anisotropic and oblate, very similar to simulated NSCs without IMBHs. \item The IMBHs decrease the relaxation time of the NSC by a factor of a few hundred, leading to the fast refilling of the loss cone for tidal disruption of stars by the central MBH. The system is in the full loss cone regime and the tidal disruption rates are two orders of magnitude larger than in an NSC without IMBHs. In addition the rate of TDEs by the IMBHs themselves is comparable to that of the central MBH. \item { The large rate of TDEs expected from IMBHs-hosting NSCs compared to the much lower TDE rate estimates from observations, suggests that typical NSCs likely do not contain IMBHs. If cluster-infall is one of the main channels for NSC formation, this also implies that the majority of massive clusters do not host IMBHs. We do caution, however, that the accuracy and reliability of current TDE rate estimates is still limited; and better estimates would be available from future surveys. } \end{itemize} | 14 | 3 | 1403.3094 |
1403 | 1403.1328_arXiv.txt | The influence of initially given small scale magnetic energy($E_M(0)$) and helicity($H_M(0)$) on the magnetohydrodynamics(MHD) dynamo was investigated. Equations for $E_M$(t), $H_M$(t), and electromotive force($\langle {\bf v}\times {\bf b}\rangle$, $EMF$) were derived and solved. The solutions indicate small scale magnetic field(${\bf b}_i$) caused by $E_M$(0) modifies $EMF$ and generates additional terms of which effect depends on magnetic diffusivity $\eta$, position of initial conditions($IC$s) $k_f$, and time ($\sim e^{-\eta k_f^2 t}$). ${\bf b}_i$ increases the inverse cascade of energy resulting in the enhanced growth of large scale magnetic field($\overline{{\bf B}}$). Simulation data show that $E_M$(0) in small scale boosts the growth rate, which also proportionally depends on $H_M(0)$. If $E_M$(0) is the same, positive $H_M(0)$ is more effective for MHD dynamo than negative $H_M(0)$ is. It was discussed why large scale magnetic helicity should have the opposite sign of the injected kinetic helicity. | 14 | 3 | 1403.1328 |
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1403 | 1403.6813_arXiv.txt | {We study dualities of the general Galileon theory in $d$ dimensions in terms of coordinate transformations on the coset space corresponding to the spontaneously broken Galileon group. The most general duality transformation is found to be determined uniquely up to four free parameters and under compositions these transformations form a group which can be identified with $GL(2,\mathbf{R})$. This group represents a unified framework for all the up to now known Galileon dualities. We discuss a representation of this group on the Galileon theory space and using concrete examples we illustrate its applicability both on the classical and quantum level.} \begin{document} | The Galileons represent a particular class of models of real scalar field $\phi $ with derivative interactions and posses number of interesting properties. It emerges in its simplest form as an effective theory of the Dvali-Gabadadze-Porrati model \cite{Dvali:2000hr,Deffayet:2001pu} as well as of the de Rham-Gabadadze-Tolley massive gravity theory \cite{deRham:2010kj} in the decoupling limit. Generalization of the Galileon Lagrangian was proposed by Nicolis, Rattazzi and Trincherini \cite{Nicolis:2008in} as the long distance modifications of General relativity. In the seminal paper \cite% {Nicolis:2008in} also the complete classification of possible terms of the Galileon Lagrangian has been made and some of the physical consequences have been studied in detail, i.a. it was demonstrated that such theories exhibit the so-called Vainshtein mechanism \cite{Vainshtein:1972sx}. In fact the general structures appearing in the Galileon Lagrangian have been already discovered in the 70' as a building blocks of the Horndeski Lagrangian \cite% {Horndeski:1974wa}, which is the most general Lagrangian built from no more than the second order derivatives of the scalar field and leading to the second-order Euler-Lagrange equations. Generalization of such Lagrangians to curved backgrounds and arbitrary $p$-form fields has been studied in \cite% {Deffayet:2009mn,Deffayet:2010zh}. From another point of view the Galileon Lagrangian can be obtained as a special non-relativistic limit of the Dirac-Born-Infeld Lagrangian describing the fluctuations of the $d$% -dimensional brane in the $d+1$ dimensional space-time \cite{deRham:2010eu}. For a pedagogical introduction into the Galileon physics as well as for the complete list of literature see e.g. \cite{Khoury:2013tda}. Putting aside very important cosmological aspects, the Galileon theory itself has an amazing structure which has been studied intensively in the literature (for pedagogical introductions into the technical aspects see e.g. \cite{Curtright:2012gx,Deffayet:2013lga}). For instance on the quantum level it exhibits the so-called non-renormalization theorem which prevents the tree-level Galileon couplings from obtaining the quantum corrections stemming from loops \cite{Luty:2003vm, Hinterbichler:2010xn, deRham:2012ew, Brouzakis:2013lla}. Another interesting feature is the existence of dualities, i.e. such transformations of fields and coordinates which preserve the form of the Galileon Lagrangian, though it changes its couplings. The duality transformations therefore interrelate different Galileon theories on the contrary to the symmetry transformations which leave the action invariant. The first such a duality has been recognized already in the paper \cite{Nicolis:2008in} where it was shown that the transformation $\phi \to \phi + \tfrac{1}{4}H^{2}x^{2}$ converts one form of the Lagrangian into another one. The latter then describes the fluctuations of the Galileon field about the de Sitter background solution. Another example of duality was mentioned and studied in \cite{Curtright:2012gx} and it corresponds to the dual Legendre transform of the field. The most interesting duality has been discovered in \cite{Fasiello:2013woa} in the context of massive gravity and bigravity and has been further studied in \cite{deRham:2013hsa}. In this paper we study these dualities from the unified point of view. We make use of the fact that the general Galileon theory can be understood as a low-energy effective theory describing the Goldstone bosons corresponding to the spontaneously broken symmetry according to the pattern $GAL(d,1)\to ISO(d-1,1)$ where $GAL(d,1)$ is the so-called Galileon group and its Lagrangian can be identified with generalized Wess-Zumino-Witten terms \cite% {Goon:2012dy}. This allows us to classify the most general duality transformation and identify it as a non-linear coordinate transformations on the coset space $GAL(d,1)/SO(d-1,1)$. As we will show such duality transformations form a four-parametric group which can be identified with $% GL(2,\mathbf{R})$ and which contains all the above mentioned dualities as special cases. We will also study the representation of this duality group on the Galileon theory space and give examples of physical applications of the duality. Namely we discuss the duality of classical covariant phase spaces and corresponding observables, the duality of fluctuations on the the classical background, the dual realization of the symmetries, the duality of the $S$ matrix and its applications on the tree and one-loop level. We also classify the Galileon theories with respect to the duality generated with specific subgroup of $GL(2,\mathbf{R})$ which leaves the $S$ matrix invariant or under which the tree-level amplitudes trivially scale. We illustrate most of the above topics by means of explicit examples. This paper is organized as follows. First, in Section~\ref{motivations} we introduce the Galileon symmetry and Lagrangian, discuss the Feynman rules and as an illustration we calculate the tree-level amplitudes up to the five-point one. In Section~\ref{sec3coset} we review the coset construction of the Galileon Lagrangian. Section~\ref{sec4galileon} and~\ref{sec5gl2} contain the main results of this work. In Section~\ref{sec4galileon} we construct the most general duality transformations and in Section~\ref% {sec5gl2} we discuss their group structure. Several applications then follow in Section~\ref{sec6applications}. Some technical details and alternative approaches are postponed in appendices. | In this paper we have studied the duality transformations of the general Galileon theories in $d$ dimensions. First we have reviewed the interpretation of the Galileon as a Goldstone boson of the spontaneous symmetry breakdown according to the pattern $GAL(d,1)\rightarrow ISO(d-1,1)$ and the identification of its action as the generalized WZW term. Then we have studied the most general coordinate transformations on the corresponding coset space $GAL(d,1)/SO(d-1,1)$. The requirement that such a general transformation acts linearly on the basic building blocks of the Galileon Lagrangian (and therefore it represents a duality transformation) constraints the form of the transformation uniquely up to four free parameters. Under composition these duality transformations form a group which can be identified with $GL(2,{\mathbf{R}})$. The explicit form of the duality transformation for general $\alpha \equiv \{\alpha _{ij}\}_{i,j=1}^{2}\in GL(2,{\mathbf{R}})$ reads% \begin{eqnarray} x_{\alpha } &=&\alpha _{11}x+\alpha _{12}\partial \phi (x) \notag \\ \phi _{\alpha }(x_{\alpha }) &=&\det \left( \mathbf{\alpha }\right) \phi (x) \notag \\ &&+\frac{1}{2}\left( \alpha _{12}\alpha _{22}\partial \phi (x)\cdot \partial \phi (x)+2\alpha _{12}\alpha _{21}x\cdot \partial \phi (x)+\alpha _{11}\alpha _{21}x^{2}\right) . \notag \end{eqnarray}% All the up to now known Galileon dualities can be identified as special elements (or one-parametric subgroups) of this duality group. We have also studied its action on the space of the Galileon theories and found a basis of the independent invariants of one of its most interesting one-parametric subgroups \begin{equation*} \mathbf{\alpha }_{D}(\theta )=\left( \begin{array}{cc} 1 & -2\theta \\ 0 & 1% \end{array}% \right) \end{equation*}% This subgroup is represented in the space of fields as a field redefinition which can be understood both as a simultaneous space-time coordinates and field transformation or as a non-local change of the fields which includes infinite number of derivative dependent terms. We have then studied the applications of the duality group. We have shown that we can relate the classical covariant phase spaces of dual theories and enlarge the duality transformation to the classical observables. In order to avoid apparent paradoxes, correct dual observables within the dual theory have to be used when we want to get the results of the original theory. The duality of phase spaces can be used to generate classical solution of the interacting Galileon theory from the solution of the more simple one even when the Galileon is coupled to the local external source. We have studied two such sources, namely the point-like and string-like ones. Here the duality appears to be an efficient tool because of the symmetries which effectively reduce the dimensionality of the problems. We have also discussed the fluctuations of the classical solutions in the linearized approximation. We have found duality transformation of the corresponding classical covariant phase spaces and corresponding observables and discussed the geometrical aspects of the problem with superluminal propagation of fluctuations. The general consideration has been illustrated by two explicit examples, namely the fluctuations of the plane wave and cylindrically symmetric classical solutions. We have also established the dual formulation of the additional symmetries of the Lagrangian. We have shown that these symmetries might be hidden within the dual theory and found e.g. that the $Z_{2}$ symmetry and space time translations are realized non-linearly and non-locally. Then we have discussed the transformation properties of the $S$ matrix and established its formal invariance within the dimensional regularization, though only the tree-level (classical) part of the complete action transforms nicely under the duality field redefinition. As a next issue we have demonstrated the usefulness of the $S$ matrix duality for calculations of the tree on-shell scattering amplitudes and for finding the relations between the contributions of the apparently very different Feynman graphs with completely different topologies. As another example we have classified the equivalence classes (with respect to the duality subgroup $\alpha _{D}(\theta )$ combined with scaling) of the Galileon theories (and at the same time of the tree-level $S$ matrices) in three and four dimensions. We found e.g. that there is up to the above dualities only one nontrivial interacting theory in three dimensions which exhibits the $Z_{2}$ symmetry. Then we have discussed the transformation properties of the $S$ matrix on the loop level. As we have discussed on a concrete example of the one-loop four-point on-shell amplitude, due to the counterterms the duality is not completely straightforward. It rather holds on the regularized level for the loop graphs with vertices from the basic tree-level Lagrangian. We have also touched the problem of the counterterms classification based on a generalization of the Weinberg formula and with help of the latter we discussed the non-renormalization theorem. \textbf{Note added}: After this work was completed two works \cite% {Creminelli:2014zxa,deRham:2014lqa} closely connected with the topic studied in this paper appeared. Both these papers concern the properties of the one-parametric duality subgroup denoted as $\alpha_D(\theta)$ in our notation and partially overlap with our results. | 14 | 3 | 1403.6813 |
1403 | 1403.2217_arXiv.txt | We present one of the best sampled early time light curves of a gamma-ray burst (GRB) at radio wavelengths. Using the Arcminute Mircrokelvin Imager (AMI) we observed GRB 130427A at the central frequency of 15.7 GHz between 0.36 and 59.32 days post-burst. These results yield one of the earliest radio detections of a GRB and demonstrate a clear rise in flux less than one day after the $\gamma$-ray trigger followed by a rapid decline. This early time radio emission probably originates in the GRB reverse shock so our AMI light curve reveals the first ever confirmed detection of a reverse shock peak in the radio domain. At later times (about $3.2$ days post-burst) the rate of decline decreases, indicating that the forward shock component has begun to dominate the light-curve. Comparisons of the AMI light curve with modelling conducted by Perley et al. show that the most likely explanation of the early time 15.7 GHz peak is caused by the self-absorption turn-over frequency, rather than the peak frequency, of the reverse shock moving through the observing bands. | The detection of the early time multi-wavelength radiation from gamma-ray bursts (GRBs) within the first day after the initial flash of $\gamma$-rays is essential for refining our understanding of these energetic events. The internal-external shock scenario \citep{piran99} suggests that along with the forward shock, which propagates into the circumburst medium to generate the classical afterglow, there is also emission associated with the reverse shock propagating into the relativistic ejecta \citep{sari99}. Evidence for the presence of a reverse shock has been demonstrated by the detection of optical flashes (within minutes after the $\gamma$-ray trigger) that are not correlated with the initial $\gamma$-ray emission from the GRB \citep{sari99grb}. Such emission can only be explained by the presence of different physical emitting regions. This same model suggests that the detection of radio flares approximately 1 day post-burst also emanate from the reverse shock \citep{kulkarni99}. Such early time radio signals, which imply a rapid rise and fall in emission within 1 day post-burst, are atypical when compared to the classical radio afterglow of long GRBs resulting from the forward shock, which slowly evolve on the time-scales of days to years \citep[for a review see][]{granot14}. The early time radio signature of GRBs has not been as well investigated as it has in the optical band. This is due to the limited number of large radio telescopes, which are required for such follow-up observations due to the faintness of GRB radio emission. This in turn makes it more difficult to acquire target-of-opportunity observations at the time of the event. Early time observations have traditionally required human intervention to activate, potentially resulting in the first radio observation of a given GRB being delayed several hours to even days post burst. As a result the radio emission emanating from the reverse shock of a GRB has only been observed in a few cases where the earliest radio detections have occurred around $1$ day post-burst \citep[e.g.,][]{kulkarni99,frail00,berger03}. Only a few robotized, rapid response, follow-up programmes of GRBs have been implemented in the radio domain. For example attempts were made with the Cambridge Low Frequency Synthesis Telescope at 151 MHz, which triggered on GRBs detected with the Burst And Transient Source Experiment onboard the \textit{Compton Gamma-Ray Observatory} \citep{green95,koranyi95,dessenne96}. More recently \citet{bannister12} conducted a robotized follow-up experiment using a 12~m radio dish at 1.4~GHz that was specifically designed to search for prompt radio emission associated with GRBs. This telescope triggered on those GRBs detected with the \swift\ $\gamma$-ray Burst Mission \citep{gehrels04} and was capable of being on target within a few minutes post-burst. In two out of the nine GRBs observed, a single dispersed radio pulse was possibly detected. In both cases the candidate's pulse was coincident with breaks in the GRB X-ray light curves. Over the past two years a new robotized follow-up programme using the Large Array (LA) interferometer of the Arcminute Microkelvin Imager \citep[AMI;][]{zwart08} has been implemented to obtain immediate observations, and conduct radio monitoring, of \swift\ detected GRBs at 15.7~GHz. This rapid GRB follow-up programme conducted with AMI is fully automated and is activated when \swift\ triggers on an event, with response times as low as 5 minutes \citep{staley13}. This programme is therefore capable of statistically constraining the radio properties of many \swift\ detected GRBs (both long and short) within the first hour post-burst, which has never been done before. One of the most recent radio bright long gamma-ray bursts is GRB 130427A, which was detected on 2013, April 27 by both the Gamma-ray Burst Monitor \citep[GBM;][]{meegan09} onboard the \textit{Fermi Gamma-ray Space Telescope} at 07:47:06.42 UT \citep{zhu13,vonkienlin13}, and the Burst Alert Telescope \citep[BAT;][]{barthelmy05} onboard the \swift\ GRB mission at 07:47:57 UT \citep{maselli13}. GRB 130427A is situated at a redshift of $0.340$ making it the closest high-luminosity (E$_{\gamma,\mathrm{iso}} \gtrsim 10^{54}$ erg) gamma-ray burst since GRB 030329 \citep{levan13gcn}. Such nearby high-energy events are very rare as $\sim80\%$ of \swift\ GRBs are located at $z > 1$ and low redshift GRBs are often under-energetic \citep[see][and references therein]{perley14}. The extreme brightness and very early detection of an optical counterpart spurred a rapid succession of multi-wavelength follow-up observations making GRB 130427A one of the best spectrally and temporally sampled GRBs to date \citep{ackermann14,kouveliotou13,laskar13,levan13,maselli14,perley14,Xu13}. Broadband modelling conducted by \citet{laskar13}, \citet{perley14}, and \citet{panaitescu13}, using multi-wavelength observations ranging from 1 GHz to 0.1 TeV conducted between 300s and 60 days post-burst, revealed that the emission from GRB 130427A is best described by synchrotron emission from the combination of a reverse and forward shock. Extremely early optical observations conducted with RAPTOR (RAPid Telescopes for Optical Response) also detected a peak in optical emission $<20$s post-burst \citep{vestrand14}. This optical flash was temporally coincident with GRB 130427A's prompt $\gamma$-ray emission but modelling by \citet{vestrand14} demonstrated that it is more likely generated by the GRB's reverse-shock. This reverse shock was also shown to dominate the radio and mm wavelength bands from the first hours to days post-burst \citep{laskar13,perley14,panaitescu13}. In this paper we present AMI observations of the energetic GRB 130427A starting at 0.36 days post burst, yielding the first early time ($<1$ day) radio detection of a GRB in the \citet{staley13} GRB follow-up observing campaign. The observations and data analysis are described in Section 2 with the resulting AMI fluxes and light curve presented in Section 3. In Section 3 we also present a basic broken power law fit to the AMI light curve of GRB 130427A and discuss how it compares to other early time radio detections of GRBs. The AMI light curve modelling is further discussed in Section 4 where we consider the implications of the different slopes, the position of the peak, and how our results compare to the Very Large Array (VLA) light curves and modelling of GRB 130427A conducted by \citet{perley14}. In Section 5 we summarise our findings. | The AMI light curve of GRB 130427A agrees well with the forward/reverse shock interpretation suggested by many authors such as \citet{perley14}, \citet{panaitescu13}, and \citet{laskar13}. These early time AMI observations (within one day post-burst) have enabled us to not only obtain one of the earliest detections of a long GRB, but also capture the peak in the reverse shock emission at 15.7~GHz. This result has allowed us to further constrain the possible models of the early time afterglow from GRB 130427A by demonstrating that $\nu_a$, rather than $\nu_m$, is the cause of the radio light curve peak. This scenario will be further investigated by combining the AMI and VLA light curves with fine time sampling observations obtained with the WSRT in van der Horst et al. (in prep.). The detection of the reverse shock radio peak in the AMI light curve of GRB 130427A clearly demonstrates the importance of rapid response radio follow-up programmes of GRBs. The AMI GRB follow-up programme \citep{staley13} is therefore crucial for exploring the early time radio signatures of GRBs and constraining the radio properties of these events within the first few hours post burst. | 14 | 3 | 1403.2217 |
1403 | 1403.5601_arXiv.txt | We analyze a 162 ks HETG \chandra\ observation of the O7.5 III(n)((f)) star $\xi$~Per, together with contemporaneous H$\alpha$ observations. The X-ray spectrum of this star is similar to other single O stars, and not pathological in any way. Its UV wind lines are known to display cyclical time variability, with a period of 2.086~days, which is thought to be associated with co-rotating interaction regions (CIRs). We examine the \chandra\/ and \ha\ data for variability on this time scale. We find that the X-rays vary by $\sim 15$\% over the course of the observations and that this variability is out of phase with variable absorption on the blue wing of the \ha\ profiles (assumed to be a surrogate for the UV absorption associated with CIRs). While not conclusive, both sets of data are consistent with models where the CIRs are either a source of X-rays or modulate them. | X-ray emission is ubiquitous in the O stars and taken as an indication of dynamic instabilities in their winds \citep{lw80, l82, ow88, fel97}. These 1-D hydrodynamical models predict a plasma with temperature $\sim 1$ --10\,MK, which is permeated with cool wind clumps. The models also predict very strong stochastic X-ray variability, on time scales of hours. However, it has been clear since early X-ray observations, that stochastic variability on such short time scales is very small, less than about 1\%. To explain this, \citet{cas83} suggested that the winds contain thousands of shocks, and \citet{fel97} speculated that models with full 2-D hydrodynamics would reduce the predicted level of stochastic X-ray variability. In the most extensive analysis to date, \citet{naz13} examined high quality {\em XMM-Newton} observations of the early O supergiant $\zeta$\,Pup. They used an {\it ad hoc} 2-D wind model and found that the lack of stochastic X-ray variability on short time scales required a highly fragmented wind with a huge number of small clumps. In this picture, a stellar wind consists of a large population of cool clumps, which contain the bulk of the stellar wind matter seen at UV, optical, IR and radio wavelengths, and a tenuous hot interclump medium responsible for the X-rays. Further evidence for small scale clumping has come from the analysis of optical and UV wind lines. \citet{hil91} found that it was necessary to introduce clumping to explain the shapes of electron scattering emission line wings in Wolf-Rayet (WR) stars. Given the 1-dimensional nature of his model, this would imply the presence of either concentric shells or random structures whose size and separation are much smaller than the Sobolev length, so that the angle integrations are meaningful. Stochastic variable features in the He\,{\sc ii} $\lambda 4686$\,\AA\ emission line in $\zeta$\,Pup were found by \citet{e98}, and explained as excess emission from the wind clumps. \citet{mar05} investigated H$\alpha$ line-profile variability in a large sample of O-type supergiants, and concluded that the observed variability can be explained by a wind model consisting of coherent or broken shells. \citet{lm08} presented direct spectroscopic evidence of clumping in O and WR star winds. Besides this small scale clumping, a different, perhaps related, aspect of O star winds is that all of them which have been observed over time scales of a day or more demonstrate temporally coherent UV wind line variability \citep[e.g.,][]{mega, k96, k99}. This phenomenon suggests the presence of large structures in their winds which {\bf may} originate from large regions on the surface of the star. These results led \citet{co96} to model wind variability by large spiral structures known as Co-rotating Interaction Regions \citep[CIRs,][]{m86}. These structures originate from large regions of enhanced wind flux on the surface of the star, although the exact cause of the enhanced wind remains unexplained. \citet{ham01} showed that the observed UV line variability cannot be simply explained as a consequence of rotation in the framework of the CIR model, and a more complex interplay between rotation and radial velocity is taking place. Further, \citet{pm10} analyzed the doublet ratios of wind lines of a large number of B supergiants. Their results demonstrated that the spectral signature of large, optically thick wind structures which cover only a portion of the line of sight to the stellar disk (most likely CIRs) is common in the B supergiants. The interplay between CIRs and small scale clumping is largely unexplored. A crude analysis by \citet{o99} suggested an intricate interaction between the two. \citet{lb08} considered 3-D hydrodynamic models of CIRs. While their models could reproduce the detailed time evolution of UV spectral features in a B-type supergiant, they expressed concern that too much small scale clumping could destroy the CIRs. Whatever the mechanism of X-ray production, the presence of large scale structures in stellar winds should leave a footprint on the X-ray emission. In particular, X-ray variability on a time scale compatible with the stellar rotation period should be present, and this X-ray variability should correlate with UV wind line variability. However, there are considerable observational obstacles to establishing a link between X-ray and UV wind line variability for single, normal O stars. First, the CIRs are thought to be confined to the equatorial plane. Consequently, variability may not be seen unless our line of sight to the star is relatively close to equator-on. Second, the stellar rotation period for most O stars is several days. This is considerably longer than the typical X-ray observation, so variability could easily be missed. Third, there are only a few stars which have been observed in the UV over long enough intervals for good wind line periods to be identified, and such a period is needed to claim a definite connection between X-ray and UV wind line variability. Fourth, binary stars with colliding stellar winds must be excluded from any investigation of the connection between X-rays and large scale wind structure. Despite these hurdles, observational evidence supporting a link between CIRs and X-ray emission from stellar winds is mounting. Recently, X-ray variability of $\sim 20$\% and on the time scale of recurrent DACs was detected in the WR star EZ CMa, and explained in the framework of the CIR model \citep{o12, i13}. More relevant to the current paper, is the work on O stars by \citet{berg96}, \citet{o01} and \citet{naz13}. \citet{berg96} found a marginal detection of a periodic X-ray variability in $\zeta$~Pup, which was not confirmed by \citet{o01} or \citet{kah01} and not in agreement with DAC period determined by \citet{how95}. \citet{naz13} analyzed a series of X-ray observations of $\zeta$~Pup which span 10\,years. They detected a slow modulation of the X-ray flux with a relative amplitude up to 15\%\ over 16 hours in the 0.3-4.0~keV band. They propose that these modulations can be attributed to CIRs. The most compelling evidence to date was given by the \citet{o01} analysis of the rapidly rotating O dwarf $\zeta$\,Oph. It is the only single O star that has been observed continuously in X-rays over a full rotational period. These observations showed that the X-ray variability occurred on a time scale similar to the UV wind line variability, but they were not long enough to show that the pattern repeated on the rotation time scale. Unfortunately, there were no simultaneous observations of the wind activity that could be used to determine the phase relation between the X-ray and wind activity. Such a relationship could provide valuable clues about the geometric relationship between the X-ray emitting plasma and the CIRs. In this paper we present X-ray observations covering nearly half of a rotation period of the O7~III(n)((f)) star ${\xi}$~Per, supplemented by optical \ha\ observations which bracket the X-ray observations. ${\xi}$~Per is an ideal candidate for attempting to make a connection between UV wind line and X-ray variability. First of all, $\xi$~Per appears to be a perfectly normal O7 giant \citep[see,][for descriptions of its optical and UV spectra, respectively]{w73,w85}. Its only distinguishing feature is that it has a somewhat high (but not abnormal) rotational velocity of $v\sin i = 204$ km~s$^{-1}$ \citep{p96}. Because of its moderate $v \sin i$, its expected rotation period is short enough to be captured by a time series of manageable duration. This was done in a detailed study of its UV and optical line variability by \citet{deJong}. They demonstrated the presence of a well characterized, distinctive 2.086~day period in the UV wind line variability. This turns out to be roughly half of the expected rotation period if the star is viewed nearly equator-on (see their Figure~4). Further, they were able to model the variability by two sets of two armed CIRs (see, their Figure~17), with one set of arms dominating the variability. They were also able to establish a relationship between the appearance of the discrete absorption components (DACs) associated with the CIRs and variations in the blue wing of \ha, which enable us to use \ha\ variability as a surrogate for DAC activity. These properties of $\xi$~Per, together with its relative brightness at X-ray wavelengths, make it an ideal candidate to determine whether the CIRs that are thought to modulate the UV wind lines in this star (and possibly all luminous OB stars with massive winds) might modulate its X-ray flux as well. To pursue this conjecture, we obtained a 162~ks (0.45 rotation periods and 0.90 periods of the strong CIR activity) \chandra\/ observation of $\xi$~Per. The X-ray and optical observations are described in \S\ref{observations}, anlyzed for variability in \S\ref{analysis} and discussed in \S\ref{discussion}. | \label{discussion} There are basically two ways to generate X-ray variability. One is {\it via} an impulsive event, and the other is by occulting the source with absorbing gas or the stellar disk, as in the case of $\zeta$~Oph \citep{o01}. However, if the variability were due to an impulsive, flare-like event, it must have occurred long before the observations (since the form of the variability is nearly linear, indicative of an exponential tail). This implies that the event would have been quite strong. Another possibility would be a flare on an unseen pre-main sequence companion. However, the X-ray luminosity of $\xi$~Per is $\simeq 1.2\times 10^{32}$ erg sec$^{-1}$ \citep{ofh2006}, while the X-ray luminosities of PMS stars are typically a factor of 10 smaller \citep[see,][]{pf05}. Further, the observed spectral change is grey and soft, very uncharacteristic for flare spectra \citep[e.g.][]{gn09}. While not impossible, either scenario seems unlikely, especially since the form of the variability and its autocorrelation function are consistent with cyclical behavior. As a result, we favor the explanation that the X-rays vary because of obscuration by intervening material. In addition to $\xi$~Per, X-ray variability has been observed in two other single, non-magnetic O stars: $\zeta$~Pup \citep[O4~If(n),] []{naz13}, and $\zeta$~Oph \citep[O9.5 Vnn,][]{o01}. Both vary on time scales which may be related to their stellar rotation periods. An analysis of several years of X-ray data for $\zeta$~Pup showed evidence for variability on a time scale similar to its DACs, but no single time series encompassed a rotation period. The best evidence for the interaction between CIRs and X-rays to date is the $\zeta$~Oph data set. It was observed for $\sim 1.2$ days and its X-ray flux was found to vary by about 20\%. Further, this variation appeared to repeat with a period of about 0.77 days, roughly half of the stellar rotation period and similar to a period previously determined from its UV wind lines ($0.875 \pm 0.167$~days) by \citet{h93}. However, no information relating the phases of the X-rays and the DACs was available, so it was impossible to constrain the geometry of the CIRs and X-rays. In this paper, we have presented three observational results which help constrain the relation of the CIRs and X-rays. These are: X-ray variability, consistent with the period observed in the DACs; a constant hardness ratio for the X-rays, suggesting that the source of the X-rays becomes partial obscured by optically thick structures along the line of sight, and; a phase lag of 124$^\circ$ between the maximum absorption on the blue wing of \ha\ and the maximum strength of the X-rays. To interpret our limited set of observations (covering roughly half of a rotation period) in terms of the geometry of CIRs requires some constraints. We follow \cite{deJong} and adopt the following 3 assumptions: \begin{enumerate} \item {\bf The wind of $\xi$ Per contains two pairs of spiral arms, with one arm in each pair being much stronger than the other. We concentrate on the two major arms, which are equally spaced.} \item The wind structures follow streak lines whose shapes are defined by the stellar rotational velocity, wind velocity law and terminal velocity. We adopt a rotational velocity equal to the observed $v\sin i = 205$~\kms, a velocity law of the form $v = v_\infty [1 -a/(r/R_\star)]$, where $v_\infty = 2420$~km~s$^{-1}$ and $1 -a = v(r = 1) = 0.01 v_\infty$. \item The \ha\ emission originates near the base of the wind and is associated with the spiral structures. Since increased emission reduces $W_B$(\ha), the \ha\ emission and X-ray variability are nearly in phase. \end{enumerate} In addition, we note that the grey nature of the variability suggests that the X-ray variations are due to occultation of the X-ray source by either the stellar disk or wind structures that are optically very thick to X-rays. Thus, we seek configurations where the X-ray emission and \ha\ emission are roughly in phase. However, even with these restrictions, there is still considerable latitude in how the geometry can be arranged. We consider two possibilities, but others may be possible. The right panel of Figure~\ref{fig:cartoons} shows the first configuration. It places the X-ray emission along the interacting edge of the spirals. This might be the case if some or most of the X-rays originate at the fast wind--slow wind interface, which creates the CIRs. Further, because the velocity differential between the fast and slow winds (and, presumably the potential to produce X-rays) is expected to drop with distance from the star, one expects most of the X-rays to originate within $\sim 5 R_\star$. Thus, viewing the Figure from the right, we see that the \ha\ and X-ray emitting regions are at or near maximum. Then, as the configuration rotates counter-clockwise, our view would be from below. In this case, one lobe of the \ha\ emission and much of the X-ray emission near the star are occulted by the stellar disk, producing a minimum in both. Finally, as the system rotates further and our view is from the left, and we have returned to a maximum once again. For this configuration to explain the observations, the arms must represent a relatively low density contrast in the wind since we do not observe spectral changes in the X-ray spectrum indicative of absorption. This is in accordance with typical CIR models, where UV DAC formation is attributed mostly to a velocity plateau. However, in contrast to the classic line driven instability (LDI) model (which predicts that the X-ray source should be distributed relatively uniformly throughout the wind) the observed variability implies that the CIRs must account for a large fraction of the X-rays. The left panel of Figure~\ref{fig:cartoons} depicts a second possible configuration. It consists of a spherically symmetric X-ray emitting region which is heavily weighted toward the inner wind (in accordance with the LDI model) and two spiral arms. As before, the \ha\ emission originates near the base of the spiral arms. To obtain the X-ray variability, we must assume that the spiral arms are optically very thick, and that their line of sight optical depth decreases with distance from the star. The latter assumption corresponds to the spiral structures expanding with a constant solid angle, causing their line of sight column density to decrease as $(r/R_\star)^2$ near the star. As a result, sight lines toward the inner wind are strongly obscured by the spiral arms. When viewed from the right, the arms present minimal absorption and both the X-ray emission and \ha\ emission are at maximum. When viewed from below, the denser, inner portion of the spiral pattern occults the X-ray source and both the \ha\ the X-ray emission are near a minimum. Finally, when viewed from the left, both once again return to a maximum. This configuration agrees with typical LDI models, since the X-rays are distributed throughout the wind. However, in contrast to normal CIR models, the spiral patterns must have large column densities to account for the grey variability, and this implies that much of the wind flow is channeled through the CIRs. Regardless of which, if either, configuration is correct, the important point is that cyclical X-ray variability is inconsistent with our current understanding of either wind structure formation, X-ray production, or both. Consequently, it may provide additional clues on how to interpret wind structure. Further, depending on how the wind material is distributed, the CIRs may account for a small fraction of the wind flow, as predicted by normal CIR theory \citep{lb08} or, if they originate from regions of localized magnetic activity \citep{cb11}, a much larger fraction. Thus, it is critical to verify our results if we are to arrive at a self-consistent understanding of stellar winds and mass loss rates in OB stars. \begin{figure} \vspace{-0.5in} \vspace{-0.1in}\epsfig{figure=try2.ps,width=3.7in} \vspace{-0.3in}\caption{ Cartoons of two different configurations of the \ha\ and X-ray emitting regions in a wind containing two spiral structures in the shape of streak lines. In each case, the \ha\ emitting region is depicted by two red dots, and the X-ray emitting region by an aqua region. The coordinates are in units of $r/R_\star$.\label{fig:cartoons}} \end{figure} | 14 | 3 | 1403.5601 |
1403 | 1403.7604_arXiv.txt | The general k-essence Lagrangian for the existence of cosmological scaling solutions is derived in the presence of multiple scalar fields coupled to a barotropic perfect fluid. In addition to the scaling fixed point associated with the dynamics during the radiation and matter eras, we also obtain a scalar-field dominated solution relevant to dark energy and discuss the stability of them in the two-field set-up. We apply our general results to a model of two canonical fields with coupled exponential potentials arising in string theory. Depending on model parameters and initial conditions, we show that the scaling matter-dominated epochs followed by an attractor with cosmic acceleration can be realized with/without the couplings to scalar fields. The different types of scaling solutions can be distinguished from each other by the evolution of the dark energy equation of state from high-redshifts to today. | The scalar fields may have played important roles for the expansion history of the Universe. The slow-rolling scalar field along a nearly flat potential drives the accelerated expansion in the early Universe--dubbed inflation \cite{inf}. The theoretical prediction of inflation for the generation of density perturbations from the quantum fluctuation of a scalar degree of freedom is consistent with the Cosmic Microwave Background (CMB) temperature anisotropies observed by the Planck satellite \cite{Planck}. The scalar fields can be also responsible for dark energy at the expense of having a very light mass of the order of the today's Hubble constant $H_0 \simeq 10^{-33}$~eV \cite{quinold,quinpapers}. The energy scales of scalar fields appearing in particle physics are usually much higher than the present cosmological density \cite{review}. The dominance of the field energy density $\rho_x$ over the background energy density $\rho_m$ in the early Universe after inflation contradicts with the successful cosmological sequence of the radiation, matter, and accelerated epochs. If there is a scaling solution where $\rho_x$ is proportional to $\rho_m$, however, the Universe can enter the regime in which the field energy density is sub-dominant to the total energy density. It is well known that a canonical field with the exponential potential $V(\phi)=V_0 e^{-\lambda \phi}$ gives rise to scaling solutions with $\rho_x/\rho_m={\rm constant}$ \cite{Ratra,clw,Joyce} (where the reduced Planck mass $M_{\rm pl}$ is set to unity). Provided that the constant $\lambda$ satisfies the condition $\lambda^2>3(1+w_m)$, where $w_m$ is the equation of state of the background barotropic fluid, the solutions approach the scaling attractor characterized by the field density parameter $\Omega_x=3(1+w_m)/\lambda^2$ \cite{clw}. In the radiation-dominated epoch there is the bound $\Omega_x<0.045$ from the Big-Bang-Nucleosynthesis (BBN) \cite{Bean}, so the primordial scaling field is compatible with the data for $\lambda>9.4$. The scaling solutions also arise for non-canonical scalar-field models with the Lagrangian $P(\phi,X)$ \cite{kinf,kes}, where $X$ is a kinetic term of the field $\phi$. For the existence of scaling solutions the Lagrangian is constrained to be in the form $P=Xg(Y)$, where $g$ is an arbitrary function in terms of $Y \equiv Xe^{\lambda \phi}$ and $\lambda$ is a constant \cite{Piazza,Sami} (see also Refs.~\cite{GB,Gomes}). This accommodates the case in which the coupling $Q$ between the field $\phi$ and the barotropic fluid is present. In the absence of the coupling, the field density parameter relevant to scaling solutions during the radiation or matter eras is given by $\Omega_x=3(1+w_m)P_{,X}/\lambda^2$, where $P_{,X}=\partial P/\partial X$ \cite{Tsujikawa06,Quartin}. For the Lagrangian $P=Xg(Y)$, there exists a scalar-field dominated solution ($\Omega_x=1$) characterized by the effective equation of state $w_{\rm eff}=-1+\lambda^2/(3P_{,X})$ \cite{Tsujikawa06}. With this solution the accelerated expansion can be realized for $\lambda^2/P_{,X}<2$. In this case, however, we do not have a physically meaningful scaling solution with $\Omega_x<1$ during the radiation and matter eras. There are several possible ways of realizing a transition from the scaling regime to the epoch of cosmic acceleration. One is to introduce a canonical single field composed by the sum of exponential potentials, e.g., $V(\phi)=V_1 e^{-\lambda_1 \phi}+V_2 e^{-\lambda_2 \phi}$ satisfying $\lambda_1^2>3(1+w_m)$ and $\lambda_2^2<2$ \cite{bcn} (see also Refs.~\cite{Varun}). The joint analysis based on the data of supernovae type Ia, CMB, and baryonic acoustic oscillations showed that the two slopes are constrained to be $\lambda_1>11.7$ and $\lambda_2<0.539$ (95\,\% CL) \cite{CDT}. Another way out is to introduce multiple canonical scalar fields $\phi_i$ ($i=1,2,\cdots, N$) with the sum of exponential potentials, i.e., $V=\sum_{i=1}^{N} V_ie^{-\lambda_i \phi_i}$ \cite{Liddle,Malik}. In this case there exists a so-called assisted inflationary solution characterized by the effective equation of state $w_{\rm eff}=-1+\lambda^2/3$, where $\lambda \equiv (\sum_{i=1}^N 1/\lambda_i^2)^{-1/2}$. Even if each field is unable to be responsible for cosmic acceleration (i.e., $\lambda_i^2>2$), multiple scalar fields can cooperate to give dynamics matching a single-field solution with $\lambda^2<2$. In this model, the dynamics of scaling solutions followed by the assisted dark energy attractor has been studied in Refs.~\cite{Guo,Blais,Kim,Ohashi,Bruck}. For the system of multiple fields with the more general Lagrangian $P=\sum_{i=1}^{N}X_i\,g(Y_i)$, where $g$ is an arbitrary function with respect to $Y_i \equiv X_i e^{\lambda_i \phi_i}$, it was shown in Ref.~\cite{Ohashi} that the scaling radiation/matter eras can be followed by the assisted inflationary attractor. However, this analysis does not cover the models in which the multiple scalar fields are coupled each other \cite{Guo2,tw,ddst}. Recently, Dodelson {\it et al.} \cite{ddst} presented a string-theoretic model described by canonical fields with the coupled potential $V=\sum_{i=1}^{2} V_i e^{\alpha_i \phi_1+\beta_i \phi_2}$. This multi-field theory follows from a torus compactification of an overall volume in the presence of a dilaton field. Dodelson {\it et al.} showed the existence of not only scaling solutions but also solutions with the accelerated expansion relevant to dark energy. In general there are many scalar fields present in string theory (dilaton, moduli, axion), so the action in the Einstein frame contains coupled scalar fields after a suitable compactification. It is also known that $\alpha'$ corrections to the tree-level string action give rise to higher-order derivative terms such as $X^2$ \cite{kinf}. Then, the general Einstein-frame action with $N$ scalar fields can be described by $P(\phi_1,\cdots,\phi_N, X_1,\cdots, X_N)$. In this paper, we study the construction of the multi-field Lagrangian $P(\phi_1,\cdots,\phi_N, X_1,\cdots, X_N)$ that possesses scaling solutions. For generality we also introduce the couplings between the scalar fields and the background barotropic fluid. The resulting Lagrangian for the existence of scaling solutions is surprisingly simple and it covers the multi-field models studied in Refs.~\cite{Guo2,tw,ddst} as specific cases. Moreover, the presence of multiple scalar fields allows the transition from the scaling radiation/matter eras to the epoch of cosmic acceleration. For the general multi-field scaling Lagrangian we derive the autonomous equations of motion and the fixed points relevant to dark energy as well as the scaling solution. We show the existence of scalar-field dominated solutions ($\Omega_x=1$) having the property of assisted inflation. For the scaling solution relevant to the dynamics during the radiation and matter eras, we also obtain analytic expressions for the field density parameter $\Omega_x$ and the field equation of state $w_x$. The stability of such fixed points will be discussed by considering linear perturbations about them. As an application of our general scaling Lagrangian, we also study the cosmology for a two-field model described by canonical fields with the potential $V=\sum_{i=1}^{2} V_i e^{\alpha_i \phi_1+\beta_i \phi_2}$ in the absence/presence of couplings between the fields and the background barotropic fluid. In such a model there exist scaling and accelerated fixed points other than those discussed above. We show how the scaling radiation or matter eras (including the $\phi$-matter-dominated-epoch \cite{Amendola}) can be followed by the scalar-field dominated attractor with cosmic acceleration. This paper is organized as follows. In Sec.~\ref{genesec} we derive the multi-field Lagrangian that possesses scaling solutions characterized by $\rho_x/\rho_m={\rm constant}$. In Sec.~\ref{autosec} we obtain autonomous equations and some physically important fixed points of the general multi-field scaling Lagrangian. In Sec.~\ref{stasec} the stability of fixed points for the scaling solution and the accelerated scalar-field dominated point will be discussed. In Sec.~\ref{example} we study the cosmological dynamics for two canonical fields with the potential $V=\sum_{i=1}^{2} V_i e^{\alpha_i \phi_1+\beta_i \phi_2}$ in detail. Sec.~\ref{consec} is devoted to conclusions. | \label{consec} In general k-essence model with multiple scalar fields $\phi_i$ ($i=1,\cdots,N$), we derived the Lagrangian for the existence of cosmological scaling solutions in the presence of a barotropic perfect fluid coupled to $\phi_i$. The resulting Lagrangian is simply given by Eq.~(\ref{scalinglag}), where $g$ is an arbitrary function in terms of $Y_i=X_i e^{\lambda_1 \phi_1}$ and $Z_i=\phi_{i+1}-\lambda_1 \phi_1/\lambda_{i+1}$. Along the scaling solution, the scalar fields evolve as Eq.~(\ref{phiso}) with $Y_i$ and $Z_i$ constant. For canonical multiple scalar fields the scaling solution behaves as an effective single field $\sigma$ with the trajectory given by Eq.~(\ref{dotsigma}). For the multi-field Lagrangian (\ref{scalinglag}) we obtained the autonomous equations (\ref{x1})-(\ref{z1}) by introducing the dimensionless variables $x_i=\dot{\phi}_i/(\sqrt{6}H)$ and $y=e^{-\lambda_1 \phi_1/2}/(\sqrt{3}H)$. There are two important fixed points with $y \neq 0$ for arbitrary functions of $g$. One of them is the scalar-field dominated point A ($\Omega_x=1$) satisfying the conditions (\ref{sca1})-(\ref{weffA}). We showed that the assisted inflationary mechanism is present for the fixed point A. Even if each field does not lead to cosmic acceleration, the multiple fields can cooperatively do so by reducing the effective slope $\lambda$ defined by Eq.~(\ref{lamdef}). Provided $\lambda^2<2$, the point A can be responsible for the accelerated expansion. Another fixed point B is the scaling solution ($\Omega_x=$\,\,constant generally different from 1) satisfying the conditions (\ref{weffscaling})-(\ref{wx}). In the presence of the couplings between the fields and the background fluid ($q \neq 0$), the effective equation of state $w_{\rm eff}$ for the point B is generally different from $w_m$. If the Lagrangian (\ref{scalinglag}) is specified to some form, we can show the existence of kinetically driven fixed points satisfying $y=0$. For the function (\ref{gchoice}), which involves the case of $N$ canonical scalar fields, there exist the $\phi$MDE point C characterized by $x_i=\sqrt{6}Q_i/[3c_i(w_m-1)]$ and the kinetic point D with $\Omega_x=1$. Unlike the fixed points A and B, the quantities $y$ and $Z_i$ are not necessarily constant for C and D. In the presence of the couplings $Q_i$, the standard matter-dominated epoch can be replaced by the $\phi$MDE with $w_{\rm eff}=\Omega_x=\sum_{i=1}^N 2Q_i^2/(3c_i)$. In Sec.~\ref{stasec} we studied the stability of the fixed points A and B in the two-field system described by the Lagrangian (\ref{lagtwo}). For the point A it is possible to carry out the general analysis without specifying any functional form of $g$. The stability of this scalar-field dominated solution is ensured under the conditions (\ref{Acon1}), (\ref{Acon2}), and (\ref{Acon3}). The stability analysis for the point B is too cumbersome to be written in a general way. In the model of two canonical fields with the function (\ref{twocanonical2}), we showed that the scaling solution B is stable under the conditions (\ref{Bsta1}) and (\ref{Bsta2}) with the additional requirement (\ref{hZre}). These conditions can be nicely interpreted as a geometric approach based on the rotation in the field space. In the case of two canonical fields with $Q_i=0$, we found that the point A is stable (unstable) when the point B is unstable (stable). In Sec.~\ref{example} we discussed the cosmological dynamics for the two-field model (\ref{conlag}) in the presence of radiation and non-relativistic matter. In doing so, it is convenient to employ the variables $y_1$ and $y_2$ defined in Eq.~(\ref{y2def}) instead of $y$ and $Z_1$. In this model there exist a few more fixed points than A, B, C, D relevant to the scaling radiation/matter eras and the epoch of cosmic acceleration. These fixed points are summarized in Table \ref{table2}. Depending on the model parameters and initial conditions of the model (\ref{conlag}), there are several qualitatively different cases: (i) the scaling radiation era (B1) followed by the stable scaling matter era (B2), (ii) the scaling radiation (E1) and matter (E2) eras followed by the point A with cosmic acceleration, and (iii) the $\phi$MDE point C followed by the accelerated expansion driven by another point F. The points E1, E2, and F, which satisfy $y_2=0$, arise when the potential energy of the last term of Eq.~(\ref{conlag}) is negligibly small relative to the third term. The cases (ii) and (iii) can be distinguished from each other by the different evolution of the dark energy equation of state $w_x$. It will be of interest to put observational bounds on the viable parameter space of the model (\ref{conlag}) and its extended model by using the data of supernovae type Ia, CMB, baryonic acoustic oscillations, and the BBN. The BBN bound should not be so restrictive in the multi-field context because the field density parameter $\Omega_x$ in the radiation era can be smaller than that in the single-field case. The study of matter density perturbations is also important to place constraints on the couplings $Q_i$ in the presence of the $\phi$MDE. We leave these topics for a future work. | 14 | 3 | 1403.7604 |
1403 | 1403.7574_arXiv.txt | We study relationships between properties of collective excitations in finite nuclei and the phase transition density $n_t$ and pressure $P_t$ at the inner edge separating the liquid core and the solid crust of a neutron star. A theoretical framework that includes the thermodynamic method, relativistic nuclear energy density functionals and the quasiparticle random-phase approximation is employed in a self-consistent calculation of $(n_t,P_t)$ and collective excitations in nuclei. The covariance analysis shows that properties of charge-exchange dipole transitions, isovector giant dipole and quadrupole resonances and pygmy dipole transitions are correlated with the core-crust transition density and pressure. A set of relativistic nuclear energy density functionals, characterized by systematic variation of the density dependence of the symmetry energy of nuclear matter, is used to constrain possible values for $(n_t,P_t)$. By comparing the calculated excitation energies of giant resonances, energy weighted pygmy dipole strength, and dipole polarizability with available data, we obtain the weighted average values: $n_t = 0.0955 \pm 0.0007$ fm$^{-3}$ and $P_t = 0.59 \pm 0.05$ MeV fm$^{-3}$. | 14 | 3 | 1403.7574 |
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1403 | 1403.1955_arXiv.txt | Three presumably young eclipsing binary systems in the direction of the Cygnus OB1, OB3 and OB9 associations are studied. Component spectra are reconstructed and their orbits are determined using light curves and spectra disentangling techniques. V443\,Cyg and V456\,Cyg have circular orbits, while the light curve of V2107\,Cyg imposes a slightly eccentric orbit ($e=0.045\pm0.03)$. V443~Cyg harbours F--type stars, and not young early--A stars as previously suggested in the literature based on photometry solely. It appears to be situated in the foreground (distance $0.6\pm0.2$~kpc) of the young stellar populations in Cygnus. V456~Cyg, at a distance of $0.50\pm0.03$~kpc consists of a slightly metal-weak A--type and an early--F star. The age of both systems, on or very near to the main sequence, remains uncertain by an order of magnitude. V2107~Cyg is a more massive system ($8.9\pm2$ and $4.5\pm1.2 M_\odot$) at $1.5\pm0.5$~kpc and, also kinematically, a strong candidate-member of Cyg\,OB1. The more massive component is slightly evolved and appears to undergo non-radial $\beta~Cep$--type pulsations. The Doppler signal of the secondary is barely detectable. A more extensive study is important to fix masses more precisely, and an asteroseismological study would then become appropriate. Nevertheless, the position of the primary in the HR-diagram confines the age already reasonably well to $20\pm5$~Myr, indicating for Cyg\,OB1 a similar extent of star formation history as established for Cyg\,OB2. | The star forming regions in the constellation Cygnus are among the most active regions of star birth in our Galaxy. An overview of these regions is presented in Reipurth \& Schneider (2008). These authors emphasized the confusion of regions as near as hundreds of parsecs with regions at 1--2 kpc and even well beyond, due to the fact that in the direction of Cygnus one looks down a spiral arm. Kinematical distances suffer from the fact that the radial velocity gradient up to 4 kpc is smaller than the typical velocity dispersion of interstellar gas. As a consequence, even the need for the nine classical subgroups OB1 to OB9 in Cygnus is still discussed, with for example a suggestion that OB1, OB8 and OB9 might just form one single subgroup (Melnik \& Efremov 1995). Hence, distances derived from young eclipsing binary systems are of value to localise the star forming regions and stellar associations. A major study of binary stars in Cygnus has focused on Cyg~OB2, the most famous and probably youngest subgroup. The Cyg~OB2 radial velocity survey encompasses now twenty-five massive binaries (Kobulnicky et al.\ 2012; Kiminki \& Kobulnicky 2012). Another very recent study (Mahy et al.\ 2013) focus on bright O--type stars. 274 spectra taken over three years reveal four spectroscopic binaries (three in Cyg~OB1 and one in Cyg~OB9; none in Cyg~OB3 and Cyg~OB8). Earlier studies discuss binary stars in open clusters superposed on Cyg~OB1. Zakirov (1999) presents three eclipsing binaries with early--B primary components which are plausible members of the open cluster IC~4996, but $UBVR$ photometry without orbital-phase resolved spectroscopy did not lead to precise fundamental parameters. Boeche et al.\ (2004) detect nine binaries out of sixteen stars in the direction of NGC~6913, including seven binary stars (one eclipsing) among the 12 presumed cluster members. They adopt a distance of 1.6~kpc to NGC~6913, but emphasize (their Fig. 2) that the cluster distance cannot be well constrained. Very recently, 60 eclipsing systems in the field of the Cygnus OB7 star forming region were identified by Wolk et al. (2013) from $JHK$ photometry in 100 nights spanning 1.5 year. A minority of them (23) are detached systems with orbital periods between $0\fd4$ and just over 13 days, and a majority (37) show continuous flux variations with periods ranging from 0.2 to 3 days. Most of the latter are thought to be W~UMa contact binaries. These stars have $J$ band magnitudes from 12.5 to 17.3, and are identified as diskless field stars. W~UMa type contact binaries are not observed in clusters younger than about 1~Gyr. We did not find any literature on later-type binary stars unequivocally identified as members of the young associations. The stars discussed in this paper are situated in the OB1--OB3--OB9 region. Two of them are suspected to be pre-main sequence systems with A-type components, based on photometric studies (Zakirov \& Eshankulova 2005; 2006). They suggest the stars are connected to Cyg OB9 (V443 Cyg) and to a group of OB-stars in the direction of Cyg OB1 (V456 Cyg). Nelson (2011) points out the lack of spectroscopic knowledge for V456 Cyg and leaves open the question whether the primary component of this system is an early--A or a late--A star. The third object, V2107 Cyg (HD~191473), is an early-B system at the end of or slightly off the main sequence with a primary component of luminosity class III to IV (B0.5\,III, Roman 1951; B0\,IV, Morgan et al. 1955; B1\,III, Walborn 1971), and thus a particularly interesting age calibrator. Their position in the Cygnus region, relative to the stellar subgroups, is shown in Fig.~1. \begin{figure} \begin{center} \plotone{fig1.eps} \caption{Region of OB1,2,3,8,9 stellar associations in Cygnus according to partitions of Blaha \& Humphreys (1989). \label{fig1}} \end{center} \end{figure} In this paper, we summarize first the characteristics of the observations, and the analysis methodology applied (\S~2 \& 3). Thereafter, each of the three binary systems is discussed in turn (\S~4 to 6) and the general conclusions are presented in \S~7. | Table~6 lists the characteristics of the components and their orbits, as derived from the combined spectroscopic--photometric analysis. The proper motions were taken from Roeser, Demleitner, \& Schilbach (2010), while the distances and systemic velocities of the systems were derived in this study. The space velocity components of the systems were calculated using the algorithm of Johnson \& Soderblom (1987) and corrected for the local standard of rest (U, V, W = 8.50, 13.38, 6.49) kms$^{-1}$ (Co{\c s}kuno{\v g}lu et al., 2011). The total space velocities of V2107 Cyg, V443 Cyg and V456 Cyg are 63, 72 and 14 kms$^{-1}$, respectively. Applying the N-body code of Dinescu et al.\ (1999) shows that the three systems move in the Galaxy in moderately eccentricity (Galactic) orbits ($e=0.17, 0.18, 0.07$ for V2107~Cyg, V443~Cyg and V456~Cyg, respectively). These eccentricities show that the three systems belong to the thin-disc population. However, V2107~Cyg is at a distance of $1.5 \pm 0.5$~kpc to us, while V443~Cyg (at $0.6 \pm 0.2$~kpc) and V456~Cyg (at $0.50 \pm 0.03$~kpc) are much nearer to us. \footnote{Bolometric corrections from Strai\v{z}ys \& Kuriliene 1981, and $M_{V,\odot} = 4.75$~mag were used.} The Mel'nik \& Efremov (1995) distance estimate of Cyg OB1 (1.4 kpc) and the Tetzlaff et al.\ (2010) space velocity components ($U, V, W = 45(2), -27(2), -7(2)$~kms$^{-1}$) are compatible with V2107~Cyg, that we consider a strong candidate member of this stellar group. The next subsection discusses shortly the implications for the star formation history of Cyg~OB1. \begin{deluxetable}{ccccccc} \setlength{\tabcolsep}{1pt} \scriptsize{} \tablecaption{Close binary stellar parameters of the program stars. Uncertainties of parameters are given in brackets. \label{table6}} \tablewidth{0pc} \tablehead{ & \multicolumn{2}{c}{V443~Cyg} & \multicolumn{2}{c}{V456~Cyg} & \multicolumn{2}{c}{V2107~Cyg} \\ Parameter & Primary & Secondary & Primary & Secondary & Primary & Secondary } \startdata Sp & F2/3V & F8V & A2hA7mA4V & F3V & B1III & \nodata \\ \emph{M}\,(M$_\odot$)& 1.2(0.2) & 1.1(0.3) & 1.86(0.06) & 1.58(0.05) & 8.9(2.0) & 4.5(1.3) \\ \emph{R}\,(R$_\odot$)& 1.3(0.2) & 1.1(0.2) & 1.68(0.02) & 1.47(0.02) & 7.4(0.6) & 2.4(0.3) \\ \emph{a}\,(R$_\odot$)& \multicolumn{2}{c}{9.40(0.33)}& \multicolumn{2}{c}{5.88(0.06)} & \multicolumn{2}{c}{26.4(1.9)} \\ \emph{P}\,(days) & \multicolumn{2}{c}{1.66220545(43)}& \multicolumn{2}{c}{0.89119559(17)} & \multicolumn{2}{c}{4.2845923(25)} \\ \emph{i}\,($^{\circ}$) & \multicolumn{2}{c}{89.9(0.5)} & \multicolumn{2}{c}{82.9(0.4)} & \multicolumn{2}{c}{86.2(0.5)} \\ \emph{q} $(M_2/M_1)$ & \multicolumn{2}{c}{0.98(0.09)} & \multicolumn{2}{c}{0.85(0.02)} & \multicolumn{2}{c}{0.50(0.06)} \\ $R_{L}$\,(R$_\odot$) & \multicolumn{2}{c}{3.54} & \multicolumn{2}{c}{2.15} & \multicolumn{2}{c}{8.57} \\ \emph{e} & \multicolumn{2}{c}{0.0} & \multicolumn{2}{c}{0.0} & \multicolumn{2}{c}{$<$0.03} \\ $\log g$\,(cgs) & 4.25(0.23)& 4.42(0.26) & 4.258(0.027)& 4.302(0.027) & 3.65(0.17)& 4.3(0.2) \\ \emph{V}\,(mag) & \multicolumn{2}{c}{12.31(0.03)} & \multicolumn{2}{c}{10.80(0.03)} & \multicolumn{2}{c}{8.63(0.02)} \\ $B-V$\,(mag) & \multicolumn{2}{c}{0.51(0.09)} & \multicolumn{2}{c}{0.31(0.02)} & \multicolumn{2}{c}{0.10(0.05)} \\ $E(B-V)$\,(mag)& \multicolumn{2}{c}{0.13(0.15)} & \multicolumn{2}{c}{0.11(0.05)} & \multicolumn{2}{c}{0.335(0.045)} \\ $A_{\rm v}$\,(mag) &\multicolumn{2}{c}{0.42(0.15)} &\multicolumn{2}{c}{0.35(0.08)} &\multicolumn{2}{c}{1.04(0.14)} \\ $(B-V)_{\rm0}$\,(mag) &\multicolumn{2}{c}{0.38(0.06)} &\multicolumn{2}{c}{0.20(0.02)} &\multicolumn{2}{c}{-0.23(0.05)} \\ $T_{\rm eff}$\,(K) & 6700(350) & 6200(550) & 7750(100) & 6755(400) & 22500(1500) & 15200(1600) \\ $\log$ \emph{L}\,(L$_\odot$)& 0.52(0.27) & 0.20(0.21) & 0.96(0.05) & 0.61(0.06) & 4.10(0.18) & 2.5(0.5) \\ $M_{\rm bol}$\,(mag) & 3.46(0.67) & 4.25(0.53) & 2.35(0.11) & 3.23(0.14) & --5.5(0.5) & --1.4(1.2) \\ $M_{\rm v}$\,(mag) & 3.44(0.65) & 4.27(0.53) & 2.32(0.12) & 3.21(0.14) & --3.4(0.6) & --0.1(1.2) \\ \emph{BC}\,(mag)& 0.02(0.02) & --0.02(0.05)& 0.03(0.01)& 0.02(0.01) & --2.2(0.2) & --1.3(0.2) \\ $K_{\rm 1,2}$\,(km\,s$^{-1}$)& 117.4(6.2) & 120.2(4.2) & 152.1(1.7) & 179.4(1.7) & 104(1.5) & 207(20) \\ $V_{\gamma}$\,(km\,s$^{-1}$) & \multicolumn{2}{c}{-19.7(0.7)} & \multicolumn{2}{c}{--2.7(1.0)} & \multicolumn{2}{c}{--6.3(0.2)} \\ V$_{synch}$\,(km\,s$^{-1}$) & 41(7) & 33(6) & 95(1) & 83(1) & 87(7) & 29(4) \\ V$_{\rm rot}$\,(km\,s$^{-1}$) & \nodata & \nodata & 100(10) & \nodata & 84(4) & \nodata \\ \emph{d}\,(pc) & \multicolumn{2}{c}{600(200)} & \multicolumn{2}{c}{500(30)} & \multicolumn{2}{c}{1500(500)}\\ $\mu_\alpha cos\delta$, $\mu_\delta$\,(mas yr$^{-1}$) & \multicolumn{2}{c}{--11.1(2.2), --15.8(2.2)} & \multicolumn{2}{c}{5.6(1.2), 1.3(1.1)} & \multicolumn{2}{c}{--4.8(1.3), --6.5(1.3)} \\ $U$\,(km\,s$^{-1}$) & \multicolumn{2}{c}{48.98(22.39)} & \multicolumn{2}{c}{--10.66(2.73)} & \multicolumn{2}{c}{49.76(19.06)} \\ $V$\,(km\,s$^{-1}$) & \multicolumn{2}{c}{--31.05(5.02)} & \multicolumn{2}{c}{--0.46(1.14)} & \multicolumn{2}{c}{--20.79(5.46)} \\ $W$\,(km\,s$^{-1}$) & \multicolumn{2}{c}{--0.57(6.22)} & \multicolumn{2}{c}{--9.02(2.84)} & \multicolumn{2}{c}{3.14(8.70)} \\ \enddata \end{deluxetable} \subsection{V2107~Cyg and the age of Cyg~OB1} The primary component is obviously near the end of the main-sequence, while the secondary is near the zero-age main-sequence. The fundamental stellar parameters of the secondary suggest a spectral type B5 V. In view of the light contribution of about six per cent to the total light, core of the diluted Hydrogen lines would be three per cent deep in the observed spectra and the strongest metal lines, Mg II\,$\lambda$4481 would be one per cent deep in case of rotation synchronized with the orbit. In order to detect these lines with confidence one requires on the order of 15 spectra with S/N of 100. Fig.~9 shows the location of the components in the $\log T_{eff}$ -- $\log L$ plane, and isochrones calculated using the web interface\footnote{http://stev.oapd.inaf.it/YZVAR/cgi-bin/form}. A solar mixture was assumed, i.e. Z=0.017, Y=0.26. Despite the present uncertainties, the age of the system can be constrained to $20 \pm 5$~Myr due to the favourable position of the primary component. The age estimate of 7.5 Myr for the group as a whole (Mel'nik \& Efremov 1995) suggests that star formation may have been on-going for at least $10^ 7$~yr. Hence, as well in Cyg~OB1 as in Cyg~OB2 (Comeron \& Pasquali 2012), there is evidence in favour of sustained star formation for well over 10 Myr. This study identifies V2107\,Cyg as an extremely interesting binary star with the potential to constrain the duration of the star formation process in the Cyg OB1 association with precision. More extensive high-resolution spectroscopy, as already obtained at the mid-eclipses, is required to constrain $K_2$ precisely (and thus the mass of the primary). Moreover, with the secondary component of V2107\,Cyg occulted for six consecutive hours, the intrinsic line profile variations of the primary component can be studied even without the (weak) contamination and dilution by the secondary star. It has been noticed, in secondary eclipse, that many metal lines show a peculiar line core with a similar asymmetry, as illustrated for the Si~{\sc iii} triplet in Fig.~5. Spectroscopic monitoring should confirm that this is due to non-radial pulsations. \begin{figure} \begin{center} \epsscale{0.5} \plotone{fig9a.eps} \\ \plotone{fig9b.eps} \\ \plotone{fig9c.eps} \\ \caption{Locations of program stars in the $\log~L$ -- $\log~T_{eff}$ plane together with evolutionary tracks and isochrone curves. Tracks and isochrones for V456~Cyg are calculated for Z=0.035 (full line) and Z=0.020 (dashed, PMS), respectively.}\label{fig9} \end{center} \end{figure} \clearpage | 14 | 3 | 1403.1955 |
1403 | 1403.4607_arXiv.txt | We develop a method for calculating the correlation structure of the Cosmic Microwave Background (CMB) using Feynman diagrams, when the CMB has been modified by gravitational lensing, Faraday rotation, patchy reionization, or other distorting effects. This method is used to calculate the bias of the Hu-Okamoto quadratic estimator in reconstructing the lensing power spectrum up to $\ord{\phi^4}$ in the lensing potential $\phi$. We consider both the diagonal noise \corr T T T T, \corr E B E B, etc.~and, for the first time, the off-diagonal noise \corr T T T E, \corr T B E B, etc. The previously noted large $\ord{\phi^4}$ term in the second order noise is identified to come from a particular class of diagrams. It can be significantly reduced by a reorganization of the $\phi$ expansion. These improved estimators have almost no bias for the off-diagonal case involving only one $B$ component of the CMB, such as \corr E E E B. | \label{sec:intro} Primary anisotropies in the Cosmic Microwave Background (CMB) were generated around $375,000$ years after the ``big-bang", when the perturbations were still in the linear regime. The CMB is characterized by its temperature and polarization. The temperature map $T(\nhat)$ describes the temperature fluctuation in the direction $\nhat$. The polarization of the CMB radiation is conventionally decomposed in terms of the polarization modes $E$ and $B$, which have even and odd parity, respectively. Primordial scalar perturbations create only $E$-modes of the CMB, while primordial tensor perturbations generate both parity-even $E$-modes and parity-odd $B$-modes~\cite{Seljak:1996gy, Kamionkowski:1996ks, Kamionkowski:1996zd}. The amplitude of primordial $B$-modes of the CMB is proportional to the energy scale at which inflation occurs; hence constraints on $B$-modes will provide valuable information about the early universe~\cite{2009AIPC.1141...10B}. Many ground-based and balloon experiments are looking for primordial $B$-modes of the CMB. The primordial $B$-mode gravitational wave signal has just been detected, constraining the ratio of tensor to scalar perturbations~\cite{Ade:aa}. Even in the absence of primordial $B$-modes, subsequent gravitational lensing by the Large Scale Structure (LSS) of the Universe converts $E$-mode polarization to $B$-mode polarization~\cite{1998PhRvD..58b3003Z}. Gravitational lensing of the CMB was first detected by cross-correlating the CMB with LSS data~\cite{2007PhRvD..76d3510S, 2008PhRvD..78d3520H}. It has since been detected using CMB data alone~\cite{2011PhRvL.107b1301D,Ade:2013tyw,2013PhRvL.111n1301H, 2013arXiv1312.6646P, 2013arXiv1312.6645P}. CMB lensing is a clean and powerful probe of several cosmological parameters~\cite{1997MNRAS.291L..33B,1997ApJ...488....1Z,1998ApJ...492L...1M,Kaplinghat:2003bh, 2006PhRvD..74j3510A,2002PhRvD..65b3003H}, see e.g.~the reviews in Refs.~\cite{Lewis:2006fu,2009AIPC.1141..121S}. Lensing depends on the integrated (time-dependent) gravitational potential along the path of the CMB photons, and provides the deepest possible measurement of the two dimensional mass distribution of the universe. It can be used to break degeneracies between cosmological parameters and to constrain the neutrino mass, independently of other probes like Lyman-$\alpha$ and galaxy clustering. It also probes the late time evolution of structure in the universe, and hence provides constraints on dark energy and the integrated Sachs-Wolfe effect. Finally, if the lensing map is well measured, then one can de-lens the observed CMB (i.e.~effectively subtract off lensing induced $B$-modes) to increase sensitivity to inflationary $B$-modes and the energy scale of inflation. Initial studies focused on the effect of lensing on the power spectra of the CMB modes~\cite{Seljak:1995ve,Zaldarriaga:1998ar,Seljak:1998nu}. Later, the emphasis switched to constraining the LSS from measurements of the CMB, by using e.g.~the fact that lensing introduces non-Gaussian fluctuations~\cite{Zaldarriaga:2000ud} which are not present in the primordial CMB. A particularly fruitful approach has been the use of quadratic estimators, built out of the convolution of two CMB modes to determine the lensing from the statistical breaking of isotropy. A quadratic estimator based on CMB modes in position space was introduced in Ref.~\cite{Seljak:1998aq}. We focus on the Hu-Okamoto quadratic estimator~\cite{Hu:2001kj,Okamoto:2003zw}, which is uniquely determined by the requirements that it is unbiased and has minimal variance. A likelihood-based approach shows that for the temperature mode this method is close to optimal, though improvement is possible for polarization modes~\cite{Hirata:2002jy,Hirata:2003ka}. The quadratic estimator has been used by \planck\ to detect lensing at $>25\,\sigma$ using only temperature information~\cite{Ade:2013tyw}. Recently the detection of lensing using CMB polarization data has been reported in Refs.~\cite{2013PhRvL.111n1301H, 2013arXiv1312.6646P, 2013arXiv1312.6645P}. The quadratic estimators provide an unbiased means for extracting the lensing potential $\phi$. However, an estimate of the lensing power spectrum depends on the two-point function of quadratic estimators. This two-point function will have a bias due to inherent noise from random fluctuations. This noise bias needs to be subtracted to get a reliable measurement of the lensing power spectrum. The calculation of noise is carried out in a small lensing expansion (unlike the simpler case of the lensed power spectra where such an expansion is unnecessary~\cite{Challinor:2005jy}). The first order correction for all estimators, and the second order correction for the estimator \est T T\ were determined in Refs.~\cite{Cooray:2002py,Kesden:2003cc,Hanson:2010rp,2014arXiv1403.2386J}. Simulations of the second order corrections for all channels were studied in Ref.~\cite{Anderes:2013jw}. The second-order corrections turned out to be unexpectedly large at small $L$ (large scales), casting doubt on the convergence of the small lensing expansion. In this paper, we develop the method for calculating the lensing of the CMB using Feynman diagrams presented in Ref.~\cite{2014arXiv1403.2386J}, which simplifies its calculation and enables us to readily identify the origin of this large second-order correction. Feynman diagrams have been used in other cosmological applications, see e.g.~\cite{Goroff:1986ep}, and were employed in Ref.~\cite{Rathaus:2011xi} to study the effect of a single gravitational lens. The computation of CMB correlations is very similar to the computation of correlation functions in statistical physics or quantum field theory, so the standard Feynman diagram approach is a very convenient and efficient way to organize the calculation. The lensing potential can be extracted using the $xy$ quadratic estimator, where $x,y \in \{T,E,B\}$, and the lensing power spectrum can be extracted by looking at the $xy-x^\prime y^\prime$ two-point function. Our formalism can be applied to study the noise not only for the \corr T T T T\ correlator, but for all of the $xy-x^\prime y^\prime$ correlators. In addition to gravitational lensing, the observed CMB can be distorted by other effects, and our diagrammatic treatment can be trivially extended to these cases. We will discuss two cases in \sec{other}, focussing on distortions due to patchy reionization, and rotations of the CMB polarization due to parity-violating Chern-Simons terms from axion fields, which exist in many theories. The outline of this paper is as follows. In \sec{lensing}, we review the basics of gravitational lensing of the CMB, which we formulate in the language of Feynman diagrams in \sec{feynman}. We derive the Feynman rules for patchy reionization and rotation of the CMB polarization in \sec{other}. The quadratic estimator for gravitational lensing and its noise terms up to order $\phi^4$ are calculated in \sec{quadratic}. We show numerical results in \sec{results} and conclude in \sec{summary}. | \label{sec:summary} We have calculated noise properties of Hu-Okamoto-based quadratic estimators of CMB lensing, using a Feynman diagram approach~\cite{2014arXiv1403.2386J}. This method allowed us, for the first time, to obtain analytical expressions for the higher order noise (up to $\ord{\phi^{4}}$) of lensing estimators based on any combination of CMB temperature and polarization channels. Previous analytical calculations were limited to the temperature channel (\corr T T T T) at this order. We have also discussed how to extend this calculation to other distorting fields like patchy reionization and cosmic rotation, deriving the relevant Feynman rules. Using this approach, it was straightforward to identify the origin of the (supposed) poor convergence of higher order noise terms. The previously noted large $\ord{\phi^4}$ term in the second order noise $N^{(2)}_L$ has been identified to come from a particular class of diagrams. By reorganizing the $\phi$ expansion, we significantly reduced the effect of higher order noise terms. We have shown results for the estimator noise up to $\ord{\phi^{4}}$ for 9 channels and for a wide range of experimental setups. From this computation we conclude that, using our re-arranged counting, the estimator is well behaved for all the channels. We also note that estimators with an odd number of $B$-fields have a very small noise. With more precision CMB polarization experiments on the way, it is extremely important to understand and improve the statistical techniques used to extract the lensing information from the data. High precision lensing maps open the door to constraining several fundamental cosmological parameters, the sum of the neutrino masses and the properties of dark energy. Perhaps most importantly, characterizing lensing opens up the possibility of ``delensing" which enhances sensitivity to measure inflationary $B$-modes induced by tensor perturbations. This is important for $L \gtrsim 100$ when the lensing contribution is no longer small compared to primordial $B$-modes. \begin{figure*} \includegraphics[bb=33 544 337 688,width=12cm]{N2_compare} \caption{Contributions to the lensing estimator noise $N^{TT,TT(2)}_L$ as a function of $L$ for a \planck-like experiment. The left panel uses the unlensed counting and the right one our lensed counting. Shown are the lensing power spectrum $C^{dd}_L$ (solid black), the connected contribution $N^{(2,c)}$ (red) and the disconnected contribution separated into $N^{(2,d)}_a$ (green) and $N^{(2,d)}_b$ (blue), corresponding to the fourth and fifth line of \eq{N2}. When the contributions to $N^{TT,(2)}$ are positive they are shown as solid lines, and when they are negative, they are shown as dashed lines. } \label{fig:N_contributions} \end{figure*} | 14 | 3 | 1403.4607 |
1403 | 1403.1813_arXiv.txt | Giant planet formation in the core accretion plus gas capture (CA) paradigm is predicated by the formation of a core, assembled by the coagulation of grains and later by planetesimals within a protoplanetary disc. As the core mass increases beyond a critical value, the hydrogen-dominated atmosphere around the core becomes self-gravitating and collapses onto the core, triggering rapid gas accretion which can lead to the formation of a gaseous planet. In contrast, in the disc instability paradigm, giant planet formation is believed to be independent of core formation: massive self-gravitating gas fragments cool radiatively and collapse as a whole independently of whether there is a core. In this paper we show that giant planet formation in the disc instability model may be also enhanced by core formation for reasons physically very similar to the CA paradigm. In the model explored here, efficient grain sedimentation within an initial fragment (rather than the disc) leads to the formation of a core composed of heavy elements. We find that massive atmospheres form around cores and undergo collapse as a critical core mass is exceeded, analogous to CA theory. The critical mass of the core to initiate such a collapse depends on the fragment mass and metallicity, as well as core luminosity, but ranges from less than 1 to as much as $\sim80$ Earth masses. We therefore suggest that there are two channels for the collapse of a gaseous fragment to planetary scales within the disc instability model: (i) H$_2$ dissociative collapse of the entire gaseous clump, and (ii) core-assisted gas capture, as presented here. We suggest that the first of these two is favoured in metal-poor environments and for fragments $\simgt 5-10$ Jupiter masses, whereas the second is favored in metal-rich environments and fragments of lower mass. | The self-gravity of a massive protoplanetary disc may, under certain conditions \citep{Toomre64}, cause small perturbations within the disc to grow in amplitude (i.e., a gravitational instability; GI) until the instabilities saturate in the nonlinear regime. GIs can manifest themselves as spiral arms and shocks, heat the disc, and drive angular momentum and mass transport \citep[see][for a review]{DurisenEtal07}, as well as create molecular abundance variations that could be used as observational telltales of active GIs (Ilee et al.~2010, Douglas et al. 2013). If the combination of disc heating and mass transport fail to saturate (self-regulate) GIs \citep[e.g.,][]{Gammie01,CossinsEtal09}, then the densest regions of spiral arms can collapse to form self-bound gaseous clumps, which may be precursors of giant planets \citep[e.g.,][]{Kuiper51b,Boss98}. When such fragments first form, their mass is of order the mass of a gas giant planet \citep[e.g.,][]{BoleyEtal10,Nayakshin10a,ForganRice11}. However, their internal properties are very far from present day gas giants. In particular, the fragments are $\sim$ 10-12 orders of magnitude less dense than Jupiter(!), dominated in mass by molecular hydrogen, have initial central temperatures of just a few hundred Kelvin, and do not have high-metallicity cores \citep[although they do contain an admixture of heavy elements and grains, perhaps significantly enriched over nebular abundances, see, e.g.,][]{BoleyEtal11a}. Removal of gas by tides \citep{BoleyEtal10,Nayakshin10c} or the accretion of gas \citep{SW08,KratterEtal10} and planetesimals \citep{HelledEtal06}, plus nonlinear planet-disc coupling through radiative heat exchange \citep{VazanHelled12,NayakshinCha13} together make it extremely difficult to say what becomes of the remaining object. The list of possibilities goes from rocky and giant planets to brown dwarfs and even low mass stars \citep{SW08,KratterEtal10,ZhuEtal12a,ForganRice13}. Rapid inward migration of first fragments \citep{VB05,VB06,ChaNayakshin11a,BaruteauEtal11,MichaelEtal11,ZhuEtal12a} is a particular concern for the survivability of giant planets formed by GIs. Initial fragment sizes are comparable to the corresponding Hill sphere of the clumps themselves, to within a factor of two or three \citep[see, e.g.,][]{BoleyEtal10}. To survive the rapid radial migration, the clumps need to cool and contract rapidly, so that central temperatures reach $\sim 2000$ K \citep[e.g.,][]{HelledEtal06} {\it before} tidal forces from the star disrupt the planet. At $T\sim 2000$~K, molecular hydrogen can dissociate, redirecting energy from pressure support into internal molecular processes. The resulting dynamical collapse of the protoplanet from sizes of $\sim 1000$ R$_J$ to a few to tens of R$_J$ \citep{BodenheimerEtal80} is analogous to the end of the first core stage in star formation \citep[e.g.,][]{Larson69}, and may allow the protoplanet to avoid tidal destruction \citep[e.g.,][]{Nayakshin10c}. In this paper we describe a second, so far insufficiently explored, channel for the collapse of giant planets faced with rapid inward migration. This second channel exists entirely due to the heavy element component within the fragments. As is well known, rapid grain growth can promote sedimentation of large grains to the centre of the clump \citep{McCreaWilliams65,Boss98,HelledEtal08}. In this paper we show that {\it if} the segregation of heavy elements to the central regions of the fragment is efficient, then dynamical collapse of the clump may be initiated by a dynamical instability in the clump's centre, next to the core, {\it before the mean temperature of the fragment rises to} $\sim 2000$~K . This instability is quite analogous to that laying the foundation of the core accretion plus gas capture theory of giant planet formation \citep[e.g.,][]{PollackEtal96} , which we refer to as simply core accretion (CA) hereafter. In CA, the core grows by accretion of planetesimals from the disc and gravitationally attracts a gaseous "atmosphere", again from the disc. The atmosphere's mass increases as a steep power of core mass \citep{Stevenson82}. When the two masses become comparable, there is no stable hydrostatic solution and the atmosphere collapses onto the core dynamically \citep{Mizuno80}, triggering a rapid gas accretion from the disc onto the protoplanet \citep[see][for a recent detailed treatment of the problem]{Rafikov06}. We find a very similar sequence of events, except that the core grows by accretion of small grains rather than planetesimals, and all the action occurs inside a self-gravitating gas fragment born from GIs. Understanding the core accretion-like instability discovered here is a key to understanding the pre-collapse evolution of the disc instability planets, and is necessary for identifying ways to distinguish observationally between planet formation modes. To differentiate between the instability discussed here and the one in the CA hypothesis, we call the collapse studied in this work a "core-assisted gas capture" (CAGC) instability, with emphasis that the core is embedded deep within a massive self-gravitating gaseous clump\footnote{Hereafter, we refer to the gaseous region immediately surrounding a sedimented core as the core's ``atmosphere'', even though it is deep within a fragment's interior.}. Global disc fragmentation simulations are not yet advanced enough to follow both the formation of a clump and the clump's subsequent evolution to planetary scales. In particular, the region near the growing core is numerically challenging to investigate due to the required resolution and the extreme changes in gas composition. For example, a core will have a radius of $\sim 10^{-4}$~AU, which is about 10,000 times smaller than the size scale of a molecular hydrogen gaseous clump. Although this core may be small compared with the total extent of the clump, we shall later see that destabilisation of gas directly surrounding the core can lead to collapse of the entire system. Here, as a first step, we present a series of simple 1D models that focus on the structure of the gas atmospheres near cores that have a range of core properties (masses and luminosities). Our primary purpose is to address the following: (a) the conditions necessary for a core within a self-gravitating fragment to prompt dynamical collapse of the core's atmosphere; (b) the critical core's mass at which this collapse happens, and (c) the likelihood of these conditions to be met by fragments born from GIs within a reasonable protoplanetary disc setting. We present our numerical methods in section 2, and show the results of a case study for a particular clump mass in section 3. We then explore in section 4 the parameter space for which cores can drive the collapse of their atmospheres. Our results are summarised and their implications are discussed in section 5. | \subsection{Main results of the paper}\label{sec:main} We used a series of 1D simple hydrostatic structure models to explore the stability of heavy element-rich atmospheres surrounding cores that are embedded deep within gaseous clumps. The clumps are envisaged to have formed through the fragmentation of gravitationally unstable protoplanetary discs, while the cores are assumed to have formed through grain sedimentation. Our calculations show that the core's atmosphere, i.e., the gas immediately surrounding the core, can become unstable and collapse for a range of fragment and core masses, as well as core luminosities. While the collapse of a core's atmosphere in and of itself does not immediately imply that the entire clump will be unstable, experiments with the 1D radiation hydrodynamical code presented in \cite{Nayakshin10b,Nayakshin10c} did lead to the whole gas clump collapsing on a dynamical time scale for a few test problems explored. Since such a clump collapse is driven by H2 molecule dissociation and H atom ionisation presenting very large energy sinks it seems reasonable to expect the whole clump collapsing once core atmosphere/envelope becomes unstable. This suggests that a disc fragment could form a gas giant planet via a novel channel that is discovered here. Namely, traditionally giant planet formation by gravitationally unstable discs is believed to occur via radiative cooling and contraction of self-gravitating gas fragments \citep[see the review by][]{HelledEtal13a}. Dust plays a passive (albeit important) role in this traditional picture by providing and regulating the opacity of the fragment \citep[e.g.,][]{HB10}. In the picture developed in our paper, however, the dust plays a dynamically important role. By accumulating into a massive dense core in the centre, dust may provide a significant destabilising effect onto the whole gas fragment. This newly discovered core-assisted gas capture instability (CAGC) is closely linked to the well known Core Accretion instability \citep{MizunoEtal78,Harris78,Mizuno80,Stevenson82,PollackEtal96,Rafikov06}. To enable a quantitative study of the CAGC, we made a number of model choices and simplifying assumptions. The gas fragment is assumed to be isolated from external influences and its structure is approximated by a very simple analytical model of \cite{Nayakshin10c}. The dense core has a fixed density and its luminosity is given by the binding energy divided by a cooling time, $t_{\rm kh}$, which is an important free parameter. The region closest to the core is metal polluted in our model since grains vaporise when $T> T_{\rm v} = 1400$~K. We argued that gas fragment collapse can be trigerred only if metallicity of the polluted layer is $z\simlt 0.5$ as opposed to the limit $1-z\ll 1$. Within these assumptions, our main quantitative results are as follows: \begin{enumerate} \item The critical core mass, $M_{\rm crit}$, that triggers the collapse of a core's atmosphere is dependent most strongly on the central density and temperature of the fragment, as well as the luminosity of the core. More luminous cores (which corresponds to shorter $t_{\rm kh}$) have more tenuous atmospheres, as expected, and hence brighter cores must be more massive to provoke the collapse (see fig. \ref{fig:param1}). As a numerical example, in a 2 M$_J$ clump of $10^4$ years of age from its birth, a core mass greater than $\sim 45 \mearth$ is required to cause collapse if the cooling time of the core is $10^4$ yr. ``Only'' $M_{\rm crit} \approx 12.5 \mearth$ is however needed for a cooling time of $t_{\rm kh} = 10^7$ yr. \item Since protoplanetary discs seem to disappear on time scale of order a few Million years, the most interesting value of $t_{\rm kh}$ to explore is of order 1 Million years or less (see \S \ref{sec:di} below on this point). The critical core mass is then (fig. \ref{fig:param1}) between $\sim 5$ and $\sim 50\mearth$, depending on properties of the gas fragment. \item Less massive and younger gas fragments require smaller $M_{\rm crit}$ at the same $t_{\rm kh}$ (cf. figures \ref{fig:param1} and \ref{fig:age}). \item Atmospheres of bright (short $t_{\rm kh}$) cores are convective, whereas atmospheres of dim (long $t_{\rm kh}$) cores are at least partially radiative. For our default opacity, the atmospheres become completely convective at cooling times of $\simlt 10^4$ yr. In this case, the critical core mass is a function of the central pressure and temperature in the clump. At a given clump age and a clump opacity law, these are a function of the total clump mass, so the critical core mass becomes insensitive to the core luminosity/cooling times in the limit of shortest $t_{\rm kh}$ (cf. the upper enveloping (red) curve in figure \ref{fig:param1}). \item $M_{\rm crit}$ is very weakly dependent on the core density and opacity in the envelope (\S\S 5.1 and 5.3). \item The collapse of gas fragments to much higher proto-planetary densities via CAGC is strongly dependent on the availability of heavy elements, as larger critical core masses can only be reached for high initial metallicity. {\it Therefore, based on these results, a strong positive correlation between frequency of giant planets and metallicity may be expected with CAGC}. Low metallicity environments may be dominated by clumps that undergo collapse through H$_2$ dissociation. \end{enumerate} \subsection{Implications for the disc instability model}\label{sec:di} In the disc instability paradigm \citep[e.g.,][]{Kuiper51b,Boss97,DurisenEtal07}, gas fragments of a few Jupiter masses are born in the outer cold and massive disc due to self-gravity instability. These fragments then cool and eventually collapse via hydrogen molecule dissociation in the fragment \citep{BodenheimerEtal80}. Rapid inward migration of clumps \citep[e.g.,][]{VB06,ChaNayakshin11a,BaruteauEtal11,MichaelEtal11,ZhuEtal12a} may however lead to their destruction due to tides or irradiation \citep{CameronEtal82,VazanHelled12} or tidal disruption \citep{BoleyEtal10,Nayakshin10c}. It is hence only those clumps that contracted before they were disrupted that may leave behind giant planets. The newly found CAGC instability is an additional channel via which gas fragments may collapse. Dust growth and sedimentation into a core can be faster than the time taken by the fragments to cool to H$_2$ dissociation, especially for fragments less massive than a few Jupiter masses \citep{HS08,HB10,Nayakshin10a,Nayakshin10c}. Furthermore, such fragments are also most vulnerable to external irradiation which slows down theor cooling even further \citep{VazanHelled12}. Therefore, CAGC may be the dominant channel through which such relatively low-mass fragments collapse to become giant planets. We caution however that much more work is needed to test these ideas. Adding a self-consistent evolutionary model for the internal structure of a fragment interacting with the disc and the protostar and migrating inward is essential. In addition, there are several physical processes such as turbulence and magnetic fields which are not included in this work. These processes could affect the formation of clumps by gravitational instability and should be included in future disc and planet formation models. Turbulence could also slow down formation of cores inside the fragments. \subsection{Implications for observations of giant planets}\label{sec:obs} \begin{enumerate} \item Presence of a massive core. Based on the older variants of disc instability models \citep[e.g.,][]{Boss97} that do not include planet migration and tidal stripping of the fragments, it is tempting to use the presence of a core in a gas giant as {\it prima facie} evidence of formation by core accretion plus gas capture, as core accretion {\it requires} the formation of a core \citep[e.g.,][]{Mizuno80,Stevenson82,PollackEtal96}, while fragmentation by disc instability does not. However, as emphasized here, core formation could occur nonetheless through grain sedimentation (e.g., Boss 1997, 1998; Helled \& Schubert 2008; Helled et al.~2008; Nayakshin 2010; Boley et al. 2010, 2011; Forgan \& Rice 2013), and more importantly, if it does, can promote the collapse of the fragment if the core's mass exceeds $M_{\rm crit}$. Sedimentation can hence erase or severely confuse using internal structure as evidence for a given formation mode, particularly if the formation of a heavy element core in a fragment can also lead to variations in final bulk composition, as discussed next. Since the presence of a core can be explained by both formation scenarios it becomes harder to discriminate among the two models based on the core's mass. While it is still widely believed that small cores are more consistent with disc instability while massive cores are a natural outcome of core accretion, we found in this paper that reality may be much more complex. Even the extreme cases (i.e., no core at all versus a very massive core) do not implicate the formation process directly. This is due to the fact that very massive cores can also form in the disc instability picture if enrichment from birth and/or planetesimal capture occur, while no or very small cores are not inconsistent with core accretion if core erosion takes place. We therefore suggest once more, that core mass should not be used as a proof for a given formation model. In addition, one has to keep in mind that core masses are inferred only indirectly from models of their interior stucture \citep[e.g.,][]{MillerFortney11,PodolakHelled12}. For extrasolar giant planets it is not currently possible to distinguish between the total heavy element enrichment and high core masses \citep[e.g.,][]{HelledEtal13a} \item Planet composition. Planets born by disc instability are often assumed to have the same composition as their host stars. However, several processes can occur that could potentially alter this \citep[e.g., see review by][]{HelledEtal13a}. As discussed above, the formation of a core in a clump must simultaneously deplete the clump's envelope of heavy elements. If core-assisted gas capture has occurred prior to the loss of the clump's envelope during disc migration, then any gas that is lost to the disc through tides would be preferentially depleted in heavy elements. Any remaining planet could thus appear to be significantly enriched in heavy elements. \item Giant planet frequency of occurence in the inner fraction of an AU was found to be a strongly increasing function of metallicity of the host star \citep{FischerValenti05}. This has been traditionally interpreted as evidence for formation of these planets via Core Accretion picture. Our results however show that metal rich fragments born by the disc instability are more likely to collapse via CAGC, and it is thus possible that the observed correlation could be recovered by this model as well (cf. \S \ref{sec:param} and point (vi) in \S \ref{sec:main}). This result is potentially exciting, but further work is required to better understand the dependence of CAGC on metallicity. On one hand, high heavy element abundances implies higher opacities and longer overall clump cooling times and perhaps larger cores \citep{Nayakshin10a}. On the other hand, high initial abundances could prompt prodigious grain growth. While this should promote a faster core growth, it could also actually decrease the overall cooling time of the clump \citep[e.g.,][]{HB11}. The uncertainty in opacity due to grain growth is however also a concern for the CA model \citep[e.g.,][]{Movshovitz2010}. Despite these uncertainties, it seems unlikely for CAGC to be important for clumps more massive than a few to $\sim 10$ Jupiter masses, because such clumps are unable to form massive solid cores quicker than they heat up to vaporise their grains \citep{HS08,Nayakshin10a}. Therefore, even under most optimistic assumptions about the CAGC contribution to forming giant gas planets, it seems highly unlikely that CAGC could explain planet-metallicity correlation for planets more massive than $\sim 10$ Jupiter masses. \end{enumerate} \subsection{Implications for observations of sub-giant planets}\label{sec:hot} ``Modern'' versions of the disc instability model, e.g., those that include core formation, fragment migration and disruption \citep[e.g.,][]{BoleyEtal10,Nayakshin10c} are multi-outcome valued due to a large number of physical processes operating. Due to these processes, results obtained here may have important implications for close in ``hot'' (planet-star separation $a\sim 0.03 0.3$ AU) sub-giant planets as well. In the context of such modern models, these close-in planets could have formed by the same process of grain sedimentation within the gas fragments, but in one of the following two distinctly different scenarios: \begin{enumerate} \item Hot rocky super-Earths or hot-neptune planets of mass $M_p$ are assembled by grain sedimentation inside gas fragments {\it that are disrupted before H$_2$ dissociation} while $M_p < M_{\rm crit}$. If $M_p \ll M_{\rm crit}$, only a tiny gas atmosphere would be retained after the disruption. If $M_p \simlt M_{\rm crit}$, a more massive atmosphere is retained, so that the planet may be neptune-like in terms of bulk properties and composition. Tidal disruption of pre-H$_2$ dissociation fragments can occur at distances $R\simgt 1$ AU to tens of AU, depending on the internal state of the fragments. For reference, tidal density around the Sun is $\approx 10^{-7}$~g~cm$^{-3}$. Cores are much denser, so they survive this disruption \citep{BoleyEtal10}. If they migrate inward via type I migration to $a\sim 0.1$~AU then they could potentially explain the population of close-in sub-giant exoplanets. \item In the other scenario, H$_2$ dissociation and fragment collapse do take place before the clump is tidally disrupted. This could be either because $M_p \geq M_{\rm crit}$ (collapse via CAGC instability studied here) or because the fragment contracted sufficiently strongly for H$_2$ dissociation to take place \citep[as in the classical picture of, e.g.,][]{BodenheimerEtal80}, when $M_p < M_{\rm crit}$. The key point here is that such a collapse still does not necessarily implies the survival of the giant planet for astrophysically interesting time scales. For this to be the case, the planet must stop migrating towards the star. In the opposite case, the whole giant planet may be completely lost when it is swallowed by the parent star, or, alternatively -- if the planet contracts slower than it migrates inward -- it may be tidally disrupted in the ``hot region'' \citep{Nayakshin11b}, $a\sim 0.1$ AU. In this last case just the core is left behind. In contrast to the picture discussed in (i) above, the migration of the planet in this case is via type II until the gas fragment is disrupted. \cite{NayakshinLodato12} showed that this close-in tidal disruption process reproduces some close in features of FU-Ori outbursts seen around some young stars \citep{HK96}. \end{enumerate} | 14 | 3 | 1403.1813 |
1403 | 1403.1036_arXiv.txt | We report radio observations of two stripped-envelope supernovae (SNe), 2010O and 2010P, which exploded within a few days of each other in the luminous infrared galaxy Arp\,299. Whilst SN\,2010O remains undetected at radio frequencies, SN\,2010P was detected (with an astrometric accuracy better than 1\,milli arcsec in position) in its optically thin phase in epochs ranging from $\sim1$ to $\sim3$\,yr after its explosion date, indicating a very slow radio evolution and a strong interaction of the SN ejecta with the circumstellar medium. Our late-time radio observations toward SN\,2010P probe the dense circumstellar envelope of this SN, and imply $\dot M [\ml{}] / v_{\rmn{wind}} [10\,\kms{}] =(3.0$--5.1)$\ee-5$, with a 5\,GHz peak luminosity of $\sim1.2\ee27$\,\lunits{} on day $\sim$464 after explosion. This is consistent with a Type IIb classification for SN\,2010P, making it the most distant and most slowly evolving Type IIb radio SN detected to date. | \label{sec:discussion} The optical spectroscopy reported in \citeauthor{kankare13} indicates that 2010O is a Type Ib SN, which are known to be bright at radio frequencies. In fact a peak luminosity $<3\ee26$\,\lunits{} (see Figure \ref{fig:LumTime}) is quite uncommon for Type Ib/c SNe \citep{soderberg07}. Additionally, these SNe evolve very rapidly and thus reach their peak luminosity within a few tens of days \citep[e.g., SN\,2008D,][]{soderberg08}, so the most likely explanation for the radio non-detection of SN\,2010O is that it reached its peak some time after the early-time MERLIN observations by \citeauthor{beswick10}, and well before our observations in 2011, so it went unnoticed owing to the lack of prompt radio observations. SN\,2010P lies in a location with larger extinction than seen toward SN\,2010O ($A_V =7$\,mag, \citeauthor{kankare13}) and at a $\approx 160$\,pc projected distance from the component C$^{\prime}$, SN\,2010P is the first radio SN detected in the outskirts of a bright radio emitting component in Arp\,299, despite the intense radio monitoring. Its strong radio emission 1.4 to 2.8\,yr after explosion indicates a strong interaction between the SN ejecta and the CSM. Whilst the early optical spectrum is consistent with the SN being either a Type Ib or IIb (\citeauthor{kankare13}), the late-time radio detection rules out a Type Ib origin. The parameters we obtained to fit the radio light curve of SN\,2010P are common for luminous Type II SNe (see Table \ref{tab:SNcomp}). Whilst the inferred peak luminosity and the time to reach the peak are more comparable to those of Type IIL's, its estimated mass-loss rate to wind velocity ratio, $\dot M/v_{\rmn{wind10}}=(3.0$--$5.1)\ee-5$, and its inferred deceleration parameter ($\sim0.74$), support better a Type IIb nature with a slow evolution, and are consistent with the early optical spectra lacking strong hydrogen features (\citeauthor{kankare13}). We note that some Type IIn SNe have comparable mass-loss rate to wind velocity ratios (see Table \ref{tab:SNcomp}), however, these SNe are also an order of magnitude more luminous than SN\,2010P at their peak. It has been suggested that Type IIb SNe can be divided into two broad groups \citep{chevalier10} under the assumption that SSA is the main absorption mechanism: i) slowly evolving SNe at radio frequencies ($\tpeak>100$\,d) with a large ($>0.1\,\msun{}$) hydrogen mass envelope and slow ejecta velocity ($\sim 10,000\,\kms{}$), expected from an extended progenitor ($R\sim10^{13}$\,cm); ii) rapidly evolving SNe ($\tpeak<100$\,d) with a small hydrogen mass envelope ($<0.1\,\msun{}$) and fast ejecta velocities (a factor of 3-5 larger than in case i), expected from a compact progenitor ($R\sim10^{11}$\,cm). According to this interpretation, SN\,2010P would come from an extended progenitor since it took hundreds of days to reach its peak (in both FFA and SSA fits). In \citeauthor{kankare13} we do not find evidence for a large hydrogen mass envelope (i.e., $>0.1\,\msun{}$) around SN\,2010P, in contradiction with the expectations for an extended progenitor as inferred from its radio behaviour. The properties of SNe 2011dh and 2011hs where more direct evidence of the progenitor star has been obtained \citep[][respectively]{bersten13, bufano14}, further suggest that the inferred radio properties of a SN are not a strong beacon of the progenitor size. This could be related to the assumption of SSA dominance for poorly sampled light curves. In the case of SN\,2001gd, the assumptions made to fit its radio light curve greatly affected the estimate for the time it took the SN to reach its peak at 5\,GHz, going from 173\,d in the FFA model \citep{stockdale03}, to 80\,d in the SSA$+$FFA model \citep{stockdale07}, thus placing this SN in the compact IIb category, following the interpretation from \citet{chevalier10}. Unlike in the SN\,2001gd case, we do not have early data for SN\,2010P describing its turn-on phase at any frequency, and thus, we cannot quantify the presence of SSA in its evolution. However, it is rather unlikely that SSA will dominate the absorption for SN\,2010P radio emission when the SN has been detected at such late ages. Thus, although the $\tpeak{}$ we estimated for SN\,2010P represents an upper limit owing to our assumption of FFA being the only absorption mechanism, we have seen that assuming an SSA model pushes the peak at different frequencies to even later times and the same interpretation holds. However, there is another possible scenario which we cannot exclude. SN\,2010P could have reached a first peak at early times not sampled by our data, and we could in fact be witnessing a re-brightening of the SN at later times. In this situation, \tpeak{} would have occurred earlier and SN\,2010P would appear closer to other Type IIb SNe in the \lpeak{} vs. \tpeak{} plot. Observationally, strong variations in the optically thin phase of different types of CCSNe are rather common \citep[see table 4 in][]{soderberg06} owing to the complexity of their CSM. For Type IIb SNe, it has been shown that variations in the mass-loss of the progenitor (e.g., luminous blue variable-like stars), can create inhomogeneities in the CSM and thus modulations in the SN radio light curve \citep[e.g.,][]{kotak06,moriya13}. The sparse sampling of the SN\,2010P light curve does not allow us to conclusively determine whether our late-time observations match with the forward shock encountering a first, or a second high-density region. Multiple high-density regions can be produced by mass-loss episodes, meaning variations in the mass-loss rate throughout the lifetime of the progenitor star. However, the early optical/NIR data (\citeauthor{kankare13}) do not show evidence for interaction with the CSM at the early stages of the SN, i.e., the interaction with the CSM corresponding to the end of the progenitor star's lifetime. Therefore, although we cannot rule out completely the possibility of the peak radio luminosity corresponding to a secondary mass-loss episode, the observations (both radio and NIR) strongly suggest that we are witnessing a slow radio evolution of SN\,2010P. | \label{sec:concl} We report radio observations towards SNe 2010O and 2010P in the LIRG Arp\,299. SN\,2010O was not detected throughout our radio monitoring of the host galaxy which lacked observations of the SN at ages between 0.6\,month and 1.4\,yr, time enough for a Type Ib SN to have have risen to a peak, then decayed. SN\,2010P was detected at various frequencies in its transition to, and in its optically thin phase from $\sim$1 to $\sim$3\,yr after explosion. Our observations favour FFA as the dominant absorption mechanism controlling the radio emission of this SN. We characterise it as a luminous, slowly-evolving Type IIb SN with $\dot M [\ml{}] / v_{\rmn{wind}} [10\,\kms{}]=(3.0$--5.1)$\ee-5$. We have also improved on the coordinates previously reported for SN\,2010P with a position accuracy better than 1 mas at $\alpha(J2000) = 11\h28\m31\fs3605$, $\delta(J2000) = 58\degr33\arcmin49\farcs315$. SN\,2010P is one of a select group of 12 Type IIb SNe detected at radio wavelengths: for nine of them a radio light curve has been obtained (see references in Figure \ref{fig:LumTime}), and three of them have only reported radio detections \citep[SNe 2008bo, PTF\,12os and 2013ak reported in][respectively]{stockdale08, stockdale12, kamble14}. SN\,2010P is also the most distant Type IIb SN detected so far, and the one that according to our data, has taken the longest time to reach its peak. However, a comprehensive, multi-wavelength study covering both rise and decline of the SN emission is needed to investigate possible progenitor star scenarios. We note that the early epochs after shock breakout are crucial, since this when when we can gather more information about the progenitors and the dominant absorption mechanism (SSA/FFA) shaping the radio light curves, thus helping us to better understand the CCSN phenomenon, and thereby the interplay between massive stars and their CSM. | 14 | 3 | 1403.1036 |
1403 | 1403.4339_arXiv.txt | We estimate the abundance of dust in damped Lyman-$\alpha$ absorbers (DLAs) by statistically measuring the excess reddening they induce on their background quasars. We detect systematic reddening behind DLAs consistent with the SMC type reddening curve and inconsistent with the Milky Way type. \textbftemp{We find that the derived dust-to-gas ratio is, on average, inversely proportional to the column density of neutral hydrogen, implying that the amount of dust is constant, irrespective of the column density of hydrogen.} \textbftemp{It means that} \textbftemp{the average metallicity is inversely proportional to the column density of hydrogen,} \textbfnew{unless the average dust-to-metal ratio varies with the hydrogen column density.} This indicates that the prime origin of metals seen in \textbftemp{DLAs} is not by {\it in~situ} star formation, with which $Z\sim N_{\rm HI}^{+0.4}$ is expected from the empirical star formation law, contrary to our observation. We interpret the metals observed in absorbers being deposited dominantly from nearby galaxies by galactic winds ubiquitous in intergalactic space. When extrapolating the relation between \textbftemp{dust-to-gas ratio and \HI\ column density to lower column density, we find a \textbftemp{value} which is consistent with what is observed for \MgII\ absorbers.} | \label{sec:introduction} Whether intervening absorbers seen in quasar spectra are aggregates of primordial material or results of the activity in galaxies is an elementary problem. In a previous publication (M\'enard \& Fukugita 2012; hereinafter MF12), it was advocated that Mg II clouds are likely to be a product of the activity of nearby galaxies with gas exported by galactic winds. This inference is based on the fact that the observed dust abundance of \MgII\ clouds relative to gas takes a value typical of galactic disks, while the star formation activity is not observed nor expected in such clouds. MF12 estimates that the global \HI\ abundance in \MgII\ clouds is $\Omega_{\rm HI}({\rm MgII})\approx 1.5\times 10^{-4}$, which is approximately 3\% of the fuel consumed by star formation by the present epoch, or roughly 6\% at $z\approx 2$. The amount of matter expelled through galactic winds in actively star forming galaxies \textbftemp{is often inferred to be comparable to} the star formation rate (e.g., Heckman et al. 2000; Veilleux et al. 2005; Weiner et al. 2009). The fraction of \HI\ in \MgII\ absorbers, which is 10 -- 20 \% of the gas re-shed by stars as a whole, is not an unreasonable amount as a product resulting from the star formation activity in galaxies. It is also shown that dust in \MgII\ clouds accounts for half the amount of dust estimated to reside outside galaxies. The analysis gives the example that the dust abundance, as explored by extinction of light rays passing through the absorbers, provides us with a useful indicator of the heavy element abundance, assuming that photometry is accurate. This suggests that more could be learned from dust studies \textbftemp{of} other classes of absorbers. In this paper we focus on damped Lyman-$\alpha$ absorbers (DLAs). The global mass density of \HI\ in DLA clouds has been estimated to be $\Omega_{\rm HI}({\rm DLA})\approx (4-10)\times 10^{-4}$ at $z\approx 2$ by Prochaska \& Wolfe (2009) using the quasar spectroscopic data of SDSS DR5, and more recently by Noterdaeme et al. (2012) using SDSS DR9. This quoted range \textbftemp{ arises from different treatments of the continuum around the damped Lyman $\alpha$ line, as well as survey path length and completeness. } Whichever value is taken, the DLA mass density is significantly larger than the \HI\ mass density in the \MgII\ absorbers, which is estimated to be $\Omega_{\rm HI}({\rm MgII})\approx 1.5\times 10^{-4}$ (MF12). \textbftemp{\HI\ mass density is clearly larger than that can be associated with stellar activity.} Studies of individual DLAs have indicated their metallicity [Fe/H] to be in the range $-0.5$ to $-2$ (e.g., Prochaska et al. 2003; Rafelski et al. 2012), an order of magnitude lower than that of galaxies. This contrasts \textbftemp{with} \MgII\ absorbers, which MF12 showed to have metallicity of the order of solar. These observations may be taken in favour of the interpretation that there would be two distinct populations of absorbers as to their origin: DLAs belonging to one and \MgII\ absorbers to another. This also induces the question as to where the metallicity of DLAs arises from. There have been a few attempts to detect reddening of quasars behind DLAs (\textbftemp{Fall et al. 1989, Pei et al. 1991, Murphy \& Liske 2004,} Vladilo et al. 2008; Frank \& P\'eroux 2010, Khare et al. 2011). The results have not always led to positive detections. The difficulty is that the reddening signal is small, of the order of a few hundredths of magnitude, whereas sample variation due to objects is an order of magnitude larger. For photometric studies, one needs to accurately define \textbftemp{the mean colour of reference quasars}. With spectroscopic work one needs accurate sensitivity calibrations over \textbftemp{a wide wavelength range}. In this paper we estimate the mean reddening effects induced by DLAs by comparing broad band flux of the quasar light showing Lyman $\alpha$ absorptions with damped wings to that without absorbers. We use accurately calibrated SDSS broad band photometry, and we limit the redshift range to minimise the scatter of fiducial quasar colours. We are concerned with a photometric accuracy smaller than 0.1 mag. It turns out that choosing the right range of \textbftemp{both quasar and absorber redshifts} is important in keeping the errors of colours small. Specifically, care must be made so that passbands are away from the Lyman edge or the Lyman $\alpha$ line for absorbers. \textbftemp{We study whether the dust-to-gas ratio of DLAs is correlated with their hydrogen column density, and also examine if any difference is seen between \MgII\ absorbers and DLAs in the metallicity \HI\ column density relation.} | \textbftemp{We have studied the dust-to-gas ratio of DLAs and used it as a metallicity indicator.} The advantage is that the dust abundance can be estimated passively from the extinction of light, provided that care is made to measure a small value of reddening. \textbftemp{To keep the accuracy controlled we limited our analysis to carefully chosen redshift ranges for both quasars and absorbers.} Using the fact that 30\% of heavy elements condense into dust grains, metallicity can be estimated from reddening in broad band photometry. This requires averaging over a large data set and well-controlled photometry that does not suffer from arbitrary errors. The advantage is that the estimate of metallicity, using only passive measurements, does not use the temperature, or does not depend on the environment such as the radiation field. Our central conclusion is that average metallicity is inversely proportional to the column density of the DLA, as $Z\sim N_{\rm HI}^{-1}$, or in other words, the metal column density stays constant independent of the hydrogen column density. We argued that {\it in~situ} star formation should lead to $Z\sim N_{\rm HI}^{+0.4}$, which contrasts to the observation of the inverse correlation. We have not observed any threshold in the hydrogen column density that is known for the law of star formation (Kennicutt 1998; Wyder et al. 2009). These aspects lead us to conclude that it is unlikely to ascribe the origin of the bulk of dust in DLAs to {\it in~situ} star formation. We argued \textbftemp{that, instead, } dust in DLAs is \textbftemp{a} deposit from intergalactic space through stellar activity \textbftemp{in the neighbourhood of} the cloud. \textbftemp{In this case, we expect}, on average, the same amount of dust deposited per surface area of intergalactic clouds irrespective of the column density. It then follows that $Z\sim N_{\rm HI}^{-1}$. We discussed that this view does not bring a problem into the dust budget consideration. About 10\% of DLAs show \MgII\ absorption features at $z\sim 2$. DLAs that do not show strong \MgII\ absorption features harbour 1/3 the amount of dust compared with those that show strong \MgII\ absorption. As a result, dust in DLAs as a whole resides more in those that do not show \MgII\ absorption features. The global amount of dust in DLAs, however, is yet 1/3 the amount in \MgII\ clouds, which are mostly at lower, LLS column densities. A corollary result from our study is that dust in DLA shows reddening consistent with the SMC type extinction curve. For the redshift we chose ($z\approx 2.2$) the Milky Way type extinction would lead to no colour excess or even blueing in $g-i$ colour in the presence of dust due to the 2175\AA~feature. We detected reddening when DLA is present in the foreground, and metallicity derived using the SMC type extinction curve agrees statistically with spectroscopic estimates. This rules out the Milky Way type extinction for the DLA. DLAs seem to be aggregates of primarily unprocessed gas with small amount of deposits from galaxy activity in their vicinity. \textbftemp{In constrast,} \MgII\ clouds are consistent with secondary products of galaxies, \textbftemp{weakly} diluted with pristine gas. However, there seems to be no clear dichotomy in the metallicity hydrogen column density relation between the two populations. \MgII\ clouds should thoroughly be contaminated, or even dominated by galactic wind, if they are of the primordial origin. There seems to be a significant population of hydrogen clouds with hydrogen column density comparable to \MgII\ clouds but do not show \MgII\ lines: they seem to be a population similar to DLAs at a lower column density extension. How \MgII\ clouds formed remains as an interesting problem. | 14 | 3 | 1403.4339 |
1403 | 1403.3343_arXiv.txt | {% The emergence of optical interferometers with three and more telescopes allows image reconstruction of astronomical objects at the milliarcsecond scale. However, some objects contain components with very different spectral energy distributions (SED; i.e. different temperatures), which produces strong chromatic effects on the interferograms that have to be managed with care by image reconstruction algorithms. For example, the gray approximation for the image reconstruction process \modif{results in a degraded} image if the total \uv-coverage given by the spectral supersynthesis is used. } {The relative flux contribution of the central object and an extended structure changes with wavelength for different temperatures. For young stellar objects, the known characteristics of the central object (i.e., stellar SED), or even the fit of the spectral index and the relative flux ratio, can be used to model the central star while reconstructing the image of the extended structure separately.} {We present a new method, called \texttt{SPARCO} (semi-parametric algorithm for the image reconstruction of chromatic objects), which describes the spectral characteristics \modif{of both the central object and the extended structure} to consider them properly when reconstructing the image of the surrounding environment. We adapted two image-reconstruction codes (\texttt{Macim}, \texttt{Squeeze}, and \texttt{MiRA}) to implement this new prescription.} {\texttt{SPARCO} is applied using \texttt{Macim}, \texttt{Squeeze} and \texttt{MiRA} on a young stellar object model and also on literature data on HR 5999 in the near-infrared with the VLTI. We obtain smoother images of the modeled circumstellar emission and improve the $\chi^2$ \modif{by} a factor \modif{9}. % } {This method paves the way to improved aperture-synthesis imaging of several young stellar objects with existing datasets. More generally, the approach can be used on astrophysical sources with similar features such as active galactic nuclei, planetary nebulae, and asymptotic giant branch stars.} | The number of aperture-synthesis images based on optical long-baseline interferometry measurements has recently increased thanks to easier access to visible and infrared interferometers. The interferometry technique has now reached a technical maturity level that opens new avenues for numerous astrophysical topics that require milliarcsecond model-independent imaging \citep{2012A&ARv..20...53B}. Image reconstruction \citep[see review][]{2013EAS....59..157T} is the key to achieve the most probable, noncommittal images following some global constraints (image positivity, size of the support, regularization, etc.). A thorough study \citep{Renard} has shown the limitations of image reconstruction, but never challenged the type of regularizations in use. The first images of the inner regions of the environment have been obtained around the young stellar objects HD\,163296 \citep{HD163296} and HR\,5999 \citep{HR5999}, revealing the structure around the central objects. However, one main caveat of the images reconstructed from spectrally dispersed instruments is that the visibilities measured at different wavelengths have been assumed to come from a \emph{gray} object. In addition, the central star has a much higher surface brightness than the surrounding emission, a contrast problem that limits the reliability and the quality of image reconstruction and complicates serious analysis of the morphology of the circumstellar material. These concerns motivate one to reconstruct an image of the envelope alone \modif{without the star in the image}, an approach considered first in self-aperture masking techniques \citep{2004ApJ...605..436M}. Some of the interferometric visibilities obtained with spectral resolution on young stellar objects (YSO), mainly in the near-infrared (NIR), have been seriously affected by strong spectral dependence. For example, the visibilities of the object MWC\,158, as measured with VLTI/AMBER in the $K$ band and in the $H$ band are spread over a broad range \citep[see Fig.~5 of][]{2011AA...528A..20B}. First believed to be \emph{quality (...) clearly lower compared to the K-band data} \citep{2011AA...528A..20B}, the increase of the visibilities of the $H$-band data has been observed not only on VLTI/AMBER, but also with VLTI/PIONIER. This indication caused us to consider an astrophysical interpretation: the change of visibility is generated by the different chromatic behavior between the central star that peaks in the visible and the circumstellar material that radiates out mainly in the near-infrared \citep{2012SPIE.8445E..0OK}. The star contribution is highest at the shortest wavelengths and becomes moderate at longer wavelengths. Classic gray algorithms assume the same brightness distribution for all wavelengths. Consequently, they cannot satisfactorily reconstruct an image of a chromatic object with spectral supersynthesis. We need to include the chromatic effect induced by the physical properties of the target in the image processing to retrieve a good intensity map of the observed target. In this paper, we present a semi-parametric approach that includes the knowledge of the \modif{relative stellar and environment spectral} properties in optical interferometry image reconstruction. As previously demonstrated in parametric modeling of optical interferometric data (e.g. \citet{2012ApJ...752...11K} ), we can directly take the object chromaticity into account in the process of image reconstruction, which improves the final fit. This approach is called SPARCO (semi-parametric approach for image reconstruction of chromatic objects) and consists of separating a well-known object (e.g. the central star of a YSO) and its complex, unknown environment (e.g. its dusty disk). The star is modeled by a parametric model (that can include hydrostatic models or binaries) and the environment by the reconstructed image. The chromatism is reproduced by changing the flux ratio between the two components across the observed bandwidth. To present the methods, we focus on YSOs because the star can be modeled at first order in the NIR by an unresolved component and the flux ratio \modif{can be represented by a power law with a good approximation}. Moreover, the environment is poorly known and is complex. The application of this technique to this type of objects is therefore important. However, this method does not apply only to YSO, but can also be used in any system where a known source is present that displays a spectral behavior different from the rest of the emitting material in the optical. For instance, the accretion disk of an active galactic nuclei is considered as unresolved in order to retrieve its environment \citep{2013ApJ...775L..36K}. A method of separation of the star from its environment was invoked for asymptotic giant branch stars \citep{2013A&A...559A.111H}, or planetary nebulae \citep{2006A&A...448..203L}. In Sect.~\ref{sec:sep} we demonstrate that the \texttt{SPARCO} approach allows proper modeling of the interferometric observables of YSOs, especially their chromatic content. We show how standard image reconstruction algorithms can be modified accordingly. In Sect.~\ref{sec:imgrec} we validate the method on the model of a realistic YSO. In Sect.~\ref{sec:dis}, we further discuss important aspects of the method in detail. We finish in Sect.~\ref{sec:HR5999} by applying \texttt{SPARCO} on actual data used by \citet{HR5999} to reveal the circumstellar environment of HR\,5999, and we compare it with a previous analysis. | \label{sec:sum} \modif{For image reconstruction based on interferometric data with spectral dispersion, visibilities are determined not only by the geometry of the object, but also by the (differential) spectral slope of its components. } We developed a method that includes the knowledge that we have on \modif{the relative spectral behaviors of two components} to reconstruct the intensity distribution of the \modif{extended} one. This method allowed us to improve the $\chi^2$ by one order of magnitude in our validation. The \texttt{SPARCO} method includes an analytical description of the stellar contribution and the chromatic ratio between the star and its environment in the image-reconstruction algorithms used in optical interferometry. The first component is modeled parametrically and an image is reconstructed for the second one. \texttt{SPARCO} was used correctly on a young stellar object with the following hypothesis: \begin{itemize} \item \textit{The environment spatial \modif{distribution} is assumed to be wavelength independent.} The chromatic dependence of visibilities only arises because the star has a different spectral dependence from the environment. This also implies that the environment has \modif{the same spectral dependence in the whole image}. \item \textit{The star is very close to the parametric description given in the method.} \item \textit{The fluxes ($F_{\lambda}$) of the components are approximated by power laws as a function of the wavelength. Since the chromatic parameter space can be degenerate when considering interferometric constraints alone, obtaining independent spectrophotometric observations helps to retrieve an image of the observed environment. Despite the low sensitivity of the method to the choice of regularization, one has to be careful when choosing the value of the hyper-parameter $\mu$.} \end{itemize} The method is simple but can be adapted to more complex models. As demonstrated in Sect.~\ref{sec:HR5999}, if the object is more complex some more information needs to be added to the model (add a parametric disk for the inner part of the object). In Sect.~\ref{sec:modif} we showed that we can easily model the parametric part by a uniform disk, a binary, and even an image. This tool is complementary to other observations of the target, especially with spectrophotometric observations. Since the chromatic parameter space can be degenerate, these types of observations are very important to correctly retrieve an image of the observed target. It is difficult to retrieve them by simultaneous fitting. The \texttt{SPARCO} method will be intensively used on a Large Program dataset on Herbig Ae/Be stars gathered by PIONIER at the VLTI. Unraveling the image of the close environment of young stars will help us to constrain the effects of the inclination of the inner regions of YSO and therefore to detect early signs of planet formation very close to the star. This method cannot be applied only to young stars, but can also be used in any system whith a point-like source and that displays a different spectral behavior from the rest of the emitting material in the optical such as active galactic nuclei. \modif{Algorithms that implement a fully polychromatic approach are currently being developped, e.g. \texttt{MiRA} \citep{Mira3D} and \texttt{Squeeze} (Baron et al. in prep.). Combining the \texttt{SPARCO} approach with polychromatic reconstruction will allow imaging any stellar environment with limited perturbation from its star, and greatly enhance our capability to study variations in the environment morphology (e.g. temperature gradient in the disk).} | 14 | 3 | 1403.3343 |
1403 | 1403.1346_arXiv.txt | We report short-time variations in the plasma tail of C/2013 R1(Lovejoy). A series of short (two to three minutes) exposure images with the 8.2-m Subaru telescope shows faint details of filaments and their motions over 24 minutes observing duration. We identified rapid movements of two knots in the plasma tail near the nucleus ($\sim$ 3$\times 10^5$ km). Their speeds are 20 and 25 km s$^{-1}$ along the tail and 3.8 and 2.2 km s$^{-1}$ across it, respectively. These measurements set a constraint on an acceleration model of plasma tail and knots as they set the initial speed just after their formation. We also found a rapid narrowing of the tail. After correcting the motion along the tail, the narrowing speed is estimated to be $\sim$ 8 km s$^{-1}$. These rapid motions suggest the need for high time-resolution studies of comet plasma tails with a large telescope. | Plasma tails of comets and their time variations potentially provide crucial information on solar winds and magnetic fields in the solar system \citep[e.g.,][]{Niedner1982,Mendis2006,Downs2013}. Short-time variations in plasma tails, however, are not yet fully understood. Indeed, most previous studies observed tails and structures at far distances ($>10^6$ km from the nucleus) with a time resolution of an order of an hour. Regarding the speed of movement along the tail, \citet{Niedner1981} studied 72 disconnection events (DEs) of various comet tails and found $\lesssim$ 100 km s$^{-1}$ at $\lesssim 10^7$ km from the nuclei. Their initial speeds before DEs are around 44 km s$^{-1}$ and the typical acceleration is 21 cm s$^{-2}$. \citet{Saito1987} analyzed a knot in the plasma tail of comet 1P/Halley and derived its average velocity of 58 km s$^{-1}$ at 4--9 $\times 10^5$ km from the nucleus. \citet{Kinoshita1996} observed C/1996 B2(Hyakutake) and measured the speed of a knot of 99.2 km s$^{-1}$ at 5.0 $\times 10^6$ km from the nucleus. \citet{Brandt2002} investigated DE of C/1995 O1 (Hale-Bopp) and obtained the speed of $\sim$ 500 km s$^{-1}$ at $\sim$ 7$\times 10^7$ km from the nucleus. \citet{Buffington2008} analyzed several knots in comets C/2001 Q4(NEAT) and C/2002 T7(LINEAR) using the Solar Mass Ejection Imager. They found the speed to be 50--100 km s$^{-1}$ around $10^6$ km from the nucleus. These previous studies did not catch the moment immediately after the formation of knots or the detachment of knots from the tail. The initial speed at these critical times was only an extrapolation from later observations relatively far away. In this letter, we report detections of knots in the plasma tail 3$\times 10^5$ km away from the nucleus of C/2013 R1(Lovejoy) and a direct measurement of their initial motions. We adopt the AB magnitude system throughout the paper. | The initial speeds of two moving knots ($\sim$ 22 km s$^{-1}$) are significantly slower than the ones measured by \citet{Niedner1981} (44 km s$^{-1}$; rms of 10.9 km s$^{-1}$). This speed is also smaller than that measured in comet Halley by \citet{Saito1987} (58 km s$^{-1}$), who suggested that the velocity was constant at 4--9 $\times 10^5$ km from the nucleus. Though it is not clear what the dominant factor of the initial speed of knots is, we can compare several parameters of the comets. As the data used by \citet{Niedner1981} include information from various comets, we compare parameters with the single case of the comet Halley by \citet{Saito1987}. At the observation by \citet{Saito1987}, the heliocentric distance to comet Halley was 1.016 au, the heliocentric ecliptic coordinate was ($\lambda$,$\beta$)=(30.1,8.8), and the heliocentric velocity was -26.5 km s$^{-1}$. Compared with C/2013 R1(Lovejoy) in this study, a part of the difference in the initial speed may be explained by the difference in the heliocentric velocity of the nucleus, -12.6 km s$^{-1}$ versus -26.5 km s$^{-1}$, if we assume that the initial speed of the knots might be comparable in heliocentric frame. It, however, does not fully explain $\sim$ 40 km s$^{-1}$ difference. Another difference is heliocentric ecliptic latitude, 30.7 vs 8.8, which may result in the difference in the speed of the solar wind at the comet position. Yet another point is that we have compared the speeds of our relatively faint and small knots with more prominent knots/kinks and DEs in the previous studies. The relevance of this comparison might be debated in light of future studies. In addition, we analyzed only one comet tail observed in a relatively short duration. A more systematic investigation is obviously needed as to the distribution of the initial speed of the tail as a function of the heliocentric velocity, the heliocentric distance, and the ecliptic position of the comet. In summary, we found two knots that were just formed at 3$\times 10^5$ km from the nucleus of C/2013 R1(Lovejoy). Their initial speed was smaller than the ones measured in previous studies, and a physical interpretation requires a more statistically significant sample at various heliocentric positions. We also found a rapid variation in the tail width in seven minutes, which implies a rapid change in ambient solar winds and magnetic field. These results strongly suggest that the variations in comet plasma tails, especially in their fine structures, require high time resolution observations with a large aperture telescope such as Subaru. | 14 | 3 | 1403.1346 |
1403 | 1403.6344.txt | We review the theory and phenomenology of neutrino electromagnetic interactions, which give us powerful tools to probe the physics beyond the Standard Model. After a derivation of the general structure of the electromagnetic interactions of Dirac and Majorana neutrinos in the one-photon approximation, we discuss the effects of neutrino electromagnetic interactions in terrestrial experiments and in astrophysical environments. We present the experimental bounds on neutrino electromagnetic properties and we confront them with the predictions of theories beyond the Standard Model. | \label{A001} The theoretical and experimental investigation of neutrino properties and interactions is one of the most active fields of research in current high-energy physics. It brings us precious information on the physics of the Standard Model and provides a powerful window on the physics beyond the Standard Model. The possibility that a neutrino has a magnetic moment was considered by Pauli in his famous 1930 letter addressed to ``Dear Radioactive Ladies and Gentlemen'' (see \textcite{Pauli:1992mu}), in which he proposed the existence of the neutrino and he supposed that its mass could be of the same order of magnitude as the electron mass. Neutrinos remained elusive until the detection of reactor neutrinos by Reines and Cowan around 1956 \cite{Reines:1960pr}. However, there was no sign of a neutrino mass. After the discovery of parity violation in 1957, \textcite{Landau:1957tp,Lee:1957qr,Salam:1957st} proposed the two-component theory of massless neutrinos, in which a neutrino is described by a Weyl spinor and there are only left-handed neutrinos and right-handed antineutrinos. It was however clear \cite{Touschek:1957,Case:1957zza,Mclennan:1957} that two-component neutrinos could be massive Majorana fermions and that the two-component theory of a massless neutrino is equivalent to the Majorana theory in the limit of zero neutrino mass. The two-component theory of massless neutrinos was later incorporated in the Standard Model of \textcite{Glashow:1961tr,Weinberg:1967tq,Salam:1968rm}, in which neutrinos are massless and have only weak interactions. In the Standard Model Majorana neutrino masses are forbidden by the $\text{SU}(2)_{L} \times \text{U}(1)_{Y}$ symmetry. Although in the Standard Model neutrinos are electrically neutral and do not possess electric or magnetic dipole moments, they have a charge radius which is generated by radiative corrections. We now know that neutrinos are massive, because many experiments observed neutrino oscillations (see the reviews by \textcite{Giunti-Kim-2007,Bilenky:2010zza,Xing:2011zza,GonzalezGarcia:2012sz,Bellini:2013wra,PDG-2012}), which are generated by neutrino masses and mixing \cite{Pontecorvo:1957cp,Pontecorvo:1957qd,Maki:1962mu,Pontecorvo:1968fh}. Therefore, the Standard Model must be extended to account for the neutrino masses. There are many possible extensions of the Standard Model which predict different properties for neutrinos (see \textcite{Ramond:1999vh,Mohapatra:2004,Xing:2011zza}). Among them, most important is their fundamental Dirac or Majorana character. In many extensions of the Standard Model neutrinos acquire also electromagnetic properties through quantum loops effects which allow direct interactions of neutrinos with electromagnetic fields and electromagnetic interactions of neutrinos with charged particles. Hence, the theoretical and experimental study of neutrino electromagnetic interactions is a powerful tool in the search for the fundamental theory beyond the Standard Model. Moreover, the electromagnetic interactions of neutrinos can generate important effects, especially in astrophysical environments, where neutrinos propagate over long distances in magnetic fields in vacuum and in matter. Unfortunately, in spite of many efforts in the search of neutrino electromagnetic interactions, up to now there is no positive experimental indication in favor of their existence. However, it is expected that the Standard Model neutrino charge radii should be measured in the near future. This will be a test of the Standard Model and of the physics beyond the Standard Model which contributes to the neutrino charge radii. Moreover, the existence of neutrino masses and mixing implies that neutrinos have magnetic moments. Since their values depend on the specific theory which extends the Standard Model in order to accommodate neutrino masses and mixing, experimentalists and theorists are eagerly looking for them. The structure of this review is as follows. In Section~\ref{B001} we summarize the basic theory of neutrino masses and mixing and the phenomenology of neutrino oscillations, which are important for the following discussion of theoretical models and for understanding the connection between neutrino masses and mixing and neutrino electromagnetic properties. In Section~\ref{C001} we derive the general form of the electromagnetic interactions of Dirac and Majorana neutrinos in the one-photon approximation, which are expressed in terms of electromagnetic form factors. In Section~\ref{D001} we discuss the phenomenology of the neutrino magnetic and electric dipole moments in laboratory experiments. These are the most studied electromagnetic properties of neutrinos, both experimentally and theoretically. In Section~\ref{E001} we discuss neutrino radiative decay in vacuum and in matter and related processes which are induced by the neutrino magnetic and electric dipole moments. These processes could have observable effects in astrophysical environments and could be detected on Earth by astronomical photon detectors. In Section~\ref{F001} we discuss some important effects due to the interaction of neutrino magnetic moments with classical electromagnetic fields. In particular, we derive the effective potential in a magnetic field and we discuss the corresponding spin and spin-flavor transitions in astrophysical environments. In Section~\ref{G001} we review the theory and experimental constraints on the neutrino electric charge (millicharge), the charge radius and the anapole moment. In conclusion, in Section~\ref{H001} we summarize the status of our knowledge of neutrino electromagnetic properties and we discuss the prospects for future research. This review has also several appendices. We highlight here Appendix~\ref{I001}, in which we clarify the conventions and notation used in the paper and we list some useful physical constants and formulae. Let us also remind that neutrino electromagnetic properties and interactions are discussed in the books by \textcite{Bahcall:1989ks,Boehm:1992nn,CWKim-book,Raffelt:1996wa,Fukugita:2003en,Zuber:2003,Mohapatra:2004,Xing:2011zza,Barger:2012pxa,Lesgourgues-Mangano-Miele-Pastor-2013}, and in the previous reviews by \textcite{Bilenky:1987ty,Dolgov:1981hv,Raffelt:1990yz,Salati:1993tf,Raffelt:1999gv,Raffelt:1999tx,Raffelt:2000kp,Pulido:1991fb,Dolgov:2002wy,Nowakowski:2004cv,Wong:2005pa,Studenikin:2008bd,Giunti:2008ve,Broggini:2012df,Akhmedov:2014kxa}. In this review we improved and extended the discussion presented in our previous reviews in order to cover in details the most important aspects of neutrino electromagnetic interactions. | \label{H001} In this review we discussed the theory and phenomenology of neutrino electromagnetic properties and interactions. We have seen that most of the theoretical and experimental research has been devoted to the study of magnetic and electric dipole moments, but there has been also some interest in the investigation of neutrino millicharges and of the charge radii and anapole moments of neutrinos. Unfortunately, so far there is not any experimental indication in favor of neutrino electromagnetic interactions and all neutrino electromagnetic properties are known to be small, with rather stringent upper bounds obtained in laboratory experiments or from astrophysical observations. The most accessible neutrino electromagnetic property may be the charge radius, discussed in Subsection~\ref{G043}, for which the Standard Model gives a value which is only about one order of magnitude smaller than the experimental upper bounds. A measurement of a neutrino charge radius at the level predicted by the Standard Model would be another spectacular confirmation of the Standard Model, after the recent discovery of the Higgs boson (see \textcite{1312.5672}). However, such a measurement would not give information on new physics beyond the Standard Model unless the measured value is shown to be incompatible with the Standard Model value in a high-precision experiment. The strongest current efforts to probe the physics beyond the Standard Model by measuring neutrino electromagnetic properties is the search for a neutrino magnetic moment effect in reactor $\bar\nu_{e}$-$e^{-}$ scattering experiments. The current upper bounds reviewed in Subsection~\ref{D064} are more than eight orders of magnitude larger than the prediction discussed in Subsection~\ref{D009} of the Dirac neutrino magnetic moments in the minimal extension of the Standard Model with right-handed neutrinos. Hence, a discovery of a neutrino magnetic moment effect in reactor $\bar\nu_{e}$-$e^{-}$ scattering experiments would be a very exciting discovery of non-minimal new physics beyond the Standard Model. In particular, the GEMMA-II collaboration expects to reach around the year 2017 a sensitivity to $\mgm_{\nu_{e}} \approx 1 \times 10^{-11} \bmag$ in a new series of measurements at the Kalinin Nuclear Power Plant with a doubled neutrino flux obtained by reducing the distance between the reactor and the detector from 13.9 m to 10 m and by reducing the energy threshold from 2.8 keV to 1.5 keV \cite{Beda:2012zz,Beda:2013mta}. The corresponding sensitivity to the neutrino electric millicharge discussed in Subsection~\ref{G012} will reach the level of $|\chg_{\nu_{e}}| \approx 3.7 \times 10^{-13} \, \elechg$ \cite{Studenikin:2013my}. There is also a GEMMA-III project\footnote{Victor Brudanin and Vyacheslav Egorov, private communication.} to further lower the energy threshold to about 350 eV, which may allow the experimental collaboration to reach a sensitivity of $\mgm_{\nu_{e}} \approx 9 \times 10^{-12} \bmag$. The corresponding sensitivity to neutrino millicharge will be $|\chg_{\nu_{e}}| \approx 1.8 \times 10^{-13} \, \elechg$ \cite{Studenikin:2013my}. An interesting possibility for exploring very small values of $\mgm_{\nu_{e}}$ in $\bar\nu_{e}$-$e^{-}$ scattering experiments has been proposed by \textcite{Bernabeu:2004ay} on the basis of the observation \cite{Segura:1993tu} that ``dynamical zeros'' induced by a destructive interference between the left-handed and right-handed chiral couplings of the electron in the charged and neutral-current amplitudes appear in the Standard Model contribution to the scattering cross section. It may be possible to enhance the sensitivity of an experiment to $\mgm_{\nu_{e}}$ by selecting recoil electrons contained in a forward narrow cone corresponding to a dynamical zero (see Eq.~(\ref{D036})). In the future experimental searches of neutrino electromagnetic properties may be performed also with new neutrino sources, as a tritium source \cite{McLaughlin:2003yg}, a low-energy beta-beam \cite{McLaughlin:2003yg,deGouvea:2006cb}, a stopped-pion neutrino source \cite{Scholberg:2005qs}, or a neutrino factory \cite{deGouvea:2006cb}. Recently \textcite{Coloma:2014hka} proposed to improve the existing limit on the electron neutrino magnetic moment with a megacurie $^{51}\text{Cr}$ neutrino source and a large liquid Xenon detector. Neutrino electromagnetic interactions could have important effects in astrophysical environments and in the evolution of the Universe and the current rapid advances of astrophysical and cosmological observations may lead soon to the exciting discovery of nonstandard neutrino electromagnetic properties. In particular, future high-precision observations of supernova neutrino fluxes may reveal the effects of collective spin-flavor oscillations due to Majorana transition magnetic moments as small as $10^{-21} \, \bmag$ \cite{deGouvea:2012hg,deGouvea:2013zp}. Let us finally emphasize the importance of pursuing the experimental and theoretical studies of electromagnetic neutrino interactions, which could open a powerful window to new physics beyond the Standard Model. \appendix | 14 | 3 | 1403.6344 |
1403 | 1403.7756_arXiv.txt | A seismic array has been deployed at the Sanford Underground Research Facility in the former Homestake mine, South Dakota, to study the underground seismic environment. This includes exploring the advantages of constructing a third-generation gravitational-wave detector underground. A major noise source for these detectors would be Newtonian noise, which is induced by fluctuations in the local gravitational field. The hope is that a combination of a low-noise seismic environment and coherent noise subtraction using seismometers in the vicinity of the detector could suppress the Newtonian noise to below the projected noise floor for future gravitational-wave detectors. In this paper, certain properties of the Newtonian-noise subtraction problem are studied by applying similar techniques to data of a seismic array. We use Wiener filtering techniques to subtract coherent noise in a seismic array in the frequency band 0.05 -- 1\,Hz. This achieves more than an order of magnitude noise cancellation over a majority of this band. The variation in the Wiener-filter coefficients over the course of the day, including how local activities impact the filter, is analyzed. We also study the variation in coefficients over the course of a month, showing the stability of the filter with time. How varying the filter order affects the subtraction performance is also explored. It is shown that optimizing filter order can significantly improve subtraction of seismic noise. | \label{sec:Intro} In the next few years, a second generation of laser-interferometric gravitational-wave (GW) detectors will start operating with the goal to directly observe GWs. At high frequencies, sources of GWs include core collapse supernovae \cite{OtEA2013,LoOt2012} or the merger of neutron stars and black holes \cite{AbEA2012}. At low frequencies, a stochastic background of GWs, most likely from cosmological origin, is possible \cite{AbEA2009}. The global network of detectors will consist of the two Advanced LIGO \cite{LSC2010} interferometers in Louisiana and Washington state, US, the Advanced Virgo \cite{Vir2011} interferometer near Pisa, Italy, GEO-HF in Hannover, Germany \cite{LuEA2010}, the KAGRA interferometer at the Kamioka mine in Japan \cite{AsEA2013}, and the IndIGO detector in India \cite{Unn2013}. Gravitational waves, which will hopefully be observed by these detectors, produce changes in distance between test masses smaller than $10^{-20}\,$m over kilometer distance scales. Isolating test masses in gravitational-wave detectors from seismic disturbances is one of the foremost challenges of the instrumental design of the detectors. Second-generation detectors will use sophisticated seismic-isolation systems to make possible the detection of GWs above 10\,Hz. For this purpose, the isolation systems consist of chains of coupled springs and pendulums with lowest resonance frequencies around 1\,Hz so that seismic noise well above these frequencies is suppressed by many orders of magnitude in all rotational and translational degrees of freedom. The goal of third-generation detectors, such as the Einstein Telescope \cite{PuEA2010}, will be to extend the detection band to even lower frequencies. This poses an even greater challenge to the design of the isolation system. The mechanical, or passive, component of the isolation system is assisted by a complex network of sensors and actuators forming the so-called active seismic isolation \cite{AbEA2004}. Sensor data recorded on some of the mechanical stages of the passive system are used to calculate a feedback force acting on the mechanical stage to suppress seismic disturbances. Another variant of the active isolation scheme is the feed-forward noise cancellation \cite{GiEA2003,DrEA2012,DeEA2012}. For example, data from a seismic sensor deployed directly on the ground close to a test mass can be used to cancel seismic noise further down the isolation chain, provided that there is correlation between motion of the ground and of some part of the isolation system. Feed-forward noise cancellation has also been implemented as part of the interferometer control using light sensors \cite{KoEA2014}. Seismic disturbances can also affect GW detector test-mass motion via gravitational coupling circumventing the entire seismic isolation system \cite{Sau1984,HuTh1998,Cre2008}. The change in mass density in the rock nearby the detector leads to changes in the local gravitational field, which introduces a force on the test-masses. This so-called Newtonian noise (NN) has never been observed, but it is predicted to be one of the limiting noise sources in second-generation detectors at frequencies between 10\,Hz and 20\,Hz, and for third-generation detectors, such as the Einstein Telescope, down to their lowest frequencies at 2\,Hz. As it is impossible to shield a test mass from NN, another method needs to be found to suppress it. One possibility, at least for future detectors, is to select a site characterized by very low levels of seismic noise \cite{BeEA2010,HaEA2010}. The construction of underground detectors has been proposed for this purpose \cite{PuEA2010}. However, this option is costly, and does not apply to any of the existing surface sites of the LIGO and Virgo detectors. Another idea is to attempt a feed-forward noise cancellation using auxiliary sensors \cite{Cel2000}. For example, gravity perturbations caused by seismic fields can be estimated in real time using data from an array of seismic sensors \cite{DHA2012}. There are also concepts for a number of future low-frequency GW detectors, called MANGO \cite{HaEA2013}, with sensitivity goals better than $10^{-19} / \sqrt{\rm Hz}$ in the 0.1\,Hz to 10\,Hz band. One class of possible detectors includes atom interferometers, which contain a source of ultracold atoms in free fall that interact multiple times with a laser. Another example is a torsion-bar antenna, which uses tidal-force fluctuations caused by GWs which are observed as differential rotations between two orthogonal bars, independently suspended as torsion pendulums \cite{AnEA2010b,ShEA2014}. A final possibility is using the existing Michelson interferometer detector design optimized to low frequencies. One of the main noise sources in this frequency band will be the NN from seismic surface fields. The study presented in this paper of coherent seismic-noise cancellation in the frequency range 0.05 -- 1\,Hz is a first step to investigate the feasibility of a seismic NN cancellation in the same band. It is therefore directly relevant to maximizing the sensitivity of MANGO detectors. There are a number of important GW signal sources in the MANGO GW detector band. Compact binaries in their inspiral and merger phase are strong possibilities \cite{HaEA2013}. Intermediate mass black hole binaries can merge in this band \cite{Aa2014}. Galactic white dwarf binaries would likely be detectable \cite{HaEA2013}. Other sources include helioseismic and other pulsation modes \cite{CuLi1996}. Although interesting in their own right, these would be a foreground for the potential detection of primordial GWs. GWs were recently possibly detected in the B-mode polarization of the CMB background, which would provide confirmation of the theory of inflation \cite{AdEA2014}. Assuming a slow roll inflationary model, this signal would correspond to a GW energy density spectrum $\Omega_{\rm GW} \approx 10^{-15}$ in the 0.1\,Hz to 1\,Hz band. Because $\Omega_{\rm GW}(f) \sim S_{\rm GW}(f) f^3$, where $S_{\rm GW}(f)$ is the detector power spectral density, a detector with strain sensitivity of $\sqrt{S_{\rm GW}} \approx 10^{-23} / \sqrt{\mathrm{Hz}}$ in this band might have sufficient sensitivity to detect the inflationary signal. Furthermore, due to the relatively limited astrophysically produced foregrounds, this band appears the most promising for a direct detection of the inflationary signal. Newtonian noise directly contributes to the noise of a gravitational-wave interferometer by introducing gravitational forces on the test masses. One of the dominant contributions to Newtonian noise is seismic Newtonian noise. In the following, we will generically refer to seismic Newtonian noise as Newtonian noise. There are two dominant contributions to Newtonian noise: the change in the surface-air boundary caused by seismic waves, and change in the rock densities also caused by seismic waves. There is a linear relationship between the amplitude of seismic waves and Newtonian noise, and thus identification and subtraction of seismic noise in an array of seismometers is useful for potential Newtonian noise subtraction in gravitational-wave interferometers. In the study that follows, we will use an array of seismic sensors to subtract seismic noise from a target seismic sensor, which imitates the time-series of a test mass in a gravitational-wave interferometer. Therefore, there is significant motivation for exploring techniques which would maximize the sensitivity of detectors in this low-frequency band, and this is the major focus of this paper. For stationary, linear systems, the optimal filter used for a feed-forward noise cancellation is the Wiener filter \cite{Vas2001}. It is calculated from correlations between data of the auxiliary sensors and data observed at the target point where noise is to be suppressed. The parameters of the optimal linear filter under ideal conditions are fully determined by the data correlations and the frequency range over which noise cancellation is to be achieved. In reality though, as will be shown in the following, filter parameters can be further optimized to account for non-stationary properties of the data and variations of the dynamics of the system; one possible example is temperature drift. With respect to slow changes in noise variance or system dynamics, one can simply update the Wiener filter regularly using the latest observed data or implement an adaptive filter technology \cite{Say2003}. However, as will be shown in this paper, in the presence of non-stationary seismic noise, improvement can also be achieved by optimizing the number of filter coefficients. Wiener filtering with seismic arrays has been performed in the past to improve signal-to-noise ratios towards weak seismic signals \cite{WaEA2008}. In this paper, we present results from a study of feed-forward noise cancellation by means of Wiener filters calculated from correlations between seismometers of an underground array. The sensors are broadband instruments sensitive to seismic noise between about 10\,mHz and 50\,Hz. The array is located at the Sanford Underground Research Facility in the Black Hills of South Dakota \cite{HaEA2010}. The facility has 8 environmentally shielded and isolated stations at 4 different depths. The seismometers are installed on granite tiles placed on concrete platforms connected to the bedrock. They are surrounded by a multi-layer isolation frame of rigid thermal and acoustic insulation panels to further stabilize the thermal environment and to achieve suppression of acoustical signals and air currents. These seismometers have been characterized using huddle tests, and they have been shown to have the same noise floor. Three stations with good data quality were active during this study using data from February and March 2012: one at 800\,ft depth, one at 2000\,ft, and one at 4100\,ft. Whereas the challenge of seismic Newtonian-noise subtraction cannot be fully represented by our study, it is explained that some key aspects such as stationarity of the noise, and scattering of seismic waves should affect both, Newtonian and seismic-noise subtraction, in similar ways. Therefore, the results of our study allow us to draw certain conclusions for Newtonian-noise subtraction. In our analysis, the Wiener filters are realized as finite-impulse response (FIR) filters. The two main parameters investigated here are the rate at which the filters are updated, and the number of filter coefficients. A brief summary on data quality issues in GW detectors in general, and specifically of the seismic data used in this study is given in section \ref{sec:DataQuality}. In section \ref{sec:WienerFiltering}, the Wiener filtering method used in this work is described. The results are given in section \ref{sec:results}, and our conclusions are summarized in section \ref{sec:Conclusion}. | \label{sec:Conclusion} In this paper, we have investigated limits of coherent seismic-noise subtraction, which serves as a first test bed for the more challenging problem of seismic NN subtraction. The main difference between the two problems is that a much larger number of seismometers will be required for seismic NN subtraction. The question of optimal array design and many technical issues to calculate Wiener filters based on a large number of reference channels are not addressed by our study. However, our results allow us to put constraints on the effect from seismic scattering on coherent NN subtraction with the conclusion that at least a factor 50 NN reduction should in principle be feasible at the Homestake site around 0.1\,Hz, provided that seismic scattering at the Homestake site is representative for seismic scattering of the entire region that needs to be included for NN estimates. We have demonstrated that we can achieve more than an order of magnitude seismic-noise cancellation between about 0.05-0.5\,Hz using Wiener filters with only a few seismometers separated by a distance of order 500\,m. We have also shown that this subtraction performance can be achieved without regularly updating the filter, indicating that the average properties of seismic fields at Homestake do not change significantly over timescales of weeks in this frequency band. This is beneficial for realizations in future GW detectors as it simplifies the application of the method to their output. However, in the attempt to optimize noise cancellation, it was found that filter order plays an important role. At frequencies below 0.1\,Hz, subtraction residuals varied almost by an order of magnitude for filter orders between 1 -- 1000. Whereas continuous optimization of filter order may not be feasible in many applications, especially in system control, the results also show that there are ranges of filter order with near-optimal subtraction performance over a broader range of frequencies. These filter orders maintain near-optimal performance over days. As shown in various publications in the past \cite{GiEA2003,DrEA2012,DeEA2012}, Wiener filters can be used to efficiently subtract noise from data off-line, or in real time by means of feed-forward noise cancellation. In this paper, we focused on subtraction of coherent seismic noise, but many other applications are conceivable and have already been applied in previous generations of GW detectors, often forming an important part of the detector design. Wiener filters can subtract broadband noise as well as narrow-band features such as noise lines, and they will therefore play an important role in future detectors, and also serve as a starting point for the development of more advanced filter technologies. These results are important as we have achieved factor of 50 subtraction in the low-frequency GW detectors band. For seismic NN, you would need another factor of 20 reduction to achieve $10^{-20}$ in strain sensitivity. Our work corresponds to the first test bed of noise subtraction in this band for these detectors. Although there is no direct connection to LIGO-like interferometers due to the low-frequency band, we are exploring pushing Wiener-filter subtraction to its limits and how to maximize its efficacy. With the installation of a larger seismic array, hopefully with good data quality above 1\,Hz, subtraction above 1\,Hz can be explored. As of now, it is not possible to say whether we can expect NN subtraction at the same levels as achieved in our analysis. This will require a test of the ability to monitor the seismic wavefield with an expanded seismic array. Because the residual spectra also contain a microseismic peak, it is evident that noise cancellation is not only limited by instrumental noise. A theory that should be tested in the future is whether the residual peak is produced by body waves instead of surface waves. It is known that both wave types contribute to the microseisms, but it is not clear how this affects noise cancellation. Alternatively, it is possible that topographic scattering of seismic waves play a role. The Homestake seismic underground array will be expanded in the future to more seismometers. This will make it possible to carry out a number of important studies relevant to seismic-noise cancellation, which were impossible with the more limited array used here. The extended array will have the capabilities to distinguish between body and surface waves, which, as explained, will be important to explore and possibly understand seismic-noise cancellation limits. It remains to be tested whether a larger array with greater variation in station distances would yield even better subtraction over a broader range of frequencies, potentially down to the instrumental noise limit. This possibility also motivates the development of improved seismometers with reduced self noise performance. | 14 | 3 | 1403.7756 |
1403 | 1403.0947_arXiv.txt | We analyze the physical conditions of the cool, photoionized (T $\sim 10^4$K) circumgalactic medium (CGM) using the COS-Halos suite of gas column density measurements for 44 gaseous halos within 160\,kpc of $L \sim L^*$ galaxies at $z \sim 0.2$. These data are well described by simple photoionization models, with the gas highly ionized (n$_{\rm HII}$/n$_{\rm H} \gtrsim 99\%$) by the extragalactic ultraviolet background (EUVB). Scaling by estimates for the virial radius, R$_{\rm vir}$, we show that the ionization state (tracked by the dimensionless ionization parameter, U) increases with distance from the host galaxy. The ionization parameters imply a decreasing volume density profile n$_{\rm H}$ = (10$^{-4.2 \pm 0.25}$)(R/R$_{\rm vir})^{-0.8\pm0.3}$. Our derived gas volume densities are several orders of magnitude lower than predictions from standard two-phase models with a cool medium in pressure equilibrium with a hot, coronal medium expected in virialized halos at this mass scale. Applying the ionization corrections to the \ion{H}{1} column densities, we estimate a lower limit to the cool gas mass M$_{\rm CGM}^{\rm cool} > 6.5 \times 10^{10}$ M$_{\odot}$ for the volume within R $<$ R$_{\rm vir}$. Allowing for an additional warm-hot, OVI-traced phase, the CGM accounts for {\emph{at least}} half of the baryons purported to be missing from dark matter halos at the 10$^{12}$ M$_{\odot}$ scale. | \label{sec:intro} Baryons account for 17\% of the gravitating mass in the universe ($\Omega_b$ = 0.17 $\Omega_m$; Blumenthal et al. 1984; Dunkley et al. 2009\nocite{cdm84, wmap05}). Yet, observational inventories reveal a shortage of baryons on both universal and galaxy-halo scales. The first `missing baryon problem' is illustrated by counting up all the baryons revealed by observations of stars, dust, and gas in galaxies and clusters ($\Omega_g$). The total is significantly less than the value expected from the widely-accepted Big Bang Nucleosynthesis model, weighing in at only 0.03 - 0.07$\Omega_b$ \citep{persic92, fhp98, bell03b}. Second, baryons are apparently missing from galaxies themselves in what is known as the galaxy halo missing baryon problem \citep{mcgaugh07,bregman07, mcgaugh10}. To explain these baryon shortages one must invoke unseen or poorly-defined components: highly photoionized intergalactic hydrogen, known as the Ly$\alpha$ forest \citep{lynds71, sargent80, cen94}, the warm-hot intergalactic medium, or WHIM, \citep{co99,dhk+99} and the circumgalactic medium, or CGM \citep[e.g.][]{bergeron86, lbt+95}. In cosmological hydrodynamical simulations, for instance, baryons are apportioned comparably between the Ly$\alpha$ forest (40\%), the CGM (25\%) and the WHIM (25\%, excluding the gas that is also CGM; Dav\'e et al. 2010\nocite{dave+10}). The present work concerns the halo missing baryon problem, which we briefly summarize here. Generally speaking, the condensed baryonic component of galaxies, which dominates the energy output of the system, is predicted to dynamically trace the underlying dark matter halo. Traditionally, baryon counting in this regime has focused on a galaxy's stars, cold ISM, and its hot X-ray halo gas \citep{bell03b, klypin, baldry08, yang09, mcgaugh10, andersonbregman10, papastergis12, gupta12}. Compared to the cosmological $\Omega_b/\Omega_m$ ratio, galaxies and their halos come up significantly short on baryons. For a Milky-Way luminosity galaxy, the various estimates of the ratio in stellar mass to the dark matter mass within the virial radius range from $M_*/M_{\rm DM} \approx 0.02-0.05$ \citep{behroozi10}; when we add the cold, neutral component from HI surveys \citep{martin10}, this fraction increases to only 0.07. Finally, when we add in the detected X-ray halo gas, the fraction is at most 0.08 \citep[but see][]{gupta12, fang13}. Such a deficiency is often expressed in terms of $(M_{\rm stars,gas}/M_{\rm DM})/(\Omega_b/\Omega_m)$. In this representation, galaxy halos appear to be missing approximately 60\% of their baryons, suggesting that they are structures nearly devoid of baryons both in mass and spatial extent. Models of the formation of galaxies like our Milky Way have long predicted that the central galaxy contains only a modest fraction of the available baryons (Klypin et al. 2011, and references therein\nocite{klypin}). Galaxies are inefficient producers that have converted a small portion of their available gas into stars. In turn, theorists have suggested a suite of physical processes to suppress star formation and/or expel gas from the galaxies \citep{ds86, sprimack99, od+10}. While evidence for outflowing gas from galaxies is common \citep[e.g.][]{wcp+09,rubin13}, its impact on the efficiency of galaxy formation is unclear. Furthermore, feedback processes are also required to explain the observed incidence of metal-line absorption along quasar sightlines \citep[e.g.][]{dodorico91,od+11,booth} and to enrich the CGM of modern galaxies \citep[e.g.][]{chen10,prochaska11, tumlinson11}. Over the past twenty years it has become increasingly apparent that galaxies also exhibit a diffuse baryonic component within the dark matter halo that extends far from the inner regions to the virial radius and beyond \citep{mwd+93, lbt+95, tripp+98,wakker09,prochaska11}. This halo gas or CGM is similar in concept to the intracluster medium revealed in X-ray emission, but the CGM is observed via UV absorption lines and has much lower temperature and density \citep{werk13}. As such, much of the CGM cannot be traced with X-ray imaging nor any other radiative emission process: it is simply too diffuse to permit direct detection with any present-day telescope. Our collaboration, COS-Halos, has been working to characterize this elusive multiphase medium \citep{tumlinson11, thom12, tumlinson13, werk12, werk13}. We have designed and executed a large program with the {\emph{Cosmic Origins Spectrograph}} (COS; Froning \& Green 2009\nocite{froning09}, Green et al. 2012\nocite{green12}) on the {\it Hubble Space Telescope (HST)} that observed halo gas of 44 galaxies, drawn from the imaging dataset of the Sloan Digital Sky Survey (SDSS), whose angular offsets from quasar sightlines and photometric redshifts implied impact parameters (R $<$ 160 kpc) well inside their virial radii. These data comprise a carefully-selected, statistically-sampled map of the physical state and metallicity of the CGM for L $\approx$ L$^*$ galaxies. Of particular relevance to the halo missing baryon problem is the total baryonic mass contained in the multiphase CGM, as traced by absorption from hydrogen and metal lines in various ionization states (e.g. MgII, SiII, CII, SiIII, CIII, SiIV, OVI).\footnote{The commonly used temperature-based nomenclature for the CGM gas phases is different from that of the ISM. The circumgalactic gas in the temperature range 10$^4$ K $\le$ T $<$ 10$^5$ K is typically referred to as cool; the gas in the temperature range 10$^5$ K $\le$ T $<$ 10$^7$ K is called warm-hot; and gas above 10$^7$ K is termed hot, and would be observed via X-ray transitions. Each of these gas phases is highly ionized. } Previous studies have attempted to estimate the total mass contribution of the CGM to a typical L$^{*}$ galaxy with varying degrees of success. Using absorber samples from HST and FUSE \citep{pss04, tripp+05, dsr+06, tripp08,tc08a, ds08, cm09}, and ground based follow-up spectroscopy to determine redshifts of galaxies along the lines-of-sight, Prochaska et al. (2011) report a strong H I-traced CGM out to 300 kpc for all galaxy types. They estimate a baryonic mass of 10$^{10.5\pm0.3}$ M$_{\odot}$ for an assumed constant total hydrogen column, N$_{\rm H}$ = 10$^{19}$ cm$^{-2}$. \cite{tumlinson11} determine the minimum mass of the highly-ionized CGM ($T$ $\approx$ 10$^{5 - 5.5}$ K) as traced by OVI absorption to be $>$ 10$^{9}$ M$_{\odot}$, based on the maximum possible value for the ionization fraction of OVI (f$_{\rm OVI}$ $<$ 0.2; but the fraction may be higher and the corresponding mass lower in some non-equilibrium scenarios; Vasiliev et al. 2013\nocite{vasiliev13}) and assuming the CGM extends to only 160 kpc. Based on HI measurements and a simple halo model that uses a power-law gas density profile exposed to a uniform ionizing background, \cite{thom12} estimate the total mass of the CGM could range from 10$^{9}$ - 10$^{11}$ M$_{\odot}$. \cite{zhu13} and \cite{lan14} use statistical techniques to assess the absorption from CaII and MgII in galaxy halos, and find an order of magnitude more cool gas in the CGM than in the interstellar medium of galaxies, implying a larger total gas mass in the CGM than in the ISM. Stocke et al. (2013) model the ionization state and metallicity of T $\sim$ 10$^{4}$K CGM clouds using absorption line data from COS and STIS. They statistically associate late-type galaxies from SDSS imaging with the COS/STIS absorbers using virial radii estimated from photometry. Based on their assumed galaxy/absorber associations, they estimate that the low-ion CGM can account for between 10\% and 15\% of the total baryonic budget of luminous spiral galaxies. This estimate is a lower limit because of saturated HI absorption lines. Here, we refine these mass calculations by modeling the photoionized gas of the CGM using a carefully selected sample of L $\approx$ L$^*$ galaxies with precise, accurate redshift measurements from Keck and Magellan spectroscopy \citep{werk12} whose 10$^{4}$K CGM is probed by {\emph{HST}}/COS spectroscopy. Our sample covers and detects a large suite of ions \citep{werk13, tumlinson13}. We rigorously determine the ionization state and metallicities for 33 of the COS-Halos sightlines that provide the best-determined measurements of HI and metal-line column densities. With the constraints imposed by the data and models, we are able to provide a conservative mass estimate for an L $\approx$ L$^*$ galaxy's CGM, and show that the CGM is a dominant reservoir of baryons on galactic scales. Section 2 summarizes the sample and data used in our analysis; in Section 3, we discuss the results of the photoionization modeling and tabulate all derived ionization parameters, metallicities, and total hydrogen columns of the individual lines of sight; in Section 4 we present our analysis of these results, including a mass estimate of the photoionized diffuse gas in the circumgalactic medium of L$\approx$L$^*$ galaxies; Section 5 presents a discussion of this result in the context of previous mass estimates, cosmological simulations, and simple hydrostatic solutions. We present a summary and conclusions in Section 6. We additionally provide an Appendix that details the photoionization modeling, explores additional sources of ionization, and discusses the results on a sightline-by-sightline basis. To maintain consistency with previous COS-Halos results, throughout this work we assume the 5-year WMAP cosmology with $\Omega_{\Lambda}$ = 0.74, $\Omega_{m}$ = 0.26, and H$_{0}$ = 72 km s$^{-1}$ Mpc$^{-1}$ \citep{wmap05}. Distances and galaxy virial radii are given in proper coordinates. We use atomic data for absorption lines from Morton (2003), and the solar relative abundances of metals from \cite{asplund09}. | We have assessed the physical conditions and mass of highly ionized, cool (T $\approx$ 10$^4$ K) CGM gas observed within 160 kpc of low-redshift, L $\approx$ L$^*$ galaxies drawn from the COS-Halos survey. The column densities of HI and low-ionization state metal absorption lines require a characteristic total hydrogen column density of N$_{\rm H}$ $>$ 10$^{19}$ cm$^{-2}$ in the CGM of these galaxies (\S~4.1; Figure 7). We have leveraged our unique dataset of 44 COS spectra of quasars selected to be within 160 kpc of the nearest L$^*$ galaxy to construct the first maps of the physical state of the CGM at low redshift. Our key findings are: \begin{enumerate} \item There is a 4$\sigma$ anti-correlation between ionization parameter and HI column density (\S~3.1; Figure 2). The low-ionization state metal line column densities also follow this trend (Figure 3). This result is qualitatively consistent with photoionized clouds in hydrostatic equilibrium where higher column density clouds have a greater total gas volume density. \item We find a 2$\sigma$ correlation between ionization parameter and the projected distance from the galaxy (Figure 4), which is driven by a declining gas volume density with impact parameter (Figure 10). Gas is more highly ionized further from the host galaxy because the gas is lower density at large radii, and thus less shielded from the EUVB. \item We construct gas surface density profiles of hydrogen (Figure 8) and metals (Figure 9), and find they decline out to 160 kpc (0.55 R/R$_{\rm vir}$) with power-law slopes of $-1.0\pm$0.5 and $-0.8\pm$0.3, respectively (\S~4.1). These 2$\sigma$ correlations are derived from a survival analysis including censoring in the HI column densities (lower limits). \item We provide a strict lower limit to the total mass of material in the CGM of low-redshift L$^{*}$ galaxies (\S~4.2.1). This limit does not allow for line saturation and truncates at 160 kpc. There is at least 2 $\times$ 10$^{10}$ M$_{\odot}$ of cool material in the CGM of these galaxies in the most conservative limit. \item We provide a more realistic lower limit to the mass of low-ionization-state material in the halos of L $\approx$ L$^{*}$ galaxies that allows for line saturation in HI (lower limits) and extends to 300 kpc: M$_{\rm CGM}^{\rm cool}$ $>$ 6.5 $\times$ 10$^{10}$ M$_{\odot}$ (\S~4.2.2). We emphasize that this mass estimate is a lower limit because of saturation in the HI absorption lines for over half of our sample. This mass of material suggests that over 25\% of the baryon budget of an L $\approx$ L$^{*}$ halo is accounted for by cool, photoionized gas in the CGM. When we sum the conservatively-estimated contributions from observed hotter, more highly ionized gas phases (OVI, X-ray) we conclude that galaxies may not be baryon-depleted at all relative to the cosmological baryon fraction (Figure 11). \item Finally, we analyze our derived gas volume densities in the context of simple hydrostatic one- and two-phase models (\S~5.3). Each of these models predicts higher gas volume densities by at least a two orders of magnitude. We conclude that the gas we observe is not in hydrostatic equilibrium with a hot gas phase at the virial temperature of the galaxy halo (Figure 12). There may be other means of supporting this gas (e.g. turbulence, magnetic fields), or else the very large amount of gas we observe has no support at all and is very short-lived in its observed state, such as might occur if it is cycling to and from galaxies on timescales that are very short compared to the dynamical times of dark matter halos. \end{enumerate} | 14 | 3 | 1403.0947 |
1403 | 1403.3425_arXiv.txt | { The DRAGON recoil mass separator at {\sc Triumf} exists to study radiative proton and alpha capture reactions, which are important in a variety of astrophysical scenarios. DRAGON experiments require a data acquisition system that can be triggered on either reaction product ($\gamma$ ray or heavy ion), with the additional requirement of being able to promptly recognize coincidence events in an online environment. To this end, we have designed and implemented a new data acquisition system for DRAGON which consists of two independently triggered readouts. Events from both systems are recorded with timestamps from a $20~${}MHz clock that are used to tag coincidences in the earliest possible stage of the data analysis. Here we report on the design, implementation, and commissioning of the new DRAGON data acquisition system, including the hardware, trigger logic, coincidence reconstruction algorithm, and live time considerations. We also discuss the results of an experiment commissioning the new system, which measured the strength of the $\ecm = 1113$ keV resonance in the $^{20}$Ne$\left( p , \gamma \right)^{21}$Na radiative proton capture reaction. \PACS{ {29.85.Ca}{Data acquisition, nuclear physics} \and {25.40.Lw}{Radiative capture} } % } % | \label{sec:intro} \subsection{The DRAGON Facility} \label{subsec:dragon} Radiative capture reactions typically involve the absorption of a light nucleus (typically a proton or an $\alpha$ particle) by a heavy one, followed by $\gamma$-ray emission. These reactions are important in a variety of astrophysical scenarios such as novae~\cite{PhysRevLett.110.262502, PhysRevLett.105.152501, PhysRevC.81.045808, PhysRevLett.96.252501, PhysRevLett.90.162501, PhysRevC.88.045801}, supernovae~\cite{PhysRevC.76.035801}, X-ray bursts~\cite{6208cab5983b4072a365e4e3c7079eed, 0004-637X-735-1-40}, and quiescent stellar burning~\cite{PhysRevLett.97.242503, Schurmann2011557}. They are often difficult to study directly in the laboratory. The cross sections are low, typically on the order of picobarns to millibarns, since the relevant energies are below the Coulomb barrier. Additionally, many interesting reactions involve short-lived nuclei and can only be studied using low-intensity radioactive beams. The \ac{DRAGON} facility at \triumf{}~\cite{Hutcheon2003190}, shown in \figref{fig:dragon}, is a recoil mass separator that was built to study radiative capture reactions using stable and radioactive beams from the \isac{}~\cite{Laxdal2003400} facility. \ac{DRAGON} experiments are typically performed in inverse kinematics with a beam of the heavy nucleus impinging on a windowless gas target containing the lighter one. Beam energies range from $E/A = 0.15$--\meas{1.5}{MeV}. The products of radiative capture (recoils) are transmitted through \ac{DRAGON} and detected in a series of charged particle detectors, while unreacted beam and other products are deposited at various points along the separator's flight path. The recoil detectors consist of a pair of \acp{MCP} to measure local \ac{TOF}~\cite{Vockenhuber2009372} and either a \ac{DSSSD}~\cite{Wrede2003619} or an \ac{IC} to measure energy loss. The $\gamma$ rays resulting from radiative capture are detected in an array of $30$ \ac{BGO} detectors surrounding the target. For beam normalization, the target chamber houses two \ac{IIS} detectors to record elastically scattered target nuclei. In a typical experiment, the scattering rates measured in the \ac{IIS} detectors are normalized to hourly Faraday cup readings of the absolute beam current. Experiments using low-intensity and possibly unpure radioactive beams may also include a pair of \ac{NaI} scintillators, a \ac{HPGe} detector, or both. These auxiliary detectors are located near the first mass-dispersed focus, and they detect the $\gamma$ rays resulting from the decay of radioactive beam deposited onto the nearby slits. This allows a continuous determination of the beam rate and composition throughout the experiment. In many experiments, unreacted, scattered, or charge-changed beam particles (``leaky beam'') are transmitted to the end of \ac{DRAGON} along with the recoils of interest. The rates vary depending on experimental conditions but can potentially be as much as a few thousand times the recoil rate~\cite{EPJNe20}. Hence, it is crucial that leaky beam be separable from recoils in the data analysis. In some cases, separation is possible using the signals from recoil detectors alone. In others, it is necessary to require (delayed) coincidences between the heavy ion and a $\gamma$ ray measured in the \ac{BGO} detectors. In such experiments, a measurement of the \ac{TOF} between the $\gamma$ ray and the heavy ion (``separator \ac{TOF}'') is useful for distinguishing genuine coincidences from random background. \begin{figure} \centering \includegraphics{dragon2.eps} \caption{The \ac{DRAGON} facility at \triumf{}.} \label{fig:dragon} % \end{figure} \subsection{Data Acquisition Requirements} \label{subsec:daq} As mentioned, identification of coincidences between the ``head'' ($\gamma$-ray) and ``tail'' (heavy-ion) detectors is important for many \ac{DRAGON} experiments. As a result, the original \ac{DRAGON} \ac{DAQ} was designed to trigger on singles events from either detector system while also identifying coincidences from hardware gating. The resulting trigger logic was rather complicated and required a moderate amount of hardware reconfiguration when changing the detector setup (for example, swapping the \ac{DSSSD} and \ac{IC}). With this system, the potential for logic problems due to human error or faulty modules was relatively high, resulting in the possibility of wasted beam time or otherwise non-optimal data sets. In order to alleviate the problems associated with the existing coincidence logic, we have designed and implemented a new \ac{DAQ} system for \ac{DRAGON} that identifies coincidences from timestamps instead of hardware gating. In the course of doing this, we have also upgraded the digital readout from a \ac{CAMAC} system to \ac{VME} and migrated part of the trigger logic from \ac{NIM} hardware to a \ac{FPGA}. In the new set\-up, the head and tail systems are triggered and read out completely independent of each other, and coincidences are identified in the analysis stage from timestamp matching. In this paper, we provide an overview of the new DRA\-GON \ac{DAQ} system and data analysis codes. We also discuss the results of the \ac{DAQ} commissioning experiment, which consisted of a measurement of the $\ecm = 1113$ keV resonance strength in the \rxnfull{20}{Ne}{p}{\gamma}{21}{Na} radiative proton capture reaction. | \label{sec:conclusions} In conclusion, we have developed a new \ac{DAQ} for the \ac{DRAGON} recoil mass separator at \triumf{}. The new \ac{DAQ} consists of two free-running acquisition systems with completely independent triggering and readout, one for the $\gamma$-ray detectors surrounding the target and the other for heavy-ion detectors at the end of the separator. Events are recorded with timestamps from a local \meas{20}{MHz} clock, with the clock frequencies and zero-point offsets synchronized between the two systems. Comparison of timestamp values allows coincidence events to be identified in the first stage of data analysis, and we have implemented and successfully employed a coincidence-matching algorithm that is suitable for both online and offline analysis. The new \ac{DRAGON} \ac{DAQ} was commissioned by measuring the strength of the \meas{\ecm = 1113}{keV} resonance in the \rxnfull{20}{Ne}{p}{\gamma}{21}{Na} radiative capture reaction. The experiment ran successfully, and the measured coincidence resonance strength, \meas{\wg = \wgC}{eV}, is in good agreement with our previous singles result~\cite{PhysRevC.88.038801}, as well as earlier publications~\cite{Engel2005491, ThomasAndTanner, PhysRevC.15.579}. All activities to date indicate that the \ac{DAQ} upgrade is successful and that the new system can be used in future \ac{DRAGON} experiments. | 14 | 3 | 1403.3425 |
1403 | 1403.4519_arXiv.txt | We present for the first time CCD SDSS $gr$ photometry, obtained at the Gemini South telescope with the GMOS attached, of stars in the field of the poorly studied star clusters NGC\,1768, HS\,85, SL\,676, NGC\,2107, NGC\,2190, and SL\,866, which are distributed in the main body of the Large Magellanic Cloud. We applied a subtraction procedure to statistically clean the cluster CMDs from field star contamination. In order to disentangle cluster features from those belonging to their surrounding fields, we applied a subtraction procedure which makes use of variable cells to reproduce the field star Color-Magnitude Diagrams (CMDs) as closely as possible. We then traced their stellar density radial profiles from star counts performed over the cleaned field stars dataset and derived their radii. Using the cleaned cluster CMDs, we estimated ages and metallicities from matching theoretical isochrones computed for the SDSS system. The studied star clusters have ages from 0.1 up to 2.0 Gyr and are of slightly metal-poor metal content ([Fe/H] $\approx$ -0.4 dex). | The Large Magellanic Cloud (LMC) harbors more than two thousand catalogued ordinary star clusters (Bica et al. 2008). Although they are prime indicators of the chemical evolution and the star formation history of the galaxy, only a very small percentage have been well studied (Chiosi et al. 2006, Glatt et al. 2010). In this sense, detailed investigations of even a handful of clusters represents a significant improvement of our knowledge of the chemical enrichment history of this galaxy. We have been intensively involved in a long-term project aimed at obtaining ages and metallicities of LMC clusters, as well as addressing other important related issues. For instance, we have discovered a new giant branch clump structure (Piatti et al. 1999), studied the infamous cluster age-gap (Piatti et al. 2002), searched for age and metallicity gradients (Piatti et al. 2009), derived ages and metallicities for some 81 LMC clusters (Piatti et al. 2011a, Piatti 2011), and investigated in detail the LMC field and cluster Age-Metallicity Relationships (Piatti \& Geisler 2013), among others. We continue here our previous work on LMC clusters by presenting results for six mostly unstudied clusters (NGC\,1768, HS\,85, SL\,676, NGC\,2107, NGC\,2190, and SL\,866) with the aim of adding them to our growing sample of well-studied LMC clusters that will allow us to assemble a much more comprehensive database with which to study the formation and evolution of LMC clusters and their parent galaxy. The paper is organized as follows. The next section describes the collected observations and the data reduction. Sect. 3 deals with the observed Color-Magnitude Diagrams (CMDs) and the procedure of disentangling cluster from field star features. We focus also on the estimation of the cluster structural parameters. The cluster fundamental parameters are derived in Sect. 4, while the analysis and discussion of the results are presented in Sect. 5. Our main findings are summarized in Sec. 6. | As far as we are aware from searching the literature, only NGC\,1768 has a previous age estimate. Glatt et al. (2010) obtained an age of log($t$) = 7.8 $\pm$ 0.4 in fairly good agreement with our present value, although their uncertainty is noticeably larger. Glatt et al. have used data from the Magellanic Cloud Photometric Surveys (Zaritsky et al. 2002) to build the cluster CMD. Although they mention that field contamination is a severe effect in the extracted cluster CMDs and therefore influences the age estimates, no decontamination from field CMDs were carried out. Their large age errors could reflect the composite stellar populations of the LMC Bar field towards which the cluster is projected. SL\,676 and NGC\,2017 resulted to be a cluster pair relatively close in age, with an age difference of (350 $\pm$ 210) Myr. These objects present an angular separation in the sky of 4.1', which is equivalent to 59.6 pc. However, since the upper separation limit for binary LMC star clusters is $\sim$ 20 pc (Bathia et al. 1991, Dieball et al. 2002) we concluded that they do not constitute a physical system. Finally, NGC\,2190 and SL\,866 resulted to be intermediate-age star clusters. According to their positions in the galaxy, the resulting ages are in good agreement with those of star clusters placed at a similar deprojected distance from the LMC center, whereas the present metallicties result slightly more metal-rich for those galactocentric distances (Piatti et al. 2009). Comparing the cluster ages and metallicities with those of their respective surrounding star fields (Piatti \& Geisler 2013), we found that the latter are older ($< t >$ $\sim$ 5 Gyr) and more metal-poor ([Fe/H] $\sim$ -1.0 dex). The remarkable different ages and metallicities of the star clusters and the dominant field stellar populations could be explained if we assume that the clusters were born in other parts of the galaxy and, because of their orbital motions, they are observed at the current locations. Notice that the ages of NGC\,2190 and SS\,866 are encompassed within the well-known star cluster bursting formation epoch (Piatti 2011), so that they could have been formed in regions where the cluster burst took place. | 14 | 3 | 1403.4519 |
1403 | 1403.1420_arXiv.txt | The outskirts of galaxy clusters are continuously disturbed by mergers and gas infall along filaments, which in turn induce turbulent flow motions and shock waves. We examine the properties of shocks that form within $\rtwo$ in sample galaxy clusters from structure formation simulations. While most of these shocks are weak and inefficient accelerators of cosmic rays (CRs), there are a number of strong, energetic shocks which can produce large amounts of CR protons via diffusive shock acceleration. We show that the energetic shocks reside mostly in the outskirts and a substantial fraction of them are induced by infall of the warm-hot intergalactic medium from filaments. As a result, the radial profile of the CR pressure in the intracluster medium is expected to be broad, dropping off more slowly than that of the gas pressure, and might be even temporarily inverted, peaking in the outskirts. The volume-integrated momentum spectrum of CR protons inside $\rtwo$ has the power-law slope of $4.25 - 4.5$, indicating that the average Mach number of the shocks of main CR production is in the range of $\left< M_s \right>_{\rm CR} \approx 3 - 4$. We suggest that some radio relics with relatively flat radio spectrum could be explained by primary electrons accelerated by energetic infall shocks with $M_s \ga 3$ induced in the cluster outskirts. | Clusters of galaxies are the largest gravitationally bound structures that emerged from hierarchical clustering during the large-scale structure (LSS) formation of the Universe. While the central part of many clusters looks relaxed into hydrostatic equilibrium, especially in X-ray observations \citep[e.g.,][]{mfsv98,vkfj06}, the outskirts around the virial radius, $\rvir$, are stirred by mergers of substructures and continuous infall of gas along adjacent filaments \citep[e.g.,][]{ryu2003,skillman2008,vazza2009a}. Observational evidence for the deviation from equilibrium in the cluster outskirts can be seen in the entropy distribution. The radial profile of $S \equiv kT/n_e^{2/3}$ obtained in X-ray observations follows roughly $\sim r^{1.1}$ in the inner part of clusters, but beyond it $S$ flattens off and turns down \citep{vkb05,gfsy09,walker2012,simionescu2013}. Moreover, according to structure formation simulations, turbulent flow motions develop during the formation of clusters; the ratio of turbulence to gas pressure increases outwards and reaches order of 10 \% in the outskirts of simulated clusters \citep[e.g.,][]{ryu2008,vazza2009b,lau2009}. Although flow motions are expected to be on average subsonic in the cluster outskirts\footnote{The ratio of turbulent to gas pressure of $\sim 30$ \%, for instance, corresponds to the turbulent Mach number of $\sim 0.42$.}, shock waves have been observed in X-ray as well as in radio. In X-ray observations, some of sharp discontinuities in the surface brightness are attributed to shocks, while others are attributed to cold fronts or contact discontinuities. The physical properties of these shocks including the sonic Mach number, $M_s$, can be determined using the deprojected temperature and density jumps \citep{markevitch07}. Since a shock was found in the so-called bullet cluster (1E 0657–56) \citep{markevitch02}, about a dozen of shocks have been detected with Chandra, XMM-Newton, and recently Suzaku \citep[e.g.,][]{russell10,akamatsu12,ob13}. The shocks identified so far in X-ray observations are mostly weak with $M_s \sim 1.5 - 3$. Shocks have been identified also in radio observations through the so-called radio relics \citep[e.g.,][for reviews]{feretti12,bruggen12}. Radio relics are the radio structures of megaparsec size observed within the virial radius. They often show elongated morphologies with sharp edges in one side, and occasionally come in pairs located in opposite sides of clusters. Radio emissions from these structures usually exhibit high polarization fractions. Radio relics are interpreted as shocks, where relativistic electrons emitting synchrotron radiation are accelerated or re-accelerated. The properties of radio relic shocks such as $M_s$, magnetic field strength, and the age can be estimated from the spectral index and spatial profile of synchrotron emissions \citep{vanweeren2010,krj2012}. So far several dozens of radio relics have been observed, and the Mach numbers of associated shocks are typically in the range of $M_s \sim 1.5 - 4.5$ \citep[e.g.,][]{clarke2006,bonafede2009,vanweeren2010,vanweeren2012}. In a few cases, shocks were detected both in X-ray and radio observations. Interestingly, however, the shock characteristics derived from X-ray observations are not always consistent with those from radio observations. For instance, the shock in the so-called sausage relic in CIZA J2242.8+5301 was estimated to have $M_s \simeq 4.6$ in the analysis of radio spectrum based on diffusive shock acceleration (DSA) model \citep{vanweeren2010}, but X-ray observations indicated $M_s \simeq 3.2$ \citep{akamatsu13}. And the shock in the so-called toothbrush relic in 1RXS J0603.3+4214 has $M_s \simeq 3.3 - 4.6$ according to the radio spectral analysis, but $M_s \la 2$ according to X-ray observations \citep{vanweeren2012,ogrean13}. In addition, the positions of shocks identified in radio are often spatially shifted from those found in X-ray (see the references above). Resolving these puzzles would require further observations as well as theoretical understandings of weak collisionless shocks in the intracluster medium (ICM), which is a high beta plasma with $\beta=P_{\rm th}/P_B\sim 100$ \citep[e.g.,][]{ryu2008}. Here, $P_{\rm th}$ and $P_B$ are the gas thermal and magnetic pressures, respectively. With relatively low Mach numbers as well as elongated morphologies and occasional parings in opposite sides of clusters, shocks observed in the outskirts are often considered to be induced by mergers. The hypothesis of {\it merger shocks} was explored in simulated clusters, especially for the origin of radio relics \citep{skillman2011,nuza2012,skillman2013}. In these studies, shocks in clusters are identified and the injection and acceleration of cosmic-ray (CR) electrons are modeled. Then, along with a model for the magnetic field in the intergalactic medium (IGM), synthetic radio maps are produced and examined. These studies suggested that merger shocks with sufficient kinetic energy flux are likely to be responsible for observed radio relics. However, it was also argued that typical mergers are expected to induce mostly weak shocks with $M_s\la 3$ and major mergers with similar masses, which are required to explain, for instance, the sausage relic \citep{vanweeren2010}, tend to generate very weak shocks with $M_s \la 2$ \citep[e.g.,][]{gabici2003}. The nature and origin of cosmological shocks have been studied extensively, using numerical simulations for the LSS formation of the Universe \citep{miniati2000,ryu2003,psej06,kang2007,skillman2008,hoeft08,vazza2009a,bruggen12}. Shocks are induced as a consequence of hierarchical clustering of nonlinear structures and can be classified into two categories. {\it External shocks} form around clusters and filaments of galaxies, when the cool ($T \sim 10^4$ K), tenuous gas in voids accretes onto them. So the Mach number of external shocks can be very high, reaching up to $M_s \sim 100$ or so. {\it Internal shocks}, which form inside nonlinear structures, on the other hand, have lower Mach numbers of $M_s \la 10$ or so, because they form in much hotter gas that was previously shocked. It was shown that while a large fraction of internal shocks have $M_s \la 3$, those with $2 \la M_s \la 4$ are most important in dissipating the shock kinetic energy into heat in the ICM. Internal shocks are induced by mergers of substructures, as well as by turbulent flow motions and by infall of warm gas from filaments to clusters \citep[e.g.,][]{ryu2003}. {\it Turbulent shocks}, induced by turbulent flow motions, are expect to be weak with at most $M_s \la 2$, because the root-mean-square (rms) flow motions are subsonic. {\it Inflall shocks}\footnote{Here we distinguish infall shocks from external accretion shocks that decelerate never-shocked gas accreting onto clusters and filaments from void regions. Infall shocks are by nature also accretion shocks that stop previously shocked gas accreting onto clusters from filaments.}, on the contrary, can have Mach numbers as large as $\sim 10$, since they form by the infall of the so-called warm-hot intergalactic medium (WHIM) with $T \approx 10^5 - 10^7$ K to the hot ICM with $T \approx 10^7 - 10^8$ K. We note that a continuous infall containing density clumps would be difficult to be differentiated from a stream of minor mergers with small mass ratios. But infall shocks clearly differ from merger shocks, which are generated by major mergers, in the sense that they do not appear as a pair in opposite sides of clusters. In addition, infall shocks should be found mostly in the cluster outskirts, since the gas infall from filaments normally stops around the virial radius and does not penetrate into the core. So it would be reasonable to conjecture that while weak shocks with $M_s \la 3$ in clusters are mostly merger or turbulent shocks, stronger shocks with $M_s \ga 3$ found in the cluster outskirts are likely to be infall shocks. Shocks that can be categorized as infall shocks were identified in observations before. For instance, the radio relic 1253+275 in the Coma cluster was interpreted as an infall shock formed by a group of galaxies along with the intra-group medium accreting into the ICM \citep{br11}. Also the radio structure of NGC 1265 in the Perseus cluster was modeled as the passage of the galaxy through a shock with $M_s \simeq 4.2$ formed by the infalling WHIM \citep{pj11}. However, the properties such as the frequency, spatial distribution, and energetics of infall shocks have not been studied in simulations before, partly because the automated distinction of infall shocks from merger shocks or other types of shocks in simulated clusters is not trivial. It is well established that CRs are produced via DSA process at collisionless shocks, such as interplanetary shocks, supernova remnant shocks, and shocks in clusters \citep{bell1978,bo78,drury1983}. Shocks in the LSS of the universe are the primary means through which the gravitational energy released during the structure formation is dissipated into the gas entropy, turbulence, magnetic field, and CR particles \citep[e.g.,][]{ryu2008,ryu2012}. Post-processing estimations with simulation data for the amount of CR protons produced in clusters showed that the CR pressure in the ICM may reach up to a few \% of the gas thermal pressure \citep{ryu2003,kang2007,skillman2008,vazza2009a}. Observationally, on the other hand, the CR-to-thermal pressure ratio in clusters was constrained to be less than a few \%, with the upper limits on $\gamma$-ray fluxes set by Fermi-LAT and VERITAS \citep{acke10,acke13,arlen12}. In some simulations for the LSS formation, the injection/acceleration of CR protons at shocks and the spatial advection of the CR pressure were followed self-consistently in run-time \citep[e.g.,][]{miniati2001,pfrommer2007,vazza2012}. \citet{pfrommer2007} and \citet{vazza2012}, adopting specific DSA efficiency models, showed that the CR acceleration occurs mostly in the cluster outskirts. Because of long lifetime and slow particle diffusion \citep[e.g.,][]{bbp97}, the CR protons accelerated in the outskirts are likely to be contained in clusters and accumulated in the ICM over cosmological time scales. But they can be advected with turbulent flows toward the central part of the cluster \citep[see, e.g.,][]{epms11}. For simplicity, let us assume that the transport of CR protons due to flow motions can be approximated by turbulent diffusion, then it could be be described by $\partial Q({\vec r},t)/\partial t = {\vec\nabla}\cdot[D({\vec r},t){\vec\nabla}Q({\vec r},t)]$, where $Q({\vec r},t)$ is the density of CR protons and $D({\vec r},t)$ is the turbulent diffusion coefficient. If only the radial diffusion is considered and the diffusion coefficient is approximated as $D(r,t) \sim r V(r)$, where $V(r)$ is the average flow speed at $r$, then the advection time scale can be estimated rather roughly as $\tau_{\rm adv} \sim r^2 / D \sim r / V(r)$. In the cluster outskirts, typically $r \sim 1 \Mpch$ and $V(r) \sim$ a few $\times\ 100$ km/s, so $\tau_{\rm adv} \sim$ a few $\times\ 10^9$ yrs. This is a substantial fraction of the age of the universe, implying that it would take a while for CR protons produced at energetic shocks in the outskirts to reach the core region. As a result, the radial distribution of the CR pressure would be broad, dropping off more slowly than that of the gas thermal pressure in the outskirts. \citet{vazza2012} also showed that the CR pressure distribution could be temporarily inverted, that is, the CR pressure can increase outwards. \citet{brunetti2012}, on the other hand, attempted to constrain the radial distributions of nonthermal components (including the CR proton energy density) in the Coma cluster by combining radio observations with recent Fermi-LAT $\gamma$-ray observations and with Faraday rotation measure (RM) data. They argued that the model based on the turbulent acceleration of secondary electrons would best reproduce the radio halo of the Coma cluster with the CR energy density that scales with the thermal energy density as $\varepsilon_{\rm CR} \propto \varepsilon_{\rm th}^{\theta}$ with $\theta\approx -0.1$ to $-0.35$, implying that $\varepsilon_{\rm CR}$ is higher at lower $\varepsilon_{\rm th}$. The outer region of the Coma cluster is strongly disturbed by ongoing mergers and infalls \citep[e.g.,][]{simionescu2013}, so it would be probably one of rare cases with this kind of inversion of the $\varepsilon_{\rm CR}$ profile. But these indicate that the partitioning of thermal and CR energies (and possibly turbulent and magnetic field energies too) could be very different in different parts of clusters. In this paper, we study shocks within the virial radius in a sample of clusters taken from LSS formation simulations. Specifically, we examine the properties of {\it energetic shocks} with relatively high Mach number and high shock kinetic energy flux that can produce large amounts of CR protons via DSA. The plan of this paper is as follows. In Section 2, numerical details are presented. In Section 3, the properties and nature of shocks in the cluster outskirts are described. In Section 4, the properties of CRs produced at energetic shocks are described. Discussion is given in Section 5, and summary follows in Section 5. | The existence of CR protons in galaxy clusters remains to be confirmed directly from observations. CR protons produce $\gamma$-ray radiation through $p-p$ collisions with background thermal protons, but so far only upper limits on $\gamma$-ray fluxes from clusters have been set by Fermi-LAT and VERITAS, as noted in the Introduction \citep{acke10,acke13,arlen12}. On the other hand, secondary electrons are also produced through $p-p$ collisions, and they along with $\mu$G-level magnetic fields emit the synchrotron radiation in radio over the cluster scale. The observed radio emission is produced typically by electrons with energy of several GeV corresponding to the Lorentz factor of $\gamma_e \sim 10^4$ \citep{krj2012}. For this energy range, the secondary electrons {\it at production} have the momentum distribution similar to the proton spectrum, that is, $f_e^p(p) \propto p^{-q_e^p}$ with $q_e^p \approx q$ \citep[e.g.,][]{dermer86}. For the proton power-law of $\bar{q} \approx 4.25 - 4.5$ estimated in the previous section, the secondary electron slope is also $q_e^p \approx 4.25 - 4.5$. Relativistic electrons suffers radiative coolings, dominantly by synchrotron and inverse Compton losses. The cooling time scale of electrons of $\gamma_e \sim 10^4$ is $\tau_{\rm cool} \sim 10^8$ yrs with cluster magnetic fields of a few $\mu$G \citep[e.g.,][]{krj2012}. The momentum distribution function of the secondary electrons {\it at the steady-state} governed by the production through $p-p$ collisions and the cooling processes is given as $f_e^{ss}(p) \propto p^{-q_e^{ss}}$ with $q_e^{ss} = q_e^p -1$ \citep{de00}. So for $\bar{q} \approx 4.25 - 4.5$, the secondary electron slope becomes $q_e^{ss} \approx 5.25 - 5.5$. Radio halos associated with some galaxy clusters are explained as diffuse synchrotron emissions over the cluster scale. The spectral index of observed radio halos is typically in the range of $\alpha_R \approx 1 - 1.5$, although in some radio halos it is much steeper \citep{feretti12}. For $\alpha_R = (q_e-3)/2$, this requires the existence of relativistic electrons with the power-slope $q_e \approx 5 - 6$, which nicely embraces the slope of steady-state secondary electrons described above. In the so-called {\it hadronic model}, for instance, the CR electrons emitting synchrotron radiation are assumed to be secondaries produced through $p-p$ collisions \citep[e.g.,][]{de00,pfrommer2004}. Our results indicate that the CR protons accelerated at shocks in the cluster outskirts may be capable of producing secondary electrons with the right energy spectral slope ($q_e^{ss} \approx 5.25 - 5.5$) for the spectral index of observed radio halos. \citet{brunetti2012}, however, argued that at least for the Coma cluster, the hadronic model that requires the CR proton energy $\ga 3 - 5 \%$ of the thermal energy may violate the $\gamma$-ray upper limit set by Fermi-LAT observations, provided that the magnetic field is not much stronger than that measured/constrained by Faraday RM. Moreover, according to a more recent Fermi-LAT paper \citep{acke13}, this limit has become even more stringent, constraining the CR proton energy down to $\la 1 \%$ of the thermal energy in clusters. In the so-called {\it re-acceleration model}, on the other hand, the secondary electrons are further accelerated by turbulence in the ICM, so the CR proton requirement is alleviated somewhat \citep[e.g.,][]{bsfg01}. The detailed implications of our results for radio halo are complicated, and addressing them properly requires studies which are beyond the scope of this paper. So far, our discussions on DSA at energetic shocks have been focused mostly on CR protons and secondary electrons resulted from $p-p$ collisions. As for supernova remnant shocks, primary CR electrons can be accelerated at ICM shocks in the same manner as CR protons, although the injection (pre-acceleration) of electrons into DSA process remains rather uncertain. We point that the projected surfaces of energetic shocks would have morphologies of partial shells or elongated arcs (see Figure 4), so diffuse synchrotron emissions from primary CR electrons accelerated at these shocks could produce radio structures that resemble radio relics discovered in the cluster outskirts \citep[e.g.][]{vanweeren2010,krj2012}. Moreover, the average radial distance of the MESs in Figure 8 is $\left< r_\mathrm{MES} \right> \sim 0.5\ \rtwo \sim 0.5 - 1.5 \Mpch$, which is comparable to the average distance of observed radio relics from the cluster center \citep[e.g.,][]{vanweeren2009}. So, for instance, radio relics with flat radio spectrum such as the sausage relic in CIZA J2242.8+5301 ($\alpha_R\approx 0.6$) could be explained by primary electrons accelerated by energetic shocks (a substantial fraction of which are infall shocks) in the cluster outskirts. In addition, pre-existing CR electrons in the ICM (previously produced at shocks or through $p-p$ collisions) can be re-accelerated at energetic shocks \citep{krj2012,pop13}. To explain the observed properties of radio relics with flat radio spectra, \citet{krj2012}, for instance, proposed a model in which pre-existing CR electrons with the momentum distribution corresponding to the observed radio spectral index are re-accelerated at weak shocks with $M_s \la 2 - 3$. The sausage relic in CIZA J2242.8+5301 then requires pre-existing CR electrons with $q_e \approx 4.2$. It is interesting to note that this is close to the slope of the secondary electrons at production ($q_e^{p} \approx 4.25 - 4.5$) due CR protons accelerated at energetic shocks in the cluster outskirts. The electrons with $\gamma_e \la 10^2$ have the cooling time longer than the age of the universe. So we may conjecture a scenario in which the secondary electrons produced with $\gamma_e \la 10^2$ are boosted to $\gamma_e \ga10^4$ at shocks in the cluster outskirts, producing radio-relic-like structures. The acceleration or re-acceleration of CR electrons at shocks in clusters, compared to those of CR protons, involve additional complications such as injection, pre-existing CR population, and cooling. | 14 | 3 | 1403.1420 |
1403 | 1403.4944_arXiv.txt | The stability of thermonuclear burning of hydrogen and helium accreted onto neutron stars is strongly dependent on the mass accretion rate. The burning behavior is observed to change from Type I X-ray bursts to stable burning, with oscillatory burning occurring at the transition. Simulations predict the transition at a ten times higher mass accretion rate than observed. Using numerical models we investigate how the transition depends on the hydrogen, helium, and CNO mass fractions of the accreted material, as well as on the nuclear reaction rates of $3\alpha$ and the hot-CNO breakout reactions $^{15}\mathrm{O}\left(\alpha,\gamma\right)\mathrm{^{19}Ne}$ and $^{18}\mathrm{Ne}\left(\alpha,p\right)\mathrm{^{21}Na}$. For a lower hydrogen content the transition is at higher accretion rates. Furthermore, most experimentally allowed reaction rate variations change the transition accretion rate by at most $10\,\%$. A factor ten decrease of the $^{15}\mathrm{O}\left(\alpha,\gamma\right)\mathrm{^{19}Ne}$ rate, however, produces an increase of the transition accretion rate of $35\,\%$. None of our models reproduce the transition at the observed rate, and depending on the true $^{15}\mathrm{O}\left(\alpha,\gamma\right)\mathrm{^{19}Ne}$ reaction rate, the actual discrepancy may be substantially larger. We find that the width of the interval of accretion rates with marginally stable burning depends strongly on both composition and reaction rates. Furthermore, close to the stability transition, our models predict that X-ray bursts have extended tails where freshly accreted fuel prolongs nuclear burning. | \label{sec:Introduction} The thin envelope of neutron stars in low-mass X-ray binaries (LMXBs) is continuously replenished by Roche-lobe overflow of the companion star. The hydrogen- and helium-rich material is quickly compressed, and after mere hours the density and temperature required for thermonuclear burning can be reached \citep{Woosley1976,Maraschi1977,Joss1977,Lamb1978}. If a thermonuclear runaway ensues, the unstable burning engulfs the entire atmosphere, consuming most hydrogen and helium within seconds. This powers the frequently observed Type~I X-ray bursts (\citealt{Grindlay1976,1976Belian}; see also \citealt{Cornelisse2003,Galloway2008catalog}; for reviews see \citealt{Lewin1993,Strohmayer2006}). For LMXBs accretion rates, $\dot{M}$, are inferred of up to the Eddington limit of approximately $\dot{M}_{\mathrm{Edd}}\sim10^{-8}\, M_{\odot}\mathrm{year^{-1}}$ (see Section~\ref{sec:methods}). At high rates close to this limit, the high heating rate from compression and nuclear burning as well as the fast inflow of new fuel allow for steady-state burning of hydrogen and helium (e.g., \citealt{Fujimoto1981,Bildsten1998}). A lower burst rate and ultimately an absence of bursts is observed at increasingly large $\dot{M}$, roughly between $0.1\dot{M}_{\mathrm{Edd}}$ and $0.3\dot{M}_{\mathrm{Edd}}$ \citep{Paradijs1988,Cornelisse2003}. When the burst rate is reduced, the presence of steady-state burning becomes evident from an increase of the $\alpha$ parameter, i.e., the ratio of the persistent X-ray fluence between subsequent bursts and the burst fluence: there is an increase in the fraction of fuel that burns in a stable manner \citep{Paradijs1988}. Understanding the burning regimes at different $\dot{M}$ allows us to accurately predict the composition of the burning ashes that form the neutron star crust, which has observable consequences for, e.g., the cooling of X-ray transients \citep[e.g.,][]{Schatz2013}. Whereas observations place the transition of stability around $0.1\dot{M}_{\mathrm{Edd}}$ to $0.3\dot{M}_{\mathrm{Edd}}$ , models predict it to occur at a mass accretion rate, $\dot{M}_{\mathrm{st}}$, close to $\dot{M}_{\mathrm{Edd}}$ \citep{Fujimoto1981}. The observed $\dot{M}$ is determined from the persistent X-ray flux. As material from the companion star falls to the neutron star, most of the rotational and gravitational energy is dissipated at the inner region of the accretion disk and at a boundary layer close to the neutron star surface. This causes these regions to thermally emit soft X-rays, and Compton scattering in a corona is thought to produce X-rays in the classical band \citep[e.g.,][]{Done2007review}. The broad-band X-ray flux is, therefore, used to infer $\dot{M}$. There is some uncertainty in the efficiency of converting the liberated gravitational potential energy to X-rays, as well as obscuration of the X-ray emitting regions by the disk. These uncertainties, however, are generally believed to be at most several tens of percents, whereas the discrepancy is close to an order of magnitude \citep[see also the discussion in][]{kee06}. This discrepancy is one of the main challenges for neutron star envelope models. At the transition, nuclear burning is marginally stable and produces oscillations in the light curve \citep{Heger2005}. This has been identified with mHz quasi-periodic oscillations (mHz QPOs) observed from hydrogen-accreting neutron stars, which typically occur at accretion rates close to $0.1\,\dot{M}_{\mathrm{Edd}}$ \citep{Revnivtsev2001,Altamirano2008,Linares2011}. In the neutron star envelope hydrogen burns through the hot-CNO cycle \citep[e.g.,][]{Wallace1981}, and helium burns through the $3\alpha$ process. At temperatures above $T\gtrsim5\times10^{8}\,\mathrm{K}$, breakout from the CNO cycle occurs through the $^{15}\mathrm{O}\left(\alpha,\gamma\right)\mathrm{^{19}Ne}$ reaction, and for $T\gtrsim6\times10^{8}\,\mathrm{K}$ through $^{18}\mathrm{Ne}\left(\alpha,p\right)\mathrm{^{21}Na}$. This is followed by long chains of $(\alpha,p)$ and $(p,\gamma)$ reactions (the \textsl{$\alpha$p}-process; \citealt{vanWormer1994}) as well as $(p,\gamma)$ reactions and $\beta$-decays (the \textsl{rp}-process). Isotopes are produced with mass numbers as high as $108$ (\citealt{Schatz2001}; for further discussion about the end point see \citealt{Koike2004,Elomaa2009}). Detailed numerical studies implement these processes in large nuclear networks \citep[e.g.,][]{Woosley2004,Fisker2008,Jose2010}. The importance of key nuclear reactions, such as $^{15}\mathrm{O}\left(\alpha,\gamma\right)\mathrm{^{19}Ne}$, has been demonstrated for the stability of nuclear burning \citep{Fisker2006,Cooper2006,Fisker2007,Parikh2008,Davids2011,Keek2012}. Especially for the two breakout reactions the rates are poorly constrained by nuclear experiment \citep{Davids2011,Mohr2013}, and experimental work to improve this is ongoing \citep[e.g.,][]{Tan2007,Tan2009,Salter2012,He2013}. In this paper we investigate the dependence of $\dot{M}_{\mathrm{st}}$ on the reaction rates of the $3\alpha$-process and the CNO-cycle breakout reactions $^{15}\mathrm{O}\left(\alpha,\gamma\right)\mathrm{^{19}Ne}$ and $^{18}\mathrm{Ne}\left(\alpha,p\right)\mathrm{^{21}Na}$, as well as on the composition of the accreted material. | Using large series of one-dimensional multi-zone simulations, we investigate the dependence of the transition of stability of thermonuclear burning on neutron stars on the reaction rates of the triple-alpha reaction and the hot-CNO cycle breakout reactions $^{15}\mathrm{O}\left(\alpha,\gamma\right)\mathrm{^{19}Ne}$ and $^{18}\mathrm{Ne}\left(\alpha,p\right)\mathrm{^{21}Na}$. Within the nuclear experimental uncertainties of the rates, a reduction of the $^{15}\mathrm{O}\left(\alpha,\gamma\right)\mathrm{^{19}Ne}$ by a factor $0.1$ produces the largest change in the mass accretion rate where stability changes: $\dot{M}_{\mathrm{st}}$ is increased from $1.08\,\dot{M}_{\mathrm{Edd}}$ to $1.46\,\dot{M}_{\mathrm{Edd}}$. The lowest value of $\dot{M}_{\mathrm{st}}=0.97\,\dot{M}_{\mathrm{Edd}}$ is obtained for an increased $^{15}\mathrm{O}\left(\alpha,\gamma\right)\mathrm{^{19}Ne}$ rate by a factor $10$. Within the current nuclear uncertainties we are, therefore, unable to explain the discrepancy with observations, which find $\dot{M}_{\mathrm{st}}\simeq0.1\,\dot{M}_{\mathrm{Edd}}$. We also study the dependence of $\dot{M}_{\mathrm{st}}$ on the accretion composition. Reducing the hydrogen mass fraction below the solar value increases $\dot{M}_{\mathrm{st}}$, leading it further away from the observed value. An additional effect is the increase of the accretion rate interval $\Delta\dot{M}_{\mathrm{st}}$ where burning is marginally stable. For several reaction rate variations $\Delta\dot{M}_{\mathrm{st}}$ increases as well. $\Delta\dot{M}_{\mathrm{st}}$ appears to have a complex dependence on the different reaction rates and the composition, which requires further study to determine. Close to the stability transition, we identify X-ray bursts with extended tails lasting over $10$ minutes, where freshly accreted material continues the nuclear burning. Our simulations yield values of $\Delta\dot{M}_{\mathrm{st}}$, of the marginally stable burning period, and of the $\alpha$-parameter that are consistent with observations. Because of the dependency of these parameters on $\dot{M}_{\mathrm{st}}$, however, quantitative comparisons are problematic as long as the observed $\dot{M}_{\mathrm{st}}$ is not reproduced. Furthermore, given the degeneracy in many of these parameters with respect to variations in reaction rates, accretion composition, as well as the effective surface gravity, it remains challenging to place constraints with current observations. | 14 | 3 | 1403.4944 |
1403 | 1403.6493_arXiv.txt | The cosmological model consisting of a nonlinear magnetic field obeying the Lagrangian $\Ln= \gamma F^{\alpha}$, $F$ being the electromagnetic invariant, coupled to a Robertson-Walker geometry is tested with observational data of Type Ia Supernovae, Long Gamma-Ray Bursts and Hubble parameter measurements. The statistical analysis show that the inclusion of nonlinear electromagnetic matter is enough to produce the observed accelerated expansion, with not need of including a dark energy component. The electromagnetic matter with abundance $\Omega_B$, gives as best fit from the combination of all observational data sets $\Omega_B=0.562^{+0.037}_{-0.038}$ for the scenario in which $\alpha=-1$, $\Omega_B=0.654^{+0.040}_{-0.040}$ for the scenario with $\alpha=-1/4$ and $\Omega_B=0.683^{+0.039}_{-0.043}$ for the one with $\alpha=-1/8$. These results indicate that nonlinear electromagnetic matter could play the role of dark energy, with the theoretical advantage of being a mensurable field. | According to Einstein's equations and assuming a Robertson-Walker (RW) geometry, the currently inferred accelerated expansion of the universe is attributed to a kind of repulsive gravity that makes fall apart spacetime. Such expansion is possible if the dominant component of the universe, the so called dark energy (DE), acts with a negative pressure that overcomes the attractive effect of ordinary matter; its corresponding energy density $\rho$ and pressure $p$ should be such that $\rho + 3 p <0$, in order to produce the mentioned acceleration. It has been shown that the effect of coupling nonlinear electrodynamics to gravity produces negative pressures that in turn accelerate the expansion \cite{Novello2,Novello3,Vollick2008,Labun2010}. In \cite{Dyadichev2002} cosmological models involving homogeneous and isotropic Yang-Mills fields were proposed as an alternative to scalar models of cosmic acceleration; while in \cite{Elizalde2003} a quantum condensate is considered as driven the accelerated expansion. In \cite{Jimenez:2008au} it is shown that a vector-tensor theory consisting of gauge fields coupled to gravity could be the origin of the accelerated expansion of the Universe. In \cite{Jimenez:2008nm} it is pointed out that an effective cosmological constant may arise from an electromagnetic mode or degree of freedom, considering that the electromagnetic field contains an additional (scalar) polarization, such that quantum fluctuations of the energy density get frozen on cosmological scales giving rise to an effective cosmological constant. In \cite{Jimenez:2008er} a timelike electromagnetic field on cosmological scales generates an effective cosmological constant; this field could be originated in primordial electromagnetic quantum fluctuations producing during the inflationary epoch. These models open the possibility that DE originates in properties of ponderable fields and matter. Unlike early universes where high energies justify the appearance of nonlinear electromagnetic effects, in late epochs, the reason to invoke nonlinear electromagnetic behaviour may be different: it can be implemented as a phenomenological approach \cite{Medeiros2012}, in which the cosmic substratum is modeled as a material media with electric permeability and magnetic susceptibility that depend in nonlinear way on the fields \cite{Pleban}. Another argument relies in the view that General Relativity is a low energy quantum effective field theory of gravity, provided that the Einstein-Hilbert classical action is augmented by the additional terms required by the trace anomaly characteristic of nonlinear electrodynamics \cite{Mottola}. Assuming that the cosmological background affects the transmission of light signals, there is another approach that considers nonlinear behaviour in the propagation of light, similar to light traveling in non vacuum spacetime \cite{Mosquera2}. This approach has its basis in the fact that the nonlinear electromagnetic Born-Infeld equations are of the same form than Maxwell's for a material media with the difference that the electric permeability and magnetic susceptibility are functions of the field strengths \cite{BI}. A technical problem arises in the coupling of an electromagnetic field to an isotropic geometry, as the electromagnetic field defines preferred directions, so an isotropization process of the energy-momentum tensor should be adopted. To this end several proposals have come up: one of them is to take a spatial average in the electromagnetic field, \cite{Tolman,Vollick2008,Elizalde2003,Novello2,Novello3}, Alternatively, it has been considered a vector triplet compatible with space homogeneity and isotropy of RW \cite{ArmendarizP}. This is a set of three equal length vectors that point in three mutually orthogonal spatial directions. While the triad guarantees the isotropy of the background, it does not automatically imply the isotropy of its perturbations that are necessary to model some observed anomalies in the CMB radiation. In fact the cosmic triad can be realized with a classical SU(2) vector field configuration \cite{ArmendarizP,Dyadichev2002}. The purpose of this work is to investigate to what extent nonlinear magnetic matter can be considered as source of the present cosmic acceleration as an alternative to the DE component. We shall consider a phenomenological model with a nonlinear magnetic field, proposed in \cite{Novello3}, associated to the nonlinear Lagrangian $\Ln=\gamma F^{\alpha}$, where $\gamma$ and $\alpha$ are two constants to be adjusted from observations. We perform a $\chi^2$ statistical analysis by using a Markov Chain Monte Carlo (MCMC) code; we probe the model with Type Ia Supernovae (SNe Ia), Long Gamma-Ray Bursts (LGRBs) and observational Hubble data (OHD). We analyze three cases, namely, $\alpha=-1$, $\alpha=-1/4$ and $\alpha=-1/8$. We could possibly think of considering a time dependent $\alpha$, which in turn, would lead to a time dependent equation of state (EoS) parameter, $w(z)$, however, a constant $w$ has the great advantage of simplicity and that is why we performed the analysis with fixed $\alpha$. In all cases, we obtain good best fits without introducing the DE component. The paper is organized as follows. In Section 2 we address the coupling of nonlinear electrodynamics (NLED) to a RW geometry. In Section 3, theoretical details of the nonlinear magnetic universe are given. In Section 4 the observational data samples and the statistical method used are presented. In Section 5 the obtained constraints and best fits are discussed, and finally the last section is for concluding remarks. | As a phenomenological approach to describe DE, it is interesting to study nonlinear magnetic scenarios with a Lagrangian of the form $\Ln= \gamma F^{\alpha}$. We performed the adjustment of $\Omega_m$ parameter with three probes: SNe Ia, LGRBs and the Hubble parameter measurements. Technical difficulties lead us to consider the parameter $\alpha$ fixed instead of depending on redshift, and it turned out that $\alpha=-1/4$ and $\alpha=-1/8$ reproduce pretty well the current observational data. The best fit for the magnetic component obtained from the combination of all observational data sets is $\Omega_B=0.562^{+0.037}_{-0.038}$ for the scenario in which $\alpha=-1$, $\Omega_B=0.654^{+0.040}_{-0.040}$ for the one with $\alpha=-1/4$ and $\Omega_B=0.683^{+0.039}_{-0.043}$ for the one with $\alpha=-1/8$ These results allow us to conclude that the nonlinear magnetic matter could play the role of DE. In general, the adjustments of $\Omega_m$ and $h$ for the scenario with $\alpha=-1/4$ and for the one with $\alpha=-1/8$ are considerably better than the one with $\alpha=-1$. In addition, although Eq. (\ref{Eq.5}) sets an upper bound for the value of $\alpha$ in order to produce accelerated expansion, from Eq. (\ref{Eq:alpha}) and our fits from the combination of all observational data sets for the scenario with $\alpha=-1/4$ without assuming any prior on $\Omega_m$, we obtain a bound for $\alpha < 0.368$. A similar bound of $\alpha<0.385$ can be calculated from the $\Lambda$CDM model. In spite that we obtained poor constraints for the $\Omega_m$ parameter from LGRBs data without assuming any prior, we should keep in mind that the use of GRBs as cosmological probes is still in debate and LGRBs data are not as reliable as SNe Ia and OHD; however they can give a general idea of the evolution and behaviour of cosmological models at high redshifts. On the other hand, regarding LGRBs, notice from the value of $\chi^2$ in Table \ref{Table:1} ($\alpha=-1$) that $\Omega_m $ is better adjusted than in Table \ref{Table:3} ($\alpha=-1/8$). Remember that a good adjustment is such that $\chi^2$ is closest to the number of data in the sample. The opposite occurs with SNe Ia: $\Omega_m $ is better adjusted for $\alpha=-1/8$ (see Table \ref{Table:3}) than for $\alpha=-1$ (Table \ref{Table:1} ). If we relate this result with the different redshift ranges that correspond to these probes, $1.547 < z < 3.57$ for LGRBs and $0.015 < z < 1.4$ for SNe Ia, the difference in the adjustments might indicate that for large redshift the EoS with $\alpha=-1$ models better the cosmic fluid than $\alpha=-1/4$. While for near epochs, a better description is accomplished with $\alpha=-1/4$. This result might point to considering the EoS parameter $w(z)$ as redshift dependent. Finally, although our analysis, that reduces to a perfect fluid one with a constant EoS-parameter, may overlap with some existing in the literature, e.g., with the presented in \cite{Medeiros2012}, in this work we have used the most recent compilation of SNe Ia released by the SCP, unlike the referred work in which it has been used the Union compilation which only includes 307 data points. Additionally, we have considered direct Hubble parameter measurements and LGRBs data which have extended the range of redshift of study. Besides, we would like to point out that our test was done employing a MCMC method which is more refined one than a standard $\chi^2$ minimization, thus leading more reliable results. | 14 | 3 | 1403.6493 |
1403 | 1403.7340_arXiv.txt | We trace the full evolution of low- and intermediate-mass stars ($1 M_{\odot} \leq M \leq 8M_{\odot}$) during the Asymptotic Giant Branch (AGB) phase in the {\it Spitzer} two-color and color-magnitude diagrams. We follow the formation and growth of dust particles in the circumstellar envelope with an isotropically expanding wind, in which gas molecules impinge upon pre--existing seed nuclei, favour their growth. These models are the first able to identify the main regions in the {\it Spitzer} data occupied by AGB stars in the Large Magellanic Cloud (LMC). The main diagonal sequence traced by LMC extreme stars in the [3.6]-[4.5] vs. [5.8]-[8.0] and [3.6]-[8.0] vs. [8.0] planes are nicely fit by carbon stars models; it results to be an evolutionary sequence with the reddest objects being at the final stages of their AGB evolution. The most extreme stars, with $[3.6]-[4.5] > 1.5$ and $[3.6]-[8.0] > 3$, are $2.5-3 M_{\odot}$ stars surrounded by solid carbon grains. In higher mass ($>3 M_{\odot}$) models dust formation is driven by the extent of Hot Bottom Burning (HBB) - most of the dust formed is in the form of silicates and the maximum obscuration phase by dust particles occurs when the HBB experienced is strongest, before the mass of the envelope is considerably reduced. | A reliable estimate of the nature and the amount of dust produced by Asymptotic Giant Branch (AGB) stars proves essential for a number of scientific issues. These stars are believed to be the dominant stellar sources of dust in the present-day Universe and their contribution to dust enrichment can not be neglected even at redshift $z > 6$ \citep{valiante11}. In order to properly include their contribution in chemical evolution models with dust, the mass and composition of dust grains released by each star as a function of its mass and metallicity need to be known. In addition, the corresponding size distribution function allows to compute the extinction properties associated with these grains, which is fundamental information required for correctly interpreting the optical-near infrared properties of high-z quasar and gamma ray burst spectra \citep{gallerani10}. Theoretical modelling of dust formation around AGBs has made considerable steps forward in the last years, owing to the pioneering explorations by the Heidelberg group \citep[see e.g.][and references therein]{fg06}, complemented by recent studies \citep{paperI, paperII, paperIII, paperIV,nanni13a, nanni13b}. The reliability of these models must be tested against the observations, given the many uncertainties affecting the AGB phase modelling \citep[see e.g.][]{herwig05}, and the description of the formation and growth of dust grains in the winds of AGBs \citep{fg06}. On the observational side, the study of evolved stars in the Galaxy is hampered by the obscuration determined by the interstellar medium, and by the unknown distances of the objects observed. This pushed the interest towards other nearby galaxies. The Large Magellanic Cloud (LMC) is an optimum target for this scope. This stems from the low average reddening ($E(B-V) \sim 0.075$) and its proximity, that allows the determination of important stellar properties, such as the absolute magnitudes. The LMC has been surveyed in the optical by the Magellanic Clouds Photometric Survey \citep[MCPS,][]{zaritsky97}, in the near-IR by the Deep Near Infrared Survey of the Southern Sky \citep[DENIS,][]{epchtein94} and the Two Micron All Sky Survey \citep[2MASS,][]{skrutskie06}. Of particular interest for the study of evolved, dust--surrounded stars, is the SAGE (Surveying the Agents of Galaxy Evolution) survey obtained by the {\it Spitzer Space Telescope} with the Infrared Array Camera (IRAC; 3.6, 4.5, 5.8 and 8.0 $\mu$m) and the Multiband Imaging Photometer for {\it Spitzer} (MIPS; 24, 70 and 160 $\mu$m) \citep[see e.g.][]{meixner06}. The analysis of evolved stars based on 2MASS and IRAC data lead to a classification of LMC stars into three main categories. \citet{cioni06}, based on their analysis of the 2MASS color-magnitude diagram (CMDs), divided the region above the tip of the Red Giant Branch (RGB) into O-rich and C-rich zones; \citet{lmcpaperII} selected a group of stars, called ``extreme'', showing the clear signature of the presence of a dusty circumstellar envelope. These stars were shown to contribute about 75\% of the overall dust from AGBs (Riebel et al. 2012). This stimulated an interesting series of papers \citep{lmcpaperI, lmcpaperIV, lmcpaperV, lmcpaperVI} aimed at refining such a classification, to interpret the observed CMDs by using grids of synthetic spectra for various chemical and physical inputs, the most relevant being the surface chemistry of the star, the size of the grains formed, the borders of the dusty shell, the surface gravity and the effective temperature of the central object. In this Letter we use models of dust formation around AGBs, based on full evolutionary computations, to interpret the distribution of LMC extreme stars in the color--color and color--magnitude diagrams (CCDs and CMDs, respectively). Our goal is to characterize the extreme stars in terms of the evolutionary phase, and of the amount and type of dust present in their surroundings. This investigation will be an important benchmark for the studies focused on dust formation around AGB stars. | For the first time, we are able to interpret the {\it Spitzer} observations of extreme stars in the LMC on the basis of evolutionary models of dusty AGB stars. We find that the main diagonal sequence traced by the observations in the $[5.8]-[8.0]$ vs. $[3.6]-[4.5]$ diagram is an evolutionary sequence of C-stars that become progressively more obscured as their surface layers dredge-up carbon from the stellar interior. We identify the reddest objects, with $[3.6]-[4.5] > 1.5$, as the descendants of $\sim 2.5-3 M_{\odot}$ stars in the latest evolutionary phases, when they develop an optically thick circumstellar envelope. The distribution of stars in the {\it Spitzer} CCD is generally determined by the optical depth, but it is also dependent on the relative percentages of carbon and SiC dust formed. Stars with a high SiC/carbon ratio populate the upper region of the diagram. This motivates the spread in $[5.8]-[8.0]$ observed for $[3.6]-[4.5] < 1.5$; the scatter disappears for more obscured objects, because the dust formed is entirely dominated by solid carbon. Massive AGBs with $M > 3 M_{\odot}$ experience HBB, which prevents them from becoming C-stars; in their circumstellar envelopes the formation of silicates occurs. The strong feature of silicates at $9.7 \mu$m favours a considerable increase in the $8.0 \mu$m flux, thus the slope traced by the evolutionary sequences in the {\it Spitzer} CCD is different from C-stars. Owing to the small extinction coefficients of silicates, M-stars do not reach extremely red colors; most of dusty massive AGBs are confined in the region $[3.6]-[4.5] < 0.5$, also populated by C-stars. The most massive AGBs experience a phase of strong HBB, when the production of silicates is strongly enhanced: both in the {\it Spitzer} CCD and CMD these stars will evolve above the main sequence, where most of the stars are detected. Follow-up (optical and near-IR) spectroscopic observations of the extreme stars in the LMC will be extremely useful to confirm our analysis. | 14 | 3 | 1403.7340 |
1403 | 1403.6146.txt | In this paper we present an exact general analytic expression $Z(sSFR)={y_Z \over \Lambda(sSFR)}+I(sSFR)$ linking the gas metallicity $Z$ to the specific star formation rate (sSFR), that validates and extends the approximate relation put forward by Lilly et al. (2013, L13), where $y_z$ is the yield per stellar generation, $\Lambda(sSFR)$ is the instantaneous ratio between inflow and star formation rate expressed as a function of the sSFR, and $I$ is the integral of the past enrichment history, respectively. We then demonstrate that the instantaneous metallicity of a self-regulating system, such that its sSFR decreases with decreasing redshift, can be well approximated by the first term on the right-hand side in the above formula, which provide an upper bound to the metallicity. The metallicity is well approximated also by $Z_{L13}^{id}=Z(sSFR)={y_Z \over 1+\eta+sSFR/\nu}$ (L13 ideal regulator case), which provides a lower bound to the actual metallicity. We compare these approximate analytic formulae to numerical results and infer a discrepancy $<0.1$ dex in a range of metallicities ($log (Z/Z_{\odot})\in [-1.5,0]$, for $y_z\equiv Z_{\odot}=0.02$) and almost three orders of magnitude in the sSFR. We explore the consequences of the L13 model on the mass-weighted metallicity in the stellar component of the galaxies. We find that the stellar average metallicity lags $\sim 0.1-0.2$ dex behind the gas-phase metallicity relation, in agreement with the data. | The evolution of the metallicity in galaxies constrains the history of the gas accretion relative to the star formation, as well as the relative importance of outflows. As such it has been extensively studied at different cosmic epochs. Whilst the full stellar metallicity distribution is available only for a few selected nearby galaxies, including the Milky Way and its components, average metallicities in the stars and in the gas of star forming regions are available for many more objects at different cosmic epochs (O'Connell, 1976, Lequeux et al., 1979, Tremonti et al., 2004, Savaglio et al. 2005, Mayer et al., 2005, 2006, Erb et al. 2006, Cid Fernandes et al. 2007, Maiolino et al., 2008, Mannucci et al., 2009, Zahid, et al., 2012, Leja, et al., 2013, Gallazzi et al., 2006, Panter et al., 2008, Sommariva et al., 2010, and references therein). These observations have established that, at any redshift $z<4$, the most massive galaxies are the most metal rich in both their gas and stellar components. Moreover, at fixed mass, the gas metallicity of star forming objects decreases with increasing redshift (Erb et al., 2006, Maiolino et al., 2008, Mannucci et al., 2009, Mannucci et al., 2011 Richard et al., 2011, Yuan et al., 2013). Among the many theoretical attempts to understand the drivers of such relations, analytical chemical evolution models appeal either to a decreasing importance of outflows (e.g. Garnett, 2002, Tremonti et al., 2004, Spitoni et al., 2010, and/or { differential winds}\footnote{Namely the outflows in which the ejection of some chemical species is enhanced with respect to others} , e.g., Dalcanton, 2007, Recchi et al., 2008) or to an increase in the star formation efficiency (Dalcanton, 2007, Spitoni et all., 2010, Peeples \& Shankar, 2010), a variation in the yield (via a flattening of the IMF, e.g. Koeppen et al., 2007) or an increase in the fraction of re-accreted metals (Dav\'e et al., 2012), with galactic mass, as possible explanations. It is worth pointing out that, in many cases, these models are adopted to interpret the data at a single epoch. On the other hand, galaxy formation numerical experiments, such as cosmological simulations and semi-analytical models, despite qualitatively matching the $z\sim 0$ relation, do not reproduce its slope (e.g. Pipino et al., 2009) and generally suffer from over-predicting the metallicity of high-redshift star forming galaxies (see, e.g., the discussion in Maiolino et al., 2008, and references therein, Sakstein et al., 2011, Yates et al., 2012). { More recently, a new dimension was added to the observational picture, with studies suggesting that, \emph{globally}, the gas metallicity $Z$ of $z\sim 0$ galaxies depends also on the (specific) star formation rate (SFR): at a fixed galaxy mass, higher metallicities correspond to a lower star formation activity (e.g. Ellison et al., 2008, Mannucci et al., 2010, but see e.g. Yates et al., 2012). Furthermore, there is empirical evidence suggesting the \emph{local} nature of such mass-SFR-Z relation (Rosales-Ortega et al., 2012), and the question becomes as to whether the $Z=Z(M,SFR)$ relation is redshift independent; that is if high redshift galaxies populate the extrapolation of the so-called $z\sim0$ fundamental-metallicity relation, that can either be a surface in the mass-SFR-metallicity space (Mannucci et al., 2010) or a plane (Lara-Lopez et al., 2010, 2013) out to $z\sim2$, or even $z=3$ (when accounting for changes in the ionization parameter, Nagajima \& Ouchi, 2013, Cullen et al., 2013). Such a debate is lively and far from being set (see e.g., Cresci et al., 2011, Richard, et al. 2011, Niino, 2012, Wuyts et al., 2012, Christensen et al., 2012, Stott et al., 2013, Henry et al., 2013, Troncoso et al., 2013, Zahid et al., 2013, Belli et al., 2013). However, irrespective of the final answer, it highlights the importance of fully and simultaneously addressing galaxy evolution in terms of mass-metallicity and mass-SFR relations and their evolution with redshift. It offers an independent test-bed to the above-mentioned analytic chemical evolution models and provides new constraints to the the increasing body of theoretical works (e.g. Bouch\'e et al., 2010, Dutton et al., 2010, Lilly et al. 2013, L13, and references therein) aimed at explaining the existence and the evolution of the SFR-mass relation (e.g., Elbaz et al., 2007, Noeske et al., 2007, Daddi et al., 2007, Pannella et al., 2009, Oliver et al., 2010) and/or the cosmic run of the specific SFR (sSFR) with simple models for the galaxy growth. In particular, L13 show that the $z<2$ evolution of the sSFR of galaxies may be controlled by the cosmological infall of gas, through the regulating action of the gas reservoir via a Schmidt (1959) linear star formation law. Such a simple model broadly explains at the same time the cosmic evolution of the sSFR (e.g. Gonzalez et al., 2010, Stark et al., 2012, and references therein) and the stellar-to-halo mass ratio (e.g., Moster et al., 2010). More importantly, the L13 model has the additional appealing property of offering an explanation both to the evolution of the gas phase metallicity and to its scatter at a given epoch by directly linking it to variations in the sSFR with just one equation. That is, the sSFR both enters as a second paramenter in setting the metallicity and gives and explanation to the epoch-invariant fundamental metallicity relation, thereby linking the epoch-dependence and the SFR-dependence of the mass-metallicity relation. At a fixed epoch, the slope of the mass-metallicity relation is then given by the variation of both the star formation and outflow efficiencies with stellar mass (see also Calura et al., 2009). In L13, however, { the instantaneous gas phase metallicity is replaced with the value derived considering the system in equilibrium (i.e. imposing $dZ/dt=0$) for both an ideal case of regulator (steady state at constant gas fraction) and a case in which the gas fraction is slowly changing}.\\ \indent Other analytic models do not either make explicit the sSFR dependence of metallicity (e.g. Dayal et al., 2013) or, despite their similarity to L13, adopt a different notion of ``steady-state'' (e.g., constant gas mass evolution, e.g., Dav\'e et al., 2012), claiming that the temporal variation in the metallicity for a given galaxy is driven by the amount of metals ejected in the surrounding medium and then re-accreted. Also in the case of Dav\'e et al. (2012), approximate values for the metallicity are adopted.\\ \indent Given the important role of metallicity as a constraint to galaxy formation theories and the progresses in the measurement of Z, SFR and stellar masses at progressively higher redshifts, it is important to derive full analytical expressions that link the metallicity evolution to the sSFR evolution of a single galaxy for generic gas accretion and outflow histories. If correct, the above mentioned approximated formulae (e.g. L13, Dav\'e et al., 2012) could be then re-derived from such general solutions as special cases and applied in suitable regimes of the galaxy growth. \begin{table*} \centering \begin{minipage}{120mm} %\scriptsize \begin{flushleft} \caption[]{Input parameters} \begin{tabular}{c | c | c | c} \hline \hline & This paper & L13 & remarks\\ \hline Baryonic accretion rate & $\dot{M}_{acc}$ & $\Phi$ & given by cosmological background \\ SF efficiency & $\nu$ & $\epsilon$ & constant (may vary with galactic mass and/or cosmic time) \\ Gas-to-total fraction & $\mu$ & $\nu_{gas}$ & - \\ Gas-to-star fraction & $f$ & $\mu$ & - \\ Infall rate-to-SFR ratio & $\Lambda$ & $1/f_{star}$ & { varies with time} \\ Outflow rate-to-SFR ratio & $\eta$ & $\lambda/(1-R)$ & { may vary with time} \\ Yield per stellar generation & $y_z$ & $y$ & constant\\ Metallicity of infalling gas & $Z_A$ & $Z_0$ & constant \\ Returned fraction & $R$ & $R$ & constant \\ \hline \end{tabular} \end{flushleft} \end{minipage} \label{t1} \end{table*} The aim of this paper is thus to validate and extend the L13 relation between $Z$ and sSFR. To this end, we revisit the L13 equations, link them to analytical models of chemical evolution (Pagel \& Patchett 1975; Hartwick 1976; Twarog 1980; Tinsley, 1980, Matteucci \& Chiosi, 1983, Clayton, 1988, Edmunds, 1990, Koeppen \& Edmunds, 1999, Matteucci 2001, Spitoni et al., 2010), and derive a more general relation in which the metallicity Z is an explicit function of the sSFR for arbitrary gas inflow and outflow histories. We then derive simplified relation for the the Closed Box Model, the steady state evolution and the L13 model as special cases of the general solution. Furthermore, we study the range of validity and the goodness of the L13 approximation by comparing these results to a direct numerical integration of the same equations as well as to the predictions of full numerical chemical evolution models, which relax some of the assumptions done to make the problem analytically tractable.\\ \indent Finally, we test the predictions of the L13 model for the evolution of the mass-stellar metallicity relation with redshift in the specific L13 case and compare it to recent observations (Sommariva et al., 2012). The L13 model and some of its equations are briefly summarized in Sec. 2 with the double aim to both set the stage, introducing the relevant physical quantities and symbols, and to link it to the standard equation of analytic chemical evolution model. In Sec. 3, general relations between gas-phase metallicity and sSFR are presented and their special cases discussed. L13 model predictions regarding the stellar average metallicity and its comparison to data are presented in Sec. 4. Finally, in Sec. 5 we summarize and discuss our main conclusions. | \subsection{The gas phase metallicity} \subsubsection{On the variation of the $Z$ as a function of the sSFR evolution} In this paper we explored the dependence of gas-phase metallicity on the specific star formation rate. In particular, we derived general analytic formulae that relate the gas phase metallicity to both the infall-to-star formation rate ratio and the sSFR, for the case of single-zone single-phase galaxies and a linear Schmidt (1959) star formation law. The derived relations take the typical form {$Z(sSFR)={y_Z \over \Lambda(sSFR/\nu)}+I(sSFR/\nu)$, where $I$ is the integral of the past enrichment history over the sSFR.} In this article, both the inflow- and the outflow-to-star formation rate ratios are functions of time (equivalently of the sSFR) and may depend on the input cosmology, the amount of gas that may penetrate the star forming regions of galaxies, and the adopted star formation law. It is important to stress that this approach is different from that adopted in many analytical chemical evolution works in the literature, and still in use to interpret the metallicity of galaxies. These studies adopt $\Lambda$ and $\eta$ constant, that is they do not take into account that realistic outflow- and inflow-to-star formation rate ratios may change with time as the result of the evolution of the galaxy. Therefore, when they are compared to data at a given redshift, they can only give a simple parameterized understanding of the mass-metallicity relation at that fixed epoch, rather than offering a comprehensive view of galaxy evolution with time. We show that \emph{in many circumstances (early evolution, quasi-steady state evolution with slowly decreasing sSFR) a good estimate of the gas-metallicity is obtained by the value $Z_{L13}^{id}$ that the system would have if in steady-state evolution with the infall-to-star formation ratio set by the current value of the sSFR} (Eq.~11), that is as in the ideal regulator case of L13. On the other hand, a metallicity obtained by current value of the infall-to-star formation rate, i.e. $Z_{L13} $(Eq.~8, L13 non ideal regulator case), would slightly overestimate the current metallicity of the system. These two values bracket with high accuracy the current metallicity of the system. Therefore, we provide the formal justification to L13 approximations and extend their validity to a larger range of cases. In particular, \emph{the formula adopted by L13 for the ideal case is exact for systems with $\Lambda-\eta>1$ and $\Lambda , \eta = const$. It is also a good approximation of the actual metallicity for galaxies with slowly decreasing accretion rates}. Also, we add that, since we did not specify anything on both the size and the geometry of what we called the ``galaxy'', the equations in principle hold at both the \emph{local} (i.e. for each star forming region) and the \emph{global} galaxy level, with the latter being a suitable weighted average of the single star forming regions. Such a theoretical expectation seems to be corroborated by very recent observations (Rosales-Ortega et al., 2012). Finally, L13 (their Fig. 7) show the predictions for the mass-metallicity relation at $z>0$ in some specific cases calibrated on either the Mannucci et al. (2010) or the Tremonti et al. (2004) $z\sim0$ relations. A qualitative agreement with the $z\sim 2$ data is achieved. In this paper we do not repeat the exercise. However, { in Fig.~6, we plot the $z<4$ tracks of the models shown in Fig.~5, compared to the full three-dimensional fundamental metallicity relation as given by Mannucci et al. (2010, their Eq.~2). In the L13 framework, galaxies evolve along the surface given by the fundamental metallicity relation.} Also, we wish to highlight the following point, which was not discussed in L13, but it is implied by the assumed $Z=f(sSFR)$ relation at a fixed epoch. It is in fact relevant to the data-model comparison to note that the galaxies in the $z=2$ and $z=3$ observational samples show a rather flat SFR-mass relation (c.f. Mannucci et al., 2009, their Fig. 6), by virtue of their selection. This relation is flatter than the typical SFR-mass relation for star forming galaxies at the same epoch (e.g. Daddi et al., 2007). If we assume that the observational samples are culled out from the star forming population at their respective redshifts, the selection of the most massive galaxies with systematically below the average SFR creates a bias such that the most massive galaxies tend to systematically have the lowest sSFR, and thus to be the most metal rich at their mass scale (at least this is the expectation in the L13 theoretical framework). By virtue of the SFR selection threshold, the low mass galaxies will have higher than average SFR, and hence a lower metallicity than the typical star forming galaxy at the same redshift and mass. This means that the observational samples might be biased in the direction of having a steeper mass-metallicity relation than the typical relation of an unbiased sample of star forming galaxies at the same redshift. Therefore any empirical conclusion on the evolution in the slope of the mass-metallicity relation must be treated with caution. \subsubsection{On the L13 metallicity formula: accuracy and comparison to full numerical chemical evolution models } In order to quantify the accuracy of the L13 approximations, we compared them to a full and direct numerical integration of the same equations, finding an \emph{excellent agreement ($<0.1$ dex) for three orders of magnitude in the sSFR and almost 2 dex in $Z$}. By comparing tracks in the $Z-sSFR$ plane given by either the L13 approximation or the Closed Box relation to the predictions of full numerical chemical evolution models which relax some of the simplifying assumptions adopted in the analytic case, we find that \emph{star forming $z<2$ (spiral) galaxies, where the sSFR slowly decreases with time, the system evolves along the locus of the steady-state solutions of decreasing gas fraction and increasing metallicity, exactly as in the L13 gas-regulated model}. In particular, these system seeks the steady state metallicity without attaining it and the current metallicity is set by the current value of the sSFR. \emph{Fast-forming (elliptical) galaxies, evolve at higher sSFR than slowly-evolving systems at the same metallicity, with a remarkable similarity to the well-known behavior in the [$\alpha$/Fe]-mass plane}. Their track in the sSFR-Z plane is better approximated by Closed Box models.\\ \subsubsection{The SFR as the second parameter in L13 and other special cases (closed box, evolution at constant gas mass)} The { actual functional form} of a mass-(s)SFR-metallicity relation is quite controversial, with empirical findings also including claims of a reversal (namely high SFR would correspond to high metallicity) at high stellar masses (e.g. Yates et al., 2012) and a lack of any SFR effects at all masses (e.g. Sanchez et al., 2013). Clearly, differences may originate from a variety of empirical issues related to the sample specifics (including redshift range and aperture effects, e.g. Sanchez et al., 2013) as well as to the methods used to derive the metallicity (and the SFR). In L13 the metallicity dependes inversely on the sSFR. To some extent we expect a smaller dependence on the sSFR as a second parameter at high masses, simply because in L13 $\nu$ becomes larger and hence the term $sSFR/\nu$ smaller than the other terms in the denominator of Eq. (8). In other words, the most massive models settle earlier on an evolutionary track where the metallicity quickly asymptotes to the yield, and the second parameter effect caused by variations in the (s)SFR become consequently small. It seems more difficult to explain a reversal of the trend above a given stellar mass scale. Moreover, L13 model is meant to reproduce the average galaxy, therefore, it does not take into account that episodic bursts and mergers may also happen and move galaxies further out of the ``average'' quasi steady-state evolution represented by our tracks.\\ As a matter of fact, in this paper \emph{we also show that an anti-correlation between $Z$ and sSFR is found also in the early evolutionary phases of the Closed Box model}. In other works (e.g. Dav\'e et al., 2012), the case $\Lambda-\eta=1$ (which is a generalization of Larson 1972's \emph{extreme infall} in the context of analytic chemical evolution) has been dubbed ``steady-state''. In other words, all the net accreted gas is used up to form stars. More specifically, it is the $\Lambda=1,\eta=0$ case which is known as the ``extreme infall''. It has the property of preserving the gas mass, rather than the gas fraction, and that the metallicity would evolve as $Z=(Z_A+y_z) \, (1-exp(1/\mu -1))$, asymptotically approaching the yield for $\mu$ approaching 0. The generalization of \emph{extreme infall} where both inflows and outflows are present ($\Lambda$, $\eta=\, const$), has the following analytical solution for the metallicity: \begin{equation} Z={(Z_A \Lambda +y_z) \over \Lambda}\biggl\lbrace 1 - e^{- \Lambda \, (1/\mu -1)}\biggr\rbrace \end{equation} or equivalently: \begin{equation} Z={[Z_A (1+\eta) +y_z] \over (1+\eta)}\biggl\lbrace 1 - e^{- (1+\eta) \, (1/\mu -1)}\biggr\rbrace \label{eq:ei} \end{equation} It is important to note here that, despite assuming the validity of the condition $\Lambda-\eta\simeq1$, Dav\'e et al. (2012) do not fully derive these solutions. In fact, they base their model on their Eq. 9. We find that their formula can be re-arranged, after discarding the trivial solution $Z(t)=0$ and assuming that $Z\ne Z_A$, as: $Z\sim {[Z_A (1+\eta) +y_z] \over (1+\eta)}$, which is only an approximation to our exact solutions (e.g., Eq.~\ref{eq:ei}), valid when the gas fraction is small (as pointed out also by Dayal et al., 2013). This latter condition ($\mu << 1$) is unlikely to be true in high redshift galaxies. When the galaxy is in its asymptotic regime at constant gas mass, that is $Z\simeq Z_A+y_z$, the only way to increase its metallicity is by acting on $Z_A\ne 0$. In the first place, in the light of our full analytic derivation, we stress that the correct solutions for $Z$ in a standard analytic chemical evolution model when the infalling gas metallicity changes with time and it is linked to the past history of the galaxy, must take into account that $Z_A=Z_A(t)$ in the formal integration (Eq.~22 in this paper, see also the implementation of galactic fountains in Recchi et al., 2008). We then note that when $Z_A$ drives the metallicity, it increases with time in a manner that is not necessarily linked to the sSFR evolution. In other words, in systems with $\Lambda(-\eta)\sim1$ the gas fraction still changes with time, leading to changes in the sSFR which are un-correlated to variations in the gas metallicity (in principle set by yield, and varied through a changing metallicity in the ``re-accretion'' of previously ejected material). This also implies that the scatter around the average $Z-sSFR$ relation cannot be described by the same equation that governs the $Z=Z(sSFR)$ evolution as in L13. On the contrary, in Dav\'e et al. (2012), the explanation of the scatter (and of the $Z-sSFR$ anti-correlation) requires stochastic events that drive the galaxies out of equilibrium, either enhancing the SFR (e.g. mergers) or momentarily suppressing it (e.g. a sudden decrease in the accretion), and causing either a decrease or a increase in $Z$, respectively. This perspective is not dissimilar to the explanation given by Mannucci et al. (2010) when they first presented the empirical results on the ``fundamental metallicity relation'', and it is further extended in other recent work (e.g. Forbes et al., 2013) which depict both the mass-SFR and the mass-Z relations as the result of statistical equilibrium in the galaxy population at a given epoch. \begin{figure} \begin{center} \includegraphics[width=3in]{treD.eps} \caption{$z<4$ evolution of the models discussed in Figs.~\ref{fig5} and~\ref{fig4diff} in the three dimensional space given by star formation rate, mass and metallicity. The shaded area is the analytical formula of the fundamental metallicity relation given by Mannucci et al. (2010). } \label{fig3d} \end{center} \end{figure} \subsection{Stellar metallicities in the L13 model} As a further extension of the L13 model, we show that \emph{it also naturally predicts a mass-metallicity relation in the stellar component which matches the current data at different epochs}. We compare our predictions to the data, and despite the encouraging qualitative agreement, no firm quantitative conclusions can be drawn due to: i) the large scatter in the high redshift data; ii) the lack of consistency among the stellar metallicity measurements at different epochs; and iii) the presence of both passive and star forming galaxies in the $z\sim0$ dataset. L13's slowly evolving galaxies therefore match both the observed cosmic evolution in the gas and in the average stellar metallicity at a given mass. This is a consequence of the fact that the \emph{average stellar metallicity systematically lags $\sim 0.1-0.2$ behind the gas metallicity of the same galaxy} (as observed). In evolved ($\mu << 1$) systems, such a small difference can be explained by the fact the both the gas and the average stellar metallicity tends to the yield. Whereas, during earlier stages of the evolution, the explanation lies in the fact that the SFR is steadily increasing in time. Therefore the youngest stellar generations (whose composition is the same of the gas-phase) have a larger weight in the computation of the average stellar metallicity. These findings imply that the evolution of the average stellar metallicity in the early phases of L13 galaxies as a function of the sSFR can be approximated by the same formula adopted for the gas phase metallicity. %\end{itemize} | 14 | 3 | 1403.6146 |
1403 | 1403.0560_arXiv.txt | We present the first ground-based CCD ($\lambda < 1\mu$m) image of an extrasolar planet. Using MagAO's VisAO camera we detected the extrasolar giant planet (EGP) $\beta$ Pictoris b in $Y$-short ($Y_S$, 0.985 $\mu$m), at a separation of $0.470 \pm 0.010''$ and a contrast of $(1.63 \pm 0.49) \times 10^{-5}$. This detection has a signal-to-noise ratio of 4.1, with an empirically estimated upper-limit on false alarm probability of 1.0\%. We also present new photometry from the NICI instrument on the Gemini-South telescope, in $CH_{4S,1\%}$ ($1.58$ $\mu m$), $K_S$ ($2.18\mu m$), and $K_{cont}$ (2.27 $\mu m$). A thorough analysis of our photometry combined with previous measurements yields an estimated near-IR spectral type of L$2.5\pm1.5$, consistent with previous estimates. We estimate $\log (L_{bol}/L_\sun ) = -3.86 \pm 0.04$, which is consistent with prior estimates for $\bPicb$ and with field early-L brown dwarfs. This yields a hot-start mass estimate of $11.9 \pm 0.7$ $M_{Jup}$ for an age of $21\pm4$ Myr, with an upper limit below the deuterium burning mass. Our $L_{bol}$ based hot-start estimate for temperature is $T_{eff}=1643\pm32$ K (not including model dependent uncertainty). Due to the large corresponding model-derived radius of $R=1.43\pm0.02$ $R_{Jup}$, this $T_{eff}$ is $\sim$$250$ K cooler than would be expected for a field L2.5 brown dwarf. Other young, low-gravity (large radius), ultracool dwarfs and directly-imaged EGPs also have lower effective temperatures than are implied by their spectral types. However, such objects tend to be anomalously red in the near-IR compared to field brown dwarfs. In contrast, $\bPicb$ has near-IR colors more typical of an early-L dwarf despite its lower inferred temperature. | In contrast to the stellar main sequence, brown dwarfs (BDs) form a true evolutionary sequence. BDs are not massive enough to maintain a constant effective temperature ($T_{eff}$) via hydrogen fusion \citep[$M\lesssim0.075$ $M_\sun$, e.g.][]{1997ApJ...491..856B}. Thus, a BD cools as it ages, radiating away the gravitational potential energy from its formation \citep{2001RvMP...73..719B}. BDs are classified into spectral types by comparison to anchor objects. Various clues to classification were judiciously chosen such that they should correspond to temperature, at least in a relative sense \citep{1999ApJ...519..802K, 2002ApJ...564..421B, 2003ApJ...594..510B}. Temperature is not a readily observable quantity, however it is well established from theory that substellar objects ranging in mass from $\sim$$1$ to $\sim$$75$ $M_{Jup}$ will have radius in a narrow range of $0.8-1.1$ $R_{Jup}$ \citep[e.g.][]{2011ApJ...736...47B, 2007ApJ...659.1661F}. This means that bolometric luminosity, $L_{bol} = 4\pi\sigma_B R^2 T_{eff}^4$, is approximately determined by temperature alone. $L_{bol}$ is observable, so with our theoretical understanding of radius we can infer $T_{eff}$ and find that field brown dwarf spectral types appear to be a well-defined temperature sequence \citep{2004AJ....127.3516G, 2009ApJ...702..154S}, except perhaps for the coolest objects \citep{2013arXiv1309.1422D}. The result is that the SpT of a brown dwarf is a function of both mass and age. The situation is even more challenging for young objects, which have not completed post-formation contraction. The radius of such an object can be significantly larger, depending on how it formed \citep{1997ApJ...491..856B, 2000ApJ...542..464C, 2003A&A...402..701B, 2007ApJ...655..541M, 2012ApJ...745..174S}. Young objects will also have lower mass than older objects of the same temperature. With lower mass and larger radius these young objects have lower surface gravity (low-g), which changes their spectral morphology \citep[e.g.][]{2001MNRAS.326..695L, 2006ApJ...639.1120K}, but even so their spectra can generally be classified within the ultracool dwarf sequence \citep{2009AJ....137.3345C, 2013ApJ...772...79A}. This population of such low-g BDs is especially interesting because they potentially serve as analogs for young extrasolar giant planets (EGPs). Many of the best studied low-g BDs are companions, such as AB Pic B \citep{2005A&A...438L..29C}, and 2M0122 B \citep{2013ApJ...774...55B}. Examples of isolated low-g objects are 2M0355 \citep{2013AJ....145....2F}, and PSO318.5 \citep{2013arXiv1310.0457L}. These objects tend to have fainter near-IR absolute magnitudes \citep{2013AJ....145....2F,2013AN....334...85L}, and have $T_{eff}$ several hundred K cooler than field BDs of the same SpT \citep{2013ApJ...774...55B,2013arXiv1310.0457L}. These low-g BDs also tend to be much redder in near-IR colors, and despite being fainter in the bluer filters and having lower $T_{eff}$, their bolometric luminosities tend to be consistent with the field for their spectral types \citep{2013arXiv1310.0457L}. The first handful of directly imaged planets show similar properties, highlighting the challenges of studying substellar objects in the new physical regime of low-g. For instance, the EGP HR 8799 b and the planetary mass companion 2M1207 b have L-dwarf like very red near-IR colors, but their luminosities and inferred temperatures (800-1000 K) are more like mid-T dwarfs \citep{2005A&A...438L..25C,2008Sci...322.1348M}. This has been interpreted as a consequence of the gravity dependence of the L-T transition \citep{2006ApJ...651.1166M}. Thick dust clouds, which cause the redward progression of the L dwarf sequence as temperature drops, persist to even lower temperatures at low-g \citep{2011ApJ...732..107S,2011ApJ...735L..39B, 2012ApJ...754..135M}. In this framework extremely red and under-luminous HR 8799 b and 2M1207 b are objects which have yet to make the transition to the cloudless, bluer, T dwarf sequence, hence they are often thought of as extensions of the L dwarf sequence \citep{2010ApJ...723..850B,2011ApJ...733...65B, 2011ApJ...737...34M}. The directly imaged EGP $\bPicb$, in contrast, is much hotter and its near-IR SED is much more typical when compared to field and low-g BDs \citep{2010Sci...329...57L,2010ApJ...722L..49Q, 2011AA...528L..15B, 2011ApJ...736L..33C, 2013arXiv1302.1160B, 2013ApJ...776...15C}. $\bPicb$ is unique among the directly imaged EGPs in that we have a dynamical constraint on its mass from radial velocity (RV). A complete orbit has not yet been observed, so RV monitoring constrains the mass depending on the semi-major axis ($a$): for $a < 8, 9, 10, 11, 12$ AU the upper mass limit is $M<10,12,15.5, 20, 25$ $M_{Jup}$ respectively \citep{2012A&A...542A..18L}. The astrometry currently favors $8 \lesssim a \lesssim 9$ AU \citep{2012A&A...542A..41C}, hence $M\lesssim12$ $M_{Jup}$, though larger values are not ruled out. We can expect to have a good dynamical understanding of $\bPicb$'s mass in the near future. $\bPicb$ is also noteworthy in that we have relatively good constraints on the age of its primary star. The age of the $\beta$ Pictoris moving group, of which $\bPicA$ is the eponymous member, has recently been revised upward to $21\pm4$ Myr \citep{2013arXiv1310.2613B} using the lithium depletion boundary technique. Though somewhat larger than the earlier age estimate of $12^{+8}_{-4}$ Myr by \citet{2001ApJ...562L..87Z}, these two estimates are consistent at the $1\sigma$ level. A well determined age and a dynamical mass constraint make $\bPicb$ a valuable benchmark for understanding the formation and evolution of both low-mass BDs and giant planets. Here we present the bluest observations of $\bPicb$ from the first light of the Magellan Adaptive Optics (MagAO) system, using its visible wavelength imager VisAO. We also present detections with the Gemini Near Infrared Coronagraphic Imager (NICI). In Section \ref{sec:magao} we describe MagAO and VisAO, and briefly discuss calibrations of this new high contrast imaging system. In Section \ref{sec:obs} we present our observations and data reduction procedures. We analyze the $0.9-2.4$ $\mu m$ SED of $\bPicb$ in Section \ref{sec:anal}, showing that this EGP looks like a typical early L dwarf. We explore the ramifications of this for the physical properties ($L_{bol}$, mass, $T_{eff}$, and radius) of the planet. Then in Section \ref{sec:discuss}, we compare these derived properties to field objects, and discuss the relationship of the measured characteristics of EGPs and BDs. Finally, we summarize our conclusions in Section \ref{sec:conclude}. \vspace{-.25in} | } We have presented the first high-contrast far-red optical observations of an EGP with MagAO's VisAO CCD camera, detecting $\bPicb$ in $Y_S$ at a contrast of $(1.63\pm0.49)\times10^{-5}$, at a separation of $0.470\pm0.010''$. The VisAO detection has S/N = 4.1, and a conservative upper-limit on false alarm probability of 1.0\%. We also present observations of $\bPicb$ in the near-IR made with the NICI instrument at the Gemini-South Telescope. Combining our VisAO $Y_S$ and NICI $CH_{4S,1\%}$, $K_S$, and $K_{cont}$ photometry with previous measurements in $J,H$ and $K$, we estimated that $\bPicb$ has a spectral type of L$2.5\pm1.5$. In color-color and color-magnitude plots, $\bPicb$ fits very well with other early-L dwarfs, perhaps being slightly redder in $H-K$. We used our spectral type estimate to evaluate the physical properties of $\bPicb$. Using field brown dwarf bolometric corrections, we estimate $\log(L/L_\sun) = -3.86\pm0.04$ dex. This is consistent with previous estimates. Using hot-start evolutionary models at an age of $21\pm4$ Myr, our $L_{bol}$ measurement yields a mass estimate of $M=11.9 \pm 0.7$ $M_{Jup}$, with an upper limit at $M\approx13$ $M_{Jup}$ due to the model treatment of deuterium burning. For temperature we find $T_{eff} = 1643 \pm 32$ K. For radius we find $R = 1.43 \pm 0.02$ $R_{Jup}$. All of these results are consistent with those of prior studies. If we instead used the field BD sequence to estimate temperature, we would find a $T_{eff}$ $\sim$$250$K hotter than expected from the evolutionary models. The population of low surface gravity ultracool dwarfs and directly-imaged EGPs likewise have low effective temperatures compared to field brown dwarfs of similar spectral type. However, these objects tend to be very red in near-IR colors, and so don't follow the field brown dwarf sequence in color-magnitude diagrams. In contrast to other directly-imaged young EGPs (such as HR 8799 b and 2M 1207 b), $\bPicb$ looks much more like a typical early L dwarf in the near-IR, both in terms of its colors and luminosity, despite its inferred low gravity and cooler temperature. \vspace{-.25in} | 14 | 3 | 1403.0560 |
1403 | 1403.5173_arXiv.txt | \noindent Neutrinoless double beta decay experiments constrain one combination of neutrino parameters, while cosmic surveys constrain another. This complementarity opens up an exciting range of possibilities. If neutrinos are Majorana particles, and the neutrino masses follow an inverted hierarchy, then the upcoming sets of both experiments will detect signals. The combined constraints will pin down not only the neutrino masses but also constrain one of the Majorana phases. If the hierarchy is normal, then a beta decay detection with the upcoming generation of experiments is unlikely, but cosmic surveys could constrain the sum of the masses to be relatively heavy, thereby producing a lower bound for the neutrinoless double beta decay rate, and therefore an argument for a next generation beta decay experiment. In this case as well, a combination of the phases will be constrained. | One of the most important questions in particle physics is whether neutrinos are Dirac or Majorana particles. If they are Dirac, then their couplings to the Higgs are extremely small, thereby exacerbating the already perplexing problem of understanding mass in the Universe. If neutrinos are Majorana particles, then lepton number is violated, and there is new physics responsible for the effective operator $(\bar{L}\tilde{H})(\tilde{H}^TL^c)/\Lambda + h.c.$, where $\tilde{H} = i\tau_2 H$ is the conjugated Higgs doublet, $L$ is the usual lepton doublet, $\Lambda$ is the scale above which the new physics manifests itself, and flavor indices are suppressed. A definitive way to resolve this question is to observe neutrinoless double beta decay (see, e.g., \cite{Bilenky:2012qi,deGouvea:2013onf} for reviews), which violates lepton number. The effective Majorana mass that governs neutrinoless double beta decay is \bea m_{\beta\beta} &=& \Big\vert m_1\cos^2\theta_{12}\cos^2\theta_{13} \,+\, m_2 e^{2i\lambda_2} \sin^2\theta_{12}\cos^2\theta_{13} \vs &&\,+\, m_3 e^{2i[\lambda_3-\delta]} \sin^2\theta_{13} \Big\vert \eql{mbb}\eea where $(m_1,m_2,m_3)$ are the masses of the three mass eigenstates; and the mixing angles ($\theta_{12},\theta_{13}$), CP violating phases $\delta$ and Majorana phases ($\lambda_2,\lambda_3$) are the elements of the unitary matrix relating the mass and flavor eigenstates: \begin{widetext} \be U = \left( \begin{matrix} c_{12}c_{13} & s_{12}c_{13} & s_{13}e^{-i\delta}\cr -s_{12}c_{23}-c_{12}s_{23}s_{13}e^{i\delta} & c_{12}c_{23}-s_{12}s_{23}s_{13}e^{i\delta} &s_{23}c_{13}\cr s_{12}s_{23}-c_{12}c_{23}s_{13}e^{i\delta} & -c_{12}s_{23}-s_{12}c_{23}s_{13}e^{i\delta} & c_{23}c_{13} \end{matrix} \right)\, \left(\begin{matrix} 1 & 0 & 0\cr 0 & e^{i\lambda_2} & 0\cr 0 & 0 & e^{i\lambda_3} \end{matrix} \right) \ee \end{widetext} where $c$ and $s$ are cos and sin. The masses are related to one another via two measured differences of mass squared: the solar mass difference $\Delta m_{12}^2\equiv m_2^2-m_1^2=\msol^2$, known to be positive and the atmospheric mass scale with \be \matm^2 = \begin{cases} \Delta m_{23}^2 & ({\rm Normal\ Hierarchy})\cr \Delta m_{31}^2 & ({\rm Inverted\ Hierarchy}) \end{cases} \ee Current and upcoming underground experiments~\cite{Barabash:2010bd,KamLANDZen:2012aa,Auger:2012gs,Ackermann:2012xja} with 10-100 kg of detector mass have the reach to explore $\mbb$ as small as 100 meV, while the ton scale experiments currently planned can push down to 10 meV. The experiments of the last decades have pinned down many of the neutrino parameters, so the allowed range of $\mbb$ -- and therefore the decay rate -- is becoming clearer. In particular there emerges a key relationship \cite{Pascoli:2005zb} between $m_{\beta\beta}$ and the sum of the neutrino masses. Here we point out that cosmic surveys, which are sensitive to the sum of the neutrino masses~\cite{Abazajian:2011dt,Abazajian:2013oma}, can further narrow the allowed range of $\mbb$ and, in the future, the two sets of experiments can work together to measure one of the Majorana phases~\cite{Minakata:2014jba}. The cosmic surveys exploit the fact that the ratio of the energy density of the cosmic neutrinos to the matter density is $f_\nu=0.008 \, \sum m_\nu/{\rm 100 \, meV}$. Even this small of a fraction disrupts the delicate balance between the gravitational accretion of cold dark matter and the expansion of the universe that would otherwise produce constant gravitational potentials. A small fraction of non-clumping matter (and neutrinos are traveling fast enough not to clump on most scales) leads to decaying gravitational potentials. For example, the power spectrum of the potential is reduced by 5\% if $\sum m_\nu=100$ meV. This suppression can be inferred by measuring the cosmic microwave background temperature~\cite{Hu:2001tn} and polarization~\cite{Hu:2001kj} on small angular scales. Following the initial detections by the Atacama Cosmology Telescope~\cite{Das:2011ak} and the South Pole Telescope~\cite{vanEngelen:2012va}, the Planck satellite has now mapped the potential with 27-sigma~\cite{Ade:2013tyw} significance. Prospects for measuring the spectrum of the potential with upcoming small scale CMB polarization experiments and with galaxy surveys~\cite{Font-Ribera:2013rwa} lead to projections that $\sum m_\nu$ can be constrained at 16 meV level. | 14 | 3 | 1403.5173 |
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1403 | 1403.3224_arXiv.txt | A relativistic framework for the description of bound states consisting of a large number of quantum constituents is presented, and applied to black-hole interiors. At the parton level, the constituent distribution, number and energy density inside black holes are calculated, and gauge corrections are discussed. A simple scaling relation between the black-hole mass and constituent number is established. | \label{sec:intro} Systems that can be characterised by a dimensionless parameter $N\gg 1$ are of considerable experimental and theoretical significance. Prominent examples include interacting Bose-Einstein condensates and baryons in the quantum theory of SU($N$)-chromodynamics. In the case of Bose-Einstein condensates, $N$ simply counts the bosons that constitute the system. In quantum chromodynamics, colour neutrality of baryons implies that $N$ can be identified with the number of valence quarks confined inside the baryons. The main amenity offered by large-$N$ systems is a natural expansion parameter given by $1/N$. In quantum chromodynamics this expansion parameter has a diagrammatic interpretation as planar dominance, which has been exploited, for instance, in the $1/N$-expansion of heavy baryons \cite{tHooft, Witten}. A generic feature of large-$N$ systems is their non--perturbative character. Even if the elementary interactions between individual constituents are consistent with a weak-coupling regime, the large number of constituents can lead to strong collective effects experienced by any individual constituent amidst the others. This suggests a mean-field description, which is well understood in the non-relativistic domain, when the Hartree approximation can be applied. Large-$N$ systems in the relativistic domain, however, are much less understood from a theoretical point of view. Attempts to describe these systems based on, for instance, the Dyson-Schwinger equations are usually too complicated to allow for a consistent approximation scheme. The purpose of this article is to provide an analytical and quantitative framework for realising a mean-field description of large-$N$ systems in the relativistic domain. In the approach presented here, the mean field is provided by a non-trivial vacuum structure causing in-medium modifications of the constituent dynamics that can be related to collective binding effects. At this level, the bound-state description is similar to the one developed by Shifman, Vainshtein \& Zakharov for using quantum chromodynamics as a predictive theory of hadrons. Besides the celebrated quark-hadron duality, certain vacuum condensates (Lorentz- and gauge-invariant compositions of fields in the normal-ordering prescription) of quarks and gluons \cite{Shifman, Shifman1, QuarkHadron} are central concepts in their approach. These condensates parametrise the non-trivial vacuum structure of quantum chromodynamics and allow to represent hadron properties at sufficiently low energies to account for confinement. In contrast, the approach presented here does not intend to model confinement effects. Rather, condensates are used as phenomenological bookkeeping devices to parametrise the mean field experienced by individual constituents in large-$N$ systems. Our main objective is to construct a solid theoretical framework which makes good use of these phenomenological ideas. As will be shown in detail, this leads to a representation of relativistic quantum bound-states qualifying as large-$N$ systems in terms of an auxiliary current. Such a representation is valid both for the asymptotic framework pertinent to the scattering matrix, as well as for the construction of kinematical states associated with large-$N$ systems. Thus, with the aide of the auxiliary current, the corresponding bound states can be reduced in the sense of Nishijima and Lehmann, Symanzik \& Zimmermann \cite{Zimmermann, Nishijima}, as well as in the usual sense of absorption and emission processes. Obviously, this is an important prerequisite for calculating static and dynamical properties of these bound states. As an application, following a recent proposal put forward in \cite{Dvali}-\cite{Dvali3}\footnote{For Schwarzschild black holes in the context of matrix models see \cite{Banks}.}, black holes will be considered as large-$N$ systems at a quantitative level strictly following the logic of the general bound-state formalism developed in this article. The key idea is to model black holes as quantum bound-states of $N\gg 1$ constituents in Minkowski space-time. Here, constituents include all graviton polarisations, in particular scalar gravitons. In this application, Minkowski space-time is not considered as a specific background geometry, rather it has the status of a distinguished space-time. Of course, Schwarzschild space-times are non-perturbative deformations of Minkowski space-time, in the sense that arbitrary many couplings between gravitons and the associated energy-momentum tensor have to be considered \cite{Duff} in order to reproduce this geometry. But the bound-state description suggested here goes beyond a purely perturbative reconstruction. From a geometrical point of view, the condensates represent non-perturbative deformations of Minkowski space-time. Furthermore, the description allows to construct observables sensitive to the constituent structure inside the black hole, such as the momentum-distribution, the number and energy density of the constituents. This leads to a description that is also complementary to the geometrical picture. It is, however, not equivalent, since typical quantum corrections are only suppressed by $1/N$, as opposed to being suppressed exponentially, thus qualifying the notion of black holes as classical entities. As has been pointed out in \cite{Dvali1}, it is exactly this feature which could shed new light on old problems like the information paradox. The outline of this article is as follows. In Section 2 we introduce the auxiliary current construction as a tool of representing bound states in terms of the fields appearing in the microscopic Lagrangian. In order to give a self-contained discussion, we first explain how these bound states can be embedded into the asymptotic framework of S-matrix theory as in-states or out-states. Secondly, we show how the construction can be generalised to situations where the bound state is not an asymptotic state. Finally, it is explained how symmetries of the bound state can be implemented directly in the auxiliary current description. We proceed by constructing the gauge-invariant constituent distribution functions of scalars inside the bound state. Although we are ultimately interested in the distribution of gravitons inside black holes, it suffices to consider scalar distribution at the parton level. Higher order corrections are, however, sensitive to the constituent polarisations, as will be show. In Section 4 we discuss the renormalisation of composite operators at parton level. We show that in the limit of infinite black hole mass, a consistent renormalisation prescription at lowest order amounts to setting all loop contributions to zero. Using this prescription, we calculate observables related to the interior structure of black holes such as the constituent distribution, energy density and total number of black hole constituents at the parton level. While composite operator renormalisation implies that all loop contributions vanish, non-triviality of our results suggests that condensation must take place. These condensates correspond to normal-ordered contributions in Wick's theorem. In the last Section we discuss how gauge corrections can be taken into account. In order to highlight the practical value of external field methods in this context, we calculate a specific diagram, leaving a systematic study of gauge corrections for future work. | \label{sum} In this article, a description of bound states consisting of $N\gg1$ constituent quanta has been given. It is based on a relativistic quantum theory of individually weakly coupled constituent fields that experience strong binding mechanisms caused by the collection of constituents. The quantum states associated with the composite objects are represented by auxiliary currents composed of the constituent fields in accordance with the quantum numbers and isometries carried by the bound states. This implies the usual reduction formalism pertinent to the asymptotic framework of scattering theory, but allows also the reduction of kinematical states representing bound states beyond an asymptotic framework. As an application, Schwarzschild black-holes have been considered as bound states of $N\gg1$ constituent gravitons (of all polarisations). Strictly following the logic of the framework presented here, we calculated the wave\-length-distribution $\mathcal{D}(\lambda)$ of constituents inside black holes at the parton level. It turns out that the distribution favours to populate black-hole interiors with constituents of maximal wavelength, $\mathcal{D}(\lambda)=\mathcal{D}(r_{\rm S}) \lambda/r_{\rm S}$, where $r_{\rm S}$ denotes an arbitrary pivot scale, for instance the Schwarzschild radius which then enters as an external quantity. We showed how gauge corrections arise and how they can be taken into account. Systematic studies of gauge corrections are, however, left for future publications. In addition, we calculated the constituent number $\mathcal{N}_{\rm c}$ and the constituents energy density $\mathcal{E}$ inside the black hole, both of which depend on the hypersurface-localisation. Integrating the energy density per constituent $\mathcal{E}/\mathcal{N}_{\rm c}$ over a spacial slice, we obtained the localisation-independent scaling law $M=\sqrt{N_{\rm c}} E$ relating the black-hole mass $M$ to the physical constituent number $N_{\rm c}$ and the average energy $E$ per condensed constituent. The derivation of this result is transparent and fully anchored in a field-theoretical context with an interesting interpretation and relation to previous works such as \cite{tHooft, Witten, Dvali}. While it is plausible to describe black holes as composite quantum-systems (they certainly allow for an asymptotic particle-like characterisation), we are convinced that the framework presented here allows to illuminate the relation between space-time geometry and quantum physics in general. This is left for future work. \appendix | 14 | 3 | 1403.3224 |
1403 | 1403.3538_arXiv.txt | {The chemically peculiar (CP) stars of the upper main sequence are mainly characterized by strong overabundances of heavy elements. Two subgroups (CP2 and CP4) have strong local magnetic fields which make them interesting targets for astrophysical studies. This star group, in general, is often used for the analysis of stellar formation and evolution in the context of diffusion as well as meridional circulation.} {In continuation of a long term study of CP stars (initiated in the 1980ies), we present new results based on photoelectric measurements for ten open clusters that are, with one exception, younger than 235\,Myr. Observations in star clusters are favourable because they represent samples of stars of constant age and homogeneous chemical composition.} {The very efficient tool of $\Delta a$ photometry was applied. It samples the flux depression at 5200\AA\, typically for CP stars. In addition, it is able to trace emission line Be/Ae and $\lambda$ Bootis stars. Virtually all CP2 and CP4 stars can be detected via this tool, and it has been successfully applied even in the Large Magellanic Cloud. For all targets in the cluster areas, we performed a kinematic membership analysis.} {We obtained new photoelectric $\Delta a$ photometry of 304 stars from which 207 objects have a membership probability higher than 50\%. Our search for chemically peculiar objects results in fifteen detections. The stars have masses between 1.7\,M$_{\sun}$ and 7.7\,M$_{\sun}$ and are between the zero- and terminal-age-main-sequence. We discuss the published spectral classifications in the light of our $\Delta a$ photometry and identify several misclassified CP stars. We are also able to establish and support the nature of known bona fide CP candidates.} {It is vital to use kinematic data for the membership determination and also to compare published spectral types with other data, such as $\Delta a$ photometry. There are no doubts about the accuracy of photoelectric measurements, especially for stars brighter than 10th magnitude. The new and confirmed CP stars are interesting targets for spectroscopic follow-up observations to put constraints on the formation and evolution of CP stars.} | More than a century ago, \citet{Maur97} detected a subclass of A-type stars with peculiar lines and line strengths, which thereafter became known as Ap stars. Later on, the spectral range was widened and the class become known as chemically peculiar (CP) stars of the upper main sequence. These stars revealed other peculiar features, for example the existence of a strong global magnetic field (CP2 and CP4 objects) with a predominant dipole component located at random with respect to the stellar rotation axis and the centre of the star as well as overabundances with respect to the Sun for heavy elements such as silicon, chromium, strontium, and europium. The peculiar surface abundances for CP stars have been explained either by diffusion of chemical elements depending on the balance between gravitational pull and uplift by the radiation field through absorption in spectral lines or by selective accretion from the interstellar medium via the stellar magnetic field \citep{Szkl13}. Therefore, the correlation of stellar magnetic field strengths with astrophysical processes like diffusion and meridional circulation as well as their evolutionary status can be very well studied with this stellar group \citep{Glag13}. \begin{table*} \caption{Fundamental parameters of the target clusters taken from \citet{Paun06} and \citet{Zejd12}.} \label{cluster_par} \begin{center} \begin{tabular}{lcccrrrcrrr} \hline \hline Cluster & & $\alpha$(2000) & $\delta$(2000) & \multicolumn{1}{c}{$l$} & \multicolumn{1}{c}{$b$} & $d_{\sun}$ & $E(B-V)$ & \multicolumn{1}{c}{$age$} & \multicolumn{1}{c}{$\mu_{\alpha}\cos\delta$} & \multicolumn{1}{c}{$\mu_{\delta}$} \\ & & & & & & [pc] & [mag] & [Myr] & [mas/yr] & [mas/yr] \\ \hline Feinstein~1 & C1103$-$595 & 11 05 56 & $-$59 49 00 & 290.03 & +0.39 & 1180 & 0.41 & 5 & $-$6.1 & +2.9 \\ NGC~2168 & C0605+243 & 06 09 00 & +24 21 00 & 186.59 & +2.22 & 830 & 0.23 & 100 & +1.5 & $-$2.9 \\ NGC~2323 & C0700$-$082 & 07 02 42 & $-$08 23 00 & 221.67 & $-$1.33 & 895 & 0.23 & 100 & +0.4 & $-$2.0 \\ NGC~2437 & C0739$-$147 & 07 41 46 & $-$14 48 36 & 231.86 & +4.06 & 1495 & 0.16 & 235 & $-$5.0 & +0.4 \\ NGC~2547 & C0809$-$491 & 08 10 09 & $-$49 12 54 & 264.47 & $-$8.60 & 430 & 0.05 & 45 & $-$7.7 & +5.0 \\ NGC~4103 & C1204$-$609 & 12 06 40 & $-$61 15 00 & 297.57 & +1.16 & 1810 & 0.25 & 30 & $-$5.6 & $-$0.5 \\ NGC~6025 & C1559$-$603 & 16 03 17 & $-$60 25 54 & 324.55 & $-$5.88 & 725 & 0.28 & 75 & $-$3.3 & $-$3.1 \\ NGC~6633 & C1825+065 & 18 27 15 & +06 30 30 & 36.01 & +8.33 & 335 & 0.18 & 505 & +0.2 & $-$1.2 \\ Stock~2 & C0211+590 & 02 15 00 & +59 16 00 & 133.33 & $-$1.69 & 300 & 0.33 & 130 & +16.6 & $-$13.5 \\ Trumpler~2 & C0233+557 & 02 37 18 & +55 59 00 & 137.38 & $-$3.97 & 605 & 0.32 & 120 & +1.0 & $-$4.6 \\ \hline \end{tabular} \end{center} \end{table*} \begin{table*} \caption{Observations log, the description of the used equipment can be found in the last column.} \label{obs_log} \begin{center} \begin{tabular}{lcccc} \hline \hline Cluster & Observatory & Telescope & Time & Reference \\ \hline Feinstein~1 & ESO & Bochum 0.61\,m & 89/04 & \citet{Mait93} \\ NGC~2168 & Hvar & 0.65\,m & 89/01, 89/02 & \citet{Mait87b} \\ NGC~2323 & ESO & 0.5\,m & 85/02, 85/03 & \citet{Mait87a} \\ NGC~2437 & ESO & 1.0\,m & 84/02 & \citet{Mait93} \\ NGC~2547 & ESO & 1.0\,m & 84/02 & \citet{Mait93} \\ NGC~4103 & ESO & 1.0\,m & 84/02 & \citet{Mait93} \\ NGC~6025 & ESO & Bochum 0.61\,m & 89/04 & \citet{Mait93} \\ NGC~6633 & ESO & Bochum 0.61\,m & 89/04 & \citet{Mait93} \\ Stock~2 & Hvar & 0.65\,m & 84/12, 85/12 & \citet{Mait87b} \\ Trumpler~2 & Hvar & 0.65\,m & 85/12, 86/10, 86/11 & \citet{Mait87b} \\ \hline \end{tabular} \end{center} \end{table*} Nearly four decades ago, \citet{Mait76} introduced the $\Delta a$ photometric system in order to investigate the flux depression at 5200\AA\, typically for CP stars. An overview of the system and its applications can be found in \citet{Paun05}. The $a$ index samples the depth of this flux depression by comparing the flux at the centre with the adjacent regions. The final intrinsic peculiarity index $\Delta a$ was defined as the difference between the individual $a$-values and the $a$-values of non-peculiar stars of the same colour (spectral type). It was shown \citep{Paun05} that virtually all CP2 and CP4 stars have positive $\Delta a$-values up to 95\,mmag. Extreme cases of the non-magnetic CP1 and CP3 objects may exhibit marginally positive $\Delta a$ values, whereas emission line Be/Ae and $\lambda$ Bootis stars exhibit significant negative values. Since the detailed study of Pleione \citep{Pavl89}, it is known that Be stars can change their $\Delta a$ values from significantly positive at their shell to negative at their emission phase. Starting with the paper by \citet{Mait81}, fourteen parts of a large photoelectric $\Delta a$ survey to detect CP stars in open clusters and stellar associations were published. Observations in star clusters are preferable because they represent samples of objects of constant age and homogeneous chemical composition, suited to the study of processes linked to stellar structure and evolution, and to fixing lines or loci in several very important astrophysical diagrams such as the colour-magnitude diagram (CMD), or the Hertzsprung-Russell diagram (HRD). In this paper, we present new photoelectric $\Delta a$ data for the ten open clusters Feinstein~1, NGC~2168, NGC~2323, NGC~2437, NGC~2547, NGC~4103, NGC~6025, NGC~6633, Stock~2, and Trumpler~2. All aggregates, with the exception of NGC~6633, are younger than 250\,Myr (Table \ref{cluster_par}). The detection of CP stars in young open clusters will help us to understand the formation and evolution of these objects and their magnetic fields. \begin{figure*} \begin{center} \includegraphics[width=165mm]{figure1a.eps} \caption{$a_{\rm corr}$ versus $(B-V)_0$ diagrams for target clusters. Filled circles denote stars with a kinematical membership probability of more than 50\%, open circles less than 50\%. Stars with a statistical significant $\Delta a$ value, are denoted with their WEBDA numbers (W no.). The solid line is the normality line and the dotted lines are the confidence intervals corresponding to 99.9\,\%.} \label{aversusbv} \end{center} \end{figure*} \addtocounter{figure}{-1} \begin{figure*} \begin{center} \includegraphics[width=165mm]{figure1b.eps} \caption{continued.} \label{aversusbv} \end{center} \end{figure*} | We presented new photoelectric $\Delta a$ photometry of 304 stars in ten open cluster fields. From a detailed kinematical analysis, we concluded that 207 stars of the sample have a membership probability higher than 50\%. The ten clusters (Feinstein~1, NGC~2168, NGC~2323, NGC~2437, NGC~2547, NGC~4103, NGC~6025, NGC~6633, Stock~2, and Trumpler~2) are young to intermediate age aggregates (5\,Myr\,$<$\,age\,$<$\,500\,Myr) with distances from 300\,pc to 1800\,pc from the Sun. Our search for chemically peculiar objects results in fifteen detections from which ten objects seem to be true member of the corresponding star cluster. Of the remaining five stars, four are probable field CP objects and one is $\lambda$ Bootis-type candidate. The objects have masses between 1.7\,M$_{\sun}$ and 7.7\,M$_{\sun}$ and are dwarf stars (luminosity class V). We discussed the already published spectral classifications in the light of our $\Delta a$ photometry and identified several misclassified CP stars. On the other hand, we were also able to establish and support the nature of known bona fide CP candidates. The newly detected CP stars, close to the zero-age-main-sequence, will help to understand the formation and evolution of this phenomenon. To this end, follow-up observations, for example, to detect and trace the local stellar magnetic fields, as well as an abundance analysis are needed. | 14 | 3 | 1403.3538 |
1403 | 1403.1198_arXiv.txt | {Primordial Black Holes (PBHs) are, typically light, black holes which can form in the early Universe. There are a number of formation mechanisms, including the collapse of large density perturbations, cosmic string loops and bubble collisions. The number of PBHs formed is tightly constrained by the consequences of their evaporation and their lensing and dynamical effects. Therefore PBHs are a powerful probe of the physics of the early Universe, in particular models of inflation. They are also a potential cold dark matter candidate.} | Primordial Black Holes (PBHs) are black holes which may form in the early Universe~\cite{Zeldovich:1967,Hawking:1971ei}. There are various formation mechanisms: the collapse of large density fluctuations (Sec.~\ref{sec-form-dens}), cosmic string loops~\cite{Hawking:1987bn} (Sec.~\ref{sec-form-cs}) or bubble collisions~\cite{Crawford:1982,Hawking:1982} (Sec.~\ref{sec-form-bubble}). In most cases the PBH mass, $M_{\rm PBH}$, is roughly equal to the horizon mass, $M_{\rm H}$, at the formation epoch (e.g. Ref.~\cite{Carr2005}): \begin{equation} M_{\rm PBH} \sim M_{\rm H} \sim \frac{c^3 t}{G} \sim 10^{15} \left( \frac{t}{10^{-23} \, {\rm s} } \right) \, {\rm g} \,. \end{equation} For instance PBHs formed at the QCD phase transition at $t \sim 10^{-6} \, {\rm s}$ would have mass of order a solar mass, $M_{\rm PBH} \sim M_{\odot}= 2 \times 10^{30} \, {\rm kg}$. As famously realised by Hawking~\cite{Hawking:1974rv}, PBHs radiate thermally and hence evaporate on a timescale, $\tau(M_{\rm PBH})$, (e.g. Ref.~\cite{Carr2005}): \begin{equation} \tau(M_{\rm PBH}) \sim \frac{\hbar c^4}{G^2 M_{\rm PBH}^3} \sim 10^{10} \left( \frac{M_{\rm PBH}}{10^{15} \, {\rm g}} \right)^3 \, {\rm Gyr} \,. \end{equation} PBHs with $M_{\rm PBH} \sim 10^{15} \, {\rm g}$ will be evaporating today and their abundance is constrained by the flux of $\gamma$-rays~\cite{Page:1976wx} (Sec.~\ref{sec-abund-evap-gamma}). Lighter PBHs evaporated in the past and are constrained by the effects of their evaporation products on Big Bang Nucleosynthesis~\cite{Vainer:1977,Zeldovich:1977} (Sec.~\ref{sec-abund-evap-nucleo}) and the present day density of any stable relic particles~\cite{MacGibbon:1987my} (Sec.~\ref{sec-abund-evap-relic}). Heavier PBHs are stable and their abundance is limited by their lensing (Sec.~\ref{sec-abund-lens}) and dynamical~\cite{Carr:1997cn} (Sec.~\ref{sec-abund-dyn}) effects and also their effects on various other astrophysical processes and objects (Sec.~\ref{sec-abund-other}). Since PBHs are matter, the fraction of the total energy density in the form of PBHs increases proportional to the scale factor, $a$, during radiation domination. Therefore the constraints on the fraction of the initial energy density in the form of PBHs, $\beta(M_{\rm PBH})= \rho_{\rm PBH}/\rho_{\rm tot}$, are very tight, lying in the range $10^{-5} - 10^{-30}$. Cosmological inflation, a period of accelerated expansion in the early Universe, may have generated the primordial fluctuations from which galaxies and large scale structure form (see e.g. Ref.~\cite{Lyth:2009zz}). The constraints on the initial fraction of the energy density in the form of PBHs can be translated into limits on the primordial power spectrum of density perturbations on small scales, and can therefore be used to constrain models of inflation~\cite{Carr:1994ar} (Sec.~\ref{sec-inf}). Finally there is extensive astronomical and cosmological evidence that the majority of the matter in the universe is in the form of non-baryonic cold dark matter (CDM) (see e.g. Ref.~\cite{Bertone:2004pz}). Since PBHs form before nucleosynthesis they are non-baryonic and therefore a candidate for the CDM (Sec.~\ref{sec-dm}). | 14 | 3 | 1403.1198 |
|
1403 | 1403.8070_arXiv.txt | {Thanks to the heroic observational campaigns carried out in recent years we now have large samples of metal-poor stars for which measurements of detailed abundances exist. In particular, large samples of stars with metallicities $-$5 $< $ [Fe/H] $ < -$1 and measured abundances of Sr, Ba, Y, and Eu are now available. These data hold important clues on the nature of the contribution of the first stellar generations to the enrichment of our Galaxy.}{We aim to explain the scatter in Sr, Ba, Y, and Eu abundance ratio diagrams unveiled by the metal-poor halo stars.} {We computed inhomogeneous chemical evolution models for the Galactic halo assuming different scenarios for the r-process site: the electron-capture supernovae (EC) and the magnetorotationally driven (MRD) supernovae scenario. We also considered models with and without the contribution of fast-rotating massive stars (spinstars) to an early enrichment by the s-process. A detailed comparison with the now large sample of stars with measured abundances of Sr, Ba, Y, Eu, and Fe is provided (both in terms of scatter plots and number distributions for several abundance ratios).} {The scatter observed in these abundance ratios of the very metal-poor stars (with [Fe/H] $<-$2.5) can be explained by combining the s-process production in spinstars, and the r-process contribution coming from massive stars. For the r-process we have developed models for both the EC and the MRD scenario that match the observations.}{With the present observational and theoretical constraints we cannot distinguish between the EC and the MRD scenario in the Galactic halo. Independently of the r-process scenarios adopted, the production of elements by an s-process in spinstars is needed to reproduce the spread in abundances of the light neutron capture elements (Sr and Y) over heavy neutron capture elements (Ba and Eu). We provide a way to test our suggestions by means of the distribution of the Ba isotopic ratios in a [Ba/Fe] or [Sr/Ba] vs. [Fe/H] diagram.} | The site for the production of the heaviest elements built via rapid neutron captures (the so-called r-process) is still unclear, and has been driving large theoretical efforts \citep[e.g.][]{Goriely13,Nakamura13, Wanajo13, Qian12, Winteler12, Arcones11, Thielemann11}. The r-process requires high neutron fluxes (i.e., the high neutron-to-seed ratios needed for the r-process to occur). The site of the r-process must also reproduce the abundance patterns seen in strongly r-process-enhanced metal-poor stars (which match the solar r-process pattern in a wide range of elements), and hence enrich the interstellar medium (ISM) on short timescales. In our latest work \citep{Cescutti13}, we studied the impact on the chemical evolution of the Galactic halo of the s-process generated by massive fast-rotating metal-poor stars (spinstars). We showed that spinstars can explain the long-standing problem of the [Sr/Ba] spread in the Galactic halo \citep[for alternative scenarios see][]{ArcoMonte11,AokiSuda13}. However, to achieve this, it was necessary to consider the contribution of an r-process to the chemical enrichment. In \citet{Cescutti13}, we followed the scenario described by \citet{WNJ09}, where the r-process occurs in a relative narrow mass range (8-10~\msun). We underline that these assumptions on the r-process did not influence our main result, which was to show how spinstars can explain the Sr/Ba ratios. In the present work we verify this by testing other r-process scenarios. First, this offers the opportunity to confirm the important role of spinstars not only in the chemical evolution of the light elements such as C and N \citep{Chiappini06,Chiappini08}, but also for the heavier elements \citep{Pigna08, Chiappini11, Frisch12}. Second, we aim to find observational constraints on the nature of the r-process by studying chemical evolution models of the earliest phases of the chemical enrichment of our Galaxy. In the present work we compute a new chemical evolution model that includes the site of production of r-process recently suggested by \citet{Winteler12}. These authors suggested that magnetorotationally driven supernovae might be the source of the r-process in the early Galaxy. These SN explode in a rare progenitor configuration that is characterized by a high rotation rate and a strong magnetic field necessary for the formation of bipolar jets. The findings of \citet{Winteler12} suggest that the second and third peaks of the solar r-process distribution can be reproduced well. Here, we test whether this site for the r-process provides an enrichment for the earliest phases of the Galactic chemical evolution consistent with the abundances observed in metal-poor halo stars. We anticipate that the results we obtain in the Galactic halo for the Winteler scenario cannot clearly be distinguished from the r-process scenario based on electron-capture SNe used in \citet{Cescutti13}, at least not before a substantial improvement in the number of stars measured in the Galactic halo has provided stronger constraints. Therefore, spinstars play a key role in this scenario as well, and the oldest halo stars are formed from an ISM enriched by both r- and s-processes. A clear prediction of both models is that EMP stars with a high [Sr/Ba] ratio should be almost entirely enriched by the s-process. This prediction is original and differs from the other possible scenarios in which the spread in [Sr/Ba] ratio is explained by a weak r-process \citep{ArcoMonte11} or a truncated r-process \citep{AokiSuda13}. A way to distinguish in this mixture between s-process and r-process in the early phase of the Galaxy formation is to examine the prediction of our models for the Ba isotopes. The s-process preferentially produces even isotopes, whereas the r-process produces approximately the same amount of odd and even isotopes. According to nucleosynthesis calculations \citep{Arla99}, it is expected that an odd fraction of Ba isotopes ($f_{odd}$) = 0.11 $\pm$ 0.01 occurs in the case of a pure s-process, and an {\it fodd} = 0.46 $\pm$ 0.06 in the case of pure r-process. \citet{Magain1995} measured for the first time the isotopic ratio of a very bright halo star, HD 140283, finding an s-process signature ([Sr/Ba]=0.9), which agrees with our theoretical results. However, his results have been challenged and still need to be confirmed \citep{Lambert2002,Collet2009,Gallagher10}. The biggest challenge is to correctly take into account the 3D effects on the line formation. More recently, \citet{Gallagher12} have again attempted to measure isotopic ratios in other metal-poor stars, but all their candidates are expected to be s-process dominated. Although the measurement of the Ba isotopic ratio is not trivial, it is feasible, and we intend to provide our results to compare them with future measurements, which will provide an important test for our models. The paper is organized as follows: in Section 2 we describe the observational data; Section 3 describes the chemical model and the adopted stellar yields. In Section 4 our results are presented, and in Section 5 we summarize our conclusions. | We have developed inhomogeneous chemical evolution models for the Milky Way halo. We adopted different hypotheses for the site of the r-process, and also included an early production of the s-process by fast-rotating massive stars (spinstars). We compared our predictions for Ba, Y, Sr, and Eu with the abundance patterns of very metal-poor stars for these elements. Our main conclusions can be summarized as follows: \begin{itemize} \item Independently of the r-process scenarios adopted, the spinstars production of s-process is needed to reproduce the spread in the light neutron capture elements (Sr and Y) over the heavy neutron capture elements (Ba and Eu). \item Our two best models based on two different r-process scenarios when coupled with the spinstars s-process production are able to reproduce the observational data for a set of four neutron capture elements (Sr, Y, Ba, and Eu). \item The r-process scenarios are not clearly distinguishable with the present data for the Galactic halo. Both scenarios reproduce the fraction of r-process-rich stars fairly well. The abundance measurements of [Eu/Fe] at intermediate metallicity $-$2$<$[Fe/H]$<-$1 tend to favor the MRD scenario. More data at this intermediate metallicity might provide an important constraint. \item We predict the contribution of spinstars to be more pronounced in specific zones in the chemical abundance ratios, for example, at high [Sr/Ba] at [Fe/H]$<-$2 or at low [Ba/Fe] for the same metallicity range. In these zones the stars are expected to present an s-process signature in the Ba isotopic ratios, leading to low odd isotopic ratios possibly measurable by the hyperfine splitting of the Ba lines (in particular, the line at 455.4 nm). \item The change of the rate of MRD SNe produces differences in the predictions of the model that can be distinguished using the distribution functions. This comparison confirms that the rate we have chosen produces results that better agree with the observational data. There is also a caveat, since we can conclude this only by assuming no bias in the observational data depending on the abundance of neutron capture elements; so not only a greater number of measured stars but also observational data without (or with known) bias are needed to proceed a step forward in the understanding of the sources of the r-process. \end{itemize} Finally, we would like to point out that a possible method to distinguish between these two scenarios might be to apply these same r-process nucleosynthesis prescriptions to another Galactic component. A different star formation history provides new constraints because a fast evolution of the metallicity can help to enhance the differences in the timescales between the two scenarios considered here. We will study this possibility in a future work by comparing our model predictions for the halo and the bulge. | 14 | 3 | 1403.8070 |
1403 | 1403.6895_arXiv.txt | Afterglow jets are Rayleigh-Taylor unstable and therefore turbulent during the early part of their deceleration. There are also several processes which actively cool the jet. In this letter, we demonstrate that if cooling significantly increases the compressibility of the flow, the turbulence collides with the forward shock, destabilizing and corrugating it. In this case, the forward shock is turbulent enough to produce the magnetic fields responsible for synchrotron emission via small scale turbulent dynamo. We calculate light curves assuming the magnetic field is in energy equipartition with the turbulent kinetic energy and discover that dynamic magnetic fields are well-approximated by a constant magnetic-to-thermal energy ratio of $1\%$, though there is a sizeable delay in the time of peak flux as the magnetic field turns on only after the turbulence has activated. The reverse shock is found to be significantly more magnetized than the forward shock, with a magnetic-to-thermal energy ratio of order 10\%. This work motivates future Rayleigh-Taylor calculations using more physical cooling models. | \label{sec:intro} Magnetized relativistic jets are important astrophysical phenomena, most notably in the context of gamma ray bursts (GRBs), but also in active galactic nuclei and tidal disruption events. As a result, the dynamics of relativistic jets have been studied extensively, often in terms of the GRB central engine \citep{1999ApJ...524..262M, 2000ApJ...531L.119A, 2007ApJ...665..569M, 2009MNRAS.397.1153K, 2012MNRAS.423.3083M, 2013ApJ...767...19L}, but also in the largely engine-independent afterglow phase when ejecta accelerated by the central engine has transferred its energy to a collimated blast wave \citep{1999ApJ...525..737R, 2000ApJ...541L...9K, 2001grba.conf..312G, 2002ApJ...571..779P, 2002bjgr.conf..146L, 2003ApJ...586..356Z, 2005ApJ...626..966P, 2007RMxAC..27..140G, 2010ApJ...722..235V, 2012ApJ...751...57D}. Given these extensive studies, there are still many fundamental questions which remain unanswered. For example, afterglow jets are thought to be magnetized, as synchrotron emission necessitates a strong magnetic field, yet no clear mechanism has been demonstrated which robustly generates such a field. Additionally, current jet models are parameterized by a small handful of parameters \citep{2012ApJ...747L..30V}, which would seem to suggest a straightforward standardization of GRB afterglow light curves. However, GRB afterglows display a great deal of variety and variability, especially at early times, hence there likely exist additional important elements missing from simplified hydrodynamical models. One avenue which potentially addresses these issues is vorticity generation behind the forward shock. Vorticity could both amplify magnetic fields via turbulent dynamo and produce variability in GRB light curves. Understanding where vorticity comes from and how much is present will help to complete the picture of how relativistic jets generate afterglow emission. The source of vorticity is still unclear, but several mechanisms have been suggested. One possibility is small-scale Weibel instabilities in the plasma particles making up the shock itself \citep{2008ApJ...673L..39S}. However, such instabilities may have a short range of influence. Alternatively, vorticity can be generated when a shock overtakes high-density clumps in the interstellar medium (ISM) \citep{2007ApJ...671.1858S, 2008JFM...604..325G}, but it is unclear whether large enough clumps exist to make this a robust mechanism. In this work, we consider the vorticity generated by Rayleigh-Taylor (RT) instability, as first suggested by \cite{2009ApJ...705L.213L}. After a GRB ejects a relativistic flow (ejecta), it expands and its thermal energy drops adiabatically until it is subdominant to the kinetic energy. The ejecta then coasts and becomes a very thin shell with width $\Delta r / r \sim 1 / \Gamma^2$, where $\Gamma$ is the Lorentz factor \citep{1999ApJ...513..669K}. When deceleration finally occurs, shocks are generated at the interface between ejecta and ISM. A forward shock pushes its way into the ISM, and a reverse shock pushes its way back into the ejecta. In the heated region between these two shocks resides the contact discontinuity, separating ejecta from ISM. This contact discontinuity is Rayleigh-Taylor unstable. Nonrelativistic RT-unstable outflows were first studied by \cite{1992ApJ...392..118C}, both analytically and numerically. \cite{1996ApJ...465..800J} later performed a two-dimensional magnetohydrodynamics calculation which demonstrated how magnetic fields tend to align themselves along RT fingers. More recently, \cite{2010AnA...509L..10F} and \cite{2011MNRAS.415...83W} have demonstrated the importance of various microphysical processes at the shock front, and \cite{2010AnA...515A.104F} has performed 3D numerical calculations. To extend the nonrelativistic results into the relativistic regime, \cite{2010GApFD.104...85L} performed a stability analysis on the two-shock solution \citep{2006ApJ...645..431N} and found linear growth rates which could potentially be large enough to impact the forward shock. In the first numerical studies of the relativistic case, \cite{2013ApJ...775...87D} found that Rayleigh-Taylor generates turbulence which could amplify magnetic fields to within a few percent of equipartition with the thermal energy density. However, in that work we found the turbulence remained confined within a region behind the forward shock and did not impact the forward shock, though turbulence did penetrate part of the energetic post-shock region. \\ In this letter we demonstrate that it is possible for the Rayleigh-Taylor turbulence to collide with the forward shock. As a result, the shock is perturbed and corrugated and significant turbulence is present everywhere behind it. This turbulence persists for a long time, until the shock becomes nonrelativistic, possibly due to the non-universality of the Blandford-McKee solution \citep{2000astro.ph.12364G}. The key ingredient allowing the turbulence to collide with the forward shock is a softer equation of state. In fact, a softened equation of state has already been invoked in the nonrelativistic case to explain how Rayleigh-Taylor fingers can catch up to the forward shock in Type 1A supernovae \citep{2001ApJ...560..244B, 2010AnA...509L..10F, 2011MNRAS.415...83W}. In this case, the collision of the ejecta with the forward shock is apparent in images of supernova remnants \citep{2005AnA...433..229V, 2011AnA...532A.114K, 2012SSRv..173..369H}. A softened equation of state can result in a reduced pressure gradient in the forward shock. This pressure gradient acts as a restoring force keeping the Rayleigh-Taylor fingers behind the forward shock, so if the pressure gradient is reduced significantly, the turbulence can collide with the forward shock. Therefore, if cooling removes a non-negligible fraction of the internal energy, it can reduce this pressure gradient, facilitating the collision of the turbulence with the shock. There are several reasons that the equation of state of GRB jets is expected to be softer than an adiabatic $4/3$ law. Cosmic ray acceleration at the forward shock can carry a significant amount of thermal energy, cooling the shock \citep{2011AnA...532A.114K, 2012ApJ...749..156O}, which may effectively result in a softer equation of state. Additionally, the shock is highly radiative, so that photon production also provides cooling. Photon cooling can potentially impact the dynamics; for example, GRB 080319B was estimated to emit $\sim 10^{51}$ ergs in X-rays \citep{2008Natur.455..183R, 2009ApJ...691..723B}, which should be a non-negligible fraction of the energy in the blastwave. If this cooling is responsible for reduced pressure in the shock front, this might require some coupling between the leptons and the baryons. The cooling from cosmic rays has less certain observational constraints in the GRB context, but it has been found to be important dynamically for the nonrelativistic case of supernova remnants \citep{2011AnA...532A.114K, 2012APh....39...33A}. In order to elucidate the effects of increased compressibility, we compare two Rayleigh-Taylor setups differing only in the adiabatic index: the usual relativistic $\gamma = 4/3$, and $\gamma = 1.1$ representing the case where cooling is dynamically important. It is straightforward to see that this reduced index should result in a lower pressure for fixed internal energy, since $P = (\gamma-1)\epsilon$. The difference between $4/3$ and $1.1$ can be envisioned as the difference between $P = \epsilon/3$ and $P = \epsilon/10$, so that for a given internal energy the pressure is reduced by about a factor of $3$. Such a change can be roughly interpreted as losing $2/3$ of the thermal energy to cooling. Therefore, this adiabatic index models a system which loses a non-negligible fraction of its internal energy. As we shall see, the reduced pressure allows Rayleigh-Taylor fingers to impact the forward shock, and generate plenty of vorticity for the entire time the shock is relativistic. Thus, the shock will continue to be corrugated as long as the relevant cooling processes are effective at softening the equation of state. The choice of $\gamma = 1.1$ to represent effects of cooling is a proof-of-concept which motivates further study using a more accurate cooling prescription. \\ | \label{sec:disc} We demonstrate that Rayleigh-Taylor instabilities can generate vorticity and magnetic fields in GRB afterglow jets. The only ingredient necessary for this mechanism is an equation of state which is softer than the usual $\gamma = 4/3$ model. This soft equation of state represents a mechanism for energy loss which reduces the pressure in the forward shock so that RT fingers can collide with it. Several processes occur at the forward shock which act to cool it; cosmic rays and radiation, for example, may carry significant energy away from the shock. Regardless of what cools the shock front, this seemingly benign change to the dynamics can completely change the structure and magnetization of the blastwave. We estimate a magnetic energy fraction of $\epsilon_B \sim 1\%$ in the forward shock, and $\sim 10\%$ in the reverse shock. We show that the choice $\epsilon_B =$ constant $= 0.01$ agrees surprisingly well with late-time afterglow calculated assuming a local value of $\epsilon_B = \epsilon_{turb}$, although this result may change when more accurate models for cooling are employed in the calculation. Finally, we show that this magnetic field does not turn on until an observer time later than the deceleration time (a factor of $5$ later in our $\Gamma = 30$ case). This occurs at a time when the shock has decelerated to a Lorentz factor lower than its original value. This could be related to the cause of observed early-time plateaus in GRB afterglows. | 14 | 3 | 1403.6895 |
1403 | 1403.5258_arXiv.txt | In this paper we study the evolution of the equation of state of viscous dark energy in the scope of Bianchi type III space-time. We consider the case when the dark energy is minimally coupled to the perfect fluid as well as direct interaction with it. The viscosity and the interaction between the two fluids are parameterized by constants $\zeta_{0}$ and $\sigma$ respectively. We have made a detailed investigation on the cosmological implications of this parametrization. To differentiate between different dark energy models, we have performed a geometrical diagnostic by using the statefinder pair $\{s, r\}$. | Recent Astronomical and astrophysical observations indicate that we live in an accelerating expanding universe (Perlmutter et al. 1997, 1999; Riess et al. 1998, 2001; Tonry et al. 2003; Tegmark et al. 2004). This fact opens a very fundamental question regarding to the source which can produce such an accelerating expansion. Since the ordinary matter (energy) generates an attractive gravitational force, there should be a kind of un-known, non-baryonic source of energy with negative pressure in order to make the expansion of the universe to be accelerating. Of course, the amount of this energy should be larger than the ordinary matter (energy) since first a fraction of this force has to counterbalance the attractive force of ordinary matter and then the rest give rise to acceleration. According to the recent observations we live in a nearly spatially flat Universe composed of approximately $4\%$ baryonic matter, $22\%$ dark matter and $74\%$ dark energy (DE). We know that the ultimate fate of our universe will be determined by dark energy but unfortunately our knowledge about its nature and properties is still very limited. It is not even known what is the current value of the dark energy effective equation of state (EoS) parameter $\omega^{X} = p^{X}/\rho^{X}$. We only know that a kind of exotic energy with negative pressure drives the current accelerating expansion of the universe; and although it dominates the present universe, it was small at early times. This is why so far many candidates have been proposed for dark energy including: cosmological constant ($\omega^{X}=-1$) (Weinberg 1989; Carroll 2001; Padmanabhan 2003; Peebles \& Ratra 2003), quintessence ($-1<\omega^{X}<-\frac{1}{3}$) (Wetterich 1988; Ratra \& Peebles 1988), phantom ($\omega^{X}<-1$) (Caldwell 2002), quintom ($\omega^{X}<-\frac{1}{3}$) (Feng et al. 2005), interacting dark energy models, Chaplygin gas as well as generalized Chaplygin gas models (Srivastava 2005; Bertolami et al. 2004; Bento et al. 2002; Alam et al. 2003), and etc. A cosmological constant (or vacuum energy) seems to be a proper candidate for dark energy which can explain the current acceleration in a natural way, but it would suffer from some theoretical problems such as the fine-tuning and coincidence problems. Quintessence and phantom dark energy models are provided by scalar fields. These models are also encounter to some problems. For example, since recent observations (Hinshaw et al. 2009; Komatsu et al. 2009; Copland et al. 2006; Perivolaropoulos 2006) indicate that $\omega^{X}<-1$ is allowed at $68\%$ confidence level, quintessence with $\omega^{X}>-1$ may not be a proper candidate as dark energy. Phantom dark energy models are also suffer from some fundamental problems, such as future singularity problem called Big Rip (Caldwell ey al. 2003; Nesseris \& Perivolaropoulos 2004) and the ultraviolet quantum instabilities problem (Carroll et al. 2003). Since recent cosmological observations mildly favor models with a transition from $\omega^{X}>-1$ to $\omega^{X}<-1$ near the past (Riess et al. 2004; Choudhury \& Padmanabhan 2005), a combination of quintessence and phantom in a unified model called quintom has been proposed (Feng et al. 2005).\\ Recently the dissipative DE models in which the negative pressure, responsible for the current acceleration, is an effective bulk viscous pressure have been proposed in order to avoid the occurrence of the big rip (McInnes 2002; Barrow 2004). The general theory of dissipation in relativistic imperfect fluid was first suggested by Eckart (1940), Landau and Lifshitz (1987). Although this is only the first-order deviation from equilibrium and may suffer from causality problem, one can still apply it to phenomena which are quasi-stationary, i.e. slowly varying on space and time characterized by the mean free path and the mean collision time. It is worth to mention that the second-order causal theory was obtained by Israel (1976) and developed by Israel and Stewart (1976). The effect of bulk viscosity on the background expansion of the universe has been investigated from different points of view (Cataldo et al. 2005; Bervik \& Gorbunova 2005; Szydlowski \& Hrycyna 2007; Singh 2008; Feng \& Zhou 2009; Oliver et al. 2011; Amirhashchi 2013a,b). There are also some astrophysical observational evidences indicate that the cosmic media is not a perfect fluid (Jaffe et al. 2005). Therefore, the viscosity effect could be concerned in the evolution of the universe. The role of viscous pressure as an agent that drives the present acceleration of the Universe has also been studied in Refs (Zimdhal et al. 2001; Balakin et al. 2003). The possibility of a viscosity dominated late epoch of the Universe with accelerated expansion was already mentioned by Padmanabhan and Chitre (1987).\\ Interaction between dark energy and dark matter (DM) is a proposal suggested as a possible solution to the coincidence problem (Setare 2007; Jamil \& Rashid 2008, 2009; Cimento et al. 2003). Moreover, DE-DM interaction provides the possibility of detecting the dark energy in a natural way. It is worth to mention that the possibility of such an interaction has been supported by the recent observations (Bertolami et al. 2007; Le Delliou et al. 2007; Berger \& Shojaei 2006). Interacting dark energy models have been widely investigated in literatures (for example see Amirhashchi et al. 2011 a, b; Amirhashchi et al. 2012; Amirhashchi et al. 2013 ; Amirhashchi 2013a,b,c; Saha et al. 2012; Yadav and Sharma 2013; Yadav 2012; Pradhan et al 2011; Setare 2007a,b,c; Setare et al 2009; Sheykhi \& Setare 2010; Jamil \& farooq 2010; Zhang 2005; Sajadi \& Vodood 2008 ). A Full dynamical analysis of anisotropic scalar-field cosmology with arbitrary potentials has been studied by Fadragas et al (2013). Recently, Long Zu et al. (2014) have investigated a class of transient acceleration models consistent with Big Bang Cosmology. In this paper, we study the behavior of the viscous dark energy EoS parameter in an anisotropic space-time namely Bianchi type III universe in the following two cases: (i) when DE and DM are minimally coupled i.e there is no any interaction between these two dark components and (ii) when there is an interaction between viscous DE and DM. We parameterize the interaction by a constant $\sigma$ and viscosity by $\zeta_{0}$, then a detailed investigation of the cosmological implications of this parametrization will be provided by assuming an energy flow from DE to DM. Finally, to discriminate the different interaction parameters, as usual, a statefinder diagnostic is also performed. | In this paper we studied dark energy in the scope of anisotropic Bianchi type III space-time. We considered two cases (i) when DE and DM do not interact with each other and (ii) when there is an interaction between these two dark components. In non-interacting as well as weak interacting ($\sigma\sim 0$) cases we observed that in absence of viscosity, dark energy EoS parameter dose not cross the phantom divided line (PDL) and hence always vary in quintessence region. However, in both cases when dark energy is considered to be viscous rather than perfect, it's EoS parameter could cross the PDL depending on the values of coupling constant $\sigma$ and bulk viscosity coefficient $\zeta_{0}$. But in this case although the dark energy EoS parameter could cross PDL and vary in phantom region ultimately tends to the cosmological constant region $\omega^{de}=-1$. This special behavior of the EoS parameter is because of our choose of bulk viscosity which is a decreasing function of time (redshift) in expanding universe. It has also been shown that in both cases according to the $\Omega^{X}$-$\Omega^{m}$ phase diagram (see figs. 3, 7, 8), deviation from flat universe ($\Omega=1$) only depends on the geometric parameter $\alpha$ not to the interaction parameter $\sigma$. | 14 | 3 | 1403.5258 |
1403 | 1403.6221_arXiv.txt | Convective dynamo simulations are performed in local Cartesian geometry. We report the first successful simulation of a large-scale oscillatory dynamo in rigidly rotating convection without stably stratified layers. A key requirement for exciting the large-scale dynamo is a sufficiently long integration time comparable to the ohmic diffusion time. By comparing two models with and without stably stratified layers, their effect on the large-scale dynamo is also studied. The spatiotemporal evolution of the large-scale magnetic field is similar in both models. However, it is intriguing that the magnetic cycle is much shorter in the model without the stable layer than with the stable layer. This suggests that the stable layer impedes the cyclic variations of the large-scale magnetic field. | A grand challenge in astrophysics is to understand a self-organizing property of magnetic fields in highly turbulent flows. The solar magnetism is the front line of this area. The solar magnetic field shows a remarkable spatiotemporal coherence even though it is generated by turbulent convection operating within its interior. Our understanding on the solar magnetism has been accelerated over the past decade in response to the broadening, deepening and refining of numerical dynamo models \citep{charbonneau10,brandenburg+12,miesch12}. However, it is still unclear what dynamo mode is excited in the solar interior and how it regulates the magnetic cycle. Various geometries have been applied to the numerical dynamo modeling; global spherical shell geometry (e.g., \cite{gilman+81,brun+04,ghizaru+10,masada+13}), spherical-wedge geometry (e.g., \cite{brandenburg+07,kapyla+10}), and local Cartesian geometry (e.g., \cite{cattaneo+91, brandenburg+96}). Among them, the local Cartesian geometry is the most simplified one and is often used for distilling the physical essence of the convective dynamo process by more accurately resolving convective eddies. A long-standing goal in the numerical dynamo modeling in the local Cartesian geometry is to realize the successful simulation of self-organized and self-sustained large-scale magnetic fields, so-called ``large-scale dynamos", by rotating convection alone without mean shear flow. The mean-field dynamo theory predicts that the rigidly rotating convection can generate net helicity and then excite the large-scale dynamo even without the mean shear effect via a stochastic process, which is known as the $\alpha$-effect (\cite{moffatt+78,krause+80}). However, no evidence of the large-scale dynamo was found in earlier studies of the rigidly rotating convection (e.g., \cite{cattaneo+06,tobias+08}). \citet{kapyla+09} brought a breakthrough in the dynamo modeling in the local system. They were the first to demonstrate that the rigidly rotating convection can excite the large-scale dynamo in the local system that consists of the convection layer and the stably stratified layers. Subsequently, the oscillatory behavior of the large-scale magnetic field was reported in \citet{kapyla+13}. Since the large-scale dynamo can be excited only when the Coriolis number is large, they concluded that the absence of the large-scale dynamo in earlier studies is caused by the slow rotation speed. However, even in the sufficiently-rapid rotating convection, Favier \& Bushby (2013) could not find the evidence for the large-scale dynamo in the local system with the convection zone alone. They suggested that the essential part for the large-scale dynamo might be the stably stratified layer assumed in the model of \citet{kapyla+09} rather than the rapid rotation. Therefore, at present, the key requirement for the large-scale dynamo is still controversial. The purpose of this work is to find the evidence of the large-scale magnetic field in the system only with the convection zone in order to demonstrate that the rigidly rotating convection is a sufficient condition for the large-scale dynamo. In addition, by comparing two convective dynamo models with and without stably stratified layers, we will discuss their effect on the large-scale dynamo. | \begin{figure}[tbp] \begin{center} \includegraphics[width=73mm]{f4.eps} \end{center} \caption{Distribution of the magnetic energy on the horizontal plane at the middle of the convection zone when (a) $t=100\tau_{\rm cv}$ (first saturation phase) and (b) $t=330\tau_{\rm cv}$ (second saturation phase) for the model A.} \label{fig4} \end{figure} We performed rigidly rotating convective dynamo simulations in the local Cartesian geometry. By comparing two models with and without stably stratified layers, their effect on a large-scale dynamo was studied. We for the first time successfully simulated an oscillatory large-scale dynamo in the local system without the stable layer, whereas it was not found in the similar earlier studies (\cite{cattaneo+06,favier+13}). The absence of large-scale dynamos in earlier studies might be a simple consequence of their relatively short integration time. For the excitation of the large-scale dynamo, we should evolve the simulation for a sufficiently long time because the large-scale magnetic component is gradually built up in an order of the ohmic diffusion time. The spatiotemporal evolution of the magnetic field was similar in two models. The large-scale magnetic component was the strongest at around the middle of the convection layer and propagated from there to the upper and lower convection zones. According to the mean-field dynamo theory, the $\alpha$-effect would be solely responsible for the large-scale dynamo in our models because the $\Omega$-effect is absent in the rigidly rotating system (c.f., \cite{kapyla+13}). However, the nonlinear properties of the $\alpha$-effect dynamo in the natural rotating convection are still veiled in mystery. We will examine quantitatively whether the $\alpha$-effect dynamo can reproduce the spatiotemporal evolution of large-scale magnetic fields observed at the nonlinear saturated phase in our simulations in a subsequent paper. An intriguing finding was the difference in the oscillation period of the large-scale magnetic field between two models. The magnetic cycle was about three-times longer in the model with the stable layer than without the stable layer, although the properties of the convective motion was similar in two models. This suggests that the stably stratified layer rather impedes the cyclic variation of large-scale magnetic fields. One possible cause making a difference in the cycle period is the ejection process of the magnetic helicity, which is known to affect nonlinear properties of dynamos (c.f., \cite{blackman+00}). In the model without stable layers, the magnetic helicity can be ejected from the system via advective transport processes because the top open boundary is placed just above the convection zone. In contrast, in the model with the top stable layer, the magnetic field must be transported throughout the stable layer for the loss of the magnetic helicity. The relatively slow ohmic diffusion dominates the transport process there. We thus speculate that the smaller magnetic helicity flux dominated by the slower ohmic diffusion process is responsible for the longer cycle period in the model with stable layers. The dynamo number and thus frequency of excited dynamo mode might be different between two models. This is because the mode with longer wavelength can be allowed in the system with conducting stable layers above and below the convection zone. This might be an another possibility to explain the cycle period difference (c.f., \cite{radler+87,rudiger+03}). In any case, further simulations with varying the thickness of the stable layers and the resistivity are necessary to % elucidate the cause and will be a target of our future work. \bigskip We acknowledge the anonymous referee for constructive comments. Computations were carried on XC30 at NAOJ, and K-Computer at RIKEN. This work was supported by JSPS KAKENHI Grant number 24740125 and the joint research project of the Institute of Laser Engineering, Osaka University. | 14 | 3 | 1403.6221 |
1403 | 1403.6198_arXiv.txt | We investigate the star formation histories (SFHs) of high redshift ($3 \la z \la 5$) star-forming galaxies selected based on their rest-frame ultraviolet (UV) colors in the CANDELS/GOODS-S field. By comparing the results from the spectral-energy-distribution-fitting analysis with two different assumptions about the SFHs --- i.e., exponentially declining SFHs as well as increasing ones, we conclude that the SFHs of high-redshift star-forming galaxies increase with time rather than exponentially decline. We also examine the correlations between the star formation rates (SFRs) and the stellar masses. When the galaxies are fit with rising SFRs, we find that the trend seen in the data qualitatively matches the expectations from a semi-analytic model of galaxy formation. The mean specific SFR is shown to increase with redshift, also in agreement with the theoretical prediction. From the derived tight correlation between stellar masses and SFRs, we derive the mean SFH of star-forming galaxies in the redshift range of $3 \leq z \leq 5$, which shows a steep power-law (with power $\alpha = 5.85$) increase with time. We also investigate the formation timescales and the mean stellar population ages of these star-forming galaxies. Our analysis reveals that UV-selected star-forming galaxies have a broad range of the formation redshift. The derived stellar masses and the stellar population ages show positive correlation in a sense that more massive galaxies are on average older, but with significant scatter. This large scatter implies that the galaxies' mass is not the only factor which affects the growth or star formation of high-redshift galaxies. | Knowledge of the physical properties --- such as stellar masses, star-formation rates (SFRs), and stellar population ages --- of galaxies is indispensable for understanding the evolution and formation of galaxies. Thanks to the last decade's boom in the panchromatic observation of remote galaxies, which has benefited from the development of powerful facilities, including the $Hubble$ $Space$ $Telescope$ ($HST$), a great advance has been made in the study of high-redshift galaxies as well as in our understanding of galaxy evolution. This advance in multi-wavelength studies of high-redshift galaxies is expected to accelerate in near future with the arrival of powerful space telescopes, such as $James$ $Webb$ $Space$ $Telescope$, as well as large ground facilities, including the Giant Magellan Telescope and the Thirty Meter Telescope. However, constraining the star-formation histories (SFHs) of galaxies from observation is not an easy task, even with these panchromatic data over a wide wavelength range. The difficulty comes from several factors : for example, the degenerate effects between SFHs and other properties of galaxies --- such as dust extinction, metallicity, and redshift --- on the overall spectral energy distributions (SEDs) of galaxies, and the fact that the observed SEDs can be easily dominated by the light from massive, young stars, readily concealing the old stellar populations. One way to derive physical inferences from the multi-band photometry is to compare the distribution of galaxy colors and magnitudes to the predictions of theoretical models of galaxy formation \citep[e.g.][]{som01,som08,som12,idz04,nag05,men06,nig06,fon09}. Another way is to derive physical parameters, such as stellar masses, SFRs, metallicities, and constraints on SFHs, by fitting the SED of each individual galaxy. This method is very useful in studying galaxies' physical parameters or stellar population properties of unresolved galaxies, and has been used in analyzing various galaxy populations over a wide redshift range \citep[][just to name a few]{saw98,pap01,sha01,sal05,guo12}. In this SED-fitting analysis, we should inevitably make assumptions on several properties of galaxies --- including their stellar initial mass function (IMF), chemical composition and its evolution, dust-attenuation law, and SFHs of galaxies. The assumption on each of these properties can affect the SED-fitting results, and has been studied by several authors --- for example, the effects of IMF are studied by \citet{pap01} and \citet{con09}, the effects of stellar evolution model are investigated by \citet{mar06}, and the effects of metallicity evolution are studied by \citet{con09}. Detailed review on this issue can be found in \citet{con13} Among these assumptions, the effects of assumed (forms of) SFHs have been studied by \citet[][hereafter L10]{lee10}. In L10, comparing with the results from \citet[][hereafter L09]{lee09}, we have extensively analyzed the effects of assumed SFHs on the SED-fitting results of $3 < z < 6$ Lyman-break galaxies (LBGs) by analyzing mock LBGs from semi-analytic models (SAMs) of galaxy formation. This analysis has revealed that the assumptions about SFHs can significantly bias the inferences about stellar-population parameters: particularly ages and SFRs. Also, the results of L09 and L10 --- combined with the prediction from SAMs of galaxy formation (represented in L10) --- suggest that the SFHs of star-forming galaxies at this redshift range ($3 < z < 6$) increase with time --- in contrast to the assumption of declining SFHs as done in some previous works. For observed galaxies, \citet{mar10} show that the SEDs of $z \sim 2$ ($BzKs$-selected) star-forming galaxies are fitted better with the exponentially increasing SFHs rather than exponentially decreasing SFHs. \citet{pap11} have analyzed the evolution of average SFRs of high-redshift galaxies, by studying the high-redshift galaxy samples with constant comoving number density, n=$2 \times 10^{-4}$ Mpc$^{-3}$, from $z=8$ to $z=3$. By following the galaxy samples with same comoving number density, they can study the evolution of SFRs and stellar masses of galaxies, which may be connected as the predecessors and their descendants. From this study, they show that the average SFR of high-redshift galaxies increases as a power law with decreasing redshift (with best-fit power $\alpha = 1.7$). Motivated by the evidence favouring increasing SFHs, we analyze the SEDs of $3 \leq z \leq 5$ LBGs observed in the southern field of the Great Observatories Origins Deep Survey (GOODS-S). Specifically, we focus on the SFHs of these observed galaxies, and examine if assuming the increasing SFHs would be more appropriate, providing better constraints on the stellar population properties of the observed high-redshift galaxies than assuming the declining SFHs in SED-fitting analysis. The relatively tight correlation between stellar masses and SFRs for star-forming galaxies at redshifts up to $z \sim 2$ has come to be known as the star-forming main sequence \citep{dad07,elb07,noe07,pan09}. Over a wide redshift range, this correlation remains tight with near unity slope while the normalization evolves quickly with redshift. This tight correlation of galaxies' star-formation activities with their already-formed stellar masses, the nearly invariant slope of this correlation, and the evolution of its normalization in a broad redshift range must impose important constraints on the formation histories of these star-forming galaxies (for example, see \citet{ren09} and \citet{saw12}). In this work, we investigate this correlation and its evolution in the redshift range between $z \sim 5$ and $z \sim 3$, and derive meaningful constraints on the formation histories of the star-forming galaxies at this redshift range. Regarding the SFHs of high-redshift galaxies, another interesting issue is the significant disagreement regarding the evolution of the specific SFR (SSFR; defined as SFR/$M_{*}$) at $z \gtrsim 2$ between the observational results and the prediction from galaxy formation models. SAMs of galaxy formation predict that the SSFR increases with redshift beyond $z \gtrsim 2$ \citep[e.g.,][]{dut10}, while observational results show a plateau over a wide range of redshift, $2 \lesssim z \lesssim 7$ \citep[e.g.,][]{gon10}. Recently, there have been several works trying to resolve this tension. For examples, \citet{bou12} have re-calculated the SSFRs based on their new estimation of UV spectral slope ($\beta$) and dust attenuation, reporting higher values of SSFRs than previously estimated. Another improvement comes from the realization of the fact that the estimation of stellar mass can be affected by the contribution from the nebular emission \citep[e.g.,][]{deb12,gon12,sta13}. Even though the derived SSFR values show some discrepancies with each other, the results of these works commonly indicate that the stellar masses can be overestimated (thus SSFRs can be underestimated) without taking into account this effect. In this work, we explore this sharp contrast between model and observation, which imposes a challenge for our understanding of galaxy evolution. Our investigation of SSFR evolution is based on the consistent estimation of stellar masses and SFRs (thus, SSFRs also) from same SED-fitting (instead of estimating stellar mass and SFR separately, applying different assumptions), especially powered by more realistic assumption about SFHs of galaxies. In Section 2, we describe the observational data and sample selection. The details of SED-fitting procedure and the results of this SED-fitting are explained in Section 3. In Section 4, we analyze the results in more detail, especially focusing on the SFHs of high-redshift galaxies and the SSFR evolution at high redshift, and we analyze the ages of these galaxies in Section 5. We summarize our results in Section 6. Throughout the paper, we adopt a flat $\rm{\Lambda}$CDM cosmology, with ($\Omega_{m}, \Omega_{\Lambda}$) = (0.3,0.7), and $H_{0}$ = 70 $km$ $s^{-1}$ $Mpc^{-1}$. All magnitudes are given in AB magnitude system \citep{oke74} | In this work, we analyze the broadband photometric SEDs of observed GOODS-S star-forming galaxies, selected by their rest-frame UV colors. Through the detailed analysis of the SED-fitting results of these observed galaxies with two different assumptions about their SFHs --- the exponentially declining SFHs and the increasing SFHs, we examine the representative SFHs of these high-redshift star-forming galaxies. We also compare our results to the theoretical predictions from SAM for the additional constraints on the SFHs of high-redshift star-forming galaxies. Our main results are summarized as follows: \begin{enumerate} \item {The comparison of $\chi^2$ values from SED-fittings with different SFHs indicates that the synthetic galaxy spectral templates with the increasing SFH assumption provide better match to the observed SEDs of GOODS-S LBGs than the templates with the exponentially declining SFH assumption do.} \item {The stellar masses and the SFRs of observed GOODS-S LBGs show a tight correlation, which is in good agreement with the predictions from SAM of galaxy formation as well as with the observational results at lower redshift, when we derive these quantities through the SED-fitting with increasing SFHs. If we derive the masses and the SFRs through the SED-fitting with exponentially declining SFHs, the SFRs and stellar masses show much weaker correlation with large scatter. This is thus an additional support for the plausibility of rising SFHs for $z \geq 3$ LBGs.} \item {From the observed tight correlation between galaxies' stellar masses and SFRs in the redshift range of $3 \leq z \leq 5$, we can also deduce the average SFH of these continuously star-forming galaxies in the given redshift range. From the evolution of the mean SSFR values, we infer the average SFH of galaxies at this redshift range as $\sim (t/1.16 \rm{Gyr})^{5.85}$ --- i.e., the average SFH of continuously star-forming galaxies, when the correction for the nebular emission included, increases steeply from $z \simeq 5$ to $z \simeq 3$ during about one billion years.} \item {Our measured SSFR values show an increase with redshift from $z \sim 3$ to $z \sim 5$, indicating continuous increase from $z \sim 2$. While contrary to the previous results, this increasing trend is well consistent with recent results, such as \citet{deb12}, \citet{gon12}, and \citet{sta13}. The SSFR values, which are estimated in a consistent manner from SED-fitting, as well as its increasing trend are in good agreement with the prediction from the SAM.} \item {While many of our observed LBGs lie on the upper envelope in the redshift-$t_f$ relation (i.e., have the largest possible $t_f$ at given redshift), there is a non-negligible fraction of galaxies whose $t_f$ values are much lower than the values at the envelope, especially, at redshifts $3.0 \lesssim z \lesssim 3.7$. This spread in $t_f$ implies that not all star-forming galaxies caught at redshifts $z \geq 3$ were formed coevally.} \item {Our sample of continuously star-forming galaxies follows a positive correlation between their masses and stellar population ages, in a sense that more massive galaxies are, on average older, but with significant scatter --- similar with the prediction from the SAM. While this positive correlation implies that the formation histories of these star-forming galaxies vary depending on their masses, the broad scatter in this $M_{*}$--age correlation might be an indication that mass is not the only property which affects the formation histories of galaxies. The correlations between the ages and other properties of mock LBGs from the SAM suggest that mergers have certain effects on the age-spread among galaxies with given stellar mass.} \end{enumerate} From this SED-fitting analysis of the broadband SEDs of GOODS-S LBGs, we can confirm the important speculations of L09 and L10, which are drawn from the analysis of mock galaxies from SAMs: namely, (1) the assumed form of SFHs affect the estimation of stellar population parameters of galaxies from broadband SEDs, and (2) we should assume the increasing SFHs, not the exponentially declining ones, in the analysis of UV-selected, star-forming galaxies within the redshift range of $3 \lesssim z \lesssim 5$. Any bias arising in the fitting procedure due to incorrect assumption on the SFHs would propagate to the inferences on galaxy evolution drawn from the SED-fitting analysis. The increasing SFHs inferred here from these SED-fitting arguments are qualitatively consistent with the prediction from hierarchical models of galaxy formation in $\Lambda$CDM \citep[e.g.][]{lee10,fin11}. This smooth increase of SFR at high redshift is expected if the gas accretion rate closely traces the growth rate of halos, which grow as $\dot{M_h} \varpropto M_h$, and the SFR traces gas accretion rate. Semi-analytic models and hydrodynamic simulations generically predict fairly smoothly rising SFHs at these redshifts \citep{lee10,fin07}. The generic rising form of the SFH is primarily due to the growth of structure \citep{her03}, but the {\it slope} of the SFH may be modulated by the overall efficiency of converting accreted gas into stars, which depends on the physics of star formation and stellar feedback, and may not be invariant with cosmic time. Also, this continuously rising SFH provides a natural explanation for the tight correlation between the stellar masses and the SFRs of star-forming galaxies over a wide range of redshift, at least up to $z \leq 5$, as shown in this work --- spanning more than 12 billion years, because galaxies would move along this relation as they grow in time. However, if high-redshift galaxies form their stars with decreasing SFRs with time, this $M_{*}$--SFR correlation should become too broad with decreasing redshift, unless all the LBGs were formed coevally --- which is not physically reasonable and is also disfavored from the results in Section 5. On the other hand, if the majority of star-forming galaxies form their stars in a bursty way with short duty cycles, it is very hard to explain the formation of this tight correlation. In other words, why should the bursty star-formation activity care about the stellar mass of host galaxy, which is mainly the result of past star-formation activity (see also \citet{wuy11} and \citet{saw12} for the supporting evidence for the continuous SFHs of high-redshift galaxies)? Accurate determination of the shape of the SF main sequence and its intrinsic scatter will be very valuable to constrain the efficiency of star formation in different galaxies and environment. In fact, \citet{ren09} suggested that the rising slope of the SFH and the subsequent evolution are related --- the more gentle the rise, the more prolonged the overall SF activity. We speculate that if this is indeed the case, then the best predictor of the future evolution of SF of a galaxy is its position in the main sequence relative to other galaxies, assuming that the main sequence is accurately measured. The results of Section 5 --- broad distribution of formation times, and the positive correlation between stellar masses and stellar population ages with large scatter --- are broadly consistent with the expectations from the SAMs. However, because the parameter $t$ or stellar population age is one of the most poorly constrained parameters in SED-fitting \citep[e.g.][]{pap01,lee09,lee10,gua11}, it is still premature to draw firm conclusions. While it is not easy to provide better constraints on the ages of these high-redshift LBGs, more samples from the possible future surveys which are as deep as and wider than the GOODS would be helpful by providing better statistics and enabling us to study the dependence on interaction or environment. Also, because our sample is selected based on the rest-frame UV, these may be biased toward star-forming galaxies with moderate amount of dust extinction. Inclusion of dusty star-forming galaxies, selected based on their rest-frame optical \citep[e.g.,][]{guo12}, would provide more complete samples of high-redshift star-forming galaxies, which will be helpful for understanding the formation histories of high-redshift galaxies. | 14 | 3 | 1403.6198 |
1403 | 1403.7966_arXiv.txt | We have investigated the gas content of a sample of several hundred AGN host galaxies at z$<$1 and compared it with a sample of inactive galaxies, matched in bins of stellar mass and redshift. Gas masses have been inferred from the dust masses, obtained by stacked Herschel far-IR and sub-mm data in the GOODS and COSMOS fields, under reasonable assumptions and metallicity scaling relations for the dust-to-gas ratio. We find that AGNs are on average hosted in galaxies much more gas rich than inactive galaxies. In the vast majority of stellar mass bins, the average gas content of AGN hosts is higher than in inactive galaxies. The difference is up to a factor of ten higher in low stellar mass galaxies, with a significance of 6.5$\sigma$. In almost half of the AGN sample the gas content is three times higher than in the control sample of inactive galaxies. Our result strongly suggests that the probability of having an AGN activated is simply driven by the amount of gas in the host galaxy; this can be explained in simple terms of statistical probability of having a gas cloud falling into the gravitational potential of the black hole. The increased probability of an AGN being hosted by a star-forming galaxy, identified by previous works, may be a consequence of the relationship between gas content and AGN activity, found in this paper, combined with the Schmidt-Kennicutt law for star formation. | \label{sec1} During the past decades several studies have revealed a connection between the mass of Black Holes hosted in galactic nuclei and the properties of their host galaxies \citep[e.g.][]{Magorrian98, Ferrarese00, Marconi03}. Moreover, the redshift evolution of the cosmic Star Formation Rate (SFR) and SMBH accretion rate density are very similar \citep{Boyle98,Granato01,Marconi04,Hopkins06,Silverman09,Aird10}. This ``co-evolution" has led various authors to investigate a possible connection between the presence of Active Galactic Nuclei (AGN, which traces the SMBH growth) and the star formation properties of galaxies. A correlation was indeed observed between star formation and nuclear activity in numerous works at high AGN luminosities \citep[e.g][]{Lutz08, Lutz10, Shao10, Rosario12}, while at low AGN luminosities this link is more debated, with \cite{Silverman09} reporting no significant difference in the SFR between active and inactive galaxies, while \cite{Santini12} found a slight enhancement for AGN-hosting galaxies. They also reported that the enhancement of star formation activity in AGN with respect to the bulk of inactive galaxies disappeared if quiescent galaxies were discarded, i.e. AGN are more likely hosted in star forming galaxies. The gas content is often regarded as a more fundamental property of galaxies, with respect to the SFR. The SFR is tightly related to the gas content through the Schmidt-Kennicutt relation \citep[SK relation hereafter]{Schmidt59,Kennicutt98}. Currently, one of the most favored scenarios is that the cosmic evolution of the star formation rate in galaxies is mostly a consequence, through the SK relation, of the more fundamental evolution of their (molecular) gas content \citep[e.g.][]{Obreschkow09,Lagos11}. Within the context of AGNs, gas is the fundamental ingredient both for nuclear activity and star formation. Possible differences in terms of star formation properties between AGNs hosts and inactive galaxies could be due to more fundamental differences in terms of gas content, as recently argued by \cite{Santini12} and \cite{Rosario12,Rosario13_2}. Therefore, it is most important to obtain information on the gas content of AGN host galaxies, possibly as a function of galaxy properties (e.g. stellar mass) and redshift. The molecular gas content can be inferred from the luminosity of the CO millimeter transitions, by assuming a proper CO-to-H$_2$ conversion factor. However, CO observations are very time consuming, and surveys of large samples are extremely difficult and time demanding. Alternatively, the total (molecular and atomic) gas mass can be derived from the dust content, inferred from the FIR-submm SED, by assuming a dust-to-gas ratio \citep[DGR; e.g.][]{Eales10, Leroy11,Magdis11}. The uncertainties on the dust-to-gas ratio and its dependence on metallicity are similar to those affecting the CO-to-H$_2$ conversion factor, making the two methods comparable in terms of accuracy, at least at metallicities 12+log(O/H)$>$8.0 \citep{Bolatto13, Remy-Ruyer13}. In this work we exploit the dust method for measuring the gas masses in AGN host galaxies. In particular, gas masses are obtained from the dust mass derived from the FIR SED of several hundred AGN host galaxies at z$<$1, along with a control sample of normal galaxies selected in the same stellar mass and redshift ranges. The aim of this work is to investigate differences in terms of gas content between AGN hosts and the bulk of the galaxy population (i.e. star forming and quiescent galaxies), in bins of stellar mass and redshift, to avoid potential biases caused by the dependency of the gas content on these two quantities. We make use of a stacking procedure to increase the luminosity completeness of the studied samples (\S~\ref{stack}). | Making use of the wide multiwavelength data available in the COSMOS, GOODS-S and GOODS-N fields, we selected a sample of AGN and galaxies at $z<1$. A stacking procedure on the Herschel maps was implemented to derive the average Herschel fluxes in bins of stellar mass and redshift. The stacked FIR fluxes were then used to derive the average dust mass and, under reasonable assumptions, the average gas mass of AGN hosts and normal galaxies. Finally we compared the average gas mass of AGN hosts and inactive galaxies in the same $z-\rmn{M_*}$ bins. We find that, at a given stellar mass and redshift, AGNs are hosted in galaxies much more gas-rich than inactive ones. The difference is strongest at low stellar masses, $\rmn{log(M_*/M_\odot)}<10.5$, where the gas mass in AGN hosts is on average ten times higher than in normal galaxies (a result significant at $6.5\sigma$). Significantly higher gas masses, relative to the normal galaxy population, are however also observed in AGN hosts with stellar masses higher than $\rmn{log(M_*/M_\odot)}>10.5$. Taken altogether, in nearly all stellar mass and redshift bins AGN host galaxies have higher gas content than normal galaxies; in almost half of the sample the gas fraction of AGN host galaxies is more than three times higher than in normal galaxies. Our result strongly suggests that the likelihood of having an AGN in a galaxy is primarily given by the amount of gas in the host galaxy, while dynamical triggering processes (bars, galaxy mergers and interactions) likely play a secondary role, at least in the luminosity range probed by us. This result can be interpreted in simple statistics terms that it is more likely that a gas cloud falls into the potential of the supermassive black hole if there are overall more gas clouds in the host galaxy. More elaborated models, in which secular fuelling of AGNs is caused by disk instabilities, which are stronger in more gas rich disks, are also supported by our results. | 14 | 3 | 1403.7966 |
1403 | 1403.5152_arXiv.txt | We use 3D hydrodynamical numerical simulations and show that jittering bipolar jets that power core-collapse supernova (CCSN) explosions channel further accretion onto the newly born neutron star (NS) such that consecutive bipolar jets tend to be launched in the same plane as the first two bipolar jet episodes. In the jittering-jets model the explosion of CCSNe is powered by jittering jets launched by an intermittent accretion disk formed by accreted gas having a stochastic angular momentum. The first two bipolar jets episodes eject mass mainly from the plane defined by the two bipolar axes. Accretion then proceeds from the two opposite directions normal to that plane. Such a flow has an angular momentum in the direction of the same plane. If the gas forms an accretion disk, the jets will be launched in more or less the same plane as the one defined by the jets of the first two launching episodes. The outflow from the core of the star might have a higher mass flux in the plane define by the jets. In giant stellar progenitors we don't expect this planar morphology to survive as the massive hydrogen envelope will tend to make the explosion more spherical. In SNe types Ib and Ic, where there is no massive envelope, the planar morphology might have an imprint on the supernova remnant. We speculate that planar jittering-jets are behind the morphology of the Cassiopeia A supernova remnant. | \label{sec:intro} One class of core collapse supernova (CCSN) explosion models is based on neutrino \citep{Colgate1966}, mainly the delayed neutrino mechanism (e.g., \citealt{bethe1985,Burrows1985,Burrows1995,Fryer2002,Ott2008,Marek2009,Nordhaus2010,Kuroda2012,Hanke2012,Janka2012,Bruenn2013}). However, recent 3D numerical studies have shown that the desired explosions are harder to achieve \citep{Couch2013, Jankaetal2013, Takiwakietal2013} than what 2D numerical simulations had suggested (for a summary of problems of the delayed neutrino mechanism see \citealt{Papishetal2014}). The problems of the delayed-neutrino mechanism can be overcome if there is a strong wind, either from an accretion disk \citep{kohri2005} or from the newly born neutron star (NS). Such a wind is not part of the delayed-neutrino mechanism, and most researchers consider this wind to have a limited contribution to the explosion. Another class of explosion mechanisms is the jittering-jet scenario \citep{Soker2010, Papish2011, Papish2012a, Papish2012b, PapishSoker2014, GilkisSoker2013}. Processes for CCSN explosion by jets were considered before the development of the jittering-jet scenario (e.g. \citealt{LeBlanc1970, Meier1976, Bisnovatyi1976, Khokhlov1999, MacFadyen2001,Hoflich2001, Woosley2005, Burrows2007, Couch2009,Couch2011,Lazzati2011}). However, most of these MHD models require a rapidly spinning core before collapse starts, and hence are limited to a small fraction of all CCSNe. The jittering-jet scenario posits that {\it all CCSNe are exploded by jets}. Recent observations (e.g. \citealt{Milisavljevic2013,Lopez2013,Ellerbroek2013}) show indeed that jets might have a much more general role in CCSNe than what is expected in the neutrino-driven mechanisms and mechanisms that require rapidly rotating cores. In the jittering-jets scenario the sources of the angular momentum for disk formation are the convective regions in the core \citep{GilkisSoker2013} and instabilities in the shocked region of the collapsing core, e.g., neutrino-driven convection or the standing accretion shock instability (SASI). Recent 3D numerical simulations show indeed that neutrino-driven convection and SASI are well developed in the first second after core bounce \citep{Hankeetal2013, Takiwakietal2013} and the unstable spiral modes of the SASI can amplify magnetic fields \citep{Endeveetal2012}. The spiral modes with the amplification of magnetic fields build the ingredients necessary for jets' launching. When the average specific angular momentum of the matter in the pre-collapse core is small relative to the amplitude of the specific angular momentum of these instabilities, intermittent jets-launching episodes with random directions occur. The two launching axes of the first two launching episodes define a plane. Using the FLASH numerical code we set a numerical study of the accretion pattern that is likely to be formed after the first two launching episodes. The code and numerical set-up for the 3D simulations are described in section \ref{sec:setup}. The accretion pattern following two jets-launching episodes, followed by a third episode, is described in section \ref{sec:accretion}. Our summery is in section \ref{sec:summary}. | In this study we assumed that CCSN explosions are driven by jittering-jets \citep{Papish2011}, an examined the pattern by which jets are launched. The angular momentum of the core is not large, such that the jets' axis is determined, at least in part, by stochastic processes such as instabilities and convection in the pre-collapse core \citep{GilkisSoker2013}. We conducted 3D numerical simulations with the FLASH code \citep{Fryxell2000}. In each jets-launching episode we launched one pair of bipolar jets (Fig. \ref{fig:dens-temp-05}). Under these assumptions, the directions of the two first jets-launching episode are more or less random. However, the first two episodes, if not along the same direction, define a plane. We set this plane to be $y=0$ in our study. The jets will eject mass outward along their propagation direction, leaving inflow in perpendicular directions. After the first episode the inflow is from a belt region perpendicular to the jets' axis, as demonstrated in Fig. \ref{fig:one_jets500}. After the second episode the inflow is concentrated in two opposite directions, as clearly seen in Figs. \ref{fig:NormalPlane}, \ref{fig:second_jets-40} and \ref{fig:second_jets-70}. A bipolar accretion flow with its axis perpendicular to the $y=0$ plane has been formed. An inflowing gas along a direction normal to the $y=0$ plane has an angular momentum direction within that plane. This implies that bipolar accretion flow makes it more likely that the axis of the jets in the third episode will be in the $y=0$ plane, as jets are launched along the angular momentum axis. Namely, the third episode jets' axis will be in about the same plane as the axes of the first two episodes. We simulated two directions for the second jets episode, and for each of these we simulated two directions in the same plane for a third jets episode. These four cases are summarized in Table \ref{table}. The asymmetry in the accretion flow is presented for these cases in Fig. \ref{fig:acc}, quantified by the the parameter $\xi$ defined in equation (\ref{eq:xi1}). The bipolar pattern becomes more prominent as the inflow becomes more concentrated along two opposite directions. This implies that the following launching of jets will be in, or near, the $y=0$ plane. The bubbles inflated by the jets grow and expand to all directions as they move toward lower density core gas. Eventually they will close on directions perpendicular to their axis as well, and expel the outer core and the rest of the star in all directions. In addition, some of the jets will not be exactly in the same plane, but will have some stochastic variations from the $y=0$ plane. This will also help in expelling the rest of the star in all directions. This later evolutionary phase will be studied in the future. Can this \emph{planar jittering pattern} have any observational consequences? We don't expect prominent signature in in Type II SNe where the shock will become more spherical as it propagate through the extended massive hydrogen envelope. In type Ic SNe the imprint might exist in the supernova remnant (SNR). We raise here the possibility that the torus morphology of a tilted thick disk with multiple jets in Cassiopeia A SNR (\citealt{DeLaneyetal2010,MilisavljevicFesen2013}) is a result of a planar jittering pattern. The last jet was more free to expand, as all core has been removed, and it is now observed as the high velocity jet like outflow along the northeast direction of the SNR. Finally, we point out that a planar jet launching patter might have taken place during galaxy formation, where a feedback between accretion of cold gas and jet activity might have taken place. {{{ {We thank the referee, Jason Nordhaus, for helpful comments.} }}} This research was supported by the Asher Fund for Space Research at the Technion, and a generous grant from the president of the Technion Prof. Peretz Lavie. OP is supported by the Gutwirth Fellowship. {{{ {The software used in this work was developed in part by the DOE NNSA ASC- and DOE Office of Science ASCR-supported Flash Center for Computational Science at the University of Chicago. This work was supported by the Cy-Tera Project (ΝΕΑ $ \rm Y \Pi O \Delta O M H / \Sigma T P A T H$/0308/31), which is co-funded by the European Regional Development Fund and the Republic of Cyprus through the Research Promotion Foundation.}}}} | 14 | 3 | 1403.5152 |
1403 | 1403.7157_arXiv.txt | \ \ \ \ \ With the latest results from the Planck experiment \cite{Ade:2013uln} and the recent BICEP2 measurement of $r$ \cite{Ade:2014xna}, cosmology and in particular inflation has finally started to become a testable field, since now we can rule out more and more models. Planck showed a preference towards inflation models with a plateau, as opposed to $\lambda_n\phi^n$ potentials, and in particular the Starobinsky model is close to the center of its allowed region in the $(r,n_s)$ plane. But the BICEP2 experiment favors instead large values of $r$, in the 0.15-0.25 range, seeming to exclude the simple Starobinsky model, with typically $r\sim 0.01$. The Starobinsky model is understood as inflation due to the presence of an $R^2$ term in the action, but it is equivalent to a usual Einstein plus scalar model with a potential that is an exponentially-corrected plateau. We will review this mechanism and generalize it to more complicated $f(R)$ actions, in the next section, in particular obtaining a generalization that allows for large $r$ compatible with BICEP2, tracing a line in the $(r,n_s)$ plane. But given that we want a plateau at large scalar field displacement values, where inflation happens, it is relevant to ask whether there is a symmetry at high energies that forces the model to have a plateau, in which case inflation would arise naturally, avoiding some non-naturalness arguments raised by \cite{Ijjas:2013vea} (see \cite{Guth:2013sya} for counter-arguments to those criticisms). If we are looking for such a symmetry, one possible answer that comes to mind is local Weyl symmetry (symmetry under rescaling the metric and the scalar fields by a local conformal factor), which in the Standard Model with masses given by VEVs of the Higgs seems to be broken at the fundamental level (not considering VEVs) only by the Planck mass $M_{\rm Pl}$. That can be remedied however by considering that $M_{\rm Pl}$ is also an effective coupling coming from the VEVs of scalars, leading to a theory with no {\em a priori} mass scales. To avoid issues with the experiments on the time variation of the Planck scale, one is led to models with 2 scalar fields and the local Weyl symmetry, as in \cite{Kallosh:2013hoa,Kallosh:2013maa} (following earlier work by \cite{Ferrara:2010in,Bars:2011mh,Bars:2011aa,Bars:2012mt,Kallosh:2013lkr,Kallosh:2013pby}; note that for \cite{Ferrara:2010in,Kallosh:2013lkr,Kallosh:2013pby} the initial motivation was superconformal symmetry, whereas for \cite{Bars:2011mh,Bars:2011aa,Bars:2012mt} the initial motivation was from 2-time physics models. The two are however related since for both we have an $SO(4,2)$ symmetry, conformal in 3+1 dimensions vs. Lorentz in 4+2). If we impose also an $SO(1,1)$ symmetry (can be thought of as a remnant of breaking of $SO(4,2)\rightarrow SO(3,1)\times SO(1,1)$) between the 2 scalars at large field value (equivalent of high energies in terms of the evolution of the Universe), and simply deform away from it, we are led to a theory with a plateau, giving naturally inflation. Moreover, one obtains the Starobinsky model predictions, basically because the scalars are related to the canonical scalar in an exponential way $\phi_i\sim e^{c\varphi}$ at large $\varphi$, and the $SO(1,1)$ symmetry cancels the leading contribution in the potential, as we will argue in detail in the current paper. The local Weyl invariance allows us to remove one of the scalars by fixing a gauge, and be left with a single physical scalar. Naturally one obtains the original Starobinsky model, so it is important to look for generalizations that allow us to go to the Starobinsky line described above. On the other hand, in \cite{Bars:2013yba,Bars:2013vba}, a similar model was considered, again with local Weyl invariance and two scalars, and the Planck mass coming from the same scalar VEVs, but without the $SO(1,1)$ symmetry between the scalars (except in the kinetic term, where it is essential!), where the motivation was different. We know of only one observed scalar so far, the recently discovered Higgs boson \cite{Aad:2012tfa,Chatrchyan:2012ufa}, so the natural question to be asked is whether we can identify the physical scalar with the Higgs. Indeed one can, and it was shown that if we want to keep only some terms natural for the Higgs in the potential, one obtains a cyclic cosmology. Basically, the reason is the same exponential relation to the canonical scalar $\phi_i\sim e^{c\varphi}$, but now because the lack of $SO(1,1)$ symmetry avoids the cancellation of the leading term and leaves us with a positive exponential potential $V\sim e^{a\varphi}$, which is one of the natural class of potentials for cyclic cosmology. Thus paradoxically, a similar set-up with or without $SO(1,1)$ at large energies leads to natural inflation or natural cyclic cosmology. In this paper we mix the two approaches by taking the best out of each one, noticing loopholes left by the two constructions, as well as generalizing the Starobinsky model by completing it to an infinite series summing to an $f(R)$, giving a line in the $(n_s,r)$ plane for CMB fluctuations. We note that allowing for $SO(1,1)$ at large field value and deforming away from it allows for a more general set-up than the one considered in \cite{Kallosh:2013hoa,Kallosh:2013maa}, in particular allowing us to naturally get the Higgs potential at low energies. Reversely, in \cite{Bars:2013yba,Bars:2013vba} one considered a quartic Higgs potential, but one can consider that the potential is modified at large field values to reach asymptotically the $SO(1,1)$ invariant form that can hold at high energies. We will see that in this case we obtain again a model giving predictions similar to the Starobinsky one in the $(r,n_s)$ plane, as in \cite{Kallosh:2013hoa,Kallosh:2013maa}, but now the model is more general, and more importantly we can also obtain the generalized Starobinsky line, allowing us to agree with the $n_s$ value of Planck but also the $r$ value of BICEP2. Note that the idea of Higgs inflation is far from new, being in fact quite popular recently. Moreover, there is in fact a model that can be thought of as a particular case of the general set-up with a Higgs coupled to the Einstein action, otherwise with only a Higgs potential, the Bezrukov-Shaposhnikov model \cite{Bezrukov:2007ep}. It was shown in \cite{Bars:2013yba} that it admits a Weyl-symmetric formulation. The paper is organized as follows. In section 2 we review the Starobinsky model, generalized to $f(R)$ actions, in particular showing that we can obtain an arbitrary value of $r$ through an $f(R)$ that goes over to $R^2$ in the inflating large $R$ region. Then we show that a class of exponentially-corrected potentials give rise to the same predictions (this observation was made in \cite{Ellis:2013nxa}, and in the context of supergravity models in \cite{Ferrara:2013rsa,Kallosh:2013yoa,Cecotti:2014ipa}; we learned about these references after the paper was posted on the arXiv, from the referee; also \cite{Ellis:2014cma} made this prediction), also generalized to a line in the $(r,n_s)$ plane. In section 3 we present the class of models we will be considering, with local Weyl invariance, $SO(1,1)$ invariance and reducing to the Higgs potential at low energies. In section 4 we discuss various particular cases and the inflationary models arising from them. In section 5 we show that with a modification of the scalar-Einstein coupling $\xi$ away from the conformal value of $-1/6$ leads nevertheless to a strong attraction to the same Starobinsky line. In section 6 we conclude. | In this paper we have studied the possibility to have the Higgs field as the inflaton in a model with local Weyl symmetry, where the Planck mass appears by fixing a gauge, and with $SO(1,1)$ invariance at large field values (in the inflationary region). We have shown that in these models, defined by two functions $f(x)$ and $g(x)$, generically we obtain the predictions of the original Starobinsky model, but we can find functions that give a generalized version of the Starobinsky model, for which we can obtain any value of the tensor to scalar ratio of CMB fluctuations $r$, including the value of $\sim 0.2 $ preferred by BICEP2. The potential is approximated in the inflationary region by a general exponentially-corrected plateau, which can be obtained from a generalized form of the Starobinsky model, with an infinite series of $R^p$ corrections that sum to a function $f(R)$. The functions $f(x)$ and $g(x)$ were analyzed from the point of view of needing to interpolate between the inflationary region and the Higgs potential, and in the inflationary region should arise in a consistent quantum gravity theory. Of course, specific functions would arise in the case of specific models, that would describe in particular how does the $SO(1,1)$ invariance get broken at low energies, but we did not attempt here to construct such models. That is left for future work. If we modify the non-minimal coupling $\xi$ of the scalar to gravity away from the conformal point, the Starobinsky line is a strong attractor, as in the original Starobinsky point, so it seems that the local Weyl invariance is not really essential to the inflationary predictions. {\bf Note}. While this paper was being written, the paper \cite{Hertzberg:2014aha} appeared, which also discusses the possibility of deforming the $SO(1,1)$ symmetry in a more general way. We would like to thank the referee for pointing out to us references \cite{Ellis:2013nxa}, where the general potential (\ref{generalpot}) was analyzed, and \cite{Ferrara:2013rsa,Kallosh:2013yoa,Cecotti:2014ipa}, where it was obtained in the context of supergravity models. We have been instead focusing on getting this potential from the conformal inflation from the Higgs. | 14 | 3 | 1403.7157 |
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1403 | 1403.5963_arXiv.txt | The phenomenology of MOND (flat rotation curves of galaxies, baryonic Tully-Fisher relation, \textit{etc.}) is a basic set of phenomena relevant to galaxy dynamics and dark matter distribution at galaxy scales. Still unexplained today, it enjoys a remarkable property, known as the \textit{dielectric analogy}, which could have far-reaching implications. In the present paper we discuss this analogy in the framework of simple non-relativistic models. We show how a specific form of dark matter, made of two different species of particles coupled to different Newtonian gravitational potentials, could permit to interpret in the most natural way the dielectric analogy of MOND by a mechanism of gravitational polarization. | \label{sec:MOND} The Modified Newtonian Dynamics (MOND) was introduced more than 30 years ago by Milgrom~\cite{Milg1,Milg2,Milg3} as an alternative to dark matter, designed to explain a variety of phenomena taking place at the scale of galaxies, which are now collectively referred to as the phenomenology of MOND (see Refs.~\cite{SandMcG02,FamMcG12} for reviews). The ability of MOND at reproducing this phenomenology is astonishing, and it is fair to say that this still represents a complete mystery today. The rotation curves of almost all spiral galaxies are reproduced in great details with a single-parameter fit --- the mass-to-luminosity ratio which is \textit{a posteriori} seen to be consistent with the expectations coming from stellar populations. The baryonic Tully-Fisher (BTF) relation~\cite{TF77,McG00}, an empirical relation between the asymptotic rotation velocity and the baryonic mass of galaxies, and valid for a large range of masses of galaxies~\cite{FamMcG12}, is naturally reproduced. In particular, for dwarf galaxies dominated by the gas there is little uncertainty on both the rotation velocity and the baryonic mass, so the evidence for the BTF relation is very strong~\cite{McG11}. The \textit{original sin} of MOND is Milgrom's law, namely that the discrepancy between the dynamical and luminous masses, \textit{i.e.} the presence of dark matter, is correlated with the involved scale of acceleration or magnitude of the gravitational field, see Fig.~\ref{fig:MD} which is taken from Ref.~\cite{FamMcG12}. \begin{figure}[htb] \vspace{-1cm} \centerline{\includegraphics[width=9.5cm]{MDacc.pdf}} \vspace{-3.5cm} \caption{\footnotesize{The mass discrepancy, defined by the ratio $(V/V_\text{b})^2$ where $V$ is the observed velocity and $V_\text{b}$ is the velocity attributable to visible baryonic matter, in spiral galaxies. No correlation is found with the distance scale (top panel); however a strong correlation is seen with the acceleration scale (middle panel) and with the gravitational field scale (bottom panel). The mass discrepancy (\textit{i.e.} the presence of dark matter) appears below the critical acceleration scale $a_0\sim 10^{-10}\,\text{m}/\text{s}^2$.}} \label{fig:MD} \end{figure} In this paper we shall adopt for MOND the modified Poisson equation\footnote{Boldface letters indicate ordinary Euclidean vectors; $G$ is Newton's gravitational constant.} of Bekenstein \& Milgrom~\cite{BekM84} \be{BMeq} \bm{\nabla}\cdot\biggl[\mu\left(\frac{g}{a_0}\right) \bm{g} \biggr] = - 4\pi G\,\rho_\text{b} \,, \ee where $\rho_\text{b}$ denotes the ordinary mass density of the baryons. The gravitational field is irrotational, $\bm{g}=\bm{\nabla}U$, with $U$ the gravitational potential, and we denote its norm by $g=\vert\bm{g}\vert$. The function $\mu$ of the ratio $g/a_0$ is the MOND interpolating function, which interpolates between the Newtonian regime $g\gg a_0$ for which $\mu\simeq 1$ (thus one recovers the usual Poisson equation of Newtonian gravity in this regime), and the MOND weak-field regime $g\ll a_0$ for which $\mu$ is linear in its argument, $\mu\simeq g/a_0$. The constant $a_0$ represents the MOND acceleration scale separating the two regimes as evidenced by Fig.~\ref{fig:MD}. Several relativistic MOND theories, extending general relativity with new fields and without the need of dark matter, have been proposed: \begin{itemize} \item The Tensor-Vector-Scalar (TeVeS) theory, which extends general relativity with a time-like vector field and one scalar field~\cite{Sand97,Bek04,Sand05}; \item Einstein-{\ae}ther theories, originally motivated by the phenomenology of Lorentz invariance violation~\cite{JM00,JM04}, involve a unit time-like vector field non-minimally coupled to the metric and with a non-canonical kinetic term~\cite{ZFS07,Halle08}; \item A bimetric theory of gravity in which the two metrics are coupled through the difference of their Christoffel symbols~\cite{bimond1,bimond2}; \item A variant of TeVeS using a Galileon field and a Vainshtein mechanism to prevent deviations from general relativity at small distances~\cite{BDgef11}; \item A theory based on a preferred time foliation labelled by the so-called Khronon scalar field~\cite{BM11,Sand11}. \end{itemize} The cosmology of these theories has been extensively investigated, notably in TeVeS and non-canonical Einstein-{\ae}ther theories~\cite{Sk06,LiMo08,Sk08,Zu10}. However, all these theories have difficulties in reproducing the CMB spectrum, even when adding a component of hot dark matter~\cite{Sk06}. | 14 | 3 | 1403.5963 |
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1403 | 1403.2258_arXiv.txt | The onset of convection in a rotating cylindrical annulus with parallel ends filled with a compressible fluid is studied in the anelastic approximation. Thermal Rossby waves propagating in the azimuthal direction are found as solutions. The analogy to the case of Boussinesq convection in the presence of conical end surfaces of the annular region is emphasised. As in the latter case the results can be applied as an approximation for the description of the onset of anelastic convection in rotating spherical fluid shells. Reasonable agreement with three-dimensional numerical results published \blue{by Jones {\em et~al.} ({\em J.~Fluid Mech.}, vol.~634, 2009, pp.~291-319)} for the latter problem is found. \red{As in those results the location of the onset of convection shifts outwards from the tangent cylinder with increasing number $N_\rho$ of density scale heights until it reaches the equatorial boundary. A new result is that at a much higher number $N_\rho$ the onset location returns to the interior of the fluid shell.} | \label{sec1} The tendency of fluid motions in rapidly rotating systems to develop nearly two-dimensional structures has often been exploited to simplify the theoretical analysis. The description of convection flows in systems where the gravity vector and the rotation axis are not parallel provides a typical example \citep{Bu70,Bu02}. In applications of convection problems to rotating planets and stars the tendency towards two-dimensionality is partly obscured by the strong variation of fluid density as a function of radius in the nearly spherical systems. It is thus of interest to investigate the extent to which the quasi-geostrophic two-dimensional description can still provide an approximation for three-dimensional convection in rapidly rotating systems with strong variations of density. In the case of the Boussinesq approximation, in which the density is regarded as constant except in connection with the gravity term, the results derived from the two-dimensional \red{quasi-geostrophic} analysis of the onset of convection in rotating spherical fluid shells \red{compare} well with the results of the three-dimensional numerical analysis \citep{Simitev03}. In this \red{paper} the two-dimensional model was based on the problem of convection in a rotating cylindrical annulus with conical end boundaries \citep{Bu70, Bu86}. \reviv{For a more detailed} {discussion of the role of the quasi-geostrophic model in relationship to more accurate three-dimensional solutions for convection in rotating spherical fluid shells we refer} \red{to the paper of \cite{Gi06}.} In recent years the anelastic approximation \citep{Go69} has been widely used to obtain more realistic descriptions of convection in the atmospheres of planets and stars with strong variations of density. In the paper by \citet[hereafter B86]{Bu86}, the analogy between the effect of changing height induced by the conical boundaries of the cylindrical annulus and the effect of a radial variation of density \reviv{has} already been pointed out. In the present paper we intend to demonstrate quantitatively that the two-dimensional analysis of the annulus model provides a reasonable approximation for the onset of convection in the presence of strong anelastic density variations in rotating spherical fluid shells. The analysis of \blue{the present} paper resembles to some extent the two-dimensional analysis of anelastic convection pursued by \citet{Evo04,Evo06}; see also \citet{Gla09}. Because of the high computational cost of three-dimensional simulations of convection in the presence of density variation over many scale heights these authors restricted their attention to two-dimensional numerical simulation of convection close to the equatorial plane. Evonuk \& Glatzmaier were interested in the nonlinear properties of two-dimensional convection including zonal flows in the presence of strong density variations. In contrast, our analytical model focuses on the linear problem of the onset of convection at high values of the rotation parameter. The main purpose of this paper is not the demonstration of a high accuracy of the two-dimensional approximation. Instead we wish to emphasise the insights into anelastic convection in rotating spheres gained from the analytical quasi-geostrophic model. In the \red{next section} we first introduce the narrow-gap cylindrical annulus with parallel ends as shown in figure \ref{f.010}(a) and derive the two-dimensional solution describing anelastic convection. In \textsection \ref{sec3} the model is modified for applications to the onset of anelastic convection in rotating spherical shells as indicated in figure \ref{f.010}(b). \blue{Detailed} comparisons with numerical solutions are evaluated in \textsection \ref{sec4}. Some \reviv{nonlinear aspects} are discussed in the final \textsection of the paper. | \label{sec5} As should be expected, the comparison of the asymptotic expressions with the numerical results does not show as good agreement in the anelastic case as in the Boussinesq case studied by \citet{Simitev03}. On the other hand, the conceptional value of the approximate asymptotic theory increases in proportion to the complexities introduced by anelastic density stratifications. In the present paper only the linear local problem of the onset of convection has been investigated. An extension could be considered in connection with the spiralling nature of convection which depends on the second derivative of the density variation in the $x-$direction. But since an analytical theory for this effect is not yet available even in the Boussinesq case of the cylindrical annulus with varying inclination of the conical end surfaces, such an analysis will be deferred to future research. The spiralling of the convection columns is an important feature since it is associated with Reynolds stresses that generate a differential rotation. Such a mechanism could eventually be described by an analytical theory based on an expansion in powers of the amplitude of motion as has been done \reviv{by \citet{Bu83}} in the case of the Boussinesq version of the problem. \blue{Finally the two-dimensional approximate analysis of the linear problem of anelastic convection could be improved through a three-dimensional {multiscale} analysis as has recently been done by \citet{Calk13} in the limit of the Boussinesq approximation.} \vspace{4mm} \noindent \textbf{Acknowledgements}\\ This research has been supported by NASA grant NNX-09AJ85G. We acknowledge the hospitality of Stanford University and UCLA. R.D.S.\ enjoyed a period of study leave granted by the University of Glasgow and the support of the Leverhulme Trust via Research Project Grant RPG-2012-600. | 14 | 3 | 1403.2258 |
1403 | 1403.2772_arXiv.txt | Investigations of the point spread functions (PSFs) and flux calibrations for stacking analysis have been performed with the far-infrared (wavelengths range of 60 to 140\,$\mu$m) all-sky maps taken by the Far-Infrared Surveyor (FIS) onboard the AKARI satellite. The PSFs are investigated by stacking the maps at the positions of standard stars with their fluxes of $0.02-10$ Jy. The derived full widths at the half maximum (FWHMs) of the PSFs are $\sim 60\arcsec$ at 65 and 90\,$\mu$m and $\sim 90\arcsec$ at 140\,$\mu$m, which are much smaller than that of the previous all-sky maps obtained with IRAS ($\sim 6\arcmin$). Any flux dependence in the PSFs is not seen on the investigated flux range. By performing the flux calibrations, we found that absolute photometry for faint sources can be carried out with constant calibration factors, which range from 0.6 to 0.8. After applying the calibration factors, the photometric accuracies for the stacked sources in the 65, 90, and 140\,$\mu$m bands are 9, 3, and 21\,\%, respectively, even below the detection limits of the survey. Any systematic dependence between the observed flux and model flux is not found. These results indicate that the FIS map is a useful dataset for the stacking analyses of faint sources at far-infrared wavelengths. | \label{intro} {\it AKARI} is the first Japanese satellite that carried out mid- to far-infrared (hereafter MIR and FIR, respectively) all-sky survey \citep{murakami07}. The {\it AKARI} all-sky survey in FIR wavelengths were performed from May 8 2006 to August 28 2007 with the Far-Infrared Surveyor (FIS, \cite{kawada07}) in four bands, covering 50-180\,$\mu$m wavelength range. The FIS scanned more than 98\,\% of the sky, and provides the FIR all-sky dataset with the highest spatial resolutions and sensitivities \citep{kawada07}. From the FIS all-sky data, the FIR bright source catalogue (BSC) is released \citep{yamamura10}. In addition, the diffuse maps will be publicly available from the data archive server. The diffuse map will enable us to perform studies of an arbitrary set of objects. However, individual objects are usually faint at FIR wavelengths, and severely contaminated with diffuse emission from interstellar dust and background unresolved sources. Stacking analysis is a useful technique to investigate FIR spectral trends and spatial structures of these faint sources. The confusion noise can be reduced by stacking the map at the positions of the sample objects. With the staked image, we can perform investigation of the FIR photometry and spatial structures of the sample with higher sensitivities than that performed with individual images. Therefore the stacking technique has been extensively adopted for the statistical FIR studies for faint sources, such as known normal galaxies and quasars \citep{kashiwagi12}, galaxy clusters \citep{giard08}, and mid-infrared extragalactic sources in deep field survey areas \citep{jauzac11}. These studies have succeeded in investigating FIR spectral and spatial features of the faint sources with fluxes even below the detection limits. This paper focuses on the investigations of point spread functions (PSFs) and flux calibrations for sources with fluxes comparable or below the detection limits, which are major objectives for the stacking analyses. In the stacking analyses, photometric accuracy and the PSFs of the images are assumed to be well calibrated, and have no systematic change throughout the considering flux range. However, the characteristics of point sources in the far-infrared data is generally very complex, mainly due to slow transient response of detectors \citep{shirahata09}, and diffraction and scattering inside the detector arrays and optical instruments \citep{arimatsu11}. These instrumental effects result in the differences between point source and diffuse calibration. The FIS all-sky map is calibrated in the earlier stage of the data processing by diffuse sources, such as zodiacal and interstellar dust emission \citep{matsuura11}. Additional investigation and calibration of the maps are required for point source studies. Furthermore, the photometric performance and the PSFs can systematically change in the FIS maps against the source fluxes, because the detector response is flux dependent \citep{shirahata09}. Therefore, we need to check the flux-dependent characteristics of the FIS maps. In the present study, we investigate the {\it AKARI} FIS diffuse maps at the N60 (65\,$\mu$m), Wide-S (90\,$\mu$m), and Wide-L (140\,$\mu$m) bands. The $5 \sigma$ detection limits for individual point sources in the maps are 2.4, 0.55, and 1.4\,Jy at 65, 90, and 140\,$\mu$m, respetively \citep{kawada07}. The N160 (160\,$\mu$m) band map is not investigated because the sensitivity is too low ($5 \sigma$ detection limit: 6.3 Jy, \cite{kawada07}) for the faint source studies. As calibration sources, we adopted infrared standard stars proposed by \citet{cohen99} with a flux range from 0.05 to 10 Jy. To investigate spatial and photometric properties of the faint stars, we stacked the FIS maps at their positions. By using the standard stars and the staking method, we present the characteristics of the point sources in the FIS maps with fluxes down to ten times lower than the detection limits for individual objects. These investigations provide not only useful information for users of the {\it AKARI}/FIS maps, but also demonstrate a investigation and calibration scheme for the general FIR stacking analyses. In this paper, section 2 describes a selection of standard stars and the stacking method used for the present study. In section 3, we report the results of the stacking analysis, trying to constrain the properties of point spread functions and calibration factors for point sources with different fluxes. Our results are discussed in section 4, and section 5 offers a short summary. | \subsection{Comparison with the Pointed Observations} In addition to the all-sky survey investigated so far, {\it AKARI}/FIS had performed slow-scan observations for the pointing mode, and their data characteristics were investigated in previous studies. In the previous studies for the slow-scan observations, the observed-to-expected flux ratios show a clear flux-dependence at the 65 and 90\,$\mu$m bands in the pointed observations \citep{shirahata09}. On the other hand, the flux dependence of the imaging performance is not confirmed for the FIS all-sky maps. The observed trend was thought to be due to the slow transient response of the FIS detector, which is also used for the all-sky survey. The previous results seem to be inconsistent with the present ones. In order to explain the discrepancy, we should firstly take the differences of the scan speeds for these observation modes into account. During the slow-scan observations, the detector sweeps the sky with a scan speed of 8 or $15\arcsec\, {\rm sec^{-1}}$, and the integration times of the point sources are 9.1 or 4.6 seconds, respectively, assuming a angular size of the point sources is comparable to the FWHM of the 90\,$\mu$m band PSF (73\arcsec, see table~\ref{taba3}). These timescales are comparable to duration time of the slow transient response (10-30 seconds), which are flux-dependent. The observed flux ratio thus shows the clear flux-dependence. On the other hand, the detector sweeps the sky with a much higher scan speed ($3\arcmin.6\ {\rm sec^{-1}}$) in the all-sky observations. A point source is scanned in only about 0.38 second, and the output signals are dominated by the initial pulse signals with duration time $\sim 0.2$ second \citep{shirahata04}. According to \citet{shirahata04}, the intensities of the pulse linearly depends on the stepwise input signal variations, and can purely be interpreted as the source fluxes. Therefore the slow transient effect can be less severe for the all-sky mode. Another possible reason for the inconsistency comes from the differences of the fluxes of the objects selected for the standards. According to the empirical model for the FIS detector established by \citet{kaneda02}, the slow transient response should depend on the total photo-current, which includes background light and signal from the target. Since the previous study used only the bright sources as calibration objects, such as asteroids and bright stars, the total fluxes are dominated by the source flux. Therefore the flux-dependence was clearly seen in the previous study. In contrast, the faint sources are used in our investigations, and the total fluxes are dominated by the background light. It is thus natural the flux ratio does not show a clear source-flux dependence in the present study. \subsection{Comparison with the Previous FIR All-sky Map} The FIS all-sky diffuse map provides an improved version of the FIR stacking study that has been performed with the previous finest FIR all-sky map obtained with IRAS. As described in section \ref{resluts1}, the average FWHMs are less than $90\arcsec$ at the three bands and significantly smaller than those of the IRAS map ($\sim 6\arcmin$, \cite{beichman88}). The higher angular resolution allows us to improve the sensitivity of faint source photometry by reducing confusion noise, which is one of the fundamental limiting factors for FIR stacking studies. Therefore the FIS maps will significantly improve the quality of the studies of faint sources performed with the IRAS data (e.g, \cite{kashiwagi12, bournaud12}). In addition to their sensitivities, the higher resolutions of the FIS maps enables us to study the spatial structures of the stacked sources, such as debris disks around main-sequence stars \citep{thompson10}, and intergalactic dust associated with galaxies \citep{menard10} and galaxy clusters \citep{giard08}. The angular scales of these objects are less than $10\arcmin$, which are difficult to resolve with the IRAS data. With the {\it AKARI}/FIS all-sky maps, we can improve, and newly develop spatial studies of these faint sources by stacking analysis. | 14 | 3 | 1403.2772 |
1403 | 1403.0913.txt | In the present work we derive an exact solution of an isotropic and homogeneous Universe governed by $f(T)$ gravity. We show how the torsion contribution to the FRW cosmology can provide a \textit{unique} origin for both early and late acceleration phases of the Universe. The three models ($k=0, \pm 1$) show a \textit{built-in} inflationary behavior at some early Universe time; they restore suitable conditions for the hot big bang nucleosynthesis to begin. Unlike the standard cosmology, we show that even if the Universe initially started with positive or negative sectional curvatures, the curvature density parameter enforces evolution to a flat Universe. The solution constrains the torsion scalar $T$ to be a constant function at all time $t$, for the three models. This eliminates the need for the dark energy (DE). Moreover, when the continuity equation is assumed for the torsion fluid, we show that the flat and closed Universe models \textit{violate} the conservation principle, while the open one does not. The evolution of the effective equation of state (EoS) of the torsion fluid implies a peculiar trace from a quintessence-like DE to a phantom-like one crossing a matter and radiation EoS in between; then it asymptotically approaches a de Sitter fate. | \label{S1} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Recent observational data suggest that our Universe is accelerating. Amongst the possible explanations for this phenomenon are modifications to gravitational theory \cite{NO7}. Cosmological constant, from an ideal fluid having different shapes of EoS with a negative pressure, a scalar field with quintessence-like or phantom-like behavior can also explain DE \cite{ENOSF}. It is not obvious what kind of DE is more suitable to explain the present epoch of the Universe. Observational data point to some type of DE having an EoS parameter which is close to $-1$, or even less than $-1$ (which is the phantom case). Modification of general relativity (GR) seems to be quite attractive possibility to resolve the above mentioned problem. Modifications of the Hilbert-Einstein action through the introduction of general functions of the Ricci scalar $R$ have been extensively explored \cite{NO07}-\cite{M11}. These $f(R)$ gravity theories can be reformulated in terms of scalar field quintessence. Moreover, it has been shown that when starting from $f(R)$ gravity, the phantom case in scalar tensor theory does not exist. However, when the conformal transformation becomes complex the phantom barrier is crossed, and therefore the resulting $f(R)$ function becomes complex. These cases are studied \cite{NOG} in more detail, in which, to avoid this handicap, a dark fluid was used to produce the phantom behavior such that $f(R)$ function reconstructed from the scalar tensor theory continues to be real. Initially the idea of teleparallelism theory has been proposed by Einstein in order to unify gravity and electromagnetism \cite{ Ea, OINC}. Later Einstein left the teleparallelism theory not because of its failure in the attempt of unification only, but also because of the vanishing of the curvature tensor of the Weitzenb\"{o}ck connection. But the non-vanishing torsion tensor aroused recently a great interest in astrophysical and cosmological applications in the so called $f(T)$ gravitational theory. The main motivations of such theory were: \begin{itemize} \item [(1)] GR can be viewed as a certain theory of teleparallelism; thus, it could be regarded at least as a different perspective that could lead to the same results \cite{TM}. \item[(2)] In such a context, one can define energy and momentum tensors of the gravitational field which are true tensors under all general coordinate transformations but {\it not under local Lorentz transformation}. \item [(3)] This theory is interesting because it can be seen as gauge theories of the translation group (not the full Poincar\'{e} group); consequently, one may provide an alternative interpretation of GR \cite{Hw}-\cite{YE13}. Most recently Teleparallel Equivalent of General Relativity (TEGR) has been generalized to $f(T)$ theory, a theory of modified gravity formed in the same spirit as generalizing GR to $f(R)$ gravity \cite{BF,Le}. A main merit of $f(T)$ gravity theory is that its gravitational field equation is of second order, the same as for GR, while it is of fourth order in metric $f(R)$ gravity. This merit makes the analysis of the cosmological expansion of the Universe in $f(T)$ gravity much easier than in $f(R)$ gravity. $f(T)$ gravity has gained significant attention in the literature with promising cosmological implications \cite{FF7}-\cite{RHSR}. \end{itemize} The main target of this work is to show how $f(T)$ gravity can be useful in explaining the flatness and acceleration at early and late phases of our Universe. For as is well known, current observations of the present Universe indicates that our Universe now is almost spatially flat. This leads one to exclude the closed and open Universe models. On the other hand the initial flat space assumption contradicts the presence of the strong gravitational field (i.e. the Riemann curvature) as it should! This contradiction might be explained as the flatness problem of the standard cosmology. Actually this problem has been overcome by the idea of an inflationary scenario during $\sim 10^{-36}-10^{-32}$ sec from the big bang. Lots of inflationary models have been proposed by using scalar fields. But to gain the benefits of both inflation and the standard cosmology the inflation should end at $\sim 10^{-32}$ sec from the big bang. This needs slow-roll conditions so that the inflationary Universe ends with a vacuum dominant epoch allowing the Universe to restore the big bang scenario. So the inflation can be considered as an add-on tool rather than a replacement of the big bang \cite{L03}. Until now there are no satisfactory reasons for the transition from inflation to big bang. Our trail here to treat these problems starts by diagnosing the core of the problem. We found that the curvature within the framework of the GR may lead to these conflicts, while introducing new qualities to the space-time, like torsion, might give a different insight of these problems. The work is arranged as follows: In Section \ref{S2}, we describe the fundamentals of the $f(T)$ gravity theory. We next show the contribution of the torsion scalar field to the density and the pressure of the FRW models and necessary modifications in Section \ref{S3}. Also, we obtain a model dependent scale factor $R(t)$ and $f(T)$ as a solution of the continuity equation. In Section \ref{S4}, we investigate the cosmological behavior of the flat, closed and open Universe models due to $f(T)$ modifications. Moreover, we give the physical descriptions for the obtained results. In the flat Universe the teleparallel torsion scalar field $T$ and the $f(T)$ appear as constant functions and the later might replace the cosmological constant, the Universe shows an inflationary behavior as the scale factor $R(t) \propto e^{Ht}$, where the Hubble parameter $H$ is a constant. The flat Universe shows no evolution with time. Moreover, we investigate the closed Universe model which shows an inflationary behavior as well. In spite of the torsion scalar field $T$ appears as a constant function similar to the flat case, but the $f(T)$ of the closed Universe appears as a function of time. This allows the cosmological parameters to evolve. In particular the evolution of the curvature density parameter $\Omega_{k}$ shows a clear tendency to vanish at late time, which explains how the Universe can start with initial curvature; then it goes naturally to flat behavior. Combining the curvature density parameter within the total density parameter $\Omega_{\textmd{Tot}}$ in addition to the matter $\Omega_{m}$ and the torsion $\Omega_{T}$ density parameters gives a very restrictive range for the total density parameters $|\Omega_{\textmd{Tot}}-1| \leq 10^{-16}$ at some early time. This is a suitable value to begin the primordial nucleosynthesis epoch. The late accelerating expansion of the Universe is also recognized as the Hubble parameter $H > 0$ and the deceleration parameter $q \rightarrow -1$. Furthermore, the investigation of the open Universe shows a behavior similar to the closed model. So both closed and open models suggest a unique source for early and late acceleration phases of the Universe. While the open model, uniquely, implies a time dependent effective EoS of the torsion fluid. Its evolution starting initially with a quintessence-like energy to asymptotical de Sitter crosses radiation-, dust- and phantom-like energies. Section \ref{S5} is devoted to summarizing and concluding the results. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | \label{S5} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \begin{itemize} \item In this work we have evaluated the matter density and pressure of the $f(T)$ field equations. We modified the FRW models due to the torsion contribution by replacing $\rho \rightarrow \rho+\rho_T$ and $p \rightarrow p+p_T$. Most of the cosmological models choose the scale factor $R(t)$ independent of the model. In this work we have got a model dependent $R(t)$ and $f(T)$ as order pairs, when applying the continuity equation to the Universe matter assuming that the torsion scalar and time are independent variables. The obtained solutions allow us to study the three world models, i.e. $k=0, \pm 1$. The calculations show that the torsion scalar (\ref{Tscalar}) can be written as a combination of the Hubble parameter $H$ and the curvature density parameter $\Omega_k$. These two parameters always combine keeping the torsion scalar a constant at all time $t$. \item The study of the \textit{flat} Universe model produces an inflationary cosmological model $R(t) \propto e^{Ht}$, $H=const$. But the Universe's compositions have no evolution where the matter density is constant during the expansion $\dot{\rho}=0$. Assuming the continuity equation for the torsion fluid leads to a constant torsion density, $\dot{\rho}_{T}=0$, during the expansion. This gives a steady state Universe. The total density of the Universe is equivalent to a constant Universe critical density. Then we conclude that the flat Universe model violates the conservation principle. \item The cosmological parameters for the \textit{closed} Universe model are found as functions of time. These parameters show a quick evolution at some early Universe, then they show a steady behavior at later time. Although the Universe in the closed model is chosen to be curved initially, the Universe's composition enforces the Universe to be flat at some late time as $\Omega_{\textmd{Tot}} \rightarrow 1$ and $\Omega_k \rightarrow 0$. Assuming the continuity equation for the torsion fluid implies a case similar to the flat model. \item In the case of the \textit{open} Universe model we have found a quick evolution of the cosmological parameters at some early time. The Universe in this model has been chosen to be initially curved, while the evolution of the cosmological parameters turns the Universe to be flat at some later time. The calculations show that the evolution of the open Universe prevents the violation of the conservation principle. This makes the open Universe model the most acceptable one. \item The inflationary Universe has been started as a speculative idea to solve some problems of the big bang cosmology. The inflation has been considered as an add-on extra tool to the standard big bang during some very early Universe. In this model, we get a built-in inflationary behavior at early time and then the model enables the big bang to be restored naturally. \item In the standard big bang cosmology it is known that the Universe becomes more and more curved very quickly, if it has been chosen to be initially curved, i.e. $\Omega_{\textmd{Tot}}$ diverges away from the unity. But the current cosmological observations show that our present Universe is almost flat. This requires a flat Universe initial condition. In our model, unlike the standard cosmology, we found that even if the Universe has started with an initial curvature, the evolution of $\Omega_{\textmd{Tot}}$ converges to unity. This tells that the Universe in the case of $k = \pm 1$ models is enforced to be flat. This solves many of the hot big bang cosmology problems. The closed Universe model shows that an extremely restrictive range for the total density parameter $|\Omega_{\textmd{Tot}}-1| \leq 10^{-16}$ at early Universe time, which is required for the nucleosynthesis epoch to begin and restore the big bang scenario. The open Universe shows almost the same restrictive range but much shorter interval of time. The result agrees with the BBN period ($\sim 1-200$ sec.) which again supports the open Universe model. See Figures 3(b) and 5(b). \item In the open model we have found that the teleparallel torsion fluid explains both early and late cosmic acceleration. This eliminates the need for the DE; in addition, it does not address the cosmological constant problem. Also, the use of the torsion scalar instead of the cosmological constant gives a conservative Universe. In addition, the torsion contribution gives a built-in inflationary behavior at a very early time; then the evolution of the total density parameter $\Omega_{\textmd{Tot}}$ shows good agreement with later stages. Moreover, the open Universe converges to a flat one, which agrees perfectly with the current observations. Furthermore, the evolution of the torsion fluid EoS, Fig. 4, shows a peculiar dynamical behavior during different phases of the cosmic expansion. There are many other details of these models that need further investigations. In particular, one would be interested in the torsion density and pressure in the open Universe model and their possible justifications as regards quantum cosmology. \end{itemize} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \subsection* | 14 | 3 | 1403.0913 |
1403 | 1403.7233_arXiv.txt | {We present the compilation of the first 221 supernovae classified during the Asiago Classification Program (ACP). The details of transients classification and the preliminarily reduced spectra, in fits format, are immediately posted on the Padova-Asiago SN group web site. The achieved performances for the first 2 years of the ACP are analysed, showing that half of all our classifications were made within 5 days from transient detection. The distribution of the supernova types of this sample resembles the distribution of the general list of all the supernovae listed in the Asiago SN catalog (ASNC$^8$, Barbon et al. 1999). Finally, we use our sub-sample of 78 core-collapse supernovae, for which we retrieve the host-galaxy morphology and $r$-band absolute magnitudes, to study the observed subtype distribution in dwarf compared to giant galaxies. This ongoing program will give its contribution to the classification of the large number of transients that will be soon delivered by the Gaia mission. } | The surveillance of the transient sky has greatly improved in the last ten years with the contribution of a growing number of surveys. Panoramic surveys of the nearby Universe like the Catalina Real-Time transient Survey (CRTS)\footnote{http://crts.caltech.edu}, the Palomar Transient Factory (PTF)\footnote{http://ptf.caltech.edu/iptf/}, the La Silla Quest\footnote{http://hep.yale.edu/lasillaquest}, the Mobile Astronomical System of theTElescope-Robots (MASTER)\footnote{http://observ.pereplet.ru}, the Optical Gravitational Lensing Experiment \linebreak (OGLE)\footnote{http://ogle.astrouw.edu.pl}, etc, have boosted the number of supernova (SN) discoveries from $\sim200$ in 2000 to $\sim1000$ in 2013. Many of these surveys post their discoveries in real time, allowing for a change of approach in SN science: for a specific SN science case the relevant events can be selected and studied. However, to really exploit the transient search efforts, we need to be able to identify, as early as possible, the event class with a prompt spectroscopic classification. Knowing the SN type and phase, one can activate a proper follow up campaign. The observing chain made of wide field searches, prompt classification and selective follow-ups has proved to be extremely productive as can be gathered from an \linebreak overview of the recent literature. As examples of the latest advances in SN researches, we recall the discovery of the class of super-luminous supernovae (SNe) whose explosion mechanism is still debated (Pastorello et al. 2010, Quimby et al. 2011, Gal-Yam 2012) and the real-time observations of the convulsions of massive stars on their path to explosion (Smith et al. 2013, Pastorello et al. 2013, Fraser et al. 2013, Margutti et al. 2014). The classification and early follow-up of bright nearby SNe can be efficiently done with small/medium size telescopes. In the past few years we conducted an European Southern Observatory (ESO) Large Program on { \em Supernova Variety and Nucleosynthesis Yields} (2009-2012) with the ESO-NTT (New Technology Telescope), complemented \linebreak with the INAF-TNG (Telescopio Nazionale Galileo of Istituto Nazionale di Astrofisica) for the Northern hemisphere, for the study of selected SNe (eg. Taubenberger et al. 2011, Patat et al. 2011, Fraser et al. 2011, Valenti et al. 2011a, Pastorello et al. 2012, Kankare et al. 2012, Pastorello et al. 2013, Tomasella et al. 2013a). Building on this successful experience, the ESO Large Program was merged into a new major international collaboration, the Public ESO Spectroscopic Survey of Transient Objects (PESSTO)\footnote{http://www.pessto.org/pessto/index.py}, that is using a major fraction of the time at the ESO NTT at La Silla (Chile). PESSTO started in 2012 and will be active for 4 (+1) years (Smartt et al. 2013). Despite the efforts of this large project and those of other groups worldwide, a large fraction of transients are not spectroscopically classified (around 50\%, based on the list of Latest Supernovae\footnote{http://www.rochesterastronomy.org/snimages/}). In this context, we decided to give a contribution to the classification of the brightest targets of the Northern hemisphere by exploiting our access at the observing facilities in Asiago, in particular the 1.82m Copernico telescope at Cima Ekar, operated by INAF Astronomical Observatory of Padova (OAPd). A parallel project is the photometric and spectroscopic follow-up of the most interesting transients (classified or not by us), which is not addressed in this paper. In fact, in the last few years the observations in Asiago have been deeply reorganised and a remarkable amount of telescope-time is allocated for two or three Large Programs. Our ongoing {\em Classification and follow-up of extragalactic transients discovered by \linebreak panoramic surveys} is one of those Large Programs. The Asiago SN Classification Program (ACP) started in 2011, with the aim to classify all transients that are accessible from our latitude and are bright enough for 1.82m Copernico telescope (apparent magnitude $\leq 19$ mag). The project is the latest evolution of Asiago SN programs, which begun in the early sixties. Among the historical achievements, we recall: the systematic SN search with Schmidt telescopes (Rosino 1964); the first identification of peculiar type~I SNe (Bertola 1964), named Ib or Ic twenty years later (Gaskell et al. 1986); the derivation of an average light curve for type~Ia (Barbon et al. 1973) and of the (different) photometric properties of type~II (Barbon et al. 1979) SNe; and the compilation of the Asiago SN Catalogue (Barbon et al. 1999), intended as a large database for statistical studies on the SN phenomenon (ASNC)\footnote{http://sngroup.oapd.inaf.it/asnc.html}. The ACP is allocated on average one week of observing nights per month at the Copernico 1.82m telescope. Typically half of the allocated time is used for classification of new targets and half to contribute to follow-up observations of selected objects. When the Copernico telescope is not available, the 1.22m Galileo telescope is used for transients brighter than magnitude 17. Both these telescopes are located in the Asiago Plateau, North-East of Italy, about one hundred km from Padova, at an altitude of 1.366 m for Copernico (Mt. Ekar, $11^\circ\,34'\,08.42\,{\rm E}$, $+45^\circ\,50'\,54.52''\,{\rm N}$) and 1.045 m for Galileo ($11^\circ\,31'\,3''\,{\rm E}$, $+45^\circ\,51'\,59''\,{\rm N}$). \linebreak This site is characterised by a continental climate, with dry winters and rainy spring and autumn. Summer time is on average favourable for observations. A statistics of the nights per month with open dome for the last years is plotted in Figure~\ref{nights}. The seeing is quite variable during the year, with an average $\sim2''$, but nights with seeing around $1''$ are frequently registered. Despite the proximity of large cities, the collaboration with Asiago Plateau Municipalities and the Regional Agency for the environmental preserve (ARPAV) has contributed to the decreasing of the local light pollution during the past ten years (S. Ortolani, private communication). \begin{figure} \includegraphics[width=\linewidth]{fig_worknights} \caption{Statistics of the open-dome nights at the Copernico 1.82m telescope (including the partially used nights) as a function of the month for the past two years (2012$-$2013) and comparison with the statistic over the past 24 years (1985-2008). June and July are mainly used for telescope and dome maintenance, tests and outreach activities. } \label{nights} \end{figure} The ACP program proceeds as follows. Suitable targets are identified among those posted by SNe searches. Soon after acquisition, the spectra are immediately reduced through a semi-automatic data reduction pipeline (Sect.~2). The spectra are then compared to SN templates aided by automatic SN classification codes (Sect.~3.) and the classifications are disseminated via the IAU Central Bureau Astronomical Telegrams circulars (CBET)\footnote{http://www.cbat.eps.harvard.edu/cbet/RecentCBETs.html} and/or the Astronomical Telegram (ATEL)\footnote{http://www.astronomerstelegram.org}. Some statistics of the program performances and the properties of the classified SN sample are reported in Sect.~4. Following what we believe is a very fruitful trend of many new projects, we made the results of our effort immediately public: within a few hours from observation, the details of transient classification and the (fast-)reduced spectra (fits format) are posted in the Padova-Asiago SN group web site \footnote{http://sngroup.oapd.inaf.it/}. | In this paper we have presented the Asiago Classification Program ACP and the results of the first two years of operation. We show that a small telescope located in a site with continental weather conditions can give a significant contribution to the rapidly growing quest for transients classification. The positive experience of these two years encouraged us to further improve the telescope instrumentation and operational mode. In particular, very recently, the re-furbishment of the telescope control system has been completed, which is now used in remote mode. This allows the implementation of a new control room in the main buildings in Asiago, providing a quicker access to the 1.82m telescope and improving security for astronomers. We are now proceeding with the procurement of a new CCD (Andor IKON L936) with higher response in the UV and strongly reduced fringing contamination red wards of 7500 \AA\/. The quest for transients classification is going to grow even higher in the next years, in particular with the large number of transients that will be delivered by the ESA mission Gaia. Different simulations agrees that Gaia alone will detect over a thousand SNe per year (Cappellaro 2012, Altavilla et al. 2012, Blagorodnova et al. 2013). The survey strategy guarantees for an unbiased sky coverage and will allow to explore with better statistics many open science topics, such as for instance the relation between SN and host galaxy types. In fact, as we have shown in Sect.~5, existing data are still insufficient. The limiting magnitude of Gaia transients ($\sim19.5-20$) appears well suited for our instrumentation and therefore we plan to devote a special effort to the classification of these targets with the coordination of the Gaia Science Alerts Working Group. | 14 | 3 | 1403.7233 |
1403 | 1403.0908_arXiv.txt | Observations with WFC3/IR on the {\it Hubble} Space Telescope and the use of gravitational lensing techniques have facilitated the discovery of galaxies as far back as $z\sim 10-12$, a truly remarkable achievement. However, this rapid emergence of high-$z$ galaxies, barely $\sim 200$ Myr after the transition from Population III star formation to Population II, appears to be in conflict with the standard view of how the early Universe evolved. This problem has much in common with the better known (and probably related) premature appearance of supermassive black holes at $z\ga 6$. It is difficult to understand how $\sim 10^9\;M_\odot$ black holes could have appeared so quickly after the big bang without invoking non-standard accretion physics and the formation of massive seeds, neither of which is seen in the local Universe. In earlier work, we showed that the appearance of high-$z$ quasars could instead be understood more reasonably in the context of the $R_{\rm h}=ct$ Universe, which does not suffer from the same time compression issues as $\Lambda$CDM does at early epochs. Here, we build on that work by demonstrating that the evolutionary growth of primordial galaxies was consistent with the current view of how the first stars formed, but only with the timeline afforded by the $R_{\rm h}=ct$ cosmology. We also show that the growth of high-$z$ quasars was mutually consistent with that of the earliest galaxies, though it is not yet clear whether the former grew from $5-20\;M_\odot$ seeds created in Population III or Population II supernova explosions. | The first stars and galaxies started forming towards the end of the cosmic ``Dark Ages," a period thought to have lasted $\sim 400$ Myr (in the context of $\Lambda$CDM) following hydrogen recombination at cosmic time $t\sim 0.4$ Myr. Over the past decade, many detailed simulations of this process have produced a comprehensive picture of how the Universe transformed from a simple initial (dark) state to the complicated hierarchical system we see today, including galaxy clusters and the many components contained within them (see, e.g., Barkana \& Loeb 2001; Miralda-Escud\'e 2003; Bromm \& Larson 2004; Ciardi \& Ferrara 2005; Glover 2005; Greif et al. 2007; Wise \& Abel 2008; Salvaterra et al. 2011; Greif et al. 2012; Jaacks et al. 2012; for more recent reviews, see also Bromm et al. 2009, and Yoshida et al. 2012). Interest in this cosmic dawn has been heightened recently with the dramatic discovery of faint galaxies at redshifts well beyond the end of the Epoc of Reionization (EoR), which observations show started at the end of the Dark Ages ($z\sim 15$, i.e., $t\sim 400$ Myr) and lasted until $z\sim 6$ ($t\sim 900$ Myr) (see, e.g., Zaroubi 2012). Stretching the imaging capabilities of WFC3/IR on the {\it Hubble} Space Telescope (HST), and using gravitational lensing techniques, investigators have apparently uncovered the earliest galaxies emerging as far back as $z\sim 10-12$, a truly remarkable achievement (Bouwens et al. 2011; Zheng et al. 2012; Ellis et al. 2012; Bouwens et al. 2012; Coe et al. 2013; Oesch et al. 2013; Brammer et al. 2013; Bouwens et al. 2013). These nascent galaxies may have contributed to the re-ionization of the intergalactic medium (IGM), perhaps even dominated this process. But this rapid emergence of high-$z$ galaxies so soon after the big bang may actually be in conflict with our current understanding of how they came to be. This problem is very reminiscent of the better known (and probably related) premature appearance of supermassive black holes at $z\ga 6$. It is difficult to understand how $\sim$$10^9\;M_\odot$ black holes appeared so quickly after the big bang without invoking non-standard accretion physics and the formation of massive seeds, both of which are not seen in the local Universe. Recent observations (e.g., Jiang et al. 2007; Kurk et al. 2009; Willott et al. 2010; Mortlock et al. 2011; De Rosa et al. 2011) have compounded this problem by demonstrating that most (if not all) of the high-$z$ quasars appear to be accreting at their Eddington limit. In our recent assessment of the high-$z$ quasar problem (Melia 2013a), we considered the possibility that this conflict may be due to our use of an incorrect redshift-time relationship, i.e., an incorrect cosmological expansion, rather than to the astrophysics of black-hole formation and growth, which are increasingly constrained by the ever improving observations. We showed that the high-$z$ quasar data may instead be interpreted more sensibly in the context of the $R_{\rm h}=ct$ Universe (Melia 2007, 2012; Melia \& Shevchuk 2012), for which standard $5-20\;M_\odot$ seeds forming after re-ionization had begun at $z\la 15$ could have easily grown into $\ga 10^9 \;M_\odot$ supermassive black holes by redshift $z\ga 6$, merely by accreting at the observed Eddington rate. The principal difference between $\Lambda$CDM and $R_{\rm h}=ct$ that eases the tension between theory and observations is simply $z(t)$. Specifically, in the $R_{\rm h}=ct$ cosmology, the EoR began at $\sim 900$ Myr ($z\sim 15$ in this cosmology) and ended at $\sim 1.9$ Gyr ($z\sim 6$), placing the birth of supermassive black holes at $\sim 1$ Gyr (roughly $z=13$), right where one would have expected them to form via the supernova deaths of Population III and II stars that presumably started the EoR. Now that the premature emergence of high redshift galaxies is creating comparable tension with established theory, it is necessary to consider whether their timing issues are similarly resolved by the $R_{\rm h}=ct$ Universe, and whether one can find consistency between their evolutionary history and the growth of supermassive black holes at $z\ga 6$. In \S~2 of this paper, we will briefly review the current thinking behind the formation of structure at $z\ga 6$, and then in \S~3 compare this with what the most recent observations are telling us. In \S~4, we will discuss a re-interpretation of high-$z$ galaxy growth in the context of $R_{\rm h}=ct$, and then present our conclusions in \S~5. | The problem with the premature formation of supermassive black holes at $z\ga 6$ has been with us for several years. In order to get around the limited time available in $\Lambda$CDM to form such enormous objects in only $\sim 500$ Myr, various `fixes' have been proposed and studied, though the latest observations suggest that even these modifications may not be consistent with the data. For example, the possibility that black holes might have grown at greatly super-Eddington rates seems to have been ruled out by the latest measurements (e.g., Mortlock et al. 2011; De Rosa et al. 2011; Willott et al. 2012), which indicate that the most distant quasars are accreting at no more than the standard Eddington value. The possibility that their seeds may have been much more massive than is typically seen ($5-20\;M_\odot$) in supernova explosions appears to be remote, at best, given that we simply don't see these forming in the local Universe. At the very least, new physics would have to be devised in order to account for such exotic events. All in all, our ever improving capability to constrain the high-$z$ quasar properties appears to be arguing against the viability of $\Lambda$CDM to account for the expansion of the early Universe. The dramatic recent discovery of high-$z$ galaxies has created renewed interest in this timing problem, because these structures appear to have formed too quickly following the end of the Dark Ages. In some instances, $\sim 10^9\;M_\odot$ of stellar material would have necessarily assembled in only $\sim 200$ Myr. This is inconsistent with the broad range of detailed simulations carried out to date. The premature emergence of high-$z$ galaxies confirms and reinforces the difficulty faced by $\Lambda$CDM in properly accounting for the formation of structure during the cosmic dawn. \begin{figure}[hp] {\centerline{\epsscale{1.0} \plotone{fig3.eps} }} \figcaption{Same as figures~1 and 2, except here highlighting the dependence of our results on the uncertainty in the inferred galaxy mass $M$. The shaded regions correspond to a mass range $10^8-10^9\;M_\odot$, the largest uncertainty quoted in Table 1 (and references cited therein), for two representative objects from this sample. The transition from Population III to Population II star formation would have occurred at $\sim 300$ Myr in both cosmologies, but the time interval corresponding to the redshift range of the EoR is different, so we do not show it here.} \end{figure} In this paper, we have explored the possibility that the evolutionary growth of high-$z$ quasars and galaxies might instead be better explained in the context of the $R_{\rm h}=ct$ Universe. We have found that not only would it have been possible for these structures to grow since the big bang consistent with standard physical principles, but that their evolution would have been mutually consistent with each other, and with what we currently believe was the history of early star formation. Having said this, there are at least two important caveats to this conclusion. The first clearly has to do with how accurately we know the mass of these high-redshift galaxies. This question was partially addressed in figure~3, which showed that even a factor 10 uncertainty in $M$ is insufficient to alleviate the timing problem in $\Lambda$CDM. We estimate that in order to remove the problem completely, the real mass of these objects would have to be at least 4 orders of magnitude smaller than currently measured. Whether this is feasible remains to be seen. The second caveat is that although none of the simulations carried out to date within the context of $\Lambda$CDM can adequately account for these early galaxies, it may be possible that a critical physical ingredient has been overlooked. With the high-$z$ quasars, proposals to fix the problem have invoked new physics to generate massive seeds, or unusual circumstances to permit super-Eddington accretion. It would be more difficult to generate such ideas for the early formation of galaxies, which are aggregates of many stars, not single objects and, at least as far as we know today, could not have started forming until Population III stars gave way to Population II. Perhaps the initial cooling that led to the formation of Population III stars was not due to the condensation of molecular hydrogen; maybe some other process permitted the gas to cool much faster than the currently believed several hundred Myr time frame. No doubt, future simulations will probe such ideas and new physics, and perhaps a solution may be found to work with the timeline afforded by $\Lambda$CDM. But with what we know today, our results demonstrate that the formation of high-redshift galaxies could be consistent with $R{\rm h}=ct$, but probably not with $\Lambda$CDM. Our understanding of this important early period in the Universe's life will improve quite rapidly in the coming years as primordial galaxies are observed with increasing numbers and greater sensitivity, and as simulations incorporate more detailed physics and new ideas. \vskip-0.2in | 14 | 3 | 1403.0908 |
1403 | 1403.0763_arXiv.txt | We study the onset of neutron drip in high density matter in the presence of magnetic field. It has been found that for systems having only protons and electrons, in the presence of magnetic field $\gtrsim 10^{15}$G, the neutronization occurs at a density which is atleast an order of magnitude higher compared to that in a nonmagnetic system. In a system with heavier ions, the effect of magnetic field, however, starts arising at a much higher field, $\gtrsim 10^{17}$G. These results may have important implications in high magnetic neutron stars and white dwarfs and, in general, nuclear astrophysics when the system is embedded with high magnetic field. \\ \\ {\it PACS Nos.:} 26.60.-c, 95.30.-k, 87.50.C-, 71.70.Di, 26.60.Kp | The neutron stars are believed to have surface magnetic field as large as $\sim 10^{15}$G (magnetar model) \cite{magnetar} and hence their interior field ($B_{int}$) could even be a few orders of magnitude higher, say $\sim 10^{18}$G \cite{lai,armen}, where the density is also higher. On the other hand, for white dwarfs with typical radius ($R\sim 5000$km), the maximum $B_{int}$ could be restricted to $\sim 10^{12}$G from the scalar virial theorem. However, for the recently proposed high density white dwarfs having highly tangled/fluctuating magnetic field with radius, e.g., $R\sim 70$km, $B_{int}\gtrsim 10^{17}$G \cite{prl}. Such smaller white dwarfs' central density ($\rho_c$) is $\sim 10^{13}$gm/cc and hence quite above the non-magnetic threshold density of neutronization. All the above facts motivate us to formulate the present work and to study the effect of (high) magnetic field on neutronization and neutron drip for the degenerate fermions at high densities. This will be helpful to interpret the detailed spectra of cooling neutron stars, radio emission and $\gamma-$ray burst from neutron stars. At such high density regions of neutron stars and white dwarfs, the mean Fermi energy and the cyclotron energy of an electron exceed its rest-mass energy; for the latter condition to hold, the magnetic field needs to be sufficiently high, making the electrons relativistic. We know that inverse $\beta-$decay can occur when the energy of electrons becomes higher than the difference between the rest mass energies of neutron and proton. Such a condition is expected to be modified in the presence of magnetic field which we plan to explore here at zero temperature. Effects of magnetic field to the neutron star matter were investigated earlier in different contexts. For example, the enhancement of electron and neutron densities in the inner and outer crusts of neutron stars was discussed in the presence of high magnetic field \cite{band1,band2}. The authors adopted the Thomas Fermi model for their calculations and one of the motivations was to understand the transport phenomena within magnetars. Based on fully self-consistent covariant density functional theory, the influence of strong magnetic fields on nuclear structure has been studied in the context of magnetars \cite{dft}. These authors showed that a field strength $\gtrsim 10^{17}$G, which presumably corresponds to the central/crust field of magnetars, appreciably modifies the nuclear ground state. Hence, the composition of the magnetar crust might be radically different from that of normal neutron stars which might have significant implications to various related properties, such as pulsar glitches, cooling etc. Other authors investigated numerically the variations of composition and pressure at the onset of neutron drip with the magnetic field \cite{chem}, which we will recall in \S3 again for comparison with the present work. The plan of the paper is the following. In the next section, we discuss the neutronization density in a neutron-proton-electron ($n-p-e$) gas. Subsequently, in \S 3 we derive the neutron drip density in a system with heavier ions. Finally, we end with a summary in \S 4. | We have studied the effects of magnetic field on the onsets of neutronization for $n-p-e$ and neutron drip for ion-electron gas systems. We have found that beyond a certain value of magnetic field, the neutronization/drip density of either of the systems increases, compared to the nonmagnetic cases, linearly with the magnetic field. This value of field for the former is $B^e_D \sim 2.7$ and for the latter $\sim 1115$. The significant change in the drip density arises only in the linear regime. For instance, the drip density for the ionic system increases by an order, compared to that of the nonmagnetic gas, for a field $\sim 6.4\times 10^{17}$G. Apart from the various applications to neutron stars, the present findings may have interesting consequences to the recently proposed high density ($\gtrsim 10^{11}$gm/cc), high magnetic field ($\gtrsim 10^{15}$G) white dwarfs \cite{prl,prd,apj,mpla,jcap}. The inner region of white dwarfs, in particular for $B_{int}\gtrsim 10^{17}$G, would have been neutronized if the drip density had not been changed with the field. Note that some authors have questioned the existence of such white dwarfs \cite{chem2}. For example, the authors have argued that for the white dwarfs to be stable against neutronization, the magnetic field at the center should be less than few times $10^{16}$G for typical matter compositions. However, as discussed in detail in \cite{mpla} (also see \cite{prd}), for the central field $\sim 10^{16}$G, the mass of the said highly magnetized white dwarfs already becomes significantly super-Chandrasekhar with the value $2.44$ solar mass. Therefore, such super-Chandrasekhar white dwarfs remain stable according to the neutronization limit given by the same authors only. Even if the white dwarfs with very large field consist of central region with density larger than the modified drip density, the present finding will help in putting a constraint on such models of white dwarfs with pure electron degenerate matter. | 14 | 3 | 1403.0763 |
1403 | 1403.7696_arXiv.txt | Recently BICEP2 found that the vanishing of the tensor-to-scalar ratio $r$ is excluded at $7\sigma$ level, and its most likely value is $r=0.2^{+0.07}_{-0.05}$ at $1 \sigma$ level. This immediately causes a tension with the Planck constraint $r < 0.11$. In addition, it also implies that the inflaton (in single field slow-roll inflation) experienced a Planck excursion during inflation $\Delta\phi /M_{Pl} \geq {\cal{O}}(1)$, whereby the effective theory of inflation becomes questionable. In this brief report, we show that the inflationary paradigm is still robust, even after the quantum effects are taken into account. Moreover, these effects also help to relax the tension on the different values of $r$ given by BICEP2 and Planck. | The inflationary scenario provides a framework for solving the problems of the standard big bang cosmology and, most importantly, provides a causal mechanism for generating structure in the universe and the spectrum of cosmic microwave background (CMB) anisotropies. In this picture, primordial density and gravitational wave fluctuations are created from quantum fluctuations during the inflation process. The former, which generates observational effects in temperature anisotropies in CMB, has been conformed by observations with spectacular precision \cite{WMAP,PLANCK}. For the primordial gravitational wave fluctuations, which generates the unique primordial B-mode polarization of CMB, has now been observed by the ground based BICEP2 experiment \cite{BICEP2}. The results of the experiment shows for the first time a non-zero value of the tensor-to-scalar ratio $r$ at $7\sigma$ C.L., and estimated \bq \lb{rvalue} r=0.2^{+0.07}_{-0.05}, \eq at $1 \sigma$ C.L. If it is confirmed, it will represent one of the most important discoveries in the field, and provides us a new diagnostic tools to probe the models of early universe, especially the models in the framework of quantum gravity. However, the BICEP2 results are in tension with the Planck constraint $r<0.11$ \cite{PLANCK}. As pointed out by BICEP2 group \cite{BICEP2}, this tension may be alleviated if a large negative running of scale spectral index is allowed. However, such a large negative running can not be produced in the framework of single field slow-roll inflation models. As the single-field slow-roll inflation has already been shown to be favored by the Planck data, it is very desirable to find some mechanisms to relieve this tension. More important, a large $r$, such that given by Eq.(\ref{rvalue}), also implies that the effective field theory of inflation is problematic. This is based on the analysis of the Lyth bound \cite{lyth}, which states that in the slow-roll approximations, the change of the inflationary field $\phi$ during inflation is given by \bqn \lb{Lyth} \frac{\Delta\phi}{M_{\text{pl}}} \simeq \sqrt{\frac{r}{8}} \Delta{N}, \eqn where $M_{\text{pl}}$ denotes the Planck mass, and $\Delta{N}$ the number of e-folds. If $r$ does not change as a function of $N$, this directly leads to \cite{lyth} \bqn \frac{\Delta \phi}{M_{\text{pl}}} \simeq \mathcal{O}(1) \sqrt{\frac{r}{0.01}}. \eqn Obviously, $\Delta \phi$ exceeds the Planck scale if $r$ is of the order of Eq.(\ref{rvalue}). As a result, the effective field theory of inflation with a potential $V(\phi)$, which consists of a derivative expansion of operators suppressed by Planck scale, becomes questionable. In this brief report, we take the point of view that $r$ is indeed large and of the order of Eq.(\ref{rvalue}), so that the gravitational quantum effects are important and necessary to be taken into account during the epoch of inflation. Once these effects are taken into account, we shall show that the inflation paradigm is still very robust, and almost scale-invariant can be easily produced. In addition, these effects also help to relax the tension between the values of $r$ given, respectively, by Planck and BICEP2. | In this brief report, we have taken the point of view that the tensor-to-scalar ratio $r$ is big, as found by BICEP2, and that the trans-Planckian effects become indeed important and need to be taken into account during the epoch of inflation, as indicated by the Lyth bound (\ref{Lyth}). Then, we have shown that, even after these effects are taken into account, almost scale-invariant perturbations can still be easily obtained \cite{WM,Wang2010PRD,Huang2013,Zhu2013}, and the inflation paradigm is robust. Moreover, these effects also helps to relax the tension between the values of $r$ given, respectively, by Planck BICEP2, whereby the two problems mentioned in Introduction are resolved. | 14 | 3 | 1403.7696 |
1403 | 1403.7143_arXiv.txt | \baselineskip = 11pt \leftskip = 0.65in \rightskip = 0.65in \parindent=1pc {\small The Solar System formed about 4.6 billion years ago from a condensation of matter inside a molecular cloud. Trying to reconstruct what happened is the goal of this chapter. For that, we put together our understanding of Galactic objects that will eventually form new suns and planetary systems, with our knowledge on comets, meteorites and small bodies of the Solar System today. Our specific tool is the molecular deuteration, namely the amount of deuterium with respect to hydrogen in molecules. This is the Ariadne's thread that helps us to find the way out from a labyrinth of possible histories of our Solar System. The chapter reviews the observations and theories of the deuterium fractionation in pre-stellar cores, protostars, protoplanetary disks, comets, interplanetary dust particles and meteorites and links them together trying to build up a coherent picture of the history of the Solar System formation. We emphasise the interdisciplinary nature of the chapter, which gathers together researchers from different communities with the common goal of understanding the Solar System history. \\~\\~\\~}% | }\label{sec:introduction} The ancient Greek legend reads that Theseus volunteered to enter in the Minotaur's labyrinth to kill the monster and liberate Athens from periodically providing young women and men in sacrifice. The task was almost impossible to achieve because killing the Minotaur was not even half of the problem: getting out of the labyrinth was even more difficult. But Ariadne, the guardian of the labyrinth and daughter of the king of Crete, provided Theseus with a ball of thread, so that he could unroll it going inside and follow it back to get out of the Minotaur's labyrinth. Which he did. Our labyrinth here is the history of the formation of the Solar System. We are deep inside the labyrinth, with the Earth and the planets formed, but we don't know how exactly this happened. There are several paths that go into different directions, but what is the one that will bring us out of this labyrinth, the path Nature followed to form the Earth and the other Solar System planets and bodies? Our story reads that once upon a time, it existed an interstellar cloud of gas and dust. Then, about 4.6 billions years ago, one cloud fragment became the Solar System. What happened to that primordial condensation? When, why and how did it happen? Answering these questions involves putting together all of the information we have on the present day Solar System bodies and micro particles. But this is not enough, and comparing that information with our understanding of the formation process of Solar-type stars in our Galaxy turns out to be indispensable too. Our Ariadne's thread for this chapter is the deuterium fractionation, namely the process that enriches the amount of deuterium with respect to hydrogen in molecules. Although deuterium atoms are only $\sim1.6\times 10^{-5}$ (Tab. \ref{tab:definitions}) times as abundant as the hydrogen atoms in the Universe, its relative abundance in molecules, larger than the elemental D/H abundance in very specific situations, provides a remarkable and almost unique diagnostic tool. Analysing the deuterium fractionation in different galactic objects which will eventually form new suns, and in comets, meteorites and small bodies of the Solar System is like having in our hands a box of old photos with the imprint of memories, from the very first steps of the Solar System formation. The goal of this chapter is trying to understand the message that these photos bring, using our knowledge of the different objects and, in particular, the Ariadne's thread of the deuterium fractionation to link them together in a sequence that is the one that followed the Solar System formation. The chapter is organised as follows. In \S \ref{sec:set-stage}, we review the mechanisms of the deuterium fractionation in the different environments and set the bases for understanding the language of the different communities involved in the study of the Solar System formation. We then briefly review the major steps of the formation process in \S \ref{sec:brief-hist}. The following sections, from \S \ref{sec:the-pre-stell} to \S \ref{sec:solar-nebula}, will review in detail observations and theories of deuterium fractionation in the different objects: pre-stellar cores, protostars, protoplanetary disks, comets and meteorites. In \S \ref{sec:summaryD}, we will try to follow back the thread, unrolled in the precedent sections, to understand what happened to the Solar System, including the formation of the terrestrial oceans. We will conclude with \S \ref{sec:conclusions}. | \end{center} \end{figure*} \subsection{\textbf{Overview of the deuterium fractionation through the different phases}}\label{sec:overview-meas-deut} Figure \ref{fig:DH-summary} graphically summarises the measured values of the D/H ratio from pre-stellar cores up to Solar System bodies, whereas Table \ref{tab:DH-summary} lists the measured D/H values and the main characteristics of the observations of Solar System bodies. \begin{table*}[bth] \begin{tabular}[h]{|lll|} \hline \hline {\bf Object} & {\bf D/H } & {\bf Note} \\ & [$\times 10^{-4}$] & \\ \hline Oort-Cloud Comets$^a$ & $\sim 3$ & Measured in seven comets.\\ Jupiter-Family Comets$^a$ & $\leq 2 $ & Measured in two comets: 1.61$\pm$0.24 and $\leq$2. \\ \hline IDPs$^a$ & 1.5--3.5 & The D/H distribution is peaked at 1.3, with a shoulder from 0.8 to 3.5\\ & & (it mimics the CC clays D/H distribution).\\ CC clays$^a$ & 1.2--2.3 & The D/H distribution is peaked at 1.4, with a shoulder from 1.7 to 2.3.\\ CC IOM$^b$ & 1.5--3.5 & The D/H distribution does not have a peak.\\ CC IOM functional groups$^b$ & 1.5--5.5 & The larger the functional group bonding energy the larger the D/H.\\ CC IOM deuterium hot spots$^b$ & $\sim$100 & Micrometer size spots in the IOM networks where D/H is highly enhanced.\\ CC SOC$^b$ & 2.1--7.1 & Branched $aa$ show the largest D/H.\\ \hline \end{tabular} \caption{\small Main characteristics of deuteration fractionation in comets and carbonaceous chondrites (CC). For explanations and references, see Secs. \ref{sec:the-comets} -- \ref{sec:meteorites}. Notes: $^a$HDO/H$_2$O; $^b$deuterium fractionation in organics; } \label{tab:DH-summary} \end{table*} Overall, the figure and the table, together with Fig. \ref{fig:iom-fig1}, raise a number of questions, often found in the literature too, which we address here. \\ \noindent {\it 1. Why is the D/H ratio systematically lower in Solar System bodies than in pre- and proto- stellar objects? When and how did this change occur?}\\ May be here the answer is, after all, simple. All the Solar System bodies examined in this chapter (meteorites and comets) were originally at distances less than 20 AU from the Sun. The D/H measurements of all pre- and protostellar objects examined in this chapter (pre-stellar cores, protostars and protoplanetary disks) refer to distances larger than that, where the temperatures are lower and, consequently, the deuteration is expected to be much larger (\S \ref{sec:the-chem-proc}). Therefore, this systematic difference may tell us that the PSN did not undergo a global scale re-mixing of the material from the outer ($\ga 20$ AU) to the inner regions. There are exceptions, though, represented by the ``hot spots'' in meteorites, which, on the contrary, may be the only representatives of this large scale re-mixing of the material in the PSN. \\ \noindent {\it 2. Why does organic material have a systematically higher D/H value than water, regardless the object and evolutionary phase?}\\ The study of the deuterium fractionation in the pre-stellar cores (\S \ref{sec:the-pre-stell}) and protostars (\S \ref{sec:prot-phase}) has taught us that the formation of water and organics in different epochs is likely the reason why they have a different D/H ratio (Fig. \ref{fig:protostar-1}). Water ices form first, when the density is relatively low ($\sim 10^4$ cm$^{-3}$) and CO not depleted yet, so that the H$_2$D$^+$/H$_3^+$ ratio is moderate (\S \ref{sec:the-chem-proc}). Organics form at a later stage, when the cloud is denser ($\ga 10^5$ cm$^{-3}$) and CO (and other neutrals) condenses on the grain surfaces, making possible a large enhancement of the H$_2$D$^+$/H$_3^+$. Can this also explain, {\it mutatis mutandis}, the different deuterium fractionation of water and organics measured in CCs and comets? After all, if the PSN cooled down from a hot phase, the condensation of volatiles would follow a similar sequence: oxygen/water first, when the temperature is higher, and carbon/organics later, when the temperature is lower. This would lead to deuterium fractions lower in water than in organics, regardless whether the synthesis of water and organic was on the gas phase or on the grain surfaces. Obviously, this is at the moment speculative but a road to explore, in our opinion. We emphasise that this ``different epochs formation'' hypothesis is fully compatible with the theory described in \S \ref{sec:fract-organ} of the origin of the organics deuteration from H$_2$D$^+$ \citep{2011GeCoA..75.7522R}. \\ \noindent {\it 3. Why do most comets exhibit higher D/H water values, about a factor 2, than CCs and IDPs?}\\ A possible answer is that comets are formed at distances, $\ga5$ AU, larger than those where the CCs originate, $\la 4$ AU (Tab. \ref{tab:definitions}; see also \S 9 ). A little more puzzling, though, is why the large majority of IDPs have D/H values lower than the cometary water, although at least 50\% IDPs are predicted to be cometary fragments (Nesvorny et al. 2010). The new Herschel observations of JFC indicate indeed lower D/H values than those found in OCC (Tab. \ref{tab:Dcomets}; \S \ref{sec:the-comets}), so that the observed IDPs D/H values may be consistent with the cometary fragments theory. In conclusion, the D/H distribution of CCs and IDPs is a powerful diagnostic to probe the distribution of their origin in the PSN. \\ \noindent {\it 4. Why does the D/H distribution of the IDPs, which are thought to be fragments of comets, mimic the CCs bulk D/H distribution? And why is it asymmetric, namely with a shoulder towards the large D/H values?}\\ As written above, the D/H ratio distribution depends on the distribution of the original distance from where CCs and IDPs come from or, in other words, their parental bodies. CCs are likely fragments of asteroids from the main belt and a good fraction is of cometary origin. IDPs are fragments from JFC and main belt asteroids. Their similar D/H distribution strongly suggests that the mixture of the two classes of objects is roughly similar, which argue for a difference between the two, more in terms of sizes than in origin. Besides, the asymmetric distribution testifies that a majority of CCs and IDPs originate from closer to the Sun parent bodies \citep[see also][]{2013Icar..223..722J}. \\ \noindent {\it 5. Why is the D/H ratio distribution of CC organics and water so different?}\\ There are in principle two possibilities: water and organics formed in different locations at the same time and then were mixed together or, on the contrary, they were formed at the same heliocentric distance but at (slightly) different epochs. The discussion in point 2 would favor the latter hypothesis, although this is speculative at the moment. If this is true, the D/H distribution potentially provides us with their history, namely when each of the two components formed. Again, being CC organics more D-enriched than water, they were formed at later stages. \subsection{\textbf{Comparison between PPD and PSN models: need for a paradigma shift?}}\label{sec:comp-models} There are several differences between the Proto Planetary disks (PPDs) and Proto Solar Nebula (PSN) models. A first and major one is that PSN models assume a dense and hot phase for the solar PPD. For example, temperatures higher than 1000 K persist for $\sim 10^5$ yr at 3 AU and $\sim 10^6$ yr at 1 AU in the Yang et al. (2013) model. Generally, models of PPDs around Solar-type stars, on the contrary, never predict such high temperatures at those distances. Second, PPD models consider complex chemical networks with some horizontal and vertical turbulence which modifies, though not substantially, the chemical composition across the disk. On the contrary, in PSN models, chemistry networks are generally very simplified, whereas the turbulence plays a major role in the final D-enrichment across the disk. However, since the density adopted in the PSN models is very high, more complex chemical networks have a limited impact \citep{2013Icar..290W}. Third, the above mentioned PSN models do not explicitly compute the dust particle coagulation and migration (if not as diffusion), and gas-grains decoupling whereas (some) disk models do (see Turner et al. and Testi et al., this volume). In our opinion, the most significant difference is the first one and the urgent question to answer is: what is the good description of the PSN? A highly turbulent, hot and diffusive disk or a cooler and likely calmer disk, as the ones we see around T Tauri stars? Or something else? The community of PSN models have reasons to think that the PSN disk was once hot. A major one is the measurement of depletion of moderately volatile elements (more volatile than silicon) in chondritic meteorites (Palme et al. 1988), which can be explained by dust coagulation during the cooling of an initially hot ($\ga 1400$ K, the sublimation temperature of silicates) disk in the terrestrial planet formation region \citep[$\la 2$ AU;][]{1974Sci...186..440L,1977Icar...30..447C,1996ApJ...472..789C}. The explanation, though, assumes that the heating of the dust at $\ga 1400$ K was global in extent, ordered and systematic. Alternatively, it is also possible that it was highly localised and, in this case, the hot initial nebula would be not necessary. For example, the so-called X-wind models assume that the hot processing occurs much closer to the star and then the matter is deposited outwards by the early Solar wind \citep{1996Sci...271.1545S,2000prpl.conf..789S,2001ApJ...548.1051G}. In addition, the detection of crystalline silicates in comets \citep{2004ApJ...612L..77W} has been also taken as a prove that the Solar System passed through a hot phase \citep{2007A&A...466L...9M}. For sure, T Tau disks do not show the high temperatures ($\ga 1400$ K at $r\sim 2$ AU) assumed by the PSN model. These temperatures are predicted in the midplane very close ($\la 0.1$ AU) to the central star or in the high atmosphere of the disk, but always at distances lower than fractions of AU \citep{2010MNRAS.407..232T}. Even the most recent models of very young and embedded PPDs, with or without gravitational instabilities \citep[see the review by][]{2012A&ARv..20...56C}, predict temperatures much lower than $\sim1400$ K. One can ask whether the hot phase of the PSN could in fact be the hot corino phase observed in Class 0 sources (\S \ref{sec:prot-phase}). Although the presently available facilities do not allow to probe regions of a few AUs, the very rough extrapolation of the temperature profile predicted for the envelope of IRAS16293-2422 (the prototype of Class 0 sources) by \cite{2010A&A...519A..65C} gives $\sim$1000 K at a few AUs. The question is, therefore: should we not compare the PSN with the hot corino models rather than the protoplanetary disks models? At present, the hot corino models are very much focused only on the $\ga 10$ AU regions, which can be observationally probed, and only consider the gas composition. Should we not start, then, considering what happens in the very innermost regions of the Class 0 sources and study the dust fate too? If so, the link with the deuterium fractionation that we observe in the hot corino phase may become much more relevant in the construction of realistic PSN models. Last but not least, the present PSN models are based on the transport and mixing of material with different initial D/H water because of the diffusion, responsible for the angular momentum dispersion. In Class 0 sources, though, the dispersion of the angular momentum is thought to be mainly due to the powerful jets and outflows and not by the diffusion of inward/outward matter. Moreover, during the Class 0 phase, material from the cold protostellar envelope continues to rain down onto the central region and the accreting disk (Z.-Y. Li et al., this volume), replenishing them by highly deuterated material.The resulting D/H gradient across the PSN may, therefore, be different than that predicted by current theories. \subsection{\textbf{Where does the terrestrial water come from?}}\label{sec:terr-water} Several reviews discussing the origin of the terrestrial water are present in the literature \citep[][van Dishoeck et al, this volume]{2009P&SS...57.1338H,2012AREPS..40..251M,2012A&ARv..20...56C}, so that we will here just summarise the points emphasising the open issues. Let first remind that the water budget on Earth is itself subject of debate. In fact, while the lithosphere budget is relatively easy to measure \citep[$\sim 10^{-4}$ M$_\earth$;][]{Lecuyer1998}, the water contained in the mantle, which contains by far the largest mass of our planet, is extremely difficult to measure and indirect probes, usually noble gases, are used for that \citep[e.g.,][]{1983Natur.303..762A}, with associated larger error bars. It is even more difficult to evaluate the water content of the early Earth, which was likely more volatile-rich than at present \citep{1996E&PSL.144..577K}. The most recent estimates give $\sim 2\times 10^{-3}$ M$_\earth$ \citep{2012E&PSL.313...56M}, namely 20 times larger than the value of the lithosphere water. If the water mantle is the dominant water reservoir, as it seems to be, then the D/H value of the terrestrial oceans may be misleading if geochemical processes can alter it. Evidently, measurements of the mantle water D/H is even more difficult. The last estimates suggest a value slightly lower than the terrestrial oceans water. In this chapter, given this uncertainty on the Earth bulk water content and D/H, we adopted as reference the evaporated ocean water D/H, the VSWOM (Tab. \ref{tab:definitions}). The ``problem'' of the origin of the terrestrial oceans rises because, if Earth was formed by planetesimals at $\sim$1 AU heliocentric distance, they would have been ``dry'' and no water should exist on Earth. One theory, called ``late veneer'', assumes that water was brought after Earth formed by, for example, comets \citep{1992AdSpR..12....5D,1995Icar..116..215O}. This theory is based on the assumptions that the D/H cometary water is the same than the Earth water D/H, but, based on observations towards comets (\S \ref{sec:the-comets}), this assumption is probably wrong. The second theory, based on the work by \cite{2000M&PS...35.1309M}, assumes that a fraction of the planetesimals that built the Earth came from more distant (2--4 AU) regions and were, therefore, ``wet''. Dynamical simulations of the early Solar System evolution have add support to this theory \citep{2000M&PS...35.1309M,2012A&A...546A..18M,2009Icar..203..644R}, challenging at the same time the idea that the flux of late veneer comets and asteroides could have been large enough to make up the amount of water on Earth. Moreover, the D/H value measured in CCs (\S \ref{sec:meteorites}) adds support to this theory. The recent findings by \citeauthor{2012Sci...337..721A} would argue for a large contribution of a group of CCs, the CI type. In summary, the origin of terrestrial water is still a source of intense debate. | 14 | 3 | 1403.7143 |
1403 | 1403.4971_arXiv.txt | We consider the status of Higgs Inflation in light of the recently announced detection of B-modes in the polarization of the cosmic microwave background radiation by the BICEP2 collaboration. In order for the primordial B-mode signal to be observable by BICEP2, the energy scale of inflation must be high, $V_{\rm inf} \approx 2 \times 10^{16} \GeV$. Higgs Inflation generally predicts a small amplitude of tensor perturbations, and therefore it is natural to ask if Higgs Inflation might accommodate this new measurement. We find the answer is essentially no, unless one considers either extreme fine tuning, or possibly adding new beyond the standard model fields, which remove some of the more attractive features of the original idea. We also explore the possible importance of a factor that has not previously been explicitly incorporated, namely the gauge dependence of the effective potential used in calculating inflationary observables, e.g. $n_S$ and $r$, to see if this might provide additional wiggle room. Such gauge effects are comparable to effects of Higgs mass uncertainties and other observables already considered in the analysis, and therefore they are relevant for constraining models. But, they are therefore too small to remove the apparent incompatibility between the BICEP2 observation and the predictions of Higgs Inflation. | \label{sec:Introduction} The theory of inflation \cite{Starobinsky:1980te, Guth:1980zm, Linde:1981mu} successfully addressed the twentieth century's greatest puzzles of theoretical cosmology. Over the past 20 years, increasingly precise measurements of the temperature fluctuations of the cosmic microwave background radiation (CMB) also confirmed the nearly scale invariant power spectrum of scalar perturbations, a relatively generic inflationary prediction. These many successes, however, underscored the inability to probe perhaps the most robust and unambiguous prediction of inflation, the generation of a background of gravity waves associated with what are likely enormous energy densities concomitant with inflation ({\it e.g.}, \cite{Krauss:1992ke}). Recently, the BICEP2 collaboration reported evidence of B-modes in the polarization pattern of the CMB \cite{Collaboration:2014fk}. The B-modes result from primordial gravity wave induced distortions at the surface of last scattering. If one assumes that these gravity waves are of an inflationary origin, then the BICEP2 measurement corresponds to an energy scale of inflation: \begin{align}\label{eq:BICEP} V_{\rm inf}^{1/4} \approx (2\pm 0.2) \times 10^{16} \GeV \end{align} for a reported tensor-to-scalar ratio of $r \approx 0.2 ^{+ 0.07} _{-0.05}$ (using also the Planck collaboration's measurement of the amplitude of the scalar power spectrum \cite{Ade:2013rta}). Such a high scale of inflation rules out many compelling models. For the purposes of this paper, we will assume that the observation $r \approx 0.2$ is valid\footnote{Note that the BICEP2 measurement is in tension with the upper bound, $r < 0.11 \atCL{95}$, obtained previously by the Planck collaboration \cite{Ade:2013rta}. }, and we assess the impact of this measurement on a particular model of inflation, known as Higgs Inflation. Higgs Inflation ({\rm HI}) postulates that the Standard Model Higgs field and the inflaton are one in the same \cite{Bezrukov:2007ep}. (See also \rref{Bezrukov:2013fka} for a recent review). This powerful assumption allows {\rm HI} to be, in principle much more predictive than many other models of inflation, as by measuring the masses of the Higgs boson and the top quark at the electroweak scale ($100 \GeV$), one might predict observables at much larger energy scales associated with inflation ($V_{\rm inf}^{1/4} \lesssim 10^{16} \GeV$). However, in practice this enhanced predictive power is elusive due to a strong sensitivity to quantum effects, unknown physics, and other technical subtleties in the model. Specifically, one connects observables at the electroweak and inflationary scales using the renormalization group flow (RG) of the SM couplings \cite{DeSimone:2008ei, Barvinsky:2008ia, Bezrukov:2010jz, Bezrukov:2012sa, Allison:2013uaa, Salvio:2013rja}. It is reasonable however to expect that there is new physics at intermediate scales, and even if the SM is extended only minimally to include a dark matter candidate \cite{Clark:2009dc} or neutrino masses \cite{Okada:2009wz, He:2010sk, Rodejohann:2012px, Kobakhidze:2013pya} this new physics can qualitatively affect the connection between electroweak and inflationary observables. Moreover, perturbative unitarity arguments require new physics just above the scale of inflation \cite{Barbon:2009ya, Burgess:2010zq}, and in addition the unknown coefficients of dimension six operators can significantly limit the predictive power of {\rm HI} \cite{Burgess:2014lza}. The {\rm HI} calculation also runs into various technical subtleties that arise from the requisite non-minimal gravitational coupling (see below) and quantization in a curved spacetime \cite{George:2012xs,George:2013iia, Prokopec:2014iya}. Finally, it is worth noting that {\rm HI} is also at tension with the measured Higgs boson and top quark masses, and an $O(2\sigma)$ heavier Higgs or lighter top is required to evade vacuum stability problems \cite{Buttazzo:2013uya}. Also, as we shall later discuss in detail, there is one additional source of ambiguity in calculations of {\rm HI} that had not been fully explored. Since the quantum corrections are significant when connecting the low energy and high energy observables, one should not work with the classical (tree-level) scalar potential, as is done in may models of inflation, but one must calculate the quantum effective potential. It is well-known that in a gauge theory the effective potential explicitly depends upon the choice of gauge in which the calculation is performed \cite{Weinberg:1973ua, Jackiw:1974cv}, and care must be taken to extract gauge-invariant observables from it \cite{Dolan:1974gu, Kang:1974yj, Nielsen:1975fs, Fukuda:1975di} (see also \cite{Aitchison:1983ns, Patel:2011th}). This fact can perhaps be understood most directly by recalling that the effective action is the generating functional for one-particle irreducible Green's functions, which themselves are gauge dependent \cite{Jackiw:1974cv}. In practice one often neglects this subtlety, fixes the gauge at the start of the calculation, and calculates observables with the effective potential as if it were a classical potential. In the context of finite temperature phase transitions, it is known that when calculated naively in this way, the predictions for observables depend on the choice of gauge used \cite{Weinberg:1974hy, Dolan:1973qd, Bernard:1974bq, Patel:2011th, Wainwright:2011qy, Wainwright:2012zn, Garny:2012cg}. Because of the extreme tension between {\rm HI} models and the data, we assess here the degree to which this gauge uncertainty might affect the observables in Higgs Inflation. We find that the gauge ambiguity introduces uncertainties that are comparable to the variation of the physical parameters, {\it i.e.} the Higgs mass. As a result, this ambiguity alone cannot resuscitate moribund models. | \label{sec:Conclusion} The recent detection of B-modes by the BICEP2 collaboration represents a profound and exciting leap forward in our ability to explore fundamental physics and the early universe. If the measurement of $r \approx 0.2$ is confirmed, then it is reasonable to expect that, in the not-too-distant future, measurements of the spectrum of primordial tensor perturbations will become possible, allowing further tests of inflation. And if the measured $r$ can unambiguously be shown to be due to inflation, then this also substantiates the quantization of gravity \cite{Krauss:2013pha}. Thus, future observations will provide significant constraints on particle physics and models of inflation. However the simple observation of non-zero $r$ already signals the death knell for low-scale models of inflation. This includes the class of models captured by the potential in \eref{eq:HI_potential}, and among these apparently Higgs Inflation. We have shown that $r \approx 0.2$ essentially excludes canonical Higgs Inflation in the absence of extreme fine tuning. The Higgs field may live on as the inflaton but only with significant non-minimal variants of {\rm HI}. In our analysis we have also drawn attention to the issue of gauge dependence in the Higgs Inflation calculation. We find that the energy scale of inflation acquires an artificial dependence on the gauge fixing parameter by virtue of the gauge dependence of the effective potential from it is extracted. However, we find this gauge dependence of the scale of inflation is comparable to the dependence on other physical parameter uncertainties, which are themselves small. While this may be important for model building purposes, it does not affect the robustness of the fact that large $r$ disfavors Higgs Inflation. | 14 | 3 | 1403.4971 |
1403 | 1403.6373_arXiv.txt | {Gravity waves (or their signatures) are detected in stars thanks to helio- and asteroseismology and they may play an important role in the evolution of stellar angular momentum. Moreover, the observational study of the CoRoT target HD\,51452 by Neiner and collaborators demonstrated the potential strong impact of rotation on the stochastic excitation of gravito-inertial waves in stellar interiors.} {Our goal is to explore and unravel the action of rotation on the stochastic excitation of gravity and gravito-inertial waves in stars.} {The dynamics of gravito-inertial waves in stellar interiors, both in radiation and in convection zones, is described with a local { non-traditional} f-plane model. Their couplings with convective turbulent flows leading to their stochastic excitation is studied in this framework.} {First, we find that, in the super-inertial regime in which the wave frequency is over twice the rotation frequency ($\sigma>2\Omega$), the evanescence of gravito-inertial waves in convective regions decreases with decreasing wave frequency. Next, in the sub-inertial regime ($\sigma<2\Omega$), gravito-inertial waves become purely propagative inertial waves in convection zones. { Simultaneously, turbulence in convective regions is modified by rotation. Indeed, the turbulent energy cascade towards small scales is slowed down and in the case of rapid rotation, strongly anisotropic turbulent flows are obtained that can be understood as complex non-linear triadic interactions of propagative inertial waves.} These { different} behaviours, due to the action of the Coriolis acceleration, strongly modify the wave couplings with turbulent flows. { On one hand, turbulence weakly influenced by rotation is coupled with evanescent gravito-inertial waves. On the other hand, rapidly rotating turbulence is intrinsically and strongly coupled with sub-inertial waves. Finally, to study these mechanisms, the traditional approximation cannot be assumed because it does not properly treat the couplings between gravity and inertial waves in the sub-inertial regime.}} {Our results demonstrate the action of rotation on stochastic excitation of gravity waves thanks to the Coriolis acceleration which modifies their dynamics in rapidly rotating stars { and turbulent flows}. As the ratio $2\Omega/\sigma$ increases, the couplings and thus the amplitude of stochastic gravity waves are amplified.} | Gravity waves propagate in stably stratified stellar radiation zones, such as the radiative core of low-mass stars and the external radiative envelope of intermediate-mass and massive stars \citep[e.g.][]{Aertsetal2010}. When such waves (or their signatures) are detected thanks to helioseismology \citep[][]{Garciaetal2007} and asteroseismology \citep[e.g.][]{Becketal2011,Papicsetal2012,Neineretal2012}, they constitute a powerful probe of stellar structure \citep[e.g.][]{TC2011,Beddingetal2011} and internal dynamics, for example of differential rotation \citep[][]{Garciaetal2007,Becketal2012,Deheuvelsetal2012,Mosseretal2012}. Furthermore, when they propagate, gravity waves are able to transport and deposit a net amount of angular momentum because of their radiative damping and corotation resonances \cite[e.g.][]{GoldreichNicholson1989,Schatzman1993,ZahnTalonMatias1997}. Therefore, they are invoked, together with magnetic torques, to explain the quasi-uniform rotation until 0.2R$_{\odot}$ in the solar radiative core \citep[][]{CharbonnelTalon2005}, the mixing of light elements in low-mass stars \citep[][]{TalonCharbonnel2005}, the weak differential rotation in subgiant and red giant stars \citep[][]{Cellieretal2012}, and the transport of angular momentum necessary to explain mass-loss in active massive stars such as Be stars \citep[][]{Huatetal2009,Neineretal2013,LeeSaio1993,Lee2013}. Thus, it is necessary to get a good understanding of their excitation mechanisms and a precise prediction of their amplitude. In single stars, two mechanisms can excite gravity waves: the $\kappa$-mechanism due to opacity bumps \citep[e.g.][]{Unnoetal1989,GastineDintrans2008} and stochastic motions both in the bulk of convective regions and at their interfaces with adjacent radiation zones where turbulent convective structures (plumes) penetrate because of their inertia \citep[e.g.][]{Press1981,Browningetal2004,RogersGlatzmaier2005,Belkacemetal2009,Cantielloetal2009,Samadietal2010,BMT2011,Rogersetal2012,Shiodeetal2013}. In this work, we focus on stochastic excitation.\\ In addition to stochastically excited mixed gravito-acoustic modes currently detected with space asteroseismology, e.g. in red giants \citep[e.g.][]{Becketal2011}, a detection has been obtained thanks to CoRoT in the hot Be star HD\,51452 \citep{Neineretal2012}. In this star, which rotates close to its critical angular velocity, authors discovered gravity modes strongly influenced by the Coriolis acceleration, i.e. gravito-inertial modes \citep[][]{DintransRieutord2000,Mathis2009,Ballotetal2010}. They propose that these modes are probably excited stochastically by turbulent convection in the core or/and in the subsurface convection zone, since this star is too hot to excite gravity modes with the $\kappa$-mechanism. Moreover, the detected gravito-inertial modes with the largest amplitudes have frequencies below the inertial one at $2\Omega$ (where $\Omega$ is the angular velocity of the star), which are the most influenced by the Coriolis acceleration. Therefore, this discovery points the potential important action of rotation on stochatic excitation of gravity and gravito-inertial waves (hereafter GIWs) in stellar interiors. This has been poorly explored until now \citep[][]{BelkacemMathisetal2009,Rogersetal2013}. To identify the associated mechanisms and unravel the related signatures, we choose here to generalise the work by \cite{GoldreichKumar1990} and \cite{LecoanetQuataert2013} taking the action of the Coriolis acceleration into account. First, in Sect.~\ref{giw}, we present the local rotating set-up in which we describe the different regimes of the dynamics of GIWs both in radiation and in convection zones. Next, in Sect.~\ref{turb}, their stochastic excitation by turbulent convective flows is studied { with a particular focus on the effects of slow and rapid rotation}. Finally, in Sect.~\ref{discuss}, we conclude on the impact of rotation on the stochastic excitation of GIWs in stars and we discuss the consequences for asteroseismology and for the study of their angular momentum evolution. | \label{discuss} In this work, we have formally demonstrated that rotation, through the Coriolis acceleration, modifies the stochastic excitation of gravity waves and GIWs, the control parameters being the wave's Rossby number $R_{\rm o}=\sigma/2\Omega$ { and the non-linear Rossby number $R_{\rm o}^{\rm NL}$ of convective turbulent flows}. On one hand, in the super-inertial regime ($\sigma>2\Omega,\hbox{ {\it i.e.} }R_{\rm o}>1$), the evanescent behaviour of GIWs above (below) the external (internal) turning point becomes increasingly weaker as the rotation speed grows until $R_{\rm o}=1$. { Simultaneously, the turbulent energy cascade towards small scales is slowed down.} The coupling between super-inertial GIWs and given turbulent convective flows is then amplified as described in Eq. (\ref{FA1}). On the other hand, in the sub-inertial regime ($\sigma<2\Omega,\hbox{ {\it i.e.} }R_{\rm o}<1$), GIWs become propagative inertial waves in the convection zone. { In the case of rapid rotation, turbulent flows, which become strongly anisotropic, result from their non-linear interactions. Sub-inertial GIWs that correspond to propagating inertial waves in convection zones are then intrinsically and strongly coupled to rapidly rotating turbulence as discussed in Sect. 3.2.3. These different regimes are summarised in Fig. \ref{Fig6MN}.} Such effects are of great interest for asteroseismic studies of rotating stars since gravity wave and GIW amplitudes are thus expected { to} be stronger in rapidly rotating stars, { a conclusion that is supported by recent numerical simulations \citep{Rogersetal2013}}. For example, until recently, stochastically excited gravity waves were thought to be of too low amplitude to be detected even with space missions such as CoRoT \citep{Samadietal2010}. The discovery of stochastically excited GIWs in the rapid rotator HD\,51452 \citep{Neineretal2012} proved that such waves can be detected. Our results show how the amplitude of these waves can be enhanced thanks to rotation. The interpretation of observed pulsational frequencies in rapid rotators should take this into account. In particular, oscillations observed in $\beta$\,Cep and Slowly Pulsating B (SPB) stars should not be systematically attributed to the $\kappa$-mechanism, as it was done until now, if the star rotates fast. For example, the GIWs observed in HD\,43317 \citep{Papicsetal2012} might be of stochastic origin and might have been enhanced by rapid rotation. This is of course especially true for Be and Bn stars. Moreover, the related transport of angular momentum, which until now was believed to become less efficient because of GIWs equatorial trapping in the sub-inertial regime \citep{Mathisetal2008,Mathis2009,MdB2012}, may be sustained thanks to the stronger stochastic excitation by turbulent convective flows. This may have important consequences for example for rapidly rotating young low-mass stars \citep{CharbonnelTalon2008,Charbonneletal2013} and active intermediate-mass and massive stars such as Be stars. For example, \cite{Neineretal2013} proposed that the outburst of the Be star HD\,49330 observed by CoRoT \citep{Huatetal2009} is due to the deposit of angular momentum by GIWs just below the surface \citep[see also][]{Lee2013}. Our prediction should now be confronted to realistic numerical simulations of stochastic excitation of GIWs in stellar interiors \citep[e.g.][]{BMT2011,Rogersetal2012,Rogersetal2013,ABM2013}, to laboratory experiments \citep{Perrardetal2013}, as well as to a larger statistical sample of observed pulsating stars with detected GIWs. Moreover, a global formalism to treat the problem in the spheroidal geometry corresponding to rotating stars must be built in a near future. | 14 | 3 | 1403.6373 |
1403 | 1403.4006_arXiv.txt | This work is dedicated to investigation of galaxies that do not fit into a common scenario of galaxy formation -- isolated lenticular galaxies. We have studied stellar populations and ionized gas content of a sample of 22 lenticular galaxies (among those 4 targets have appeared to be of erroneous morphological classification) by undertaking deep long-slit spectroscopy with the Russian 6-m telescope and with the Southern African Large Telescope (SALT). The obtained average ages of the stellar populations in bulges and discs covers a wide range between 1.5 and $>15$ Gyr, that indicates the absence of distinct epoch of their stellar content formation. In contrast to galaxies in groups and clusters, the stellar population ages in bulges and discs of isolated lenticulars tend to be equal, that supports the inefficiency of the bulge rejuvenation in sparse environment. Almost all the lenses and rings possess intermediate ages of the stellar populations, within the range of $2-5$ Gyr. By analyzing the emission-line spectra of galaxies, we have found that 13 out of 18 ($72\pm11$ \%) objects of our sample possess extended emission-line structures; among those, 6 galaxies ($46\pm14$ \%) demonstrate decoupled gas kinematics with respect to their stellar discs. We have found starforming off-nuclear regions in 10 galaxies; their gas oxygen abundances are nearly solar that implies tidal gas accretion from gas-rich dwarf satellites rather than accretion from cosmological large-scale structure filaments. | One of the central topics in current extragalactic astronomy is how galaxies form and how their properties change through cosmic times. This is a particularly difficult task if concerning lenticular galaxies since this class of objects show a diversity in their properties that goes beyond present day simulations. The standard scenario of formation of lenticular galaxies relates to transformation of spiral galaxies into lenticulars by dense environment effect: ram pressure in hot intracluster/intragroup medium \citep{gunn_gott_1972,quilis_2000}, gravitational tides and harassment \citep{byrd_1990,moore_1996}, direct encounters of galaxies \citep{spitzer_1951, icke_1985}, starvation \citep{larson_1980}. However 15\% of nearby field galaxies are lenticulars \cite{naim_1995}, and there are examples of strictly isolated S0 \cite{sulentic_2006}. The study of isolated lenticulars by using deep optical spectroscopy methods provided information not only for central part of galaxies but also for disc components, is important issue because it provides a crucial point for the testing of scenario formation of lenticulars at all. | In this work we have presented the results of deep long-slit spectroscopy (total exposures per target 0.3-3 hour) with the 6-m Russian telescope and with Southern African Large Telescope (SALT) for sample of 18 isolated lenticular galaxies. By using full spectral fitting technique as well as Lick indices approach and by analysing of the pure emission line spectra we have obtain that: \begin{itemize} \item The obtained average ages of the stellar populations in bulges and discs covers a wide range between 1.5 and $>15$ Gyr, that indicates the absence of distinct epoch of their stellar content formation. \item In contrast to galaxies in groups and clusters, the stellar population ages in bulges and discs of isolated lenticulars tend to be equal, that supports the inefficiency of the bulge rejuvenation in sparse environment. \item Almost all the lenses and rings possess intermediate ages of the stellar populations, within the range of $2-5$ Gyr. On average, the chemical abundances ([Z/H], [Mg/Fe]) in the lenses and rings are the same as in discs, also dynamically the stellar population in lenses and rings do not sufficiently stand out against a discs. \item More than half of the sample of isolated lenticulars ($72\pm11$ \%) possess extended emission-line structures; among those, 6 galaxies ($46\pm14$ \%) demonstrate decoupled gas kinematics with respect to their stellar discs. \item We have found starforming off-nuclear regions in 10 galaxies; their gas oxygen abundances are nearly solar that implies tidal gas accretion from gas-rich dwarf satellites rather than accretion from cosmological large-scale structure filaments. \end{itemize} \renewcommand{\baselinestretch}{0.7} {\small | 14 | 3 | 1403.4006 |
1403 | 1403.6832_arXiv.txt | Emission line diagnostic diagrams probing the ionization sources in galaxies, such as the Baldwin-Phillips-Terlevich (BPT) diagram, have been used extensively to distinguish AGN from purely star-forming galaxies. Yet, they remain poorly understood at higher redshifts. We shed light on this issue with an empirical approach based on a $z\sim0$ reference sample built from $\sim$300,000 SDSS galaxies, from which we mimic selection effects due to typical emission line detection limits at higher redshift. We combine this low-redshift reference sample with a simple prescription for luminosity evolution of the global galaxy population to predict the loci of high-redshift galaxies on the BPT and Mass-Excitation (MEx) diagnostic diagrams. The predicted bivariate distributions agree remarkably well with direct observations of galaxies out to $z\sim1.5$, including the observed stellar mass-metallicity ($MZ$) relation evolution. As a result, we infer that high-redshift star-forming galaxies are consistent with having {\it normal} ISM properties out to $z\sim1.5$, after accounting for selection effects and line luminosity evolution. Namely, their optical line ratios and gas-phase metallicities are comparable to that of low-redshift galaxies with equivalent emission-line luminosities. In contrast, AGN narrow-line regions may show a shift toward lower metallicities at higher redshift. While a physical evolution of the ISM conditions is not ruled out for purely star-forming galaxies, and may be more important starting at $z\gtrsim2$, we find that reliably quantifying this evolution is hindered by selections effects. The recipes provided here may serve as a basis for future studies toward this goal. Code to predict the loci of galaxies on the BPT and MEx diagnostic diagrams, and the $MZ$ relation as a function of emission line luminosity limits, is made publicly available. | Nebular emission lines can reveal crucial information on the ionized gas content in galaxies. In particular, several optical emission line diagnostics have been developed to probe gas properties such as metallicity, ionization parameter, electron density and temperature \citep{ost06}, which can in turn provide additional insights on the source of ionization of the gas. An important application is thus the identification of active galactic nuclei (AGNs), which leave strong signatures on nebular line ratios such as \oiiilam/\hb\ and/or \niilam/\ha. These two line ratios form the most traditional version of the BPT diagram \citep{bpt,vei87}. The latter has been calibrated with both a theoretical approach \citep{kew01,sta06,kew13a} and empirically with low-redshift galaxies \citep{kau03c}. There are now questions about the applicability of low-redshift nebular line diagnostics to higher-redshift objects. A number of studies suggest that high-redshift galaxies are offset from the locus of low-redshift reference samples on the standard BPT diagram (\oiiilam/\hb\ vs. \niilam/\ha) \citep[e.g.,][]{sha05,erb06,tru13,new14,hol14}. While there are a few hypotheses, the cause of this offset is not yet fully explained. For example, it was suggested that high-redshift galaxies may have had different \hii\ region conditions (such as electron densities, temperatures, pressures, etc.) relative to the bulk of star-forming galaxies \citep{bri08,liu08,hai09,leh09,rig11,ly13}. It was suggested \citep{leh13,shi13} that this may be due to galaxies globally forming their stars with a higher surface density in the past, which has been inferred from infrared luminosity surface densities \citep{red12} and galaxy infrared SED fitting \citep{mag12}. However, other studies claim that the offsets on excitation diagrams are instead caused by an increased contribution from AGN \citep{gro06,wri10,tru11}, which would shift the galaxies in a similar way \citep{kew13a}. If there were a higher incidence of AGN in galaxies in the past we may expect a steeper ionization profile and thus varying emission line strengths. It is crucial to disentangle the source of ionization in galaxies (young stars vs. AGN) in order to interpret and derive important quantities in galaxy evolution studies like star formation rates (SFRs), metallicities, and gas dynamics, but also to understand the interplay between black hole growth and stellar growth in galaxies. Furthermore, another complication arises because intermediate- and high-redshift galaxy samples used thus far may suffer from strong selection biases due to the emission line detection limits. Relative to existing large spectroscopic sample at low redshifts (e.g., SDSS), only galaxies with intrinsically luminous lines can be detected at intermediate to high redshifts. These potential selection biases have been mostly neglected thus far, and will be explored in this Paper along with genuine evolutionary trends. As we will show, emission line detection limits add complexity to the problem, but not taking them into account can yield misguided interpretations of how galaxy properties evolve with redshift. In addition to the traditional BPT-\nii\ diagnostic diagram, we revisit an alternative diagram using stellar mass in place of \nii/\ha\ \citep[the Mass-Excitation (MEx) diagnostic diagram from][hereafter J11]{jun11}. The MEx diagram has the advantage of requiring only the \oiii/\hb\ emission lines, which are more widely separated in wavelength and therefore easier to resolve spectroscopically than \ha\ and \nii. Furthermore, they be can observed to higher redshift in any given wavelength regime. In optical spectra, \nii/\ha\ are available out to $z\sim0.45$ whereas \oiii/\hb\ can be observed out to $z\sim0.9$. Similarly, NIR spectra in the $K$ band cover \nii/\ha\ out to $z\sim2.5$ but \oiii/\hb\ out to $z\sim3.7$. Another advantage of the MEx diagram is its probabilistic approach. For a given location on the MEx plane, and given the measurement errors, the MEx diagram yields the probability that the galaxy hosts an AGN. This method has a built-in uncertainty in the sense that ambiguous cases will have a low or intermediate AGN probability, and is well suited for statistical studies because the AGN probabilities can be used as statistical weigths to weigh for (or against) AGN. On the other hand, one might expect the MEx diagram to be more sensitive to evolution of the stellar-mass metallicity ($MZ$) relation \citep{sav05,sha05,erb06,yab12,zah13}, than the traditional BPT. We will show that an improved treatment of emission-line detection limit mitigates such bias by directly accounting for appropriate gas-phase metallicities when building a tailored low-redshift comparison sample for each survey. The aim of this Paper is twofold. First, we provide improved AGN diagnostics that account for redshift-dependent effects. More specifically, we revisit both the original BPT diagram and more recent MEx diagram in order to disentangle selection and evolution effects, and to improve their applicability to a broad range of redshifts. In addition, this work reveals insight into the ISM conditions in higher redshift galaxies, once the selection effects are taken into account. The Paper is organized as follows. We describe the galaxy samples used for low-redshift calibration and higher redshift applications in Section~\ref{sec:sample}, followed by the low-redshift revision of the MEx demarcations in Section~\ref{sec:revised}. The results (Section~\ref{sec:result}) include empirical predictions of the redshift evolution of the BPT and MEx diagrams including both genuine evolution and selection effects due to line detection limits (Section~\ref{sec:detlim}). These predictions are confronted with observations out to $z\sim2$ (Section~\ref{sec:hiz}), and compared to theoretical predictions from \citet{kew13a} in Section~\ref{sec:theo}. The implications for the high-redshift application of emission line diagnostic diagrams are discussed in Section~\ref{sec:discu}, including the stellar mass-metallicity relation, before the main findings are summarized in Section~\ref{summ}. Throughout this paper, we assume a flat $\Lambda$CDM cosmology ($\Omega_m = 0.3$, $\Omega_{\Lambda} = 0.7$, and $h = 0.7$) and a \citet{cha03} initial mass function (IMF). | \label{sec:discu} \subsection{Mass-Metallicity Relation}\label{sec:MZ} Several studies have reported evolution of the stellar mass-metallicity ($MZ$) relation for galaxies \citep{sav05,erb06,zah13}, or along a plane in the $M_{*}-SFR-Z$ space \citep{man10,lar10,yat12}. There are variations in the details of this evolution and how it may vary with galaxy stellar masses, but the general sense is that at a given stellar mass, lower redshift galaxies have higher gas-phase metallicities than their higher redshift counterparts. This could indicate the global enrichment in galaxies as they evolve and form stars. \begin{figure*}[bht] \epsscale{1.0} \plotone{f7.eps} \caption{ $MZ$ relation from SDSS observations at low-redshift (this work) and from observations compiled by \citet{zah13}, who fitted DEEP2 $z\sim0.8$ and Y12 $z\sim1.4$ results (thick orange curves). In all panels, the blue logarithmically-spaced contours show the full SDSS prior galaxy sample, with filled circles marking the median metallicity in stellar mass bins of 0.1~dex. The blue line shows the fitted relation using the functional form introduced by \citet{mou11}. Green and red contours correspond to a SDSS subsample with \ha\ and \oiii\ lines above the labeled luminosity. The median metallicity in stellar mass bins of 0.1~dex is shown with filled symbols. The choice of the threshold luminosity is made according to the fiducial scenarios, where the luminosity detectability threshold is lowered by the same amount as $L^*$(\ha) fades between the redshifts of interest: (a) $z\sim0.8$ and (b) $z\sim1.4$, and the prior sample at $z\sim0.09$. In this case, the predicted contours and median values agree remarkably well with the observed relations compiled by Z13. Alternative scenarios and more details are included in Appendix~\ref{app:altern}. Contours are logarithmically spaced (0.5~dex) with the outermost contour corresponding to 50 galaxies per bin of 0.15\,dex$\times$0.15\,dex. }\label{fig:MZ} \end{figure*} The revised MEx demarcations and probabilities now include metallicity evolution indirectly through the use of line luminosities to build a prior sample. As we show below, together with the stellar mass, emission line luminosities can trace metallicity to some extent. This is similar, though not identical, to previous work defining the Fundamental Metallicity Relation \citep{man10,lar10,yat12,cre12}. In the latter, the SFR is used as a third parameter to define a plane while here we use both \ha\ luminosity, a tracer of the SFR \citep{ken98}, and \oiii, which depends more directly on the gas-phase metallicity. We also {\it fade} the line luminosity threshold by the corresponding fading of the knee of the \ha\ luminosity function. We test directly our selection method against observed $MZ$ relations compiled by \citet[hereafter Z13]{zah13}. These authors fitted the functional form defined by \citet{mou11} to datasets in five redshift slices, including an intermediate-redshift DEEP2 sample ($z\sim0.8$), and the Y12 sample ($z\sim1.4$), which overlap with the present study. These $MZ$ relations are compared to our empirical predictions from SDSS prior samples. We compute metallicities following the same method as Z13 to facilitate a direct comparison. As in that study, we apply the KK04 calibration \citet{kob04} to the required \oii, \hb, and \oiii\ emission lines. Line fluxes were first corrected for dust attenuation by measuring the Balmer decrement and applying the \citet{cal00} dust attenuation curve with $R_V=4.05$, assuming an intrinsic ratio of \ha/\hb=2.86 \citep{ost06}. In Figure~\ref{fig:MZ}, we show predicted contours using our fiducial approach (\ha\ and \oiii\ selection above the luminosity threshold that includes $L^*_{\ha}$ evolution). There is a good global agreement between predicted contours and analytical fits to observations reported by Z13, indicating that the observed $MZ$ evolution is a built-in feature of our method. Empirical predictions were also made for two alternative scenarios (Appendix~\ref{app:altern}), namely to compare with samples selected from a single line: \oiii\ or \ha, but also to consider null evolution scenarios where the line luminosity detection limits are used as is. We note that the empirical prediction of the Y12 sample (red contours on panel b) show a cutoff at high masses ($>10^{10.5}$~\Msun). While galaxies with these high masses and low metallicities may be more common at higher redshifts. This is likely a true evolutionary trend, and could indicate that most massive galaxies have had time to enrich their ISM by $z<0.2$ and that the galaxies that remain with low metallicities at this more recent epoch have lower stellar masses. \subsection{Is there evolution in the emission-line diagnostics?}\label{sec:shift} In Section~\ref{sec:detlim}, we investigated the consequences of increasing emission-line luminosity limits to select subsamples from $z\sim0$ SDSS emission-line galaxies, and the resulting bivariate distributions on the BPT and MEx diagrams. We found that the division line developed by \citet{kau03c} still appears to be applicable to divide the two branches defined by the data at all the luminosity thresholds tested, albeit a small shift of $<$0.2dex toward higher \nii/\ha\ and/or \oiii/\hb\ could be possible at the highest line luminosities. This suggests that the luminosity dependence would more strongly affect the locus of the sample than the definition of the dividing lines. In Section~\ref{sec:hiz}, we compared $z\sim1.5-2.5$ observations with empirical predictions from SDSS subsamples selected to have equivalent line luminosity limits. We found a general but imperfect agreement between the observations and matched SDSS prior samples on the BPT diagrams. Namely, some galaxies from the T13 and N14 samples are located between the two empirically predicted BPT branches on the right-hand side panels of Figure~\ref{fig:hizbpt}. This feature could be due to lower NLR metallicity, such as predicted by K13 with their Scenario~2 (second row of Figure~\ref{fig:allsce}). Alternatively, it could be due to AGN host having higher SFRs at higher redshifts, which would boost \ha\ more strongly than \nii\ and dilute the AGN signatures. Lower NLR metallicites and/or higher SFRs at higher redshift imply a potentially greater fraction of AGNs in the BPT-composite region (between the \citealt{kau03c} and \citealt{kew01} lines) and even toward the top of the star-forming branch. This trend could explain the presence of X-ray (or otherwise securely-identified) AGNs on the boundary between the composite and star-forming regions of the BPT-\nii\ diagram in the T13 and N14 samples (red symbols in Figure~\ref{fig:hizbpt}) and the presence of data points to the left of the AGN branch on the BPT diagram. As a consequence, AGN samples should include BPT-composites to improve their completeness \citep[also see][]{tro11}. Conversely, purely star-forming galaxy samples may be harder to obtain given the higher risk of including AGN contaminants when NLR metallicities are lower. A potential solution may be to combine both the BPT and MEx classification schemes in such ambiguous cases. How do these trends compare on the MEx diagram? In contrast to the BPT diagram, the splitting of the two branches traced by $z\sim0$ SDSS samples on the MEx diagram shows an obvious offset with increasing line luminosity limits (Figure~\ref{fig:mex}). This offset corresponds to changes in the AGN probabilities as a function of line luminosity limits, with the transition region shifting toward higher stellar masses with increasing line luminosity limits (fitted in Appendix~\ref{sec:offset}). The physical interpretation depends on the nature of the main ionizing source. For purely star-forming galaxies, brighter lines correspond to lower metallicity and/or higher SFRs, and therefore higher \oiii/\hb\ ratios. However, on the AGN side, brighter lines imply higher accretion rates onto black holes (traced by \oiii) as well as higher SFR of the hosts (traced by \ha). Both of these criteria will favor high-mass hosts. First, massive galaxies are more likely to host high-mass black holes \citep[e.g.][]{mag98,geb00,fer00}, which will be brighter than lower mass black holes for a given Eddington ratio distribution \citep[also see][]{air12}. Second, high mass galaxies are more likely to have high SFRs, according to the $M_{*}-$SFR sequence for star-forming galaxies\footnote{The emission line selection will favor star-forming or active galaxies over truly passive systems which do not follow the $M_{*}-$SFR sequence. This bias will be increasingly important as the emission line luminosity limit increases, therefore increasing the likelihood that the selected galaxies are star-forming and not passive.} \citep{bri04,noe07,elb11}. Therefore, these trends likely combine to shift the star-forming galaxies/AGN division toward higher stellar masses. Regardless of the underlying interpretation, the empirical offset has been calibrated, allowing one to trace the MEx demarcation lines for various line sensitivity limits. The MEx AGN probabilities can be calculated taking into account both the survey detection limit, and the individual redshift of each galaxy (which will determine the amount of $L^*$ evolution). While the $z>1$ samples on the BPT diagram suggested offsets of AGN NLRs toward lower metallicities (i.e., lower \nii/\ha\ ratios), the MEx diagram is largely insensitive to such an effect, as the \oiii/\hb\ would vary in the opposite direction, \emph{helping} the AGN selection. Instead, the MEx diagnostic has the caveat of being incomplete to select AGNs at low stellar masses ($<10^{9.5-10}$~$M_{\sun}$). This limitation depends slightly on line luminosities with this new approach (galaxies with more luminous emission liness can only be recognized as AGNs in higher mass hosts relative to galaxies with fainter lines). This limitation may not be very severe if AGNs in low-mass hosts are rare \citep{bel11,tan12}, and/or if their relative importance was lower at higher redshift, as would be the case if AGN activity followed the downsizing phenomenon \citep[e.g.,][]{bar05,kel13,hir13}. The situation is different if one is particularly interested in low-mass AGN hosts, low-mass black holes and black hole seeds \citep{gre07,bar08,dia13,rei13}. \subsection{Lower NLR metallicity at higher redshifts: gas-poor hosts or metal dilution?}\label{sec:interp} NLRs tend to be metal-rich in nearby galaxies, and while there is evidence for low-metallicity AGNs, those systems are rare \citep{gro06}. In this study, we find tentative evidence that higher redshift AGNs may be less chemically enriched than their local counterparts. We discuss a few physical mechanisms that could explain this trend, in light of our recent understanding of high-redshift star-forming galaxies. As also mentioned by \citet{kew13a}, emission lines from NLRs may trace gas closer to galaxy nuclei than global galaxy-scale spectra. Negative metallicity gradients could therefore play a role in explaining that NLRs are typically more metal-rich than the surrounding galaxy ISM at larger scales. In this view, if disk galaxies start with flatter gradients at higher redshifts, the NRLs in these hosts would also exhibit more metal-poor characteristics. Conversely, if the nuclei of galaxies enrich on very short timescales, then the NLRs would be metal-rich already in higher redshift systems, meaning that their hosts had steeper metallicity gradients. Thus, NLR metallicities could trace whether host galaxies have had time to enrich at least their central regions, and whether disks grow inside-out \citep[e.g.,][]{jon13}. However, this simple picture may not hold during galaxy mergers because metal-rich nuclei could be diluted by inflows of more pristine gas brought in from the outskirts or from satellite galaxies. Galaxy mergers and interactions have indeed been reported to exhibit flattened or inverted metallicity gradients \citep[e.g.][]{kew10,rup10,que12}, though major merger events only account for a small fraction of the total star-forming and AGN galaxy population at a given time. Here, we discuss one more possibility related to violent disk instabilities \citep[VDIs; e.g.,][]{bou07,dek09}. High-redshift disk galaxies have been observed to have high gas fractions at $z>1.5$ \citep{dad10,tac10}, and clumpy appearances \citep[e.g.][]{elm07,elm09} that distinguish them from typical star-forming disk galaxies observed at low redshifts, and that are interpreted as observational signatures of VDIs. These instabilities are predicted to be ubiquitous at $z>1-2$, and to generate inflows toward the central regions \citep{kru10,bou11}. These inflows can bring metal-poor gas in the viscinity of active BHs and result in lower metallicity NLRs. While VDI clumps are themselves star-forming and produce metals, they undergo outflows \citep{gen11} as they migrate inwards but also accrete gas from their diffuse surrounding \citep{dek13,bou14}. This could potentially maintain somewhat lower metallicities for these clumps, or the general turbulence could contribute to erase or flatten metallicity gradients \citep[e.g.,][for an example case]{que12}. The details and timescales are still uncertain, but we speculate that VDIs could play a role in determining the observed metallicities of NLR gas at higher redshifts ($z>1$). \subsection{Comparison with previous MEx diagram results}\label{sec:comp} In this work, we have revised the MEx demarcation and probability calculations to use the SDSS DR7 sample (Section~\ref{sec:revised}). We have then implemented changing MEx demarcation and probabilities as a function of the effective luminosity threshold for emission line detection (applied to \ha\ and \oiii). How do these revisions compare to other studies of the MEx at high-redshift from the literature? On the one hand, T13 had found a good agreement between the BPT classification and the original MEx dividing curves from J11, and concluded that the original MEx classification were valid for their sample at $z\sim1.5$. The revised demarcations are now shifted slightly upwards at low masses based on the low-redshift calibration with SDSS DR7 (Figure~\ref{fig:newmex}). This feature increases the number of low-mass BPT-AGNs that lie in the MEx-SF region at the top left of the star-forming branch and slightly worsens the agreement. On the other hand, N14 noted that the original MEx demarcations should be displaced to higher stellar masses in order to improve the agreement between the BPT and MEx diagrams at $z\sim2$. In this case, our revised MEx dividing curves improve the agreement and mitigate the need for such a large shift. Instead, the revised MEx demarcation lines introduced here appear to be applicable out to $z\sim2$ given the current observational constraints, and given the limitations of the uncertain BPT classes for galaxies with measurement errors spanning the AGN/SF demarcation, which were not taken into account by N14. More recently, \citet{hen13} also suggested that the MEx demarcation should shift to higher stellar masses at higher redshifts, and suggested a 1~dex shift for their galaxy sample at $1.3<z<2.3$. The empirical approach presented in this Paper lies between these two conclusions, i.e., it includes a line luminosity dependence that mimic high-redshift galaxies on line ratio diagnostics, and accounts for $MZ$ relation evolution (Section~\ref{sec:MZ}) but the mass offsets on the MEx doagnostic diagram are generally not as extreme as those suggested by N14 and \citet{hen13}. A direct consequence of using the revised MEx diagnostic at higher redshifts will be to yield slightly lower AGN fractions compared to the original J11 version. This was noted by \citet{mig13} in their work comparing AGNs identified from \nevlam\ and X-rays in zCOSMOS, with the MEx diagnostic among others. The MEx demarcation lines can be further tested in future investigations by using the publically available MEx probability calculation code\footnote{https://sites.google.com/site/agndiagnostics/home/mex}, with optional emission line detection limits tailored to each survey, and with or without prescriptions for $L^*$ evolution to the redshift of each individual galaxy. It will also be informative to push the analysis to yet higher redshift, with new samples becoming available \citep[e.g.][]{hol14}, which will be a follow-up work to this article. In Appendix~\ref{app:stac}, we present a example application to intermediate-redshift ($0.2<z<0.8$) galaxy samples from the stacked spectral analysis of \citet{vit13}. We have also implemented alternative selections based on a single emission line, \ha\ or \oiii, in the public distribution of the code (Appendix~\ref{app:altern}). | 14 | 3 | 1403.6832 |
1403 | 1403.1589.txt | We use the moving mesh code \arepo coupled to a time-dependent chemical network to investigate molecular gas in simulated spiral galaxies that is not traced by CO emission. We calculate H$_{2}$ and CO column densities, and estimate the CO emission and CO-H$_2$ conversion factor. We find that in conditions akin to those in the local interstellar medium, around 42\% of the total molecular mass should be in CO-dark regions, in reasonable agreement with observational estimates. This fraction is almost insensitive to the CO integrated intensity threshold used to discriminate between CO-bright and CO-dark gas. The CO-dark molecular gas primarily resides in extremely long ($>100$ pc) filaments that are stretched between spiral arms by galactic shear. Only the centres of these filaments are bright in CO, suggesting that filamentary molecular clouds observed in the Milky Way may only be small parts of much larger structures. The CO-dark molecular gas mainly exists in a partially molecular phase which accounts for a significant fraction of the total disc mass budget. The dark gas fraction is higher in simulations with higher ambient UV fields or lower surface densities, implying that external galaxies with these conditions might have a greater proportion of dark gas. | \label{intro} The nature and distribution of molecular gas in the interstellar medium (ISM) is of great interest to astrophysicists because all local star formation occurs within clouds of molecular gas. Unfortunately, the main constituent of this gas, molecular hydrogen (H$_{2}$), is extremely difficult to observe directly, owing to the large energy separation between its rotational levels, as well as its absence of a dipole moment. Most studies of H$_{2}$ in the ISM therefore involve observing a tracer species, most commonly carbon monoxide (CO), and then inferring the properties of the molecular gas as a whole from the behaviour of this tracer. A crucial assumption here is that the tracer does indeed trace all of the H$_{2}$. However, there is growing observational evidence that this is not the case for CO, and that there is a component of the dense ISM that is rich in H$_{2}$ but that produces little or no CO emission \citep{Tielens85,vanDishoeck88}. Evidence for the existence of this component, termed ``dark molecular gas'' by \citet{Wolfire10}, comes from several different sources. Gamma-ray observations of the local ISM probe the total hydrogen column density, regardless of whether the hydrogen is in the form of H$^{+}$, H or H$_{2}$. Comparisons between maps of the diffuse gamma-ray emission and maps of the H{\sc i} 21~cm emission and CO emission show that there is a considerable amount of hydrogen surrounding molecular clouds that is not well traced by either form of emission, and that has therefore been interpreted as being CO-dark molecular gas \citep{Grenier05,Abdo10}. Measurements of total gas column densities using dust extinction or dust emission allow a similar comparison to be made and also show evidence for a significant fraction of CO-dark molecular material \citep[see e.g.][]{RomanDuval10,Ade11,Paradis12}. Further corroboration comes from \citet{Allen2012}, who find bright OH emission that is uncorrelated with atomic hydrogen or CO emission, and from \citet{PinedaJ13}, who show that a significant CO-dark molecular gas component is needed in order to explain the distribution of [C{\sc ii}] emission in the Galaxy. Finally, this ``dark'' molecular component can be probed directly by UV absorption line measurements \citep[see e.g.][]{Burgh07,Sonnentrucker07,Sheffer08}, albeit only along lines-of-sight where there happens to be a UV-bright background source. These UV absorption line studies confirm that significant H$_{2}$ can be present in regions without much CO, but do not allow this gas to be mapped, or its total mass to be determined. Numerical simulations can provide valuable assistance for interpreting this observational data, and allow us to better understand how ``dark'' H$_{2}$ is distributed in the ISM. For example, \citet{Wolfire10} use photodissociation region (PDR) models of a series of clouds to quantify the amount of dark H$_{2}$ that one would expect to find in their envelopes, finding values of around 30--40\% of the total H$_{2}$ content of the clouds, with little dependence on cloud properties other than metallicity and mean extinction. Although illuminating, this kind of 1D PDR model must inevitably make use of a highly approximate treatment of cloud structure. Moreover, the \citet{Wolfire10} models also assume chemical steady-state, and hence do not allow one to examine the effects of the dynamical history of the gas on its H$_{2}$ content, even though 3D dynamical models suggest that this may be highly significant \citep{Glover07a,Glover07b,Dobbs08a,Glover10}. Ideally, what one would like to do is to model the dynamical evolution of the ISM together with its chemical evolution in a self-consistent fashion, in order to better quantify how much dark molecular gas may be found around CO-bright molecular clouds, and also how much is distributed in the form of diffuse clouds with high H$_{2}$ fractions but little or no CO. However, to do this requires one to model both the H$_{2}$ and CO chemistry within a high dynamical range 3D simulation of the evolution of the ISM. Until very recently, this has not been possible. A number of groups have performed large-scale dynamical simulations of the ISM that account for H$_{2}$ formation \citep[see e.g.][]{Dobbs08a,Gnedin09,Christensen12}, but these models typically have insufficient resolution to allow one to study individual GMCs in detail, and they are also missing any treatment of the CO chemistry. On the other hand, small-scale models have for some time been able to model both H$_{2}$ and CO formation in the turbulent ISM \citep[see e.g.][]{Glover10,Glover12b}, but have not had the dynamical range necessary to model cloud formation self-consistently, or have started from highly artificial initial conditions \citep{Clark12}. In the present paper, we present initial results from a study that attempts to combine a large-scale 3D model of the dynamically evolving ISM with a detailed treatment of the small-scale cooling physics and chemistry. Our goals in this paper are to quantify how much CO-dark gas one would expect to find in the ISM, to explore how the fraction of molecular gas that is CO-dark varies as we change the mean surface density of the galaxy or the strength of its interstellar radiation field, and to better understand the spatial distribution of the CO-dark molecular gas. In Section \ref{methods}, we introduce our numerical methods and discuss the initial conditions that we use for our simulations. In Section \ref{MW}, we analyse in depth the dark gas fraction, morphology, molecular abundance and estimated CO emission for our fiducial case based on the Milky Way galaxy. In Section \ref{comparison}, we then compare to three simulations with different mean surface densities and ambient radiations fields. Finally, in Section \ref{discussion} we discuss our results, and in Section \ref{conclusions} we give our conclusions. | \label{conclusions} We have performed high-resolution simulations of a significant portion of a galactic disc with a version of the \arepo moving mesh code that includes a self-consistent treatment of the chemistry of the ISM in order to investigate the fractional amount and morphology of CO-dark molecular gas. In the most highly refined portion of our simulation we reach a mass resolution of only four solar masses. To the best of our knowledge, this is the highest resolution treatment to date of the ISM on galactic scales that has followed the chemical evolution of the gas. This approach enables us to resolve substructure within the disc sufficiently accurately to match the observed H{\sc i}--H$_2$ transition without having to resort to an ad-hoc ``clumping factor'' to represent unresolved density fluctuations. We considered four different models: a fiducial `Milky Way' model designed to be a reasonable match to the properties of our own Galaxy in the solar neighbourhood, plus three comparison simulations in which we explored the effects of varying the mean surface density of the gas and the strength of the interstellar radiation field. We evolved the simulations for 261.1 Myr (corresponding to six spiral arm passages) and then calculated the amount and morphology of CO-bright and CO-dark molecular gas. Our major findings are outlined below: \begin{enumerate} \item In conditions typical of the Milky Way disc, 46\% of the molecular gas resides below CO column densities of $10^{16}$ \cms. In terms of emission, 42\% of the molecular gas is found in regions with integrated intensities in the $J = 1 \rightarrow 0$ line of $^{12}$CO that are less than \WCO = 0.1~K~kms$^{-1}$. If we take this integrated intensity to mark the boundary between CO-bright and CO-dark gas, then we derive a dark gas fraction $f_{\rm DG} = 0.42$. Changing the value of our intensity threshold by factors of a few changes this number by only a few percent. \item Our predicted dark gas fraction of $f_{\rm DG} = 0.42$ is in excellent agreement with estimates derived from gamma rays, slightly lower than estimates from dust emission and absorption, and is slightly higher than estimates derived from [C{\sc ii}] emission. Overall our estimate of the CO-dark gas fraction is consistent with the observational literature. \item Although some CO-dark molecular gas is found in the vicinity of CO-bright clouds located in spiral arms, there is also a substantial amount in inter-arm regions in the form of extremely long filamentary clouds, with lengths of hundreds of parsecs. This filamentary geometry makes the CO more susceptible to dissociation due to the increased surface-to-volume ratio making it more difficult to shield. Observations have shown that filamentary molecular clouds are a feature of the Milky Way disc, but our simulations suggest that these might be just the observable parts of much larger structures. \item A non-negligible fraction of the total disc mass exists in a mixed phase that is neither fully atomic nor fully molecular. In the Milky Way simulation we find that 25\% of the gas has an H$_2$ abundance relative to the total number density of hydrogen nuclei of between 0.1 and 0.4 (where a value of 0.5 corresponds to fully molecular gas). This partially molecular phase is typically at lower column densities than the fully molecular gas and therefore is more likely to be CO-dark. However the majority of the H$_2$ mass still resides in regions where the fractional abundance is greater than 0.45. \item In simulations with higher radiation fields or lower mean surface densities, the dark gas fraction is much higher. However, the dark gas is still located primarily in the inter-arm regions of the disc. Higher radiation fields reduce the total amount of molecular gas and truncate the length of the filaments. Changes in the surface density of the disc have a much larger effect on $f_{\rm DG}$ than changes in the radiation field. These results suggest that external galaxies with lower surface densities or more intense radiation fields may have a greater dark gas fraction than the Milky Way. \item The globally-averaged CO-to-H$_2$ conversion factor that we measure in the CO visible gas in the Milky Way simulation, $X_{\rm CO} = 2.2\E^{20}$ cm$^{-2}$K$^{-1}$km$^{-1}$s, is in good agreement with the value inferred for the actual Milky Way. \XCO increases when the column density of the disc is reduced or the radiation field is increased, due to lower H$_2$ columns and brighter CO emission respectively. In all cases the value of \XCO rises if the diffuse gas is included in the estimate. \end{enumerate} | 14 | 3 | 1403.1589 |
1403 | 1403.4230_arXiv.txt | { Using new spectroscopic observations obtained as part of the VIMOS Ultra-Deep Survey (VUDS), we performed a systematic search for overdense environments in the early universe ($z>2$) and report here on the discovery of Cl J0227-0421, a massive protocluster at $z=3.29$. This protocluster is characterized by both the large overdensity of spectroscopically confirmed members, $\delta_{gal}=10.5\pm2.8$, and a significant overdensity in photometric redshift members. The halo mass of this protocluster is estimated by a variety of methods to be $\sim3\times10^{14}$ $\mathcal{M}_{\odot}$ at $z\sim3.3$, which, evolved to $z=0$ results in a halo mass rivaling or exceeding that of the Coma cluster. The properties of 19 spectroscopically confirmed member galaxies are compared with a large sample of VUDS/VVDS galaxies in lower density field environments at similar redshifts. We find tentative evidence for an excess of redder, brighter, and more massive galaxies within the confines of the protocluster relative to the field population, which suggests that we may be observing the beginning of environmentally induced quenching. The properties of these galaxies are investigated, including a discussion of the brightest protocluster galaxy, which appears to be undergoing vigorous coeval nuclear and starburst activity. The remaining member galaxies appear to have characteristics that are largely similar to the field population. Though we find weaker evidence of the suppression of the median star formation rates among and differences in the stacked spectra of member galaxies with respect to the field, we defer any conclusions about these trends to future work with the ensemble of protostructures that are found in the full VUDS sample.} | Large associations of galaxies provide an excellent laboratory for investigating astrophysical phenomena. The most massive of these associations, galaxy clusters and superclusters (i.e., clusters of clusters), while rare, are useful not only to constrain the dynamics and content of the universe (e.g., Bahcall et al.\ 2003; Reichardt et al.\ 2013), but also to study the evolution of galaxies, since the core of galaxy clusters are the regions of the universe where galaxy maturation occurs most rapidly (e.g., Dressler et al.\ 1984; Postman et al.\ 2005). This rapid maturation is a result of the large number of transformative mechanisms that a cluster galaxy experiences, mechanisms that are less effective or non-existent in regions of typical density in the universe (e.g., Moran et al.\ 2007). The number of processes a cluster galaxy is subject to is, however, both a virtue and a complication for studying their evolution. While the signs of transformation and evolution are prevalent among galaxies in clusters that have not already depleted their galaxies of gas, the large number of physical processes that are effective in overlapping regimes complicates interpretation. Furthermore, the effectiveness of such mechanisms appears to have a complex relationship with the halo mass of the host cluster and the dynamics of the galaxies that comprise it, the density and temperature of the intracluster medium (ICM), local galaxy density, mass of the individual galaxies, and cosmic epoch (e.g., Fujita \& Nagashima 1999; Poggianti et al.\ 2010; Lemaux et al.\ 2012; Muzzin et al.\ 2012; Dressler et al.\ 2013). The lower mass counterparts to galaxy clusters, galaxy groups, also suffer the same ambiguities. As such, despite nearly a century of study into such associations, the role that environment plays in galaxy evolution and the dominant process or processes that serve to transform cluster or group galaxies is still unclear. In the local universe, the relationship between environment and galaxy evolution has been revolutionized over the past decade with the advent of the Sloan Digital Sky Survey (SDSS). Observations from this survey have been used to great effect to study the properties of both groups and clusters and their galaxy content (e.g, G{\'o}mez et al.\ 2003; Hansen et al.\ 2009; von der Linden et al.\ 2010) and have led to insight into the nature of environmentally driven evolution in the local universe. However, these studies alone provide only a baseline for studies of cluster and group galaxies in the higher redshift universe because, in general, the galaxies populating structures in the low-redshift universe have come to the end of their evolution. Initial investigations of clusters beyond the local universe found that the fraction of galaxies that displayed a significant gas content, bluer colors, and late-type (i.e., spiral) morphologies increased rapidly with decreasing cosmic epoch (Butcher \& Oemler 1984). Yet, thirty years later, the cause or causes of such a trend have not been identified definitively. In intermediate-density environments, such as galaxy groups and pairs or small associations of galaxies, significant progress has been made in the past decade to understand the relative effect of such processes on galaxy evolution due to the emergence of spectroscopic surveys covering large portions of the sky in the intermediate-redshift universe ($z\sim1$, e.g., DEEP2, VVDS, zCOSMOS). While such surveys are typically devoid of massive clusters, a testament to their relative scarcity, the large number of spectroscopic redshifts, wide field coverage, and quality of both spectroscopic data and associated ancillary data have led to a variety of insights into the nature of galaxy evolution in intermediate-density environments (e.g., Cooper at al.\ 2006, 2007, 2008; Cucciati et al.\ 2006, 2010a, 2010b, 2012, Tasca et al.\ 2009; Peng et al.\ 2010; Presotto et al.\ 2012; George et al.\ 2012; Knobel et al.\ 2013; Kova{\v c} et al.\ 2014). At similar redshifts, systematic spectroscopic studies of clusters and cluster galaxies are somewhat rare. Surveys of clusters extending to several times the virial radius at $z\sim$0.5 (e.g., Treu et al.\ 2003; Dressler et al.\ 2004; Poggianti et al.\ 2006; Ma et al.\ 2008, 2010; Oemler et al.\ 2009, 2013) and of massive groups and clusters at $z\sim1$ (e.g., Lubin et al.\ 2009; Jeltema et al.\ 2009; Balogh et al.\ 2011; Muzzin et al.\ 2012; Hou et al.\ 2013; Mok et al.\ 2013, 2014) have begun to provide a somewhat coherent picture at these redshifts in which galaxy evolution has a complicated dependence on secular (i.e., mass-related) processes, as well as on both the global and the local environment. However, even at such redshifts, the effect of residing in the harsh cluster environment for several Gyr is evident among member galaxies, because the fraction of both red and quiescent galaxies is observed to be in excess of that of the field at similar redshifts (e.g., Patel et al.\ 2011; Lemaux et al.\ 2012; van der Burg et al.\ 2013). Going to higher redshifts, the effect of the environment should be reversed, inducing rather than suppressing star formation as gas-rich galaxies coalesce in the primeval universe. Indeed, tentative evidence for the reversal of the correlation between star formation rate (SFR) and galaxy density has already been found at slightly higher redshifts (Tran et al.\ 2010; Santos et al.\ 2014, though see also Santos et al.\ 2013; Ziparo et al.\ 2014). Observing the reversal of the SFR$-$density relation, as well as contextualizing the massive, red-sequence galaxies (RSGs) observed at $z\sim1$ in cluster and group environments, has motivated recent searches for high-redshift ($z\ga1.5$) clusters (e.g., Henry et al.\ 2010; Gobat et al.\ 2011; Papovich et al.\ 2010; Stanford et al.\ 2012; Zeimann et al.\ 2012; Newman et al. 2014) or other overdensities (i.e., protoclusters or protostructures) in the early universe (e.g., Steidel et al.\ 2005; Doherty et al.\ 2010; Toshikawa et al.\ 2012; Hayashi et al.\ 2012; Koyama et al.\ 2013; Hodge et al.\ 2013). One of the main difficulties in such searches, beyond the extreme faintness of the bulk of the member populations of such structures, is the failure of search techniques that are widely used at lower redshifts. Traditional techniques, such as searching for overdensities of RSGs or the presence of a hot ICM, are predicated on the assumption that a sufficiently long time scale has elapsed over which cluster galaxies can be processed. While these techniques can be used to find the most massive and oldest structures at any given redshift, such searches are biased against exactly the types of structures where the reversal of the SFR$-$density relation should be most apparent. One way of circumventing this bias is to search for overdensities of galaxies lying at the same redshift as estimated by broadband photometry (i.e., photometric redshifts), which have now largely supplanted searches for high-redshift overdensities of red galaxies. However, the nature of such overdensities cannot be be characterized well without dedicated spectroscopic followup. An alternative technique, which is employed especially for searches of the high-redshift universe, is to perform narrow-band imaging or photometric redshift searches around massive radio-loud quasars or other types of powerful active galactic nuclei (e.g., Kurk et al.\ 2004; Miley et al.\ 2004; Venemans et al.\ 2004, 2005; Zheng et al.\ 2006; Overzier et al.\ 2008; Kuiper et al.\ 2010, 2011, 2012). Such phenomena are typically associated with massive galaxies, which are, in turn, typically associated with galaxy overdensities. While this technique has been successful in observing large numbers of structures or protostructures in the high-redshift universe, it is not at all clear whether such environments are typical progenitors of lower redshift clusters or are exceptional in some way, which limits their usefulness in contextualizing results at lower redshifts. Additionally, narrow-band and spectroscopic searches of Lyman alpha emitter (LAEs) populations in (somewhat) random regions of the sky have revealed protostructures in the very high-redshift universe (e.g., Shimasaku et al.\ 2003; Ouchi et al.\ 2005; Lemaux et al.\ 2009; Toshikawa et al.\ 2012). However, such surveys cover rather limited portions of the sky and are only effective at observing overdensities of emission line objects, a population that, while being readily observed at high redshift because of the \emph{relative} ease of obtaining redshifts of emission line objects, is the subdominant population in the early universe (see, e.g., Shapley et al.\ 2003). As such, the structures (or protostructures) found by such searches are wildly inhomogeneous (see the recent review in Chiang et al.\ 2013). This inhomogeneity, combined with a lack of large, comparable samples of galaxies at more moderate (i.e., field) densities at similar redshifts makes interpreting such structures difficult. Ideally then, one would require a spectroscopic census of galaxy populations residing in both high- and lower-density environments in the high-redshift universe, representative in some way of the overall galaxy population at those epochs. With such a census it should be possible to make distinctions between evolution due to environmental processes and those driving overall trends observed in galaxy populations as a function of redshift and to properly connect these galaxy populations to their lower redshift descendants. The recently undertaken VIMOS Ultra-Deep Survey (VUDS; Le F{\`e}vre et al.\ 2014), an enormous 640-hour spectroscopic campaign with the 8.2-m VLT at Cerro Paranal targeting galaxies over 1 $\Box^{\circ}$ in three fields at $z>2$, for the first time provides the possibility of undertaking such a search at these redshifts. Like its predecessors at lower redshift, the fields targeted in the VUDS survey are random, albeit well-known, patches of the sky. As mentioned earlier, owing to the magnitude limited nature of field surveys (e.g., AEGIS, Davis et al.\ 2007; Newman et al.\ 2013; VVDS, Le F{\`e}vre et al.\ 2005, 2013; zCOSMOS, Lilly et al.\ 2007, 2009), the scarcity of red galaxies relative to bluer galaxies, and the rarity of massive clusters, environmental studies in field surveys, like VUDS, typically suffer the problem of limited dynamic ranges in local densities. Indeed, despite extensive spectroscopy from various surveys in the COSMOS (Scoville et al.\ 2007), CFHTLS-D1, E-CDF-S (Lehmer et al.\ 2005) fields, the three fields targeted by VUDS, only a few massive spectroscopically confirmed clusters have been found in these fields at $z<1.5$ (Gilli et al.\ 2003; Valtchanov et al. 2004; Guzzo et al.\ 2007; Silverman et al.\ 2008). However, there are several distinct differences between these surveys and VUDS in the way that they relate to a study of the effect of environment on galaxy evolution due to the nature of galaxies being probed. LAEs and other star-forming galaxies at high redshift, both of which are selected in VUDS by virtue of a photometric redshift selection, are known to be highly clumpy populations (e.g., Miyazaki et al.\ 2003; Ouchi et al.\ 2003, 2004, 2005; Lee et al.\ 2006; Bielby et al.\ 2011; Jose et al.\ 2013), making it possible to observe a wide dynamic range of local densities. In the high-redshift universe, protostructures comprised of such populations are observed (e.g., Steidel et al.\ 1998; Ouchi et al.\ 2005; Capak et al.\ 2011; Tashikawa et al.\ 2012; Chiang et al.\ 2014) and found in simulations (e.g., Chiang et al.\ 2013; Zemp 2013; Shattow et al.\ 2013) to be large in transverse extent. This large extent on the sky allows for sampling a larger number of members in a single VIMOS pointing than in traditional multi-object spectroscopic surveys of lower-redshift overdense environments. In addition, as a result of a photometric redshift selection, galaxies that have more distinguishing features in their SED, i.e., both a continuum break at $\sim$4000 \AA\ and the typical continuum break observed at the Lyman limit and Lyman$\alpha$, will be more likely to be assigned a accurate photometric redshift and are thus more likely to be targeted. Such a sample will be comprised of a mix of quiescent, post-starburst, and starburst populations. These populations are instrumental in the investigation the effect of environment on galaxy evolution. With this in mind, we performed a systematic search for overdensities of galaxies with secure spectroscopic redshifts in all three VUDS fields. The full results of this search will be published in a future work. In this paper, we focus on the discovery and study of the most significantly detected \emph{spectroscopic} overdensity in the CFHTLS-D1 field, Cl J0227-0421, a massive forming cluster at $z\sim3.3$. The structure of the paper is as follows. \S\ref{obsnred} provides an overview of the spectroscopic and imaging data available in the CFHTLS-D1 field, as well as the derivation of physical parameters of galaxies in our sample, with particular attention paid to new observations from the VUDS survey. \S\ref{analysis} describes the search methodology employed and the subsequent discovery of Cl J0227-0421, along with the estimation of its global properties. In \S\ref{memprop} we describe the investigation of the properties of the spectroscopically confirmed members of Cl J0227-0421 and compare those properties to galaxies in lower-density environments. Finally, \S\ref{conclusions} presents a summary of our results. Throughout this paper all magnitudes, including those in the IR, are presented in the AB system (Oke \& Gunn 1983; Fukugita et al.\ 1996). We adopt a standard concordance $\Lambda$CDM cosmology with $H_{0}$ = 70 km s$^{-1}$, $\Omega_{\Lambda}$ = 0.73, and $\Omega_{M}$ = 0.27. | \label{conclusions} In this paper we have described a systematic search for overdensities at high redshift ($z>2$) in the CFHTLS-D1 field using newly obtained VUDS spectroscopic data in conjunction with the wealth of other imaging and spectroscopic data available for this field. We then described the discovery and characterization of the most significant of these overdensities, the Cl J0227-0421 protocluster at $z\sim3.3$. Here we briefly outline the main conclusions of this study. \begin{itemize} \renewcommand{\labelitemi}{$\bullet$} \item With 19 confirmed spectroscopic members and six potential spectroscopic members, Cl J0227-0421 is significantly overdense relative to the field at these redshifts. Using a large field coeval population from VUDS and VVDS along with the 19 confirmed spectroscopic members, we estimated the significance of the spectroscopic overdensity of Cl J0227-0421 to be $\sigma=13.5$ (or $\delta_{gal}=10.5\pm2.8$). After accounting for spurious peaks, we found that Cl J0227-0421 is also overdense in its photometric redshift members, with a significance of $\sigma=8.0$. \item Four different methods were used to estimate or place limits on the halo mass of Cl J0227-0421 at $z\sim3.3$ or $z=0$ (or both). These were member galaxy dynamics, stellar-to-halo mass, X-ray hydrostatic equilibrium, and spectroscopic galaxy overdensity. Though the errors, uncertainties, and number assumptions used for each method were large, a consistent picture emerged in which Cl J0227-0421 has already assembled a large amount of mass in the early universe ($\mathcal{M}_{z\sim3.3}\sim3\times10^{14}$ $M_{\odot}$) and will evolve into a cluster with a halo mass rivaling or exceeding that of Coma ($\mathcal{M}_{z=0}\sim4\times10^{15}$ $M_{\odot}$). \item The properties of the spectroscopic member galaxies of Cl J0227-0421 were investigated. In the brightest protocluster galaxy, we found evidence of a powerful active galactic nuclei, as well as tentative evidence of vigorous star formation activity ($\sim750$ $M_{\odot}$ yr$^{-1}$). Within the protocluster environment, a significant excess of brighter, redder, and more massive galaxies appeared relative to a similarly selected field population at similar redshifts. This excess was quantified both absolutely, $\delta_{pRSG}=25.1\pm15.2$, and relatively, with a fractional excess of such galaxies within the protocluster of around three. Based on comparisons with models, the last major star-formation event in these galaxies was estimated to be in excess of 300 Myr prior to $z\sim3.3$, indicating that we may be witnessing the onset of environmentally-driven quenching processes. \item The remaining protocluster members had properties that were broadly similar to those of field galaxies. While we found weak evidence of suppression of the star formation rates among the general protocluster member population and subtle differences between the stacked spectra of the two populations, these differences were not significant enough to be conclusive. \end{itemize} Despite the massive nature of Cl J0227-0421, the relatively small number of protocluster members statistically limited the conclusions that could be drawn. Still, the results of several lines of analysis presented in this paper were tantalizingly suggestive of the effect of environment at $z\sim3.3$. These lines of analysis will be continued with the $\sim$40 overdensities found within the entire VUDS sample to search for definitive signs of environmentally-driven evolution and transformation in the high-redshift universe. | 14 | 3 | 1403.4230 |
1403 | 1403.1623_arXiv.txt | The scaling relation for early type galaxies in the 6dF galaxy survey does not have the velocity dispersion dependence expected from standard stellar population models. As noted in recent work with SDSS, there seems to be an additional dependence of mass to light ratio with velocity dispersion, possibly due to a bottom heavy initial mass function. Here we offer a new understanding of the 6dF galaxy survey 3D gaussian Fundamental Plane in terms of a parameterized Jeans equation, but leave mass dependence of M/L and mass dependence of structure still degenerate with just the present constraints. Hybrid models have been proposed recently. Our new analysis brings into focus promising lines of enquiry which could be pursued to lift this degeneracy, including stellar atmospheres computation, kinematic probes of ellipticals at large radius, and a large sample of one micron spectra. | Scaling relations for galaxies, like the Tully-Fisher relation for disks and the Faber-Jackson relation for ellipticals, are fundamental and powerful. They challenge theories of galaxy formation and they allow us to measure galaxy distances. Local galaxy distances provide us with maps of the mass distribution to compare with the light distribution. \cite{TF77} and \cite{FJ} pioneered scaling relations, and the latter was soon replaced with the fundamental plane \citep{L87}, or FP. The virial theorem offered a partial explanation, \eg \cite{AM79,F87}. Nearly two decades ago \cite{W96} opined that an understanding of scaling relations was within reach. But hydrodynamic models of galaxy formation and semi-analytic models have not led to a full understanding.% Nor has the comparison of mass maps and the galaxy distribution yet led to a satisfying resolution. On the one hand, the distribution of peculiar velocities in x-ray clusters of the 6dF galaxy survey (Magoulas 2012; PhD thesis\footnote{http://dtl.unimelb.edu.au} ) and of Planck's kinetic Sunyaev Zeldovich peculiar velocities \citep{A13} is comparable; on the other, the bulk flow measured locally \citep{F10, M12} is on the high side of expectations from the $\Lambda$CDM model. In this presentation we outline the scaling relation problem for early type galaxies from the perspective of the 6dF galaxy survey \citep{J05, LC}. Our findings parallel the SDSS result of \cite{C13}. We consider what this means for the elliptical scaling relation. And we suggest what needs to be done to test the notion that a varying bottom-heavy initial mass function (IMF) is what we are missing in ellipticals. | The ratio of dynamical mass to standard stellar population mass in the 6dF galaxy sample of early type galaxies is approximately proportional to velocity dispersion in the range 100 $<~\sigma~<$ 300 km s$^{-1}$. A bottom heavy IMF is a simpler explanation of this trend than the notion, for example, that the baryonic fraction of these galaxies has a peak at $\sim$ 200 km s$^{-1}$ velocity dispersion. However, a greater dark matter mass fraction in large halo potentials is an alternative hypothesis that cannot at present be ruled out. Hybrid models are very possible \citep{O13}. Stellar atmospheres \citep {A97, AF13}, stellar populations, globular cluster dynamics, and galaxy formation theory can all play valuable roles in tying down what we are missing in elliptical galaxies. The question posed in the title of this Letter deserves a simple answer. Our answer is that we are missing $\delta$ and $\epsilon$. Unambiguous determination of the level of late M dwarf light in ellipticals will measure $\epsilon$. Kinematic probes of the outer gravitational potentials of ellipticals will measure $\delta$, as neutral hydrogen did for disk galaxies back when scaling relations were first proposed. | 14 | 3 | 1403.1623 |
1403 | 1403.2370_arXiv.txt | Galaxy clusters can efficiently convert axion-like particles (ALPs) to photons. We propose that the recently claimed detection of a 3.55--3.57 keV line in the stacked spectra of a large number of galaxy clusters and the Andromeda galaxy may originate from the decay of either a scalar or fermionic $7.1$ keV dark matter species into an axion-like particle (ALP) of mass $m_{a} \lesssim 6\cdot 10^{-11}~{\rm eV}$, which subsequently converts to a photon in the cluster magnetic field. In contrast to models in which the photon line arises directly from dark matter decay or annihilation, this can explain the anomalous line strength in the Perseus cluster. As axion-photon conversion scales as $B^2$ and cool core clusters have high central magnetic fields, this model can also explain the observed peaking of the line emission in the cool cores of the Perseus, Ophiuchus and Centaurus clusters, as opposed to the much larger dark matter halos. We describe distinctive predictions of this scenario for future observations. | 14 | 3 | 1403.2370 |
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1403 | 1403.5766_arXiv.txt | {In this paper we show that the Germani-Kehagias model of Higgs inflation (or New Higgs Inflation), where the Higgs boson is kinetically non-minimally coupled to the Einstein tensor is in perfect compatibility with the latest Planck and BICEP2 data. Moreover, we show that the tension between the Planck and BICEP2 data can be relieved within the New Higgs inflation scenario by a negative running of the spectral index. Regarding the unitarity of the model, we argue that it is unitary throughout the evolution of the Universe. Weak couplings in the Higgs-Higgs and Higgs-graviton sectors are provided by a large background dependent cut-off scale during inflation. In the same regime, the W and Z gauge bosons acquire a very large mass, thus decouple. On the other hand, if they are also non-minimally coupled to the Higgs boson, their effective masses can be enormously reduced. In this case, the W and Z bosons are no longer decoupled. After inflation, the New Higgs model is well approximated by a quartic Galileon with a renormalizable potential. We argue that this can unitarily create the right conditions for inflation to eventually start.} \subheader{RESCEU-8/14\\LMU-ASC 13/14} \begin{document} | The discovery of polarized B-modes in the CMB by BICEP2 has completely changed our prospective of inflation since the release of the Planck results. With this new data, if inflation ever occurred, it must be chaotic-type, i.e. the excursion of the canonically normalized inflaton must be trans-Planckian. With the discovery of the Higgs boson at the LHC, it is very tempting to seriously consider the minimal scenario in which the Higgs boson is the inflaton. Here we showed that this would be compatible with Planck and BICEP2 data if the Higgs boson is non-minimally kinetically coupled to curvatures, as in the New Higgs Inflationary scenario of \cite{new}. In particular, we show that the mild tension between Planck and BICEP2 data can be released by a negative running of the New Higgs inflation spectral index. Finally, we have argued that our model is unitary throughout the whole evolution of the Universe. In particular, in the original model of Higgs inflation the gauge sector can be considered decoupled from the low energy effective theory. A non-minimal interaction of the gauge bosons to the Higgs and the double-dual Riemann tensor will, on the other hand, significantly lower the gauge boson masses during inflation. In this case, we showed that the gauge bosons can be treated within our effective field theory even during inflation. | 14 | 3 | 1403.5766 |
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1403 | 1403.7280_arXiv.txt | We report the detection of eight vibronic bands of C$_3$, seven of which have been hitherto unobserved in astrophysical objects, in the translucent cloud towards HD~169454. Four of these bands are also found towards two additional objects: HD~73882 and HD~154368. Very high signal-to-noise ratio ($\sim$1000 and higher) and high resolving power ($R=80,000$) UVES-VLT spectra (Paranal, Chile) allow for detecting novel spectral features of C$_3$, even revealing weak perturbed features in the strongest bands. The work presented here provides the most complete spectroscopic survey of the so far largest carbon chain detected in translucent interstellar clouds. High-quality laboratory spectra of C$_3$ are measured using cavity ring-down absorption spectroscopy in a supersonically expanding hydrocarbon plasma, to support the analysis of the identified bands towards HD~169454. A column density of N(C$_3$) = $(6.6 \pm 0.2) \times 10^{12}$ cm$^{-2}$ is inferred and the excitation of the molecule exhibits two temperature components; $T_{exc}= 22 \pm 1$ K for the low-$J$ states and $T_{exc}= 187 \pm 25$ K for the high-$J$ tail. The rotational excitation of C$_3$ is reasonably well explained by models involving a mechanism including inelastic collisions, formation and destruction of the molecule, and radiative pumping in the far-infrared. These models yield gas kinetic temperatures comparable to those found for $T_{exc}$. The assignment of spectral features in the UV-blue range 3793-4054 \AA\ may be of relevance for future studies aiming at unravelling spectra to identify interstellar molecules associated with the diffuse interstellar bands (DIBs). | Currently, some 180 different molecules have been detected in dense inter- and circumstellar clouds, largely in radio- and submilimeter surveys, with this number growing with several new species per year. However, only about ten simple molecules are observed in the visible part of the electromagnetic spectrum as absorption features originating in translucent clouds, transparent for optical wavelengths. Among them are homonuclear species, such as H$_2$, C$_2$ and C$_3$, which are not accessible to radio observations. Bare carbon chains do not exhibit pure rotational transitions, because of the lack of a permanent dipole moment, and thus only their electronic or vibrational spectral features can be observed. The latter cover the spectral range from the vacuum UV until the far infrared. Determination of the abundances of simple carbon molecules in interstellar clouds is important, as they are considered building blocks for many already known interstellar molecules with a carbon skeleton. After the $19^{th}$ century discovery of the 4052 \AA\ band in the spectrum of comet Tebbutt \citep{Huggins1881} and the assignment of this blue absorption feature to the C$_3$ molecule \citep{Douglas1951}, the triatomic carbon chain radical was detected in the circumstellar shell of the star IRC+10216 \citep{Hinkle1988} and subsequently, albeit tentatively, in the interstellar medium \citep{Haffner1995}. \citet{Cernicharo2000} detected nine lines of the $\nu_2$ bending mode towards Sgr B2 and IRC+10216, observed in the laboratory by~\citet{Giesen2001}. This spectral range is, however, not applicable to observations of translucent clouds. High resolution observations using the Herschel/HIFI instrument detected transitions of C$_{3}$ originating in the warm envelopes of massive star forming regions~\citep{Mookerjea2010, Mookerjea2012}; in these environments high densities of $10^5-10^6$ cm$^{-3}$ prevail, whereas in diffuse clouds densities are limited to $10^3$ cm$^{-3}$. The presence of C$_3$ in the diffuse interstellar medium was proven by \citet{Maier2001}, based on the detection of the \~{A}$^{1}\Pi_{u}$--\~{X}$^{1}\Sigma_{g}^{+}$ 000--000 band close to 4052~\AA\ towards three reddened stars. Up to now the highest resolution study of C$_3$ in such environments was reported by~\citet{Galazutdinov2002b}, although restricted to a few objects only. These observations provided the only reliable estimates so far of the abundance of this molecule, yielding values an order of magnitude below that of C$_2$. Other observations of C$_3$ in translucent clouds~\citep{Roueff2002, Adamkovics2003, Oka2003} suffered from lower signal-to-noise ratios, giving rise to large uncertainties in the deduced column densities. The existing data for linear carbon molecules longer than C$_3$, like C$_4$ \citep{Linnartz2000}, or C$_5$ \citep{Motylewski1999b}, fail to provide firm evidence for their existence in the interstellar medium \citep{Galazutdinov2002a, Maier2002, Maier2004}. For a systematic investigation of the conditions under which carbon-based molecules are produced in the interstellar medium, a larger class of targets exhibiting a variety of optical properties and physical conditions in the intervening clouds has to be observed. The targeted objects should be selected for translucent clouds with the carbon-bearing molecules producing a single Doppler-velocity component, thus allowing for an unambiguous analysis of the spectrum, and resulting in accurate column densities. The identification of the carriers of diffuse interstellar bands (DIBs) remains, since their discovery by \citet{Heger1922}, one of the persistently unresolved problems in spectroscopy. The current list of unidentified interstellar absorption features contains more than 400 entries \citep{Hobbs2008}. The presence of substructures inside DIB profiles, discovered by \citet{Sarre1995} and by \citet{Kerr1998}, supports the hypothesis of their molecular origin. The established relation between profile widths of DIBs and rotational temperatures of linear carbon molecules \citep{Kazmierczak2010} makes the latter interesting targets for observations. Both C$_2$ and C$_3$ may show different rotational temperatures along different lines of sight, as shown by \citet{Adamkovics2003}. This is associated with the fact that their rotational transitions are forbidden and thus cooling of their internal degrees of freedom is inefficient. For this reason, an accurate determination of rotational excitation temperatures of short carbon chains may help to shed light on the origin of the mysterious carriers of the DIBs \citep{Kazmierczak2010b}. Observationally, the spectral features originating from either C$_2$ or C$_3$ typically turn out to be rather shallow and thus high S/N ratios and high spectral resolution are required to establish accurate values for the excitation temperature. The aim of the present investigation is to use the superior capabilities of UVES-VLT to obtain high quality spectra of C$_3$. All previous studies were based solely on the strongest 000--000 band of the \~{A}$^{1}\Pi_{u}$--\~{X}$^{1}\Sigma_{g}^{+}$ electronic system. Here additional vibronic bands in the \~{A}$^{1}\Pi_{u}$--\~{X}$^{1}\Sigma_{g}^{+}$ electronic system of C$_3$ are identified along sight lines towards objects {HD~169454}, {HD~73882}, and {HD~154368}. A detailed analysis of eight vibronic bands detected towards the object {HD~169454} is presented. The astronomical observations are supported by a high-quality laboratory investigation, using cavity ring-down laser spectroscopy, producing fully rotationally resolved C$_3$ spectra of the vibronic bands in the \~{A}--\~{X} system. The combined information of laboratory and observed spectra is used to deduce column densities and a rotational temperature of C$_3$ in {HD~169454} The results are interpreted in terms of excitation models for C$_3$ \citep{Roueff2002} and a chemical model of a translucent cloud towards HD~169454. | We report on a high-quality absorption spectrum, in terms of both resolution and signal-to-noise ratio, of the C$_3$ molecule towards HD~169454. Besides a fully resolved spectrum of the \~{A} -- \~{X} 000 - 000 band at $\sim 4052$~\AA\ seven further vibronic C$_3$ bands were identified in the range 3793-4000 \AA. These absorptions had been suspected previously \citep{Gausset1965}, but now these bands are unambiguously detected for the first time with UVES-VLT. Four of those vibronic bands have been observed as well along the sight lines to two heavily reddened stars: HD~154368 and HD~73882. The observations are supported by laboratory measurements of all eight bands under high resolution using cavity ring-down laser spectroscopy. The accurate value for the C$_3$ column density in the translucent cloud towards HD~169454 was included in an excitation model, applying the Meudon PDR-code by \citet{Roueff2002}, yielding good agreement for the column densities of rotational levels. The calculations do not allow to determine gas density and destruction rate uniquely. As far as the rate coefficients and photodissociation rates are established, the model depends on the ratio of gas density to intensity of radiation field. When modifying the model with recently updated collisional rates~\citep{Abdallah2008} the model results in high destruction rates required for C$_{3}$, inconsistent with the present understanding of the destruction process. Observation of carbon chain molecules in optical spectra of diffuse clouds was up to now unsuccessful, except for C$_3$. All complex organic molecules identified in the interstellar space have been found in dense protostellar clouds. The detection and identification of a number of very weak absorptions of the C$_3$ molecule raises hope that the forest of narrow absorptions in the blue-violet part of the spectrum towards reddened stars may uncover assignments of heavier carbon species. Many weak as yet unidentified spectral features appear to be molecular lines, rather than noise. The regular pattern of C$_3$ bands allows for an unambiguous identification, as shown in the present study. Heavier species like C$_4$ and C$_5$ would produce more compact patterns due to the smaller rotational constant or due to a strong Q-branch compared to R- and P-branches characteristic for $\Sigma$--$\Pi$ transitions in case of C$_5$, particularly for low excitation temperature, which may lead to their detection. The assignment of absorption features in the UV-blue spectral range to the C$_3$ molecule is of relevance in the context of searches for carriers of the diffuse interstellar bands (DIBs). While most of the absorption features detected in translucent clouds in sight-lines toward reddened stars are ascribed to DIBs of unknown origin, the presently observed features can be excluded from the DIB-lists for which carriers are sought. | 14 | 3 | 1403.7280 |
1403 | 1403.2977_arXiv.txt | Astroparticle physics and cosmology allow us to scan the universe through multiple messengers. It is the combination of these probes that improves our understanding of the universe, both in its composition and its dynamics. Unlike other areas in science, research in astroparticle physics has a real originality in detection techniques, in infrastructure locations, and in the observed physical phenomenon that is not created directly by humans. It is these features that make the minimisation of statistical and systematic errors a perpetual challenge. In all these projects, the environment is turned into a detector medium or a target. The atmosphere is probably the environment component the most common in astroparticle physics and requires a continuous monitoring of its properties to minimise as much as possible the systematic uncertainties associated. This paper introduces the different atmospheric effects to take into account in astroparticle physics measurements and provides a non-exhaustive list of techniques and instruments to monitor the different elements composing the atmosphere. A discussion on the close link between astroparticle physics and Earth sciences ends this paper. | \label{intro} Recent years have seen the development of major infrastructure around the Earth in order to increase considerably the performances of experiments in astroparticle physics and cosmology. Unlike other fields in science where measurements are made on a physical phenomenon created in laboratory, research in astroparticle physics has originality in detection techniques and infrastructure locations. Experiments are operated over large desert areas as the Cherenkov Telescope Array (CTA)~\cite{CTA_web}, the Pierre Auger Observatory~\cite{PAO_web} or very soon the LSST telescope~\cite{LSST_web}, in oceans or ice with ANTARES~\cite{ANTARES} and IceCube~\cite{IceCube}, respectively, or even in space with projects as the AMS-02 experiment~\cite{AMS} or soon the JEM-EUSO telescope~\cite{JEM-EUSO}. In all these projects, and more than any other experience in subatomic physics, minimising statistical and systematic errors is a challenge because the physical phenomenon observed is not produced by man himself: scientists are just observers. Thus, scientists build ever larger detectors to go further in the knowledge. However, owning a large detector is not a necessary and sufficient requirement to push the limits of our knowledge: the systematic error still lurks and demands from scientists an excellent understanding of their detector. The temptation to increase the duty cycle of the detector in order to reduce still more the statistical error should not obscure the need to control the associated increase in systematic error. Therefore there is a point where these two errors become inseparable and where the optimisation of detector performance can be compared to the concept of \emph{yin} and \emph{yang}. In all these projects, the environment is turned into a detector medium or a target. The atmosphere is probably the environment component the most common in astroparticle physics, usually used as a giant calorimeter in cosmic ray experiments or as an irreducible detection volume in the case of ground-based astrophysics surveys. To minimise as much as possible the systematic errors associated to the atmosphere evolution in time, its properties have to be continuously monitored. It is to this end that extensive atmospheric monitoring programs have been developed by different collaborations in astroparticle physics. Section~\ref{sec:astropart_exps} will list briefly the different experiments where the atmosphere is a part of the detector. In all cases, at some point, photons propagate into the atmosphere and they are affected by the medium before being detected. Section~\ref{sec:atmo_effects} will describe the different physics phenomena affecting photon propagation in the atmosphere in order to remove their effect in measurements. Then, in Section~\ref{sec:atmo_facilities}, the main instruments used to monitor the atmospheric properties or the atmosphere components will be presented. Astroparticle physics experiments, equipped with such infrastructures and located in unusual places, provide an opportunity to develop interdisciplinary activities, especially in atmospheric science and geophysics: this will be the purpose of Section~\ref{sec:interdisciplinary}. | \label{sec:summary} Astroparticle physics experiments require still greater precision in measurements to answer questions in particle physics, astrophysics and cosmology. Some of these experiments use the atmosphere as a part of their detector. In order to reduce as much as possible systematic uncertainties related to the atmosphere, extensive atmospheric monitoring programs have been, are or will be developed by collaborations. It has been shown that all the astroparticle physics experiments are not at the same stage in atmospheric monitoring: whereas collaborations in ultra-high energy cosmic rays use already techniques developed in atmospheric sciences to probe atmospheric properties and correct its effect in their measurements, scientists in very-high energy gamma rays or ground-based astronomical surveys are still in a stage where atmospheric measurements are used only as a quality cut on data selection. However, this would not be true anymore in the next major projects where the challenge of environment monitoring will be a key element for developing instrumentation. In order to better carry out these future projects, it makes sense to collaborate with scientists in Earth sciences to choose the best methods and techniques to reach scientific goals. Concerning atmospheric measurements, it has been shown that many instruments and techniques developed in atmospheric sciences are available. Depending on the atmospheric component monitored -- molecular, aerosol or cloud -- same instruments will not be used. Before installing any instrument on site, it is necessary to know the effect of the chemical component planned to be measured in the wavelength range studied and its time and spatial variations. Since in the near future astroparticle physics experiments will require extensive atmospheric monitoring programs with many instruments, the idea to join a worldwide atmospheric network has to be considered. Indeed, experiments are sometimes situated in places with a few weather stations available and joining such networks could represent an opportunity for both research fields. | 14 | 3 | 1403.2977 |
1403 | 1403.7555_arXiv.txt | We analyzed several basic correlations between structural parameters of galaxies. The data were taken from various samples in different passbands which are available in the literature. We discuss disc scaling relations as well as some debatable issues concerning the so-called Photometric Plane for bulges and elliptical galaxies in different forms and various versions of the famous Kormendy relation. We show that some of the correlations under discussion are artificial (self-correlations), while others truly reveal some new essential details of the structural properties of galaxies. Our main results are as follows: (1) At present, we can not conclude that faint stellar discs are, on average, more thin than discs in high surface brightness galaxies. The ``central surface brightness -- thickness'' correlation appears only as a consequence of the transparent exponential disc model to describe real galaxy discs. (2) The Photometric Plane appears to have no independent physical sense. Various forms of this plane are merely sophisticated versions of the Kormendy relation or of the self-relation involving the central surface brightness of a bulge/elliptical galaxy and the \ser\ index $n$. (3) The Kormendy relation is a physical correlation presumably reflecting the difference in the origin of bright and faint ellipticals and bulges. We present arguments that involve creating artificial samples to prove our main idea. | 14 | 3 | 1403.7555 |
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1403 | 1403.6553_arXiv.txt | We present a model for the in situ assembly of planetary systems around a 0.5~$M_{\odot}$ star, and compare the resulting statistics to the observed sample of cool Kepler planet candidates from Dressing \& Charbonneau (2013). We are able to reproduce the distribution of planetary periods and period ratios, although we once again find an underabundance of single transit systems relative to the observations. We also demonstrate that almost every planetary system assembled in this fashion contains at least one planet in the habitable zone, and that water delivery to these planets can potentially produce a water content comparable to that of Earth. Our results broadly support the notion that habitable planets are plentiful around M~dwarfs in the solar neighbourhood. | The search for potentially habitable planets around other stars has advanced significantly in recent years. As the precision of radial velocity searches has improved, the detection limits have reached into the regime where planets are thought to be too small to be genuine gaseous planets. Furthermore, estimates of the incompleteness of such searches suggest that planetary systems composed of low mass planets may be significantly more common than those with Jovian class planets in short period orbits (Howard et al. 2010; Mayor et al. 2011). The launch of the Kepler satellite has also unearthed a substantial catalogue of planetary candidates with radii indicative of small planets (Borucki et al. 2011; Batalha et al. 2013), including several candidates at sufficiently long periods to potentially qualify as habitable (Borucki et al. 2013; Barclay et al. 2013). The origins of such plentiful systems of lower mass planets is still somewhat uncertain. There are a significant number of planets in the Neptune-size category, which were predicted to be rare in population synthesis models based on the concept of planetary migration (Ida \& Lin 2008; Mordasini et al. 2009). Furthermore, the period ratios in the multiple planet systems (Lissauer et al. 2011; Fabrycky et al. 2012) are broadly distributed, contrary to predictions of migration models (Alibert et al. 2006; Terquem \& Papaloizou 2007; Raymond et al. 2008; Mordasini et al. 2009; Ida \& Lin 2010), which anticipated a preference for commensurabilities. This has led to an alternative proposal that such systems assemble in situ (Hansen \& Murray 2012; 2013; Chiang \& Laughlin 2013), although adjustments to the migration scenario have also been proposed (Rein 2012; Goldreich \& Schlichting 2014). The bulk of the attention thus far has been focussed on solar-type stars, both for the obvious anthropocentric reasons and because the observational samples are dominated by such stars. Lower mass stars, M-dwarfs, are disfavoured by virtue of their lower luminosities, high activity levels, and redder spectra (which reduce radial velocity accuracy). On the other hand, M-dwarfs are smaller and therefore are attractive transit survey targets, as smaller planets obscure a larger fraction of the star than of a G-dwarf during transit. Furthermore, the lower stellar luminosity also implies that the habitable zone is closer to the star and so easier to study via transits because of the shorter orbital period and greater transit depths. As a result, several groups have recently begun to quantify the sample of transiting systems that orbit M-dwarfs in the Kepler data (Muirhead et al. 2012; Mann et al. 2012; Dressing \& Charbonneau 2013). Therefore, in this paper we explore the predictions of an in situ assembly model for M-dwarf planetary systems, adapting the model of Hansen \& Murray (2013)--hereafter HM13--to the environs of a 0.5~$M_{\odot}$ star. We compare these to the sample of transiting planet candidates defined by Dressing \& Charbonneau (2013)--hereafter DC13. In \S~\ref{Observe} we will quantify the observational sample and use this to estimate the initial conditions necessary for our theoretical model. In \S~\ref{Model} we then describe the construction of our theoretical model for in situ assembly and its comparison for the data. | M~dwarfs have a history of confounding pessimistic theory predictions regarding their planetary inventories. Laughlin et al. (2004) predicted that the mass and timescale requirements for giant planet formation via core accretion would make giant planets relatively rare around M~dwarfs. Although this is partially born out by the relative rarity of giant planets with short orbital periods around M dwarfs (Johnson et al. 2007), some have been discovered (Rivera et al. 2010; Haghighipour et al. 2010; Johnson et al. 2010) and microlensing surveys suggest that many M~dwarfs may indeed possess giant planets at larger distances (Gould et al. 2010). Similarly, Raymond et al. (2007) predicted that protoplanetary disks which scaled linearly with stellar mass would produce terrestrial planet systems with planets too small ($< 0.3 M_{\oplus}$) to be detected by Kepler in significant quantities, although Montgomery \& Laughlin (2009) demonstrated that more massive disks could indeed produce planetary systems of observable size via in situ accretion. Furthermore, planetary systems may also be constructed via a number of other evolutionary pathways beyond simple in situ accretion (Raymond et al. 2008, Ogihara \& Ida 2009), which could potentially alleviate the mass constraints. The Kepler mission has now revealed that planetary systems around cool stars follow a similar trend to those around hotter stars, in the sense that compact systems of low mass planets appear to be more numerous than those containing gas giants, at least at short orbital periods. Our goal in this paper has been to demonstrate that these low mass planetary systems are indeed broadly consistent with a population of planets that assembled in situ, suggesting a commonality with the planetary systems observed around sun-like stars (Hansen \& Murray 2012, HM13; Chiang \& Laughlin 2013). This is somewhat at odds with the conclusion reached by Swift et al. (2013), who favour inward migration of planets. However, that study relied on a detailed study of a single system (Kepler-32), allied to blithe generalities about the rest of the Kepler sample. The Kepler-32 system is one of the few observed systems that actually corresponds to the commensurable chain expected from migration models, and therefore cannot really be used as a representative of the sample as a whole (consider the distribution of period ratios shown in Figure~\ref{PratM}). Furthermore, the particular arrangement of the Kepler-32 system may simply be a statistical fluctuation. Figure~\ref{Zeta} shows the distribution of the four Kepler-32 period ratios when characterised by the statistic $\zeta_i$ introduced by Lissauer et al. (2011), to characterise the proximity to a resonance of order $i$. The open circles indicate the four Kepler-32 pairs, while the filled circles show the rest of the pairs in the DC13 sample. We see that the Kepler sample as a whole is spread out over the full observed range of $\zeta$, and could easily represent a generic sampling of the fuller distribution. Also shown is the distribution of $\zeta_1$ expected from our Monte-Carlo model, which reproduces the broad distribution seen in the observations. In short, in situ assembly reproduces the observed broad distribution of period ratios quite well, even with the commensurabilities of the Kepler-32 system included in the sample. Migration may still be required to explain systems with giant planets (Haghighipour 2013) however. The amount of mass used in our rocky nebula model is $6 M_{\oplus}$ between 0.05~AU and 0.5~AU. If we estimate the same amount of mass in this range using the surface density profile assumed in HM13, we get 12.4~$M_{\oplus}$, around a star of twice the size. Thus, our two models nicely match the approximate linear scaling of disk mass with respect to stellar mass inferred by Andrews et al. (2013) from observations of infrared excesses around stars in Taurus. Furthermore, this amount of rocks requires only $\sim 10^{-3} M_{\odot}$ of gas for a solar metallicity protoplanetary disk, and is thus quite possible around a young M-star. Our model also has implications for the detectability of habitable planets around M~dwarfs. We have shown in Figure~\ref{Water} that the disk mass appropriate to explaining the observed distribution of tranets usually produces one, and often two, planets in the nominal HZ. We can generalise this assertion by making use of the same semi-analytic estimate of planetary spacings outlined in HM13. If we assert that a given surface density profile results in a chain of planets that are characterised by mass, energy and angular momentum conservation, and the requirement that the final planets be spaced according to the $\Delta$ distribution shown in Figure~\ref{DeltaM}, then we can relate the $\Delta$ between two neighbouring planets to the quantity $y=a_2/a_1$, where $a_1$ and $a_2$ are the semi-major axes of the inner and outer members of the pair. Using the same $\Sigma \propto R^{-2}$ profile, we derive \begin{equation} \Delta = 115 \frac{y-1}{y+1} \left( \frac{M_*}{0.5 \rm M_{\odot}} \frac{6 \rm M_{\oplus}}{M_{disk}} \frac{1}{\ln y} \right)^{1/3} \end{equation} where $M_{disk}$ is the total planetesimal disk mass, integrated from 0.05~AU to 0.5~AU. If the habitable zone extends from 0.23 to 0.44~AU, this corresponds to y=1.91, and $\Delta \sim 42$. A comparison with Figure~\ref{DeltaM} shows that this value is larger than almost all of the pairs in our simulations, and indicates that is very difficult to form a system with this model without at least one planet in the HZ. We can extend this discussion by asking how large a disk do we need before it becomes likely to not have any planets in the HZ? If we set $\Delta=25$, the peak of the distribution, and keep $y=1.9$, we can use this to derive a disk mass (between 0.05 and 0.5~AU) of $M_{disk}>28 M_{\oplus}$. A more useful comparison is if we convert this disk mass into the summed mass of the pair of planets that might straddle the HZ. For our surface density profile, the mass fraction of the total 0.05--0.5~AU disk that is contained in the HZ is 0.28, so that a pair of planets that straddles the habitable zone and has a combined mass $>7.8 M_{\oplus}$ may not have any further planets between them. Our discussion has thus far focussed on transiting systems, because of the larger sample size and prospects for eventual atmospheric characterisation. However, several planetary systems have also been discovered around M~dwarfs, and these are also well matched by in situ assembly models. Figure~\ref{Stat2} shows the comparison of the simulation results with the observed planetary systems around the stars GJ~581 and GJ~667C (Udry et al. 2007; Mayor et al. 2009; Vogt et al. 2010; Bonfils et al. 2011; Forveille et al. 2011; Delfosse et al. 2012; Vogt, Butler \& Haghighipour 2012; Anglada-Escude' et al. 2012, 2013). We compare the relative spacings ($S_s$) and mass concentration indices ($S_c$) shown in Table~\ref{OutTab} with the mass and periods measured from radial velocities. We see that both the observed systems fall within the parameter range probed by the simulations. The star GJ~667C hosts several potentially habitable planets (Anglada-Escude' et al. 2013), as expected from the above discussion. We can also adapt this rationale to the controversy regarding the existence of habitable planets in the GJ581~system (Vogt et al. 2010, 2012; Forveille et al. 2011). The existence of GJ~581c and GJ~581d are not disputed, and they exhibit a period ratio of 6.8, corresponding to $y=3.59$. Their estimated minimum combined mass is 11.2 $M_{\oplus}$, which yields $\Delta \sim 34$, for a host star mass $\sim 0.32 M_{\odot}$. The issue is whether the claimed planet GJ581g exists between them. Comparing to Figure~\ref{DeltaM}, we see that $\Delta \sim 34$ is larger than most pairs, but does occur in a few cases. Thus, it is not impossible that there is a true gap between GJ~581c and GJ~581d, although there should be another planet between them in the majority of planetary systems of this type. If GJ~581g did exist with the nominal $3.1 M_{\oplus}$ mass and 36.6~day period, it would yield $\Delta \sim 23$ and $\Delta \sim 12$ with respect to the inner and outer companions, which is also not ruled out on dynamical grounds, although it would be comparable to the most compact systems that emerged from the simulations. It is possible to locate a planet in this gap in a more dynamically stable location, in terms of maximising the $\Delta$ values with respect to both GJ~581c and GJ~581d. Planets with periods $\sim 29$~days would have $\Delta \sim 20$ with respect to both neighbours, assuming a $1 M_{\oplus}$ planet in the gap. Although the in situ assembly model provides a good description of the distribution of period ratios and overall period distribution of the observed multiple systems, we must also note that the model once again underpredicts the number of single tranets observed, just as in the case of the solar-like hosts (HM13). Furthermore, Figure~\ref{MultiComp} shows that the ratio of singles to doubles is similar to the case of the solar hosts. However, unlike in the solar case, it does appear as though there is some difference in the distribution of single tranets, in that there is a great overabundance of single tranet systems for orbital periods $< 2.2$~days. Once again, this discrepancy suggests that our description is not complete, and that a comprehensive model for the formation of these planetary systems will require additional elements to either reduce the multiplicity of some systems, or to increase the spread in inclinations. Indeed, if we postulate a second process that produces only single planet systems, then, in order to match the observed multiplicities, we require that $\sim 58\%$ of observed tranet systems form via this alternative process. Our model would then be responsible for the remaining 42\% systems, and for 56\% of all observed tranets, including all multiple tranet systems. This has implications for the estimates of the frequency of habitable planets around M~dwarfs. The analysis of the frequency of tranet detections (DC13, Kopparapu 2013) indicate that such planetary systems around M~dwarfs are quite common, and that the nearest planet in the HZ of an M~dwarf may be within only a few parsecs. DC13 estimate that 25\% of early M-dwarfs host a planet with radius between 0.5-1.4 $\rm R_{\oplus}$. If we assume that these planets result from our in situ model plus the second, singles-only process in the proportions above, this implies that 10\% of such stars host high multiplicity systems. Given that every planetary system formed by the above model produces at least one habitable zone planet, this suggests that the incidence of habitable planets is $\sim 0.10$ per star, which is quite similar to the number obtained by DC13. It would also increase if we include those systems with radius between 1.4--4 $\rm R_{\oplus}$, which have an equivalent occurence rate. Our results therefore only strengthen the conclusion of DC13, in that we demonstrate that the bulk of the observed systems are well reproduced by an in situ assembly model, and that the model suggests that a large fraction of planetary systems observed to have a transiting earth-mass planet at any period should also have at least one planet within the HZ. | 14 | 3 | 1403.6553 |
1403 | 1403.6309_arXiv.txt | {The Kuiper belt is formed of planetesimals which failed to grow to planets and its dynamical structure has been affected by Neptune. The classical Kuiper belt contains objects both from a low-inclination, presumably primordial, distribution and from a high-inclination dynamically excited population.} {Based on a sample of classical TNOs with observations at thermal wavelengths we determine radiometric sizes, geometric albedos and thermal beaming factors for each object as well as study sample properties of dynamically hot and cold classicals.} {Observations near the thermal peak of TNOs using infra-red space telescopes are combined with optical magnitudes using the radiometric technique with near-Earth asteroid thermal model (NEATM). We have determined three-band flux densities from \emph{Herschel}/PACS observations at 70.0, 100.0 and $160.0\ \mathrm{\mu m}$ and \emph{Spitzer}/MIPS at 23.68 and $71.42\ \mathrm{\mu m}$ when available. We use reexamined absolute visual magnitudes from the literature and ground based programs in support of \emph{Herschel} observations. } {We have analysed 18 classical TNOs with previously unpublished data and re-analysed previously published targets with updated data reduction to determine their sizes and geometric albedos as well as beaming factors when data quality allows. We have combined these samples with classical TNOs with radiometric results in the literature for the analysis of sample properties of a total of 44 objects. We find a median geometric albedo for cold classical TNOs of $0.14_{-0.07}^{+0.09}$ and for dynamically hot classical TNOs, excluding the Haumea family and dwarf planets, $0.085_{-0.045}^{+0.084}$. We have determined the bulk densities of Borasisi-Pabu ($2.1_{-1.2}^{+2.6}$ g cm$^{-3}$), Varda-Ilmar\"e ($1.25_{-0.43}^{+0.40}$ g cm$^{-3}$) and 2001 QC$_{298}$ ($1.14_{-0.30}^{+0.34}$ g cm$^{-3}$) as well as updated previous density estimates of four targets. We have determined the slope parameter of the debiased cumulative size distribution of dynamically hot classical TNOs as $q$=2.3$\pm$0.1 in the diameter range 100$<$$D$$<$500 km. For dynamically cold classical TNOs we determine $q$=5.1$\pm$1.1 in the diameter range 160$<$$D$$<$280~km as the cold classical TNOs have a smaller maximum size. } {} | Transneptunian objects (TNO) are believed, based on theoretical modeling, to represent the leftovers from the formation process of the solar system. Different classes of objects may probe different regions of the protoplanetary disk and provide clues of different ways of accretion in those regions (\cite{Morbidelli2008}). Basic physical properties of TNOs, such as size and albedo, have been challenging to measure. Only a few brightest TNOs have size estimates using direct optical imaging (e.g. Quaoar with \emph{Hubble}; \cite[2004]{Brown2004}). Stellar occultations by TNOs provide a possibility to obtain an accurate size estimate, but these events are rare and require a global network of observers (e.g. Pluto's moon Charon by \cite{Sicardy2006}; and a member of the dynamical class of classical TNOs, 2002 TX$_{300}$, by \cite{Elliot2010}, 2010). Predictions of occultations are limited by astrometric uncertainties of both TNOs and stars. Combining observations of reflected light at optical wavelengths with thermal emission data, which for TNOs peaks in the far-infrared wavelengths, allows us to determine both size and geometric albedo for large samples of targets. This {\it radiometric method} using space-based ISO (e.g. \cite{Thomas2000}), \emph{Spitzer} (e.g. \cite[2008]{Stansberry2008}, \cite[2009]{Brucker2009}) and \emph{Herschel} data (\cite[2010]{Muller2010}, \cite[2010]{Lellouch2010}, \cite[2010]{Lim2010}, \cite[2012]{SantosSanz2012}, \cite[2012]{Mommert2012}, \cite[2012]{Vilenius2012}, \cite{Pal2012}, \cite[2013]{Fornasier2013}) has already changed the size estimates of several TNOs compared to those obtained by using an assumed albedo and has revealed a large scatter in albedos and differences between dynamical classes of TNOs. Observations at thermal wavelengths also provide information about thermal properties (\cite[2013]{Lellouch2013}). Depending on the thermal or thermophysical model selected it is possible to derive the thermal beaming factor or the thermal inertia, and constrain other surface properties. Ground-based submillimeter observations can also be used to determine TNO sizes using the radiometric method (e.g. \cite{Jewitt2001b}), but this technique has been limited to very few targets so far. TNOs, also known as Kuiper belt objects (KBO), have diverse dynamical properties and they are divided into classes. Slightly different definitions and names for these classes are available in the literature. Classical TNOs (hereafter CKBO) reside mostly beyond Neptune on orbits which are not very eccentric and not in mean motion resonance with Neptune. We use the \cite{Gladman2008} classification: CKBOs are non-resonant TNOs which do not belong to any other TNO class. The eccentricity limit is $e$\,$\lesssim$\,$0.24$, beyond which objects belong to {\it detached objects} or {\it scattering/scattered objects}. Classical TNOs with semimajor axis $39.4$\,$<$\,$a$\,$<$\,$48.4\ \mathrm{AU}$ occupy the {\it main classical belt}, whereas {\it inner} and {\it outer} classicals exist at smaller and larger semi-major axis, respectively. Apart from the Gladman system, another common classification is defined by the Deep Eplictic Survey Team~(DES, \cite{Elliot2005} 2005). For the work presented here, the most notable difference between the two systems is noticed with high-inclination objects. Many of them are not CKBOs in the DES system. In the inclination/eccentricity space CKBOs show two different populations, which have different frequency of binary systems (\cite{Noll2008}), different luminosity functions (LF; \cite{Fraser2010} 2010), different average geometric albedos (\cite{Grundy2005, Brucker2009} 2009) and different color distributions (\cite[2008]{Peixinho2008}). The low-inclination ``cold'' classicals are limited to the main classical belt and have a higher average albedo, more binaries and a steeper LF-derived size distribution than high-inclination ``hot'' classicals. Some amount of transfer between the hot and cold CKBOs is possible with an estimated maximum of 5\% of targets in either population originating from the other than its current location (\cite{Volk2011}). The \emph{``TNOs are Cool'': A survey of the trans-Neptunian region} open time key program (\cite{Muller2009}) of \emph{Herschel Space Observatory} has observed 12 cold CKBOs, 29 hot CKBOs, and five CKBOs in the inner classical belt, which are considered to be dynamically hot. In addition, eight CKBOs have been observed only by \emph{Spitzer Space Telescope}, whose TNO sample was mostly overlapping with the \emph{Herschel} one. This paper is organized in the following way. We begin by describing our target sample in Section~\ref{Targetsample}, followed by \emph{Herschel} observations and their planning in Section~\ref{Hobs} and \emph{Herschel} data reduction in Section~\ref{dataredux}. More far-infrared data by \emph{Spitzer} are presented in Section~\ref{Sobs} and absolute visual magnitudes in Section~\ref{auxobs}. Thermal modeling combining the above mentioned data is described in Section~\ref{model} and the results for targets in our sample in Section~\ref{resultsection}, comparing them with earlier results when available (Section~\ref{resultscomp}). In Section~\ref{discussions} we discuss sample properties, cumulative size distributions, correlations and binaries as well as debiasing of the measured size distributions. Conclusions of the sample analysis are given in Section~\ref{conclude}. | \label{conclude} The \emph{Herschel} mission and the cold phase of \emph{Spitzer} have ended. The next space mission capable of far-infrared observations of CKBOs will be in the next decade. Occultations can provide very few new size estimates annually, and the capabilities of the \emph{Atacama Large Millimeter Array} (ALMA) to significantly extend the sample of measured sizes of TNOs already presented may be limited by its sensitivity\footnote{\cite{Moullet2011} estimated 500 TNOs to be detectable by ALMA, based on assumed albedos commonly used at that time.}. In this work we have analysed 18 classical TNOs to determine their sizes and albedos using the radiometric technique and data from \emph{Herschel} and/or \emph{Spitzer}. We have also re-analysed previously published targets, part of them with updated flux densities. The number of CKBOs with size/albedo solutions in literature and this work is increased to 44 targets and additionally three targets have a diameter upper limit and albedo lower limit. We have determined the mass density of three CKBOs and updated four previous density estimates. Our main conclusions are: \begin{enumerate} \item The dynamically cold CKBOs have higher geometric albedo ($0.14$) than the dynamically hot CKBOs ($0.085$ without dwarf planets and Haumea family, 0.10 including them), although the difference is not as great as reported by \cite{Vilenius2012} (2012). \item We do not confirm the general finding of \cite{Vilenius2012} (2012) that there is an anti-correlation between diameter and albedo among all measured CKBOs as that analysis was based on a smaller number of targets. \item The cumulative size distributions of cold and hot CKBOs have been infered using a two-stage debiasing procedure. The characteristic size of cold CKBOs is smaller, which is compatible with the hypothesis that the cold sub-population may have formed at a larger heliocentric distance than the hot sub-population. The cumulative size distribution's slope parameters of hot CKBOs in the diameter range 100$<$$D$$<$500 km is $q$=2.3$\pm$0.1. Dynamically cold CKBOs have an infered slope of $q$$=$5.1$\pm$1.1 in the range 160$<$$D$$<$280. \item The bulk density of Borasisi is $2.1_{-0.59}^{+0.58}$ g cm$^{-3}$, which is higher (but within error bars) than other CKBOs of similar size. The bulk densities of % Varda and 2001 QC$_{298}$ are $1.25_{-0.43}^{+0.40}$ g cm$^{-3}$ and $1.14_{-0.30}^{+0.34}$ g cm$^{-3}$, respectively. Our re-analysis of four targets ($D$$<$400 km) has decreased their density estimates and they are mostly between 0.5 and 1 g cm$^{-3}$ implying high macroporosity. \end{enumerate} | 14 | 3 | 1403.6309 |
1403 | 1403.6664_arXiv.txt | The hot Jupiter HD 189733b is probably the best studied of the known extrasolar planets, with published transit and eclipse spectra covering the near UV to mid-IR range. Recent work on the transmission spectrum has shown clear evidence for the presence of clouds in its atmosphere, which significantly increases the model atmosphere parameter space that must be explored in order to fully characterise this planet. In this work, we apply the NEMESIS atmospheric retrieval code to the recently published \textit{HST}/STIS reflection spectrum, and also to the dayside thermal emission spectrum in the light of new \textit{Spitzer}/IRAC measurements, as well as our own re-analysis of the \textit{HST}/NICMOS data. We first use the STIS data to place some constraints on the nature of the cloud on HD 189733b, and explore solution degeneracy between different cloud properties and the abundance of Na in the atmosphere; as already noted in previous work, absorption due to Na plays a significant role in determining the shape of the reflection spectrum. We then perform a new retrieval of the temperature profile and abundances of H$_2$O, CO$_2$, CO and CH$_4$ from the dayside thermal emission spectrum. Finally, we investigate the effect of including cloud in the model on this retrieval process. We find that the current quality of data does not warrant the extra complexity introduced by including cloud in the model; however, future data are likely to be of sufficient resolution and signal-to-noise that a more complete model, including scattering particles, will be required. | Since its discovery in 2005 \citep{bouchy05}, the hot Jupiter HD 189733b has been repeatedly observed as it transits and is eclipsed by its parent star, leading to excellent coverage in both its transmission and eclipse spectra from the visible to mid-infrared. This has resulted in HD 189733b being probably the best characterised of all the known exoplanets; it is known from the observed slope of the reflectance spectrum that this unresolved planet would appear a deep shade of blue \citep{evans13}, a fact that testifies to the power of the transit spectroscopy technique. The transmission spectrum presented by \citet{pont13} shows clear evidence for haze or cloud in the atmosphere of this planet, making it the first transiting planet outside the solar system that is known to be cloudy. \citet{lee12} and \citet{line12} analyse the dayside spectrum from secondary eclipse observations, and retrieve the temperature-pressure profile and abundances of H$_2$O, CO$_2$, CO and CH$_4$. \citet{knutson12} use \textit{Spitzer}/IRAC phase curves to investigate the longitudinal temperature variability, and \citet{dewit12} combine the phase curves with an analysis of the ingress/egress shape in eclipse to place further constraint on spatial variability. The recent albedo spectrum from \textit{HST}/STIS obtained by \citet{evans13} provides an opportunity to investigate the cloud structure on the dayside. Unlike the transmission spectrum investigated by \citet{pont13} and \citet{lee13} which probes the limb of the exoplanet amosphere, the dayside reflection spectrum has near-nadir geometry, and so it is sensitive to deeper regions of the atmosphere. If the cloud is similar at the terminators and on the dayside, we can use this to further constrain its properties. Alternatively, we may see an entirely different cloud layer on the hotter dayside from that observed at the terminator. The albedo spectrum is useful for placing constraint on the cloud, as scattering particles have a significant effect on the optical reflectivity of an atmosphere. We expect there to be fewer gaseous absorbing species in the visible part of the spectrum than in the infrared; distinct absorption features of Na and K are seen in the transmission spectrum, but otherwise cloud appears to be the dominant opacity source in this region \citep{pont13}. We also expect little, if any, thermal contribution from the planet itself at these wavelengths due to its temperature. In this work, we use the Non-linear optimal Estimator for MultivariateE spectral analySIS (NEMESIS) software \citep{irwin08} to calculate synthetic spectra using a simple cloudy model atmosphere for HD 189733b, including multiple scattering. We vary cloud parameters and Na volume mixing ratio (VMR) and compare the resultant spectra to the STIS measurement; calculating the $\chi^2$ goodness-of-fit parameter allows us to determine the region of model parameter space that provides the best match. We then use a subset of our best-fitting models to examine the effect of clouds on our ability to accurately retrieve temperature and molecular abundances from the infrared dayside emission spectrum, following the work of \citet{lee12}. \subsection{NEMESIS} We use the NEMESIS spectral retrieval tool to produce forward models (predicted spectra for a range of model atmospheres) for comparison with the \textit{HST}/STIS spectra, and also to retrieve the atmospheric state from the thermal emission spectrum as in \citet{lee12}. NEMESIS was developed by \citet{irwin08} for atmospheric retrieval of solar system planets, and has since been extended to enable the same analysis for observations of transiting and directly imaged extrasolar planets \citep{lee12,lee13a,barstow13b}. NEMESIS is an Optimal Estimation retrieval model \citep{rodg00} and uses a correlated-k radiative transfer model \citep{lacis91,goodyyung}. NEMESIS is not a radiative equilibrium model; instead, it simply uses the atmospheric model provided to compute the incoming and outgoing radiative flux. In an irradiated case, it will compute the incoming and scattered/reflected flux from the star, but it will not take into account the heating effect of the incoming stellar flux on the atmosphere. For the multiple scattering runs, NEMESIS uses beams over a user-specified number of zenith angles; azimuthal dependence is accounted for using Fourier decomposition. For more details of the scattering calculation in this work, see Section~\ref{multscat}. We use the same line and collision-induced absorption data as \citet{lee12}, \citet{lee13} and \citet{barstow13}. A list of sources for absorption line data is given in Table~\ref{linedata}. H$_2$-He collision-induced absorption data are taken from the models of \citet{borysow89,borysowfm89,borysow90,borysow97} and \citet{borysow02}. The reference stellar spectrum is taken from the model set made available by Kurucz\footnote{http://kurucz.harvard.edu/stars/HD189733/}, and the stellar radius is taken from \citet{ellyn08}. We use the same planetary mass and radius as \citet{lee13}, and a H$_2$:He ratio of 9:1. As found by \citet{lee12}, the precise value of this does not have a large effect on secondary eclipse retrievals. \begin{table} \centering \begin{tabular}[c]{|c|c|} \hline Gas & Source\\ \hline H$_2$O & HITEMP2010 \citep{roth10}\\ CO$_2$ & CDSD-1000 \citep{tash03}\\ CO & HITRAN1995 \citep{roth95}\\ CH$_4$ & STDS \citep{weng98}\\ Na & VALD \citep{heiter08}\\ K & VALD \citep{heiter08}\\ \hline \end{tabular} \caption{Sources of gas absorption line data.\label{linedata}} \end{table} \subsubsection{Multiple scattering calculations}\label{multscat} In order to reproduce an accurate reflection spectrum in the presence of optically thick cloud, it is necessary to include multiple scattering as many scattering events are likely. NEMESIS uses the matrix operator algorithm of \citet{plass73} for multiple scattering calculations, where the zenith angle integration is achieved using a 5-point Gaussian-Lobatto quadrature scheme and azimuth angle integration is achieved through Fourier decomposition, with the necessary number of Fourier components determined by the stellar and emission zenith angles. The analytical disc-averaged integration scheme used in \citet{lee12} for eclipse spectra is not used here as it is not applicable to scattering situations; instead, we represent the disc average for the synthetic STIS spectra by running multiple scattering calculations with the stellar zenith angles set to each of the 5 Gaussian-Lobatto quadrature angles and the azimuth angle set for back-scattering, since during secondary transit the observer is located in the same direction as the star and the stellar zenith angle is equal to the emission angle. The disc-average is then determined using a weighted average of these 5 different calculations, assuming that the atmospheric conditions are the same at all points on the disc. Retrievals including multiple scattering are computationally expensive. Therefore, we anticipate that the majority of the thermal emission retrieval calculations in the future will still be performed with an extinction-only approximation using the disc integration described by \citet{lee12}. However, we also tested the effect of including multiple scattering, to test the sensitivity of the retrieval to differences in our modelling approach. The 5-angle approach used to calculate the STIS synthetic spectra is still too time-consuming for the emission spectrum retrieval, so in this case we approximate further by calculating the spectrum for a stellar zenith angle of 45$^{\circ}$ only, which represents the average angle of weighting function used to compute the disc-averaged spectrum since $\bar{R}=\int_0^{\frac{\pi}{2}}\!2R(\theta)\mathrm{sin}(\theta)\mathrm{cos}(\theta)\, \mathrm{d}\theta=\int_0^{\frac{\pi}{2}}\!R(\theta) \mathrm{sin}(2\theta) \,\mathrm{d}\theta$. The scattering phase function of the particles is calculated using Mie theory and approximated by the Henyey-Greenstein parameterisation \citep{hg41}, as in previous work on planetary clouds (e.g. \citealt{irwin09}, \citealt{barstow12}); expected deviations from the true phase function are small compared with the errors on the observed spectrum, so we consider this approximation to be valid. We use the double-peaked version of the phase function: $P(\theta)= \frac{1}{4\pi}\left[f\frac{1-{g^2_1}}{(1+{g^2_1}-2{g_1}\mu)^{3/2}}+(1-f)\frac{1-{g^2_2}}{(1+{g^2_2}-2{g_2}\mu)^{3/2}}\right]$, which represents the phase function as the sum of forward and backward scattering peaks. Here, $\mu$ is the cosine of the scattering angle, $g_1$ and $g_2$ are scattering asymmetry parameters for the forward and backward peaks respectively, and $f$ is the fractional contribution of the forward peak to the total phase function. For smaller particles approaching the Rayleigh scattering limit, the parameter $f$ is close to 0.5 and there are approximately equal contributions from the forward and backward scattering peaks; as the particle size increases relative to the wavelength of light the value of $f$ increases and the scattering becomes more asymmetric. However, in no case is the scattering completely isotropic. We use the enstatite refractive index values of \citet{scott96} and the MnS values from \citet{huff67} in our calculations of the scattering parameters. It is worth noting that, as we model realistic particles, they absorb as well as scatter incident radiation. The fraction of light that is absorbed rather than scattered is dependent on the particle size and also on the composition of the particles. In general, more of these realistic particles would be required than idealised, perfectly scattering particles to produce an equivalent planetary albedo in an atmospheric model, since for realistic particles some of the incident radiation is absorbed by the particles rather than scattered back to space. The single scattering albedo is also wavelength dependent, and this can therefore affect the shape of modelled planetary albedo and thermal emission spectra. | We find that a large range of enstatite cloud models can fit the measured \textit{HST}/STIS spectrum. Cloud-free model atmospheres are also acceptable solutions, although we consider this to be less plausible due to the clear evidence from the transmission spectrum that HD 189733b is cloudy \citep{pont13,lee13} and the likelihood that small particles are evenly distributed in hot Jupiter atmospheres \citep{parmentier13}. Small enstatite particles ($<$0.1 $\upmu$m) and 50 ppmv of Na provide the best fit to the STIS data of the examples we test; we find an overlap with the models of \citet{lee13} for a uniformly-distributed 0.1 $\upmu$m cloud with an optical depth of 0.5 at 0.25 $\upmu$m. However, the problem is extremely degenerate and we cannot exclude solutions with larger cloud particles. The retrieval of temperature and atmospheric composition from the thermal emission spectrum is relatively insensitive to the inclusion of cloud in the model atmosphere, for our best-fitting models. This suggests that for the case of HD 189733b accurate retrieval of temperature and gaseous abundances from the thermal emission spectrum is possible, even without detailed knowledge of the cloud properties. Solution degeneracy prevents firm conclusions from being drawn about the nature of the cloud on HD 189733b; the current quality and coverage of spectroscopic data for most exoplanets is therefore insufficient to simultaneously constrain temperature structure, gaseous abundances and multiple cloud properties. Given the additional complexity (and therefore number of degenerate solutions) introduced to the retrieval problem when clouds are included, this implies that the best approach with the currently available secondary eclipse data is to use cloud-free model atmospheres for temperature retrieval. As data quality improves, alternative, more detailed cloud models must be explored and their effect on the emission spectrum reassessed. | 14 | 3 | 1403.6664 |
1403 | 1403.1875_arXiv.txt | We present the fiducial main sequence stellar locus traced by 10 photometric colors observed by SDSS, 2MASS, and WISE. Median colors are determined using 1,052,793 stars with $r$-band extinction less than 0.125. We use this locus to measure the dust extinction curve relative to the $r$-band, which is consistent with previous measurements in the SDSS and 2MASS bands. The WISE band extinction coefficients are larger than predicted by standard extinction models. Using 13 lines of sight, we find variations in the extinction curve in $H$, $K_s$, and WISE bandpasses. Relative extinction decreases towards Galactic anti-center, in agreement with prior studies. Relative extinction increases with Galactic latitude, in contrast to previous observations. This indicates a universal mid-IR extinction law does not exist due to variations in dust grain size and chemistry with Galactocentric position. A preliminary search for outliers due to warm circumstellar dust is also presented, using stars with high signal-to-noise in the W3-band. We find 199 such outliers, identified by excess emission in $K_s-W3$. Inspection of SDSS images for these outliers reveals a large number of contaminants due to nearby galaxies. Six sources appear to be genuine dust candidates, yielding a fraction of systems with infrared excess of $0.12\pm0.05$\%. | The Two Micron All Sky Survey \citep[2MASS;][]{2mass} and Sloan Digital Sky Survey \citep[SDSS;][]{york2000} have provided revolutionary improvements in our understanding of the stellar populations within our Galaxy at near infrared and optical wavelengths, respectively. For example, Milky Way halo substructures have been discovered in both SDSS and 2MASS photometry \citep{ibata2001,ibata2002}. Normal stars have been separated from more exotic objects with much greater accuracy by matching sources between these surveys \citep{finlator2000}. This wide-field dataset has continued to set the standard for multi-wavelength studies, enabling science not possible with either survey individually, and providing astrometric and flux standards to calibrate future surveys. \citet[][hereafter C07]{covey2007} used a sample of $\sim$600,000 stars matched between SDSS and 2MASS to measure the fiducial stellar color locus in $ugrizJHK_s$ passbands as a function of $g-i$ color. This parameterization provides colors of main sequence stars as a function of their effective temperature. The C07 main sequence locus has been used to search for color outliers due to being, for example, white dwarf binaries, quasars, or post-main sequence stars. The C07 locus has also provided a robust method to identify and classify normal stars given any combination of SDSS and 2MASS colors. The Wide-field Infrared Survey Explorer \citep[WISE;][]{wise} has produced a modern census of the entire sky with unprecedented accuracy and depth in four bandpasses, ranging from 3.4 to 22 $\mu$m. Combining this survey with the well studied SDSS and 2MASS datasets will enable the discovery of new classes of both Galactic and extra-Galactic objects. Furthermore, this will provide improved understanding of the cool solar neighborhood, young stellar populations, dust content, and substructure within our Galaxy. Already, WISE and SDSS data have been used to survey infrared excesses around white dwarfs \citep{debes2011} and to discover many new brown dwarf candidates \citep{aberasturi2011}. The first confirmed Y0 dwarf \citep{cushing2011} has also been discovered with WISE, probing the stellar mass function at its lowest extrema for the first time. WISE photometry will also provide the best dataset for mapping asymptotic giant branch, with the ability to trace them to well beyond the Galactic center (Hunt-Walker 2014 in prep). All of these studies will critically depend on the proper understanding of photometric behavior for ``normal'' stars simultaneously in all 10 colors, as was done with the seven colors from SDSS and 2MASS in C07. In this paper we present the first detailed study of the stellar locus for nearby stars as observed by 2MASS, SDSS, and WISE. In \S\ref{sec:data} we describe the creation of a matched sample of low-extinction point sources. A detailed measurement of the stellar locus is given in \S\ref{sec:locus}. Using this fiducial stellar color sequence as a set of ``standard crayons'' \citep[termed by][]{peek2010}, we measure the relative dust extinction coefficients from the $u$-band to 22$\mu$m in \S\ref{sec:extinct}. We search for warm dust disks from WISE color outliers in \S\ref{sec:dust}. A summary of our work is given in \S\ref{sec:conclusion}. | \label{sec:conclusion} We have presented a study of the fundamental properties of stars across a wide range in wavelength. Our stellar locus, derived from a million low-extinction sources, will be of great utility to many future studies with the powerful multi-wavelength combination of SDSS - 2MASS - WISE. Spectroscopically confirmed low-mass stars and brown dwarfs from SDSS, matched to WISE photometry, will also provide an extension of our color locus to lower mass objects (S. J. Schmidt 2014 in prep). A brief summary of our work is as follows: \begin{enumerate} \item A measurement of the 10-dimensional color locus was presented, from SDSS $u$-band to WISE $W3$-band, using 1,052,793 stars with low extinction $(A_r<0.125)$. This locus contains the best characterization of stellar colors in WISE passbands to date. \item We have empirically measured the $r$-band relative dust extinction coefficients, $A_\lambda/A_r$, for each of the photometric bands in our sample, providing strong constraints for dust composition models in the infrared. \item Variations in the infrared dust extinction have been shown for different lines of sight. Coherent trends with both Galactic latitude and longitude were seen. Increasing relative infrared extinction with increasing Galactic latitude was found to be in opposition to previous observations. A detailed follow-up investigation of the properties of dust extinction and emission in the infrared with WISE is strongly motivated. \item From a subset of our sample we recovered 199 infrared excess candidates that span a wide range of optical colors. The majority of these were found to be contaminants from neighboring stars or background galaxies. Six objects appear to be bona fide infrared excess systems, possibly indicative of dust disks. Higher resolution and longer wavelength followup is required to verify these systems. \end{enumerate} | 14 | 3 | 1403.1875 |
1403 | 1403.3088_arXiv.txt | *{xxx} \abstract{The interpretation of observed spectra of stars in terms of fundamental stellar properties is a key problem in astrophysics. For FGK-type stars, the radiative transfer models are often computed using the assumption of local thermodynamic equilibrium (LTE). Its validity is often questionnable and needs to be supported by detailed studies, which build upon the consistent framework of non-LTE. In this review, we outline the theory of non-LTE. The processes causing departures from LTE are introduced qualitatively by their physical interpretation, as well as quantitatively by their impact on the models of stellar spectra and element abundances. We also compile and analyse the most recent results from the literature. In particular, we examine the non-LTE effects for 24 chemical elements for six late-studied FGK-type stars.} | Local thermodynamic equilibrium (LTE) is a common assumption when solving the radiative transfer problem in stellar atmospheres. The reason for adopting LTE is that it tremendously simplifies the calculation of number densities of atoms and molecules. However, the trouble is that the assumption essentially implies that stars do not radiate. Quoting \citet{1973ARA&A..11..187M}, 'Departures from LTE occur simply because stars have a boundary through which photons escape into space'. In what follows, we will take a short excursus into why this happens and provide a brief overview of the recent developments in the field. We start with reviewing the theoretical foundations of radiative transfer as it happens in reality, i.e. including its impact on the properties of matter in stellar atmospheres. This approach is known as non-local thermodynamic equilibrium (non-LTE). We discuss the non-LTE 'mechanics' and summarise the impact on spectral line formation. Finally, we present a list of non-LTE abundance corrections for the chemical elements, for which detailed statistical equilibrium calculations are available in the literature. | Non-local thermodynamic equilibrium is a framework, which describes consistently the propagation of radiation in a stellar atmosphere and its coupling to matter. LTE is the boundary case of non-LTE, the approximation in the limit of infinitely large collision rates. In this review lecture, the focus is on non-LTE conditions in the atmospheres of cool late-type stars. The impact of non-LTE on stellar parameter and abundance determinations can be gained from detailed theoretical and observational analyses, and these are now available for the most important chemical elements observed in the spectra of cool stars. So far, most studies in the literature focussed on deviations from LTE in the spectral lines of neutral or singly-ionised atoms. Generally, the results obtained by independent groups and methods are consistent and they can be summarised as follows. There are several well-defined types of non-LTE effects, and consequently several groups of species which behave similarly under the same physical conditions in a stellar atmosphere (i.e. given the same $\teff$, $\log g$, and $\feh$). The first group is formed by species sensitive to over-ionisation (e.g. \ion{Mg}{I}, \ion{Si}{I}, \ion{Ca}{I}, \ion{Fe}{I}), the second group are collision-dominated species (\ion{Na}{I}, \ion{K}{I}), and the rest are mixed-type ions (e.g. \ion{Li}{I}, \ion{O}{I}, \ion{Ba}{II}), which may show positive or negative non-LTE effects depending upon different factors. In some cases, the non-LTE abundance corrections vary by an order of magnitude between the groups of elements. In other cases, cancellation may occur such that even the LTE analysis may provide the correct abundance \emph{ratios} of two elements. Such pairs are difficult to establish, although, to first approximation, one may form pairs from elements in the same group, taking into account the excitation potential of the line as well as the elemental abundances. We want to stress, however, that a widely-spread 'rule of thumb' that the lines of the same ionisation stage can be safely modelled in LTE is a misconception. Generally, any element analysed under the assumption of LTE should be regarded with caution (at least) until careful non-LTE studies have been conducted. Even then, second-order non-LTE effects related to line strengths and excitation potentials may turn out to fundamentally alter the outcome. With these potential pitfalls, we urge the reader to take non-LTE corrections into account whenever possible. \begin{acknowledgement} Fig. 3 has been kindly provided by Rob Rutten. \end{acknowledgement} | 14 | 3 | 1403.3088 |
1403 | 1403.6722_arXiv.txt | The continuous stream of data available from the \corr{\emph{Atmospheric Imaging Assembly}} (AIA) telescopes onboard the \corr{\emph{Solar Dynamics Observatory}} (SDO) spacecraft has allowed a deeper understanding of the Sun. However, the sheer volume of data has necessitated the development of automated techniques to identify and analyse various phenomena. In this \corr{article}, we describe the Coronal Pulse Identification and Tracking Algorithm (\textsf{CorPITA}) for the identification and analysis of coronal ``EIT waves''. \textsf{CorPITA} uses an intensity-profile technique to identify the propagating pulse, tracking it throughout its evolution before returning estimates of its kinematics. The algorithm is applied here to a data-set from February~2011, allowing its capabilities to be examined and critiqued. This algorithm forms part of the SDO Feature Finding Team initiative and will be implemented as part of the Heliophysics Event Knowledgebase (HEK). This is the first fully automated algorithm to identify and track the propagating ``EIT wave'' rather than any associated phenomena and will allow a deeper understanding of this controversial phenomenon. | \label{sect:intro} Solar eruptions are the most energetic events in our solar system, releasing large bursts of radiation as solar flares and ejecting plasma into the heliosphere as coronal mass ejections (CMEs). On the Sun, these eruptions are often associated with large-scale disturbances that propagate across the solar atmosphere at typical speeds of $\approx$200\,--\,400~km~s$^{-1}$ \citep{Thompson:2009il} although more recently velocities of up to $\sim$1500~km~s$^{-1}$ have also been measured \citep[\emph{e.g.}][]{Olmedo:2012ff,Shen:2013tw}. Initially observed by \citet{Moses:1997qa}, \citet{Dere:1997fk} and \citet{Thompson:1999zt} using the \corr{\emph{Extreme ultraviolet Imaging Telescope}} \citep[EIT:][]{Delaboudiniere:1995ve}, these disturbances (commonly called ``EIT waves'') have been studied in detail for more than $\approx$15~years. However, they remain a source of debate with conflicting observations of their properties leading to a myriad of theories proposed to explain the phenomenon. Initial studies of the ``EIT wave'' feature interpreted it as a fast-mode magnetoacoustic wave using the theory originally proposed by \citet{Uchida:1968gb,Uchida:1970jl}, with the disturbance initiated by the same eruption producing the associated solar flare and CME. This is consistent with observations of refraction \citep{Veronig:2006cq} and reflection \citep{Gopalswamy:2009kl} at coronal-hole boundaries and observed pulse properties such as pulse dispersion and dissipation \citep[\emph{e.g.}][]{Long:2011fv,Muhr:2011pi}, although some authors have suggested alternate wave interpretations such as solitons \citep{Wills-Davey:2007mw} or slow-mode MHD waves \citep{Podladchikova:2010fu}. However, observations of stationary bright fronts at coronal-hole boundaries and low pulse speeds have lead to the proposal of ``pseudo-wave'' theories. These interpretations see the disturbance not as a true wave, but as a bright feature produced by Joule heating at the boundary between the erupting CME and the background coronal magnetic field as the CME propagates into the heliosphere \citep{Delannee:2000bh,Delannee:2008qf}. A third alternative, originally proposed by \citet{Chen:2002zr}, combines aspects of both the wave and pseudo-wave theories to interpret ``EIT waves'' as a hybrid of both. In this case, the erupting CME drives a fast-mode wave \corr{that} then propagates freely, while magnetic reconnection driven by the restructuring magnetic field as the CME erupts is seen as a second, slower propagating feature. This theory was built upon in a subsequent \corr{article} by \citet{Chen:2005ys} and has been the focus of further simulations performed by \citet{Cohen:2009ly} and \citet{Downs:2011nx,Downs:2012cr}. There has also been some observational evidence for a second propagating front, particularly \corr{in work} by \citet{Zhukov:2004if}, \citet{Chen:2011vn}\corr{,} and \citet{Harra:2011hc}. A detailed discussion of the different proposed theories and the evidence for and against them may be found in the recent reviews by \citet{Wills-Davey:2009qc}, \citet{Gallagher:2011oq}\corr{,} and \citet{Zhukov:2011ud}. The rich variety of theories proposed to explain the ``EIT wave'' phenomenon can be explained by the methods typically used to study them. As relatively rare events\corr{,} ``EIT waves'' are often studied in isolation, with single-event studies used to infer general phenomenological properties. Despite more than 15~years of analysis, this approach remains the primary technique for investigating them and can explain the level of uncertainty that still surrounds their true physical nature. However, several authors have proposed larger statistical surveys of the phenomenon, with the work of \citet{Thompson:2009il} in particular the benchmark for statistical analysis of ``EIT waves''. \citet{Thompson:2009il} manually identified 176 events observed by EIT onboard the \emph{SOlar and Heliospheric Observatory} \citep[SOHO:][]{Domingo:1995dq} spacecraft between 24~March~1997 and 24~June~1998, finding speeds ranging from $\approx$\,50\,--\,700~km~s$^{-1}$, with values most typically being $\approx$\,200\,--\,400~km~s$^{-1}$. This catalogue has since been utilised by multiple authors, with \citet{Biesecker:2002uq} using it to show an unambiguous correlation between ``EIT waves'' and CMEs while more recent work by \citet{Warmuth:2011kh} used it to show evidence for three distinct kinematic classes of ``EIT waves''. A recent \corr{article} by \citet{Nitta:2013kc} can be seen as the spiritual successor to the work of \citet{Thompson:2009il}, producing a catalogue of ``EIT waves'' observed by the \corr{\emph{Atmospheric Imaging Assembly}} \citep[AIA:][]{Lemen:2012bs} onboard the \emph{Solar Dynamics Observatory} \citep[SDO:][]{Pesnell:2012lh} spacecraft. In this case, 171 disturbances were manually identified between April~2010 and January~2013, with the observations used to examine the relationship between ``EIT waves'', solar flares and CMEs although no relationship was found between the wave speed and flare intensity or CME magnitude. Despite their breadth and impact, the catalogues compiled by both \citet{Thompson:2009il} and more recently \citet{Nitta:2013kc} consist of manually identified events, with the consequence that identification is entirely user-dependent. This was noted by \citet{Thompson:2009il}, who assigned one of six quality ratings to each event as ``an indicator of the observability of the wave in the data''. However, the rating is entirely subjective and dependent on consistent application by the authors. This approach means that the same parameters may not necessarily be applied consistently to events identified by both catalogues, despite best efforts. Several authors have proposed automated approaches to the identification and analysis of ``EIT waves'' to try \corr{to} overcome these issues. The Novel EIT wave Machine Observing (NEMO) catalogue was developed by \citet{Podladchikova:2005ye} to identify and analyse the coronal dimmings associated with ``EIT waves'' using data from SOHO/EIT (\url{sidc.oma.be/nemo/}). This technique was successfully implemented and operated at the Solar Influences Data Analysis Centre at the Royal Observatory of Belgium from 1997 to 2010 and primarily identified coronal-dimming regions, although a strong correlation was noted between the dimming regions and the propagating ``EIT wave'' feature. An alternative technique using Huygens tracking was proposed by \citet{Wills-Davey:2006ss}. This approach employs percentage base-difference \citep[PBD; cf.][]{Wills-Davey:1999fc} images, with the pulse identified by finding the line of peak intensities corresponding to the peak of the Gaussian cross-section of the pulse. Once the disturbance has been identified in each image, its path of propagation is found using a reverse-engineered Huygens tracking approach. While highlighting several of the issues associated with the manual identification of ``EIT waves'' and attempting to rectify them, this technique has not been implemented on a larger scale to the best of the authors' knowledge. As noted by \citet{Aschwanden:2010fk}, the sheer volume of data available from SDO/AIA has underlined the need for automated algorithms to identify and characterise events of interest to the wider solar community. The SDO \corr{\emph{Feature Finding Team}} \citep[FFT:][]{Martens:2012kl} has produced and implemented a number of automated techniques designed to identify and analyse features ranging from active regions and coronal holes to flares, coronal bright points, waves, CMEs, and filaments. In this \corr{article}, we discuss the Coronal Pulse Identification and Tracking Algorithm (\textsf{CorPITA}), the ``EIT wave'' detection module for the FFT. This algorithm is designed to identify, track, and analyse ``EIT waves'' using the continuous data stream available from SDO/AIA. The algorithm is outlined in Section~\ref{sect:alg}, before being applied to a sample data set in Section~\ref{sect:app}. The performance and caveats of the algorithm are then discussed and conclusions drawn in Section~\ref{sect:conc}. | \label{sect:conc} In this \corr{article} we have presented a new automated algorithm for identifying, tracking and analysing ``EIT waves'' in data from SDO/AIA. \corr{\textsf{CorPITA}} uses an intensity-profile technique applied to percentage base-difference images to identify the propagating pulse, tracking it for as long as possible before estimating the variation in kinematics across 360 overlapping \corr{area masks} of 10$^\circ$ width. This allows the variation of the pulse propagation across the entire Sun to be studied and analysed, providing an indication of variations in the magnetic-field strength and density of the low solar corona. The systematic, automated approach of \corr{\textsf{CorPITA} allows reproducible detections of an ``EIT wave'' following a specified series of analytical steps along with a statistically significant} estimate of the pulse kinematics, providing an opportunity to determine their true physical nature across a large sample of events. \begin{figure*}[!t] \begin{center} \includegraphics[keepaspectratio,width=0.97\textwidth,trim=0 0 0 0]{CorPITA_detect_211_20110215} \includegraphics[keepaspectratio,width=0.97\textwidth,trim=0 10 0 0]{CorPITA_detect_211_20110216} \end{center} \caption{\corr{Example of the \textsf{CorPITA} output} for the events from 15~February~2011 \citep[panels a\,--\,d, previously studied by][]{Schrijver:2011qo,Olmedo:2012ff} and 16~February~2011 \citep[panels e\,--\,h, previously studied by][]{Harra:2011hc,Veronig:2011mb,Long:2013fu}. \corr{For the 15~February~2011 (16~February~2011) event, panel a (e) shows the pulse propagation with time while panels b\,--\,d (f\,--h) show the fitted kinematics and Savitsky--Golay derived velocity and acceleration, respectively, for the highest-rated arc. \textsf{CorPITA}} has successfully identified ``EIT waves'' in both cases indicating that the algorithm is robust.} \label{fig:prev_wave} \end{figure*} The period chosen to provide a sample output from \corr{\textsf{CorPITA}} included two large-scale ``EIT wave'' events that have been previously studied. The \corr{\textsf{CorPITA}} analysis of both of these events is shown in Figure~\ref{fig:prev_wave}. In both cases, ``EIT waves'' were successfully identified by the algorithm, allowing a direct comparison with the work of \citet{Schrijver:2011qo} and \citet{Olmedo:2012ff}, who studied the 15~February~2011 event, and \citet{Harra:2011hc}, \citet{Veronig:2011mb} and \citet{Long:2013fu}, who studied the 16~February~2011 event. For the 15~February~2011 event (shown in Figure~\ref{fig:prev_wave}a\,--\,d), \corr{\textsf{CorPITA}} returned an average initial velocity of $406\pm1$~km~s$^{-1}$, although this varied from $\approx$\,0\,--\,1\,000~km~s$^{-1}$ across the arcs where a pulse was detected and tracked. Figure~\ref{fig:prev_wave}a also shows the anisotropic nature of the pulse propagation, with the pulse tracked primarily across quiet regions of the solar corona and not through the different adjacent active regions. This is consistent with the work of \citet{Olmedo:2012ff} who noted the variation in pulse velocity with direction, although the current version of \corr{\textsf{CorPITA}} cannot detect the reflection that they observed while the transmission through coronal holes and active regions would produce a jump in pulse position that cannot currently be tracked. Despite this, \corr{\textsf{CorPITA}} successfully identifies the pulse and the variation in its propagation through the quiet solar corona. The 16~February~2011 event (shown in Figure~\ref{fig:prev_wave}e\,--\,h) was also successfully identified by \corr{\textsf{CorPITA}}, with an average initial velocity of $331\pm6$~km~s$^{-1}$ and fitted initial velocities varying from $\approx$\,100\,-\,975~km~s$^{-1}$ across the different arcs. The anisotropic nature of the pulse is again shown by Figure~\ref{fig:prev_wave}e, with the pulse in this case strongly directed towards solar North away from the erupting active region. The fitted initial velocities measured by \corr{\textsf{CorPITA}} are in good correspondence with those measured by \citet{Harra:2011hc}, \citet{Veronig:2011mb} and \citet{Long:2013fu}, indicating that the same pulse is identified. However, \corr{\textsf{CorPITA} is designed to identify the forward motion of a single propagating pulse and} cannot yet identify multiple \corr{propagating} pulses associated with a single event\corr{, although this may be implemented in a future iteration of the code. As a result, it} does not observe the second propagating feature identified by \citet{Harra:2011hc}. The results shown here indicate that \corr{\textsf{CorPITA}} offers a fully automated, robust approach for identifying, tracking and analysing ``EIT waves''. This will allow \corr{a systematically reproducible} analysis of the entire SDO/AIA data-set as well as a near-real-time analysis of these events when fully incorporated into the SDO feature analysis pipeline. Automating the identification of these events from SDO/AIA data will allow an improved statistical analysis which has implications for our understanding of the physical processes involved in the eruption and propagation of this phenomenon. The ability of \corr{\textsf{CorPITA}} to identify anisotropies in the propagation of these pulses also suggests that it may be used to investigate the large-scale structure of the solar corona using coronal seismology to study density and magnetic field variations. \begin{acks} The authors wish to thank Rebecca Feeney-Barry for useful discussions\corr{, and the anonymous referee whose comments helped to improve the paper}. The SDO feature finding team effort is supported by NASA. Data from SDO/AIA are courtesy of NASA/SDO and the AIA science team. DML received funding from the European Commission's Seventh Framework Programme under the grant agreement No. 284461 (eHEROES project), while DSB was funded by the European Space Agency Prodex programme. \end{acks} | 14 | 3 | 1403.6722 |
1403 | 1403.2930_arXiv.txt | {} {We derive physical properties such as the optical depths and the column densities of $^{13}$CO and C$^{18}$O to investigate the relationship between the far ultraviolet (FUV) radiation and the abundance ratios between $^{13}$CO and C$^{18}$O.} {We have carried out wide-field (0.4 deg$^2$) observations with an angular resolution of 25.8$\arcsec$ ($\sim$ 0.05 pc) in $^{13}$CO ($J$=1--0) and C$^{18}$O ($J$=1--0) toward the Orion-A giant molecular cloud using the Nobeyama 45 m telescope in the on-the-fly mode.} {Overall distributions and velocity structures of the $^{13}$CO and C$^{18}$O emissions are similar to those of the $^{12}$CO ($J$=1--0) emission. The optical depths of the $^{13}$CO and C$^{18}$O emission lines are estimated to be \tauAmin\ $<$ $\tau_{\rm ^{13}CO}$ $<$ \tauAmax\ and \tauBmin\ $<$ $\tau_{\rm C^{18}O}$ $<$ \tauBmax, respectively. The column densities of the $^{13}$CO and C$^{18}$O emission lines are estimated to be \NAmin\ $\times$ 10$^{\NAorder}$ $<$ $N_{\rm ^{13}CO}$ $<$ 3.7 $\times$ 10$^{17}$ cm$^{-2}$ and \NBmin\ $\times$ 10$^{\NBorder}$ $<$ $N_{\rm C^{18}O}$ $<$ 3.5 $\times$ 10$^{16}$ cm$^{-2}$, respectively. The abundance ratios between $^{13}$CO and C$^{18}$O, $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$, are found to be \Xmin -- \Xmax. The mean value of $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ in the nearly edge-on photon-dominated regions is found to be 16.47 $\pm$ 0.10, which is a third larger than that the solar system value of 5.5. The mean value of $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ in the other regions is found to be 12.29 $\pm$ 0.02. The difference of the abundance ratio is most likely due to the selective FUV photodissociation of C$^{18}$O. } {} | Far ultraviolet (FUV: 6 eV $<$ $h$$\nu$ $<$ 13.6 eV) radiation emitted from massive stars influences the structure, chemistry, thermal balance, and evolution of the neutral interstellar medium of galaxies \citep{Hollenbach97}. Furthermore, stars are formed in the interstellar medium (ISM) irradiated by the FUV radiation. Hence, studies of the influence of FUV are crucial for understanding the process of star formation. Regions where FUV photons dominate the energy balance or chemistry of the gas are called photon-dominated regions (PDRs). The FUV emission selectively dissociates CO isotopes more effectively than CO because of the difference in the self shielding \citep{Glassgold85, Yurimoto04, Liszt07, Rollig13}. The FUV intensity at the wavelengths of the dissociation lines for abundant CO decays rapidly on the surface of molecular clouds, since the FUV emission becomes optically very thick at these wavelengths. For less abundant C$^{18}$O, which has shifted absorption lines owing to the difference in the vibrational-rotational energy levels, the decay of FUV is much lower. As a result, C$^{18}$O molecules are expected to be selectively dissociated by UV photons, even in a deep molecular cloud interior. In the dark cloud near a young cluster (IC 5146, for example), the ratio of the $^{13}$CO to C$^{18}$O fractional abundance, $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$, considerably exceeds the solar system value of 5.5 at visual extinction ($A_V$) values of less than 10 mag \citep[Fig. 19 in][]{Lada94}. This trend indicates the selective UV photodissociation of C$^{18}$O. A variation of the abundance ratios of the isotopes is also reported \citep{Wilson99,Wang09}. For example, the values of $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ in the solar system, local ISM, Galactic center, and Large Magellanic Cloud (LMC) are measured to be 5.5, 6.1, 12.5, and 40.8. In the Milky way, the isotopic ratio is proportional to the distance from our Galactic center \citep{Wilson99}. The Orion-A giant molecular cloud (Orion-A GMC) is the nearest GMC ($d$ = 400 pc; \cite{Menten07,Sandstrom07,Hirota08}) and is one of the best studied star-forming regions (e.g., \citet{Bally87, Dutrey93, Tatematsu99,Johnstone99,Shimajiri08,Shimajiri09, Takahashi08,Davis09,Berne10,Takahashi13,Lee13}). In the northern part of the Orion-A GMC, there are three HII regions, M 42, M 43, and NGC 1977 \citep{Goudis82}. From the comparison of the AzTEC 1.1 mm and the Nobeyama 45 m $^{12}$CO ($J$=1--0), and $Midcourse\ Space\ Experiment$ ($MSX$) 8 $\mu$m emissions (from polycyclic aromatic hydrocarbons, PAHs) maps, \citet{Shimajiri11} have identified seven PDRs and their candidates in the northern part of the Orion-A GMC: 1) Orion Bar; 2) the M 43 Shell; 3) a dark lane south filament (DLSF); 4-7) the four regions A-D. Since the stratification among these distributions can be recognized, the PDR candidates are likely to be influenced by the FUV emission from the Trapezium star cluster and from NU Ori in nearly edge-on configuration. Thus, the Orion-A GMC is one of the most suitable targets for investigating the PDRs. Recently, \citet{Shimajiri13} carried out wide-field (0.17 deg$^2$) and high-angular resolution (21.3$\arcsec$ $\sim$ 0.04 pc) observations in [CI] line toward the Orion-A GMC. The mapping region includes the nearly edge-on PDRs and the four PDR candidates of the Orion Bar, DLSF, M 43 Shell, and Region D. The overall distribution of the [CI] emission coincides with that of the $^{12}$CO emission in the nearly edge-on PDRs, which is inconsistent with the prediction by the plane-parallel PDR model \citep{Hollenbach99}. The [CI] distribution in the Orion-A GMC is found to be more similar to those of the $^{13}$CO ($J$=1--0), C$^{18}$O ($J$=1--0), and H$^{13}$CO$^+$ ($J$=1--0) lines rather than that of the $^{12}$CO ($J$=1--0) line in the inner part of the cloud, suggesting that the [CI] emission is not limited to the cloud surface, but is tracing the dense, inner parts of the cloud. \begin{table*} \centering \caption{Parameters of our observations \label{observations_parameter}} \begin{tabular}{lcc} \hline \hline Molecular line & $^{13}$CO ($J$=1--0) & C$^{18}$O ($J$=1--0) \\ \hline Rest Frequency [GHz] & 110.201354 & 109.782176\\ Observation & 2013 May & 2010 March -- 2013 May\\ Scan mode & OTF & OTF \\ Mapping size [deg$^{2}$] & 0.4 & 0.4 \\ Effective beam size [$\arcsec$] & 25.8 & 25.8 \\ Velocity resolution [km s$^{-1}$] & 0.3 & 0.3 \\ Typical noise level in $T_{\rm MB}$ [K] & 0.7 & 0.2 \\ \hline \end{tabular} \end{table*} \begin{figure} \begin{center} \includegraphics[width=90mm, angle=0]{./co_peak_all.eps} \caption{Peak intensity map in the $^{12}$CO ($J$=1--0) line in units of K ($T_{\rm MB}$). The data are from \citet{Shimajiri11} and \citet{Nakamura12}. A dashed box shows the $^{13}$CO ($J$=1--0) and C$^{18}$O ($J$=1--0) observing region. The $^{12}$CO data in FITS format are available at the NRO web page via http://www.nro.nao.ac.jp/~nro45mrt/html/results/data.html.} \label{co_peak_map} \end{center} \end{figure} This paper is organized as follows: In Sect. 2, the Nobeyama 45 m observations are described. In Sect. 3, we present the $^{13}$CO and C$^{18}$O maps of the Orion-A GMC and estimate the optical depths of the $^{13}$CO and C$^{18}$O gas and the column densities of these molecules. In Sect. 4, we discuss the variation of the ratio of the $^{13}$CO to C$^{18}$O fractional abundance in terms of the FUV radiation. In Sect. 5, we summarize our results. Detailed distributions of the filaments and dense cores and their velocity structure and mass will be reported in a forthcoming paper. | We have carried out wide-field (0.4 deg$^2$) observations with an angular resolution of 25.8$\arcsec$ ($\sim$ 0.05 pc) in the $^{13}$CO ($J$=1--0) and C$^{18}$O ($J$=1--0) emission lines toward the Orion-A GMC. The main results are summarized as follows: \begin{enumerate} \item The overall distributions and velocity structures of the $^{13}$CO and C$^{18}$O emission lines are similar to those of the $^{12}$CO ($J$=1--0) emission line. The $^{13}$CO velocity channel maps show a new shell-like structure, which is not obvious in the $^{12}$CO and C$^{18}$O maps. \item We estimated the optical depths and column densities of the $^{13}$CO and C$^{18}$O emission lines. The optical depths of $^{13}$CO and C$^{18}$O are estimated to be \tauAmin\ $<$ $\tau_{\rm ^{13}CO}$ $<$ \tauAmax\ and \tauBmin\ $<$ $\tau_{\rm C^{18}O}$ $<$ \tauBmax, respectively. The column densities of the $^{13}$CO and C$^{18}$O gas are estimated to be \NAmin\ $\times$ 10$^{\NAorder}$ $<$ $N_{\rm ^{13}CO}$ $<$ 3.7 $\times$ 10$^{17}$ and \NBmin\ $\times$ 10$^{\NBorder}$ $<$ $N_{\rm C^{18}O}$ $<$ 3.5 $\times$ 10$^{16}$ cm$^{-2}$, respectively. \item The abundance ratio between $^{13}$CO and C$^{18}$O, $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$, is found to be \Xmin\ -- \Xmax. The mean value of $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ of the the nearly edge-on PDRs such as the Orion Bar and DLSF are 16.5, which is a factor of three larger than the solar system value of 5.5. On the other hand, the mean value of $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ in the other regions is 12.3. The difference between the abundance ratios in nearly edge-on PDRs and the other regions are likely due to the different intensities of the FUV radiation that cause the selective photodissociation of C$^{18}$O. \item In the low column density regions ($N_{\rm C^{18}O}$ $<$ 5 $\times$ 10$^{15}$ cm$^{-2}$), we found that the abundance ratio exceeds 10. These regions are thought to be influenced by the FUV radiation from the OB stars embedded in the Orion-A GMC such as the NGC 1977, M 43, and Trapezium cluster as well as the interstellar UV radiation. \item To examine the influence of the beam filling factor in our observations on the abundance ratio of $^{13}$CO to C$^{18}$O, we estimated the beam filling factors for the $^{13}$CO and C$^{18}$O gas to exceed 0.8. After taking into consideration the uncertainties in the beam filling factor, we also found the high abundance ratio $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ over the Orion-A cloud, particularly toward the nearly edge-on PDRs. \item Even if we consider the lower excitation temperatures of $T_{\rm ex\ ^{13}CO}$ = 30 K and $T_{\rm ex\ C^{18}O}$ = 20 K, we come to the same conclusion that the abundance ratio $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ becomes high toward the nearly edge-on PDRs. \item We checked the robustness of our conclusions in the Orion-A GMC by varying $\phi_{\rm ^{13}CO}$, $\phi_{\rm C^{18}O}$, $T_{\rm ex\ ^{13}CO}$, $T_{\rm ex\ C^{18}O}$, $\tau_{\rm ^{13}CO}$, and $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$. To explain the mean value of 11.4 for the intensity ratio $R_{13,18}$ in our observed region, the lower limit of the $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ value should be (5.8 -- 11.5) times larger than the solar system value of 5.5. In addition, the $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ values in the nearly edge-on PDRs are most likely larger than those in the other regions. \item When studying the range of possible values of the beam filling factors and of the excitation temperatures, the conclusion remains valid that the $X_{\rm ^{13}CO}$/$X_{\rm C^{18}O}$ values are higher than solar throughout Orion A, and larger in the PDRs than in the diffuse medium. \end{enumerate} | 14 | 3 | 1403.2930 |
1403 | 1403.0618_arXiv.txt | The Magellanic Bridge is the nearest low-metallicity, tidally stripped environment, offering a unique high-resolution view of physical conditions in merging and forming galaxies. In this paper we present analysis of candidate massive young stellar objects (YSOs), i.e., {\it in situ, current} massive star formation (MSF) in the Bridge using {\it Spitzer} mid-IR and complementary optical and near-IR photometry. While we definitely find YSOs in the Bridge, the most massive are $\sim10 M_\odot$, $\ll45 M_\odot$ found in the Large Magellanic Cloud (LMC). The intensity of MSF in the Bridge also appears decreasing, as the most massive YSOs are less massive than those formed in the past. To investigate environmental effects on MSF, we have compared properties of massive YSOs in the Bridge to those in the LMC. First, YSOs in the Bridge are apparently less embedded than in the LMC: 81\% of Bridge YSOs show optical counterparts, compared to only 56\% of LMC sources with the same range of mass, circumstellar dust mass, and line-of-sight extinction. Circumstellar envelopes are evidently more porous or clumpy in the Bridge's low-metallicity environment. Second, we have used whole samples of YSOs in the LMC and the Bridge to estimate the probability of finding YSOs at a given \hi\ column density, N(HI). We found that the LMC has $\sim3\times$ higher probability than the Bridge for N(HI) $>10\times10^{20}$~cm$^{-2}$, but the trend reverses at lower N(HI). Investigating whether this lower efficiency relative to HI is due to less efficient molecular cloud formation, or less efficient cloud collapse, or both, will require sensitive molecular gas observations. | The dependence of star formation on environment is fundamental in both the nearby and distant universe. Star formation is often enhanced in galaxies undergoing interaction or merger \citep{LT78}. This enhancement can produce global starbursts such as Arp 220, though more frequently observed are local concentrations of star forming regions \citep{BNetal03}. Although the overall star formation rate (SFR) can be enhanced, it is not clear how such physical conditions affect the star or cluster mass distribution. The higher pressure and density environment of mergers might encourage the preferential formation of {\it massive} stars and clusters \citep[e.g.,][]{EE97}. On the other hand, the increased turbulence in the interstellar medium (ISM) can result in larger gas dispersions, hampering the formation of giant molecular clouds and hence the massive clusters formed within \citep[e.g.,][]{SC97}. As massive stars and clusters are the energy source of the ISM and in turn affect the evolution of their host galaxies, it is important to understand their formation in a variety of environments that are different from our Galaxy. In addition to the influence of dynamical interactions, the formation of massive stars may depend on metallicity. A lower dust abundance and greater permittivity to ultraviolet radiation of the ISM is expected to affect pre-formation gas dynamics, as well as cooling and feedback from massive young stellar objects (YSOs) \citep{Poetal95}. It is thus critical to understand massive star formation in a low-metallicity, dynamically disturbed environment, to interpret similar situations in the early universe. Furthermore, to understand the detailed physics of the process; the geometric, morphological and temporal relationships of the molecular and dust components with the forming stellar population, we must spatially resolve the relevant structures, i.e., molecular clouds and individual young massive stars. Located between the Large and Small Magellanic Clouds (LMC and SMC) at a distance of $\sim$ 50--60 kpc \citep[e.g.,][]{HHH03,TVetal10}, the Magellanic Bridge (hereafter the Bridge) is the closest tidal system, and one of the few where clusters and interstellar structures can be resolved and studied in detail. The Bridge was first identified in an \hi\ survey \citep{Hietal63} and its production has been suggested to be the result of a recent close encounter between the LMC and SMC $\sim$ 200 Myr ago \citep[e.g.,][see also \citealt{BGetal10}]{GSF94}. The Bridge's tidal environment together with its significantly low metallicity $\sim$ 1/5--1/8 $Z_\odot$ \citep{RWetal99,LJetal05} provide an excellent laboratory to study massive star formation under such physical conditions. Furthermore, the Bridge's high \hi\ mass $\sim 1.5\times10^8 M_\odot$ and substantial \hi\ column density (N(\hi )) up to $\sim 3\times10^{21}$ cm$^{-2}$ \citep{Muetal03b} make it a promising site to search for newly formed massive stars. This high \hi\ mass also qualifies the Bridge as a potential region to develop into a dwarf galaxy \citep[\hi\ mass ranging from $10^6-10^9 M_\odot$,][]{MF99}, providing insight into their development and evolution as well. Several studies have found evidence of stars in the Bridge less than $\sim$100 Myr old, which if the 200--300 Myr formation timescale is correct must have formed {\em in situ} in tidal gas. \citet{HJ07} analyzed optical color-magnitude diagrams to derive a star formation history beginning $\sim 200-300$ Myr ago, with two distinct episodes $\sim$ 160 and 40 Myr ago. Studies specifically targeting massive stars provide more compelling evidence of {\em in situ} star formation. The handful of large \ha\ shells and several small \hii\ regions \citep{MJ86,MP07} in the Bridge are most likely formed by massive stars. One of the large \ha\ shells, DEM\,S\,171, has been suggested to be ionized by one or more O-type stars or blown by a supernova explosion \citep{MJ86,Gretal01}, though for other \ha\ shells and regions the underlying stellar population is not known. A population of blue stars are also identified in the Bridge using broadband $BV$ or $BVR$ photometry, with estimated ages ranging from as young as $\sim$ 10--25 Myr to $\sim 100$ Myr \citep{Iretal90,DI91,DB98}. Nevertheless, sufficient uncertainty in the timescales, and no clear association with natal gaseous material, leave open the possibility that these massive stars actually formed in the SMC body and were stripped out along with the gas. Pre-main-sequence stars identified in the Bridge would require an order of magnitude shorter timescale and it would be very hard to argue against {\em in situ} formation. A recent near-infrared (IR) $JHK_s$ survey of the Bridge finds Herbig Ae/Be (HAeBe) candidates with ages possibly down to $\sim 2$ Myr, but this color-selected candidate list is of moderate reliability \citep[only $\sim 40$\% likely bona-fide HAeBe,][]{Nietal07} and requires spectroscopic confirmation. Recent {\it Spitzer Space Telescope} imaging observations in the mid-IR have enabled the detection of individual massive YSOs in the LMC and SMC \citep[e.g.,][]{CYetal05,WBetal08,GC09,BAetal07,Sewiloetal13}. Follow-up {\it Spitzer} spectroscopic observations further confirms a $> 95$\% reliability rate in identifying massive YSOs in the LMC using our method based on examination on multi-wavelength spectral energy distributions (SEDs) and images \citep{CCetal09,GC09,Seetal09}. In this paper we present a similar inventory of massive YSOs in the Bridge, compare their properties and distribution to the molecular clouds, and probe a causal relationship between the initial condition (gas) and the end product (stars) in the most direct way. With the knowledge of current and recent stellar content and expected stellar energy feedback, it is then possible to assess if the star formation is triggered and its relative strength to that from spontaneous processes, and further estimate how the star formation efficiency (SFE) of a molecular cloud varies with time. Comparisons between the SFEs of clouds in the Bridge to those in a variety of metallicities and galactic environments such as the Galaxy or the LMC then allow us to probe the effect of environment on massive star formation. As part of the the {\it Spitzer} survey of the SMC \citep[SAGE-SMC,][]{GKetal11}, the high N(\hi ) portion of the Bridge (where N(\hi ) = 2--27$\times10^{20}$ cm$^{-2}$ with an average = 10$\times10^{20}$ cm$^{-2}$) has been mapped fully in the {\it Spitzer} bands from 3.6 to 160 \um\ (Figure~\ref{fig:hiimg})\footnote{The region of our study is located in the western part of the continuous stellar bridge between the LMC and SMC \citep{Iretal90}. It also appears to be extending from the SMC Body and Wing and hence has been referred as ``the SMC Tail'' \citep[e.g.,][]{GKetal09}. For simplicity we call this region the Bridge.}. Molecular clouds have been detected via CO J=1-0 emission \citep{Muetal03a,MNetal06}, providing an excellent opportunity to investigate if the formation of massive YSOs depends on physical conditions of the clouds. To study the current massive star formation in the Bridge, we have used {\it Spitzer} mid-IR observations and archival catalogs and data in the optical and near-IR wavelengths. The paper is organized as follows: the observations and data reduction are described in Section 2, the identification of YSO and HAeBe candidates is reported in Section 3, the derivation of physical properties of these candidates is detailed in Section 4, the properties of massive star formation in the Bridge is discussed in Section 5, and a summary is given in Section~6. | We have analyzed the distribution of young stellar objects (YSOs) in the Magellanic Bridge, the nearest tidally disturbed low-metallicity system. We selected YSO candidates from {\em Spitzer} photometry obtained for the SAGE-SMC Legacy program \citep{GKetal09}, using a combination of color and SED cuts, building on the techniques of \citet{WBetal08} and \citet{GC09}. We present a list of high-confidence brighter massive YSOs, and a list of somewhat fainter YSO candidates which may suffer from some contamination by unresolved background galaxies. We fit each source's SED with dust radiative transfer models from \citet{RTetal06} to constrain physical properties including central mass, envelope mass, and accretion rate. We find only weak correlation between the evolutionary ``Stage'' implied by the SED fit (high accretion rate normalized to central source mass) and our ``Type'' classification, in which sources with redder SEDs and fainter parsec-scale interstellar diffuse emission are presumed to be less evolved. The differences are likely due to a combination of a not extremely tight relation between circumstellar environment and protostellar evolutionary stage (which affects the ``Type''), and the fact that not all warm circumstellar dust contributing to the FIR SED may actually accrete onto the protostar (which affects the ``Stage''). We also analyzed {\em Spitzer} photometry for a list of NIR-derived candidate HAeBe stars \citep{Naetal05} in the Bridge, and find that about half of the HAeBe candidates show no significant MIR excess, and the other half we model as Stage II, or more evolved, YSOs. All of these findings are consistent with our similar analysis of LMC star formation regions. YSOs in the Bridge do show one particular contrast to those with similar mass in the higher metallicity LMC. Sources at the same evolutionary stage (derived primarily from the FIR flux and derived circumstellar dust mass) are brighter at optical wavelengths in the Bridge compared to the LMC. One explanation is that the circumstellar envelopes are more permeable or clumpy at low metallicity, allowing more short-wavelength radiation to escape for the same mass of warm clumps. Another particularity of the Bridge arises from the comparison of MIR-selected YSOs and NIR-selected HAeBe candidates -- the most massive YSOs are less massive than the most massive HAeBe candidates. This is consistent with the star formation in the Bridge becoming less vigorous with time, i.e. forming less massive stars than in the past. Of particular interest is star formation activity at the locations where CO 1-0 has been detected by \citet{MNetal06}. All CO 1-0 clouds except their lowest signal-to-noise detection are associated with MIR YSOs. Many of the clouds contain multiple YSOs, and several show evidence of multiple generations of star formation and possible triggering, for example a large HII region with MIR YSOs on its rim. Many of the YSOs in clouds are also closely associated with UV-bright main-sequence intermediate to massive stars, further supporting heterogeneity and non-coeval formation. We compared the overall star formation rate and efficiency in CO-detected molecular clouds to the atomic and molecular gas column. Global measures of star formation (integrated 24$\mu$m luminosities) fall below the rates predicted by the total gas column density and the S-K relation. However, detailed analysis of the YSO content brings the measured rates in closer agreement with gas-based predictions -- this agrees with results in LMC regions and may be a result of incomplete sampling of the stellar mass function, which systematically reduces integrated H$\alpha$ and 24$\mu$ because of the steep stellar mass-luminosity relation. When the entire Bridge area is considered (in contrast to only the molecular clouds), the star formation rate predicted from the relatively high HI mass is significantly higher than the total star formation that we detect. This is consistent with very inefficient molecular cloud formation efficiency in this environment, but relatively efficient star formation in molecular clouds that can form. | 14 | 3 | 1403.0618 |
1403 | 1403.2047_arXiv.txt | {Recently, more than 100 Wolf-Rayet and OB stars were identified in the Galactic center. \nnn{About a third of these sources} are not spatially associated with any of the known star clusters in this region. We probe the distribution of drifted sources in numerical models of the massive clusters in the Galactic center and compare it to the observed distribution of isolated massive sources in this region. We find that stars as massive as $100\,M_\odot$ drift away from the center of each cluster by up to $\sim\,60$ pc using the cluster models. \nnn{Our best model reproduces $\sim60\%$ of the known isolated massive stars out to $80 \,\mathrm{pc}$ from the center of the Arches cluster. This number increases to $70-80\%$ when we only consider the region of $\sim 20$ pc from the Arches cluster. }} | Galactic nuclei are ideal laboratories to investigate star formation in extreme conditions such as a strong tidal field, high UV radiation and a strong magnetic field. The only galactic nucleus where we can resolve the stellar population into individual stars is the center of our Galaxy at a distance of $\sim8.0\,\mathrm{kpc}$ (\citealt{ghez2008}, \citealt{gillessen}). However, the conditions for star formation and the dynamics of this region yet need to be understood as it harbors dense molecular clouds, a high star formation rate per unit volume, and the largest concentration of massive stars and star clusters in the Milky Way (e. g. \citealt{1996ARA&A..34..645M}; \citealt{2002ASPC..285..381F}; \citealt{2007A&A...467..611F}). The GC region hosts three starburst clusters with masses in excess of $\sim 10^4\,M_{\odot}$ and core radii of $\sim 0.15-1\,\mathrm{pc}$ (\citealt{1993ApJ...407L..77E}; \citealt{figer2002}; \citealt{Espinoza}). These three compact and massive clusters are the Young Nuclear Cluster surrounding the supermassive black hole, as well as the Quintuplet and the Arches clusters. Recent observations of isolated sources in the GC region revealed that, similar to these three clusters, the field stars in this area encompass many massive sources (\citealt{2011MNRAS.417..114D}; \citealt{Mauerhan2010_main}; \citealt{Mauerhan_matching}; \citealt{Wang2010}). A population of distributed very massive Wolf-Rayet stars with initial masses in excess of $20-40\,M_{\odot}$ were detected within a few dozen of pc from the super-massive black hole, Sgr~A*, by X-ray observations, accompanied by spectroscopic studies and Paschen-$\alpha$ (Pa$\alpha$) narrow-band imaging with HST (\citealt{Wang2010}; \citealt{2011MNRAS.417..114D}). Up to now, more than 100 Wolf-Rayet stars and O supergiants have been spectroscopically identified in the Galactic center region \citep{Mauerhan2010_main}, including the known cluster members. As about a third of these sources are located outside the three massive starburst clusters, they were suggested to provide evidence for {\sl isolated} high-mass star formation in the GC (\citealt{2011ASPC..439..104D}, \citealt{2013arXiv1309.7651O}). Observations of massive stars in the solar neighborhood show that generally massive stars form in groups and associations (\citealt{2003ARA&A..41...57L}; \citealt{2007ARAA..45..481Z}; \citealt{2008AA...490.1071G}). But it is not clear if we can generalize these findings to different galactic environments. On the other hand, dynamical evolution of stellar populations in the GC region can become dramatic under the strong effect of the GC tidal field. If so, dense and massive clusters like the Arches and Quintuplet can shape the distribution of the field stars in the region. In fact, the Arches and Quintuplet clusters are observed to be already mass segregated at ages of $2-6\,\mathrm{Myr}$. A recent study by \cite{benjamin} showed that the Quintuplet cluster at an age of 3-5$\,\mathrm{Myr}$ exhibits a flat mass function slope of $ -1.68 \pm 0.1$ in the cluster center compared to the standard \cite{Salpeter} initial mass function (IMF) of $-2.3$. A similarly flat mass function was found in the central region of the Arches cluster, but the slope increases substantially towards larger radii (\nnn{\citealt{1999ApJ...525..750F}}, \citealt{stolte_2005}, \citealt{Espinoza}, \citealt{habibi}). A study by \cite{Harfst2010} implemented N-body simulations of the Arches cluster to investigate the internal dynamical evolution of the cluster. By comparing their models to the observational data from the central 0.4 pc of the Arches cluster, they could constrain the initial conditions and construct a dynamical model of the Arches cluster that best represented the central stellar mass distribution, hereafter the best-fitting model of the Arches cluster. From this model, a steep increase of the stellar mass function slope as a function of the cluster center distance was predicted. The Arches cluster was later studied by \cite{habibi} to a larger radius of $\sim 1.5\,\mathrm{pc}$. Observations of the high-mass part of the mass spectrum in this study ($M >\sim 10M_{\odot}$) revealed a depletion of massive stars in the cluster outskirts. % In this previous study, we compared the measured slope of the mass function to the slopes predicted by the N-body simulations in different annuli out to the tidal radius of the Arches cluster. % This comparison showed that the Arches cluster exhibits characteristics of a normal, i. e. with a normal initial mass function, but dynamically evolved cluster. The dynamical evolution of the Arches and Quintuplet clusters, however, not only changes the distribution of stars inside the clusters, it can also change the distribution of field stars in the GC region. Through gravitational interactions between stars in dense and compact clusters stars can accelerate to become runaways (e.g. \citealt{1967BOTT....4...86P}; \citealt{1986ApJS...61..419G}). Moreover, the dynamical evolution of clusters under the influence of the Galactic tidal field leads to the formation of tidal arms. These tidal structures are mostly observed for globular clusters which evolve for many Gyrs (e.g. \citealt{2001ApJ...548L.165O}). As massive clusters in the GC dissolve within a few Myrs (\citealt{2000ApJ...545..301K}; \citealt{2001ApJ...546L.101P}), dynamical evolution under the effects of the strong tidal field of the GC leads to the formation of extended tidal structures during shorter timescales. These tidal structures, in turn, can significantly contribute to the field stars in the GC region. In this study, we analyze the best-fitting Arches model presented by \cite{Harfst2010}, extended to incorporate the effect of the Galactic center tidal field, to investigate the contribution of the Arches and Quintuplet clusters to the observed population of isolated massive stars detected by \citep{Mauerhan2010_main}. This paper is organized as follows: In Sect. \ref{sec: data}, a summary of observational studies to detect massive stars in the GC region, together with our criteria to construct an observational reference sample, is presented. In Sect. \ref{sec: simul}, we describe the computational methods, a grid of different models based on distinct physical assumptions, and a method to find the best-matching model to reproduce the observations. In Sect. \ref{sec: results}, we analyze this best-matching model to derive the spatial distribution of the massive drifted sources compared to the observed population. The velocity distribution of the drifted sources from the modeled clusters and their expected spatial and mass distributions are also predicted in this section. A summary of our findings is presented in Sect.\ref{sec:conc}. | \label{sec:conc} In this study we present N-body simulations of the Arches cluster to create combined models of the Arches and Quintuplet clusters. The population of ejected and drifted sources from the two clusters is compared to the HST/NICMOS Paschen-$\alpha$ (Pa$\alpha$) survey of the Galactic center which detects the distribution of young massive stars in the GC region \citep{Mauerhan2010_main}. Our study can be summarized as follows: \begin{enumerate} \item We construct different combined models of massive sources outside the Arches and Quintuplet clusters assuming different ages for the Quintuplet cluster and distinct values for the initial mass of a WR progenitor (see Table \ref{h_d}). We compare these models to the observed population of massive stars presented by \cite{Mauerhan2010_main} employing a method which calculates the histogram difference between spatial distributions of stars. Among all the constructed models, the model which assumes an age of 5 Myr for the Quintuplet and an initial mass of $40 M_{\odot}$ for a WR progenitor is the most similar to the spatial distribution of the observed isolated high-mass stars. \item The strong tidal field of the GC potential results in extended tidal arms for both the Arches and Quintuplet clusters at their current age. In the best-matching model, tidal arms of the Arches cluster stretch out to 20 pc in each direction on the plane of the sky and along the Galactic plane, while the tidal arms of the Quintuplet cluster extend out to 65 pc. The observations of massive stars in the GC region \citep{Mauerhan2010_main} reveal two strips along the Galactic plane with a prominent gap along the direction of the galactic poles. The projected positions of the tidal structures of the two clusters closely reproduce this observed distribution (see Fig. \ref{dens_map}). \item The observed massive sources outside the three clusters, including the young Nuclear cluster in the Pa$\alpha$ survey area is compared to the models \nnn{ using two different methods. First, comparing histograms of the spatial distributions of observed and simulated stars shows that the best-matching model reproduces $80\%$ of the observed population out to $21\,\mathrm{pc}$ and $67\%$ of the observed population out to $80 \,\mathrm{pc}$ distance with reference to the center of the Arches cluster. Second, we create a density map of observed isolated massive stars using Voronoi diagrams. The constructed density map allows us to probe the probability of observing one or more stars in each Voronoi cell assuming our best-matching model. For 62\% of the observed isolated massive stars, at least one of the ten random realizations of our model predicts a star that can explain the observed star. This number increases to 72\% when we only consider the Voronoi cells within the central 20 pc from the center of the Arches cluster.} The sources that cannot be explained as originating from the Arches and Quintuplet are located at large distances from the tidal tails. On the other hand, the best-matching model predicts $20\%$ more massive stars outside the clusters, when we perform the comparison between the best-matching model and the list of observed WR stars only, i.e., excluding the less complete sample of OB sources (see Fig \ref{histo}). \item The best-matching model predicts 26 massive stars outside the clusters, compared to the 35 observed massive stars outside the three clusters in the Pa$\alpha$ survey region. According to our best-matching models the majority of the simulated massive sources are located close to the tidal structure of the clusters, while the list of 35 observed massive stars includes 27 sources close to the tidal arms and 8 sources which are not close to the tidal structure of the Arches and the Quintuplet cluster. Histograms of the spatial distribution of the observed sources display a major characteristic peak at $\sim 15 \, \mathrm{pc}$ from the center of the Arches cluster. This peak is well reproduced in the best-matching model (see Fig. \ref{histo}). However, two minor peaks at the distances of $\sim 40\, \mathrm{pc}$ and $\sim 60\, \mathrm{pc}$ from the Arches center are absent in the model. Possible origins of these sources are supernova kicks or dynamical ejections involving tight initial binary systems. In this case, these additional WR stars could also have emerged from the Arches and Quintuplet clusters. Currently we can also not exclude the possibility that these high-mass stars might have formed in smaller clusters or in isolation. Finally, a deviation in particular in the orbit of the Quintuplet cluster could also give rise to the remaining differences between the observed and the simulated spatial distribution of high-mass stars in the GC. \item According to our models the projected tidal arms of the Quintuplet cluster at the age of 5 Myr extends out to $60\,\mathrm{pc}$ and reaches to the Sagittarius B2 region. This implies that the evolved massive stars observed in projection toward this region might originate from the tidally drifted sources from the Quintuplet cluster. \item The tidal structure of both clusters form as a result of velocity variation in the cluster. The velocity variation along the tidal arms of the Arches cluster reaches $50\,\mathrm{km\,s}^{-1}$. This value is as high as 140 $\mathrm{km\,s}^{-1}$ for the Quintuplet cluster, that will be detectable in a proper motion diagram of high precision astrometric studies of the Quintuplet cluster. \item The trajectories of the sources in different mass ranges in our models show that the tidal drifting of the cluster stars by the GC potential is an effective process which causes the stars to recede out to $70\,\mathrm{pc}$ from the center of the Arches cluster. Both massive, $M>40 \,M_{\odot}$, and intermediate mass stars, $10 M_{\odot}< M < 20 \,M_{\odot}$, follow a similar pattern; they gain energy in the center of the cluster which causes these sources to exceed the cluster's escape velocity. This suggests that the massive and intermediate-mass stars are evolved from the clusters by the same dynamical processes into tidal tails. The extended radial coverage of the high-mass stars inside the tidal tails implies that up to $80\%$ of the isolated observed WR population can be explained by cluster stars. \end{enumerate} | 14 | 3 | 1403.2047 |
1403 | 1403.4088_arXiv.txt | Variability in emission lines is a characteristic feature in young stars and can be used as a tool to study the physics of the accretion process. Here we present a study of H$\alpha$ variability in 15 T Tauri and Herbig Ae stars (K7\,-\,B2) over a wide range of time windows, from minutes, to hours, to days, and years. We assess the variability using linewidth measurements and the time series of line profiles. All objects show gradual, slow profile changes on time-scales of days. In addition, in three cases there is evidence for rapid variations in H$\alpha$ with typical time-scales of 10\,min, which occurs in 10\% of the total covered observing time. The mean accretion-rate changes, inferred from the line fluxes,are 0.01\,-\,0.07\,dex for time-scales of $<1$\,hour, 0.04\,-\,0.4\,dex for time-scales of days, and 0.13\,-\,0.52\,dex for time-scales of years. In \cite{2012MNRAS.427.1344C} we derived an upper limit finding that the intermediate (days) variability dominated over longer (years) variability. Here our new results, based on much higher cadence observations, also provide a {\it lower} limit to accretion-rate variability on similar time-scales (days), thereby constraining the accretion rate variability physics in a much more definitive way. A plausible explanation for the gradual variations over days is an asymmetric accretion flow resulting in a rotational modulation of the accretion-related emission, although other interpretations are possible as well. In conjunction with our previous work, we find that the time-scales and the extent of the variability is similar for objects ranging in mass from $\sim 0.1$ to $\sim$5$\,M_{\odot}$. This confirms that a single mode of accretion is at work from T Tauri to Herbig Ae stars -- across a wide range of stellar masses. | Accretion is a vital and central process in the formation of low mass stars, controlling the flow of angular momentum and mass from the interstellar medium (ISM) onto the star. The accretion history can have a large effect on the long-term properties of the stellar system, such as its luminosity, radius, mass, and rotation rate \citep{2009ApJ...702L..27B}. It is also thought to be strongly connected to the mechanisms of wind and jet launching \citep{2007prpl.conf..231R}, and of disc clearing and evolution \citep{2005ApJ...625..906M}. In order to understand these processes we need to gain a full understanding of accretion. The current model for this process is magnetospheric accretion where the stellar magnetic field threads the disc, and the material in the disc falls along the field lines onto the surface of the star \citep{1991ApJ...370L..39K}. However, this simple model does not explain one of the defining features: its variability. The accretion flows and shocks emit continuum emission from the ultraviolet (UV) to the infrared (IR) as well as a number of emission lines (e.g., H$\alpha$, Ca\,II, He\,I), which, when observed are all found to be variable on time-scales from hours, to weeks, and to years \citep{2005ApJ...626..498M,Nguyen09,2006ApJ...638.1056S}. This variability can and has been used to probe the inner regions of these accreting systems, to find the source of the variations and to provide stringent constraints on the nature of the accretion process itself \citep{2012ApJ...750...73D}. The time-scale of accretion variations and the magnitude of the variations will depend on their source, so by monitoring these accreting systems over different time-scales, we can identify where these variations originate. This has been done very successfully for individual objects such as AATau \citep{2003A&A...409..169B} and V2129 \citep{2012A&A...541A.116A}, but it has yet to be done systematically for a larger sample. Our previous observing program LAMP (Long-term Accreting Monitoring Program) was designed to systematically explore the possible origin of accretion variability by monitoring H$\alpha$~and the Ca\,II triplet. Two successful runs took place over two months, with a total baseline of 15 months. Over this period 25 young stellar objects (YSOs) in Chameleon I were observed, each object 12 times. Among these targets we found 10 accretors, which covered a spectral range of G2\,-\,M5.75. All the accretors showed variations in their accretion signatures. The average accretion spread as calculated from the H$\alpha$ equivalent width (EW), was found to be 0.37\,dex \citep{2012MNRAS.427.1344C}. Although there are long term accretion rate variations in the sample, the amplitude of variations reached a maximum or are within 70\% of the maximum after 25 days. This was the shortest time-scale in our sample, which indicates the dominant cause of this variability is occurring on times scales of 2\,-\,3 weeks, or less. This result is backed up by other studies, such as a short-term monitoring covering time-scales up to $\sim$ 48 hours of $\sim$ 29 sources in Chameleon II, which defined the spread in accretion rates as derived from H$\alpha$ to be 0.2\,-\,0.6 dex \citep{2012A&A...547A.104B}. Also, \cite{Nguyen09} found a very similar accretion-rate spread over time-scales of days and months. These and our previous results rule out origins of variability such as those due to a wind form the disc or large-scale instabilities in the accretion disc, as these would occur on much longer time-scales of months and years. These time-scales of $\lesssim$~ 2\,-\,3 weeks are close to the rotation period of these kinds of objects (1\,-\,10 days \cite{2007prpl.conf..297H}), which suggests that what we are observing could be a rotational modulation of the accretion flow. If there is even a slight offset between the rotation and the magnetic field axis (2\,-\,5\degree) it has been shown that an asymmetric accretion flow will form \citep{2003ApJ...595.1009R}. As this rotates with the star, different parts of the accretion flow will become visible, which will change the accretion signatures that we observe. A second possible explanation for variations on these time-scales concerns the existence of instabilities in these systems, either in the magnetic field, or in the inner disc \citep{2013MNRAS.431.2673K, 1998ApJ...492..323G}. These will cause short term, stochastic variations occurring on the time-scales as short as hours. If we wish to distinguish between these two possible causes of accretion variations on short time-scales, we need to monitor these systems on time-scales close to the rotation period (multiple days) as well as on time-scales of hours. Just such a test can be performed using the high signal-to-noise, high-cadence linear spectropolarimetry data-set of Vink et al. (2005) on a sample of T Tauri and Herbig Ae stars. This data-set will allow us to constrain whether the mass-accretion rate variations indeed involve a slow and gradual variation in accretion emission across the rotation period, or if they concern stochastic and rapid variations. Our combined sample of T Tauri and Herbig Ae stars probing the time-scales of minutes and hours, with LAMP probing the time-scales of days, months, and years, we can constrain the dominant accretion variations. Using the H$\alpha$ emission as an accretion indicator, and to measure accretion rates, we compare the variations found on {\it all time-scales}. Furthermore, we can extend our work to higher stellar masses (up to 2-3 $M_{\sun}$ for Herbig Ae stars). This way, we can test whether the accretion process in Herbig Ae stars might be similar to that in lower mass T Tauri stars, as suggested by \citep{2001A&A...371..186N,2002MNRAS.337..356V,2003A&A...406..703V,2006A&A...459..837G,2004ApJ...613.1049E}. Whilst magnetic fields have indeed been claimed in Herbig Ae stars \citep{2004A&A...428L...1H,2005A&A...442L..31W}, the fact that the field incidence in Herbig Ae stars is similarly low to that in post-main sequence objects (e.g. \cite{2012sf2a.conf..401A}) may cast doubt on an extension of the magnetospheric accretion scenario towards higher masses. Therefore the issue of the fundamental accretion process in Herbig Ae stars remains open, and can be tested in this paper. The paper is laid out as follows: In Sect.\ref{sec:sample_observations} we discuss the chosen sample, the observations and reduction, Sect. \ref{sec:behaviour} presents the data and the variations present in the H$\alpha$ emission, Sect. \ref{sec:origin} addresses the nature of the H$\alpha$ emission, Sect. \ref{sec:accretion_rates} discuses the derivation of accretion rates and Sect. \ref{sec:discusion} deals with possible causes of the variations observed. Individual sources and their variations are presented in the Appendix. | \label{sec:discusion} This paper sets out to constrain the variations in the H$\alpha$ emission of accreting stars. The aim here was to isolate the relative variability of the mass accretion rate as derived from the emission in each star. Thus, we do not take into account any systematic uncertainties from stellar parameters or line luminosity conversions (which will effect the comparison of the accretion values measured for different stars). Only the accuracy in our line measurements are considered in order to ascertain whether the observed variations we measure are above the measurement errors. Though some of the variations may come from sources other than accretion such as winds or from secondary accretion effects such as continuum changes (see Sect.\,\ref{sec:origin}), these measurements allow us to put an upper-limit on the accretion rate changes as observed in the H$\alpha$ emission line. As previously mentioned, there are two distinct types of behaviour observed in this sample, \emph{slow variations} and \emph{rapid events}. In the following section possible origins of each are discussed. \subsection{Slow Variations: Rotational Modulation of the Accretion Rate} In the majority of the cases in this sample, small changes are seen in the profiles across the time-scale of our observation blocks. These variations are referred to as \emph{slow variations}. These occur in discrete wavelength ranges, and take the form of a gradual change in the profile emission. These variations do not translate to large accretion rate changes and on average they result in changes of the accretion rate derived from H$\alpha$ EW of 0.01\, - \,0.07\,dex. When the changes between different nights of observations are examined in this sample, the spread in accretion rate variations increases slightly to 0.04\,-\,0.4\,dex (see Table \ref{tab:accretion_rates2001and2003}). Comparing the difference in derived accretion rates for the objects with observations in 2001 and 2003 this spread does not increase from the individual observation periods remaining at 0.13\,-\,0.52\,dex. This is a strong indication that the time-scales of days are the dominant time-scales for these variations. Fig.\,\ref{fig:accretion_timescales_mean} shows a comparison of accretion rate variations on all the time-scales in the sample. For each object every accretion rate measurement is compared with every other accretion rate measurement for that object. In this way all the time-scales available in the sample can be exploited. The mean accretion rate is then plotted for each time bin for that object. In all cases but two, the variations reach a maximum after the first few days of observations. It shows that the dominant time-scale of variations in this sample is on the order of days. Two objects do show a rise in variations on the year time-scales, GW Ori and RY Tau. In the case of GW Ori, the sharp rise could be due to the lack of observations over consecutive days, and is probably not a real increase in accretion variations on the longer time-scales (see Fig.\,\ref{fig:accretion_timescales_mean}). For RY Tau, it could be a real rise in accretion variations, but it is more likely to be a single event observed at the end of the time-series. (See Fig.\,\ref{fig:accretion_timescales_all_pointsA} and Fig.\,\ref{fig:accretion_timescales_all_pointsB} for un-binned comparison of accretion variations and time-scales for all objects.) \begin{figure*} \centering \centering \begin{tabular}{ccc} \includegraphics[scale=0.4]{RYTAU_Mdot_mean_intervals.eps} & \includegraphics[scale=0.4]{GWORI_Mdot_mean_intervals.eps} \\ \includegraphics[scale=0.4]{ABAUR_Mdot_mean_intervals.eps} &\includegraphics[scale=0.4]{MWC480_Mdot_mean_intervals.eps} \\ \includegraphics[scale=0.4]{DRTAU_Mdot_mean_intervals.eps} &\includegraphics[scale=0.4]{TTAU_Mdot_mean_intervals.eps} \\ \includegraphics[scale=0.4]{RWAUR_Mdot_mean_intervals.eps} & \includegraphics[scale=0.4]{SUAUR_Mdot_mean_intervals.eps} \\ \end{tabular} \caption{ Mean differences in accretion rates~\lbrack Log(M$_{\odot}$ yr$^{-1}$)\rbrack~versus the time difference \lbrack days\rbrack. In this case we compare all accretion rates for each object, in order to cover all time-scales within the sample. The difference between all accretion rates was calculated, and the mean is plotted for each time bin. The time bins vary from object to object depending on the number of observations blocks. The same plot for all objects where the accretion rate differences are not binned can be seen in Fig.\,\ref{fig:accretion_timescales_all_pointsA} and Fig.\,\ref{fig:accretion_timescales_all_pointsB}.} \label{fig:accretion_timescales_mean} \end{figure*} The time-scales over which the accretion rate variations reach a maximum are comparable to the rotation periods within this sample, which all lie between 1\,$\sim$\,5 days (see Table \ref{tab:stellar_parameters}). These \emph{slow variations} in the profiles are consistent with what one would expect from a rotational modulation of the accretion rate. In the case of an asymmetric accretion flow, as the star rotates, different parts of the accretion flow will come into view, changing how much of the accretion flow is visible and what velocities within the flow are visible. Apart from three cases, no major changes are seen in the profiles within the single blocks of observations. These blocks represent time-scales of about an hour or less. In this sample it appears that the small, gradual changes we see within these observation blocks accumulate to larger variations between the different nights of observations. Our previous work on a sample of T Tauri stars in Chameleon supports this result. We monitored the H$\alpha$ variations over time periods of weeks\,-\,15 months, and also found that it was the short time-scales, of a few weeks or less, that were the dominant time-scales within the sample \citep{2012MNRAS.427.1344C}. In Chameleon, we found the average spread in accretion rates to be 0.37\,dex, which is very close to the variations we find in this sample of both T Tauri and Herbig Ae stars. Simulations of magnetospheric accretion have found that even a slight offset (2\,-\,5\degree) will result in an asymmetry in the accretion flow \citep{2003ApJ...595.1009R}. This is thought to be in the case for V2129 Oph \citep{2012A&A...541A.116A}. The authors used MHD simulations of the observed magnetic octupole and dipolar fields of V2129 Oph, and radiative transfer codes to reproduce the observed spectral line profiles. Earlier observations of the magnetic field found an offset between the octupole and dipole fields (15 and 25\degree) and the rotation axis \citep{2011MNRAS.412.2454D}. The modelling of these fields result in two ordered flows of material onto the star very close to the poles. The derived profile variations are similar in magnitude to the observed profiles, however the changes in the profile shape are not. At an inclination angle of 60\degree~to the viewer, the models result in a change in 8\,\AA~in H$\alpha$ EW across the rotation period. The changes that are observed in this sample over multiple days, are on the same order of magnitude as those modelled for V2129 Oph. For example AB Aur, RY Tau, SU Aur, BP Tau all have EW ranges of 5\,-\,15\,\AA~between multiple nights observations. RW Aur has a much larger spread of 35\,\AA, which may mean that this model of two rotating flows is probably not sufficient to explain all of the variations seen in this object. We can expect that rotationally-induced apparent accretion rate change will depend on the inclination of the systems to our line of sight. \citet{2006MNRAS.370..580K} showed through MHD simulations that as the inclination of an accreting system increases, the EW of the H$\alpha$ emission decreases. This is due to the fact that we see less of the accretion flow at higher angles. In these models, as the inclination increases from 10\degree~to 80\degree~the EW changes from 32\,\AA~to 21\,\AA. One can also expect to see more changes in the H$\alpha$ EW if the system is inclined to our line of sight, making it more likely that the accretion flow will move in and out of view. In the 2003 observations the three stars with the largest range in accretion rates in the sample are three of the most inclined systems. RY Tau, SU Aur, and RW Aur all have inclinations of 45$^{\circ}$ or over. DR Tau is also highly inclined, but does not show very large variations. However since the system is close to edge on, we may not have a full view of the accretion flow, and the disc may obscure a lot of the light coming from the accretion flows. The inclination angles are given in Table \ref{tab:stellar_parameters}. This assertion no longer holds true when the 2001 observations are considered. The three objects (RY Tau, SU Aur, and RW Aur) show much less variability than the other objects, whereas DR Tau shows some of the largest accretion-rate variations, and T Tau and MWC 480 with disc inclinations of $\sim$ 30\degree~also show large variations. This suggests that a picture of a stable rotating asymmetric flow is probably too simplistic to describe the full variations in these objects. Long term photometric monitoring of accreting objects also support the scenario of a more complicated accretion flow as irregular light curves have been observed in many accreting T Tauri stars \citep{1994AJ....108.1906H,2009MNRAS.398..873S,2004A&A...419..249S,2007A&A...461..183G}. Over the time-scales of years multiple different types of variations can been seen in one object. In many of the cases were variations occur on the time-scales of days, the simple explanation of a rotational modulation cannot explain the full behaviour. In their spectro-polarimetric observations of BP Tau, \citet{2008MNRAS.386.1234D} found strong signatures of rotational modulation in the accretion related emission lines. The narrow emission lines associated with accretion (He\,I, Fe\,II and narrow component of Ca\,II IR triplet) varied strongly with rotation period. However, the broad emission lines H$\alpha$, H$\beta$ (and also the broad component of the Ca\,II IR triplet) were found to vary on the level of 10\%\,-\,20\% with rotation, with the remainder of the variations coming from seemingly other sources. This could be explained by the narrow emission component originating close the base of the accretion flow, with the H$\alpha$ emission originating in the bulk of the accretion flow, which may be more sensitive to instabilities. Excess contributions in the H$\alpha$ emission from outflows could also play a role. The authors also suggest, that changes in the rotational modulation of the longitudinal magnetic field between the two observation periods in February and December 2006 implies that the large-scale field topology was distorted by variability in the system between the two periods. This suggests though rotational modulation accounts for the majority of the variations, there are other ongoing processes. There are many kinds of instabilities that can occur in the disc and accretion flow that may account for the changes in the variations observed over time, e.g.: the build up of material in the circumstellar disc \citep{2012MNRAS.420..416D}. \citet{2013MNRAS.431.2673K} show that in the case where Rayleigh-Taylor instabilities exist in the inner disc, unstable accretion flows can form. These flows change in size, shape, and numbers, meaning there is always a accretion column visible. This results in a constantly observed red-shifted absorption in the Balmer lines, more particularly in the higher Balmer lines such as H$\gamma$ and H$\delta$. This has been observed in RW Aur \citep{1994AJ....108.1056E}, where spectroscopic monitoring also confirmed the presence of an asymmetric accretion flow \citep{2001A&A...369..993P}, as well as for SU Aur \citep{1995ApJ...449..341J,1996A&A...314..821P}. It is possible that these instabilities exist in some if not all of the objects in this sample. Comparing the 2001 and the 2003 sample, some changes are seen in the H$\alpha$ profiles, but also differences in the derived accretion rates and their variations. These instabilities could account for changes in the H$\alpha$ emission if they change the form of the accretion flows. \subsection{Rapid Events}\label{sec:rapid_events} The \emph{rapid events} observed in this sample do not fit into the frame of rotational modulation. AB Aur, and to a lesser extent RY Tau and RW Aur, show significant variations in the profile over the time-scale of 1 hour. In the case of each star, these changes only occur during a single night of observations (For AB Aur see Fig.\,\ref{fig:rapid_changes}). A number of short term rapid variations have previously been discovered in objects in our sample and these are presented in the following paragraphs. Short term striking variations have been observed in the H$\alpha$ profile of AB Aur previously, where changes occurred in both the intensity and profile shape. \citet{1995A&A...298..585B} described the changes as occurring during each observing night across the emission line but especially in the absorption feature of the P-Cygni profile and the emission peak. They attribute these variations to the motion of circumstellar inhomogeneities. This is similar to the changes observed in the H$\alpha$ profile of AB Aur on night 3, where it oscillates between broad wings and low emission, and strong emission and with narrow wings. However in the case of \citet{1995A&A...298..585B}, their observations take place with hour separations, so they do not have the short term cadence that we have in this sample. Rapid variations have been seen in the H$\gamma$ profile of RW Aur, occurring on time scales as short as 10 mins. Over the course of the observations the central absorption component of the H$\gamma$ line increased in EW by a factor to 2 within an hour \citep{1982ApJ...256..156M}. These variations were not reflected in the H$\alpha$ emission, which is thought to be a result of the high optical thickness in the surrounding envelope. \citet{1982ApJ...256..156M} found these variations to be consistent with both flaring and accretion events. \begin{figure} \centering \begin{tabular}{cc} \includegraphics[scale=0.22]{profile_comaprison_day2.ps}& \includegraphics[scale=0.22]{profile_comaprison_day3.ps}\\ \end{tabular} \caption{A comparison between a sub-set of profiles on two separate nights of the 2003 observations for AB AUR. On the second night of observations in 2003, no significant change within the profile was observed. The minute of the hour in which the spectrum was observed is given to the right of the profiles in the corresponding colour. During the third night of observations there were large changes. Shown here is a sequence of spectra where the emission line falls in strength, and broadens, before returning to the initial strength.} \label{fig:rapid_changes} \end{figure} Variations in H$\alpha$ profile of RY Tau were found on the time-scales of 10\,-\,20 mins without variations in the star's brightness \citep{1975PZ.....20..153K}. Three separate nights observations took place covering time scales of a few hours each night. The profiles changed between nights, but in two out of three occasions they were stable throughout the night. The brightness of the star was lower on the night of the variations than on the other two nights, which is an indication that these variations came from magnetic activity \citep{1997A&A...324..155G}. With the short wavelength coverage in this sample and no simultaneous alternative observations, the behaviour seen in AB Aur, RY Tau or RW Aur cannot be properly defined, and the origin of these rapid variations is not clear. \citet{1994A&A...287..131G} argue that in the case of magnetospheric accretion short term variability is to be expected. The in-fall time-scale of gas towards the pole is less than one hour. Any instabilities that occur in the disc or the magnetic field at the point of their interaction will lead to a clumpy flow of material onto the surface. The changes in RY Tau and RW Aur take the form of a drop in emission across both lines, but not within the line centre, which would be more indicative of a flare (see Sect. \ref{sec:chromo_accretion}). However it is probable that the rapid events that are observed in RW Aur and RY Tau are due to a flare event. This is not the case with AB Aur. The \emph{rapid variations} that occur in the profile of AB AUR are unique in the sample. In no other object do we see these changes in emission line strength, width and surrounding absorption. It is possible that all the targets have these periods of \emph{rapid variations}. Out of the total 22.6 hours of observations, these \emph{rapid events} only take place with 3 observation blocks, which constitutes 2.4 hours or $\sim$ 10\% of our total observing time. This suggests they are not very common events and it rules out stochastic processes as the primary source of variations within the sample. However these stochastic events could be the cause of the \emph{rapid variations} we see in a small number and probably take the form of instabilities in the magnetic field \citep{1998AIPC..431..533G}, or inner disc \citep{2013MNRAS.431.2673K}. \subsection{Comparison between Herbig Ae and T Tauri sample} Herbig Ae stars are the intermediate mass equivalent of T Tauri stars and are thought to go through a similar process of accretion as T Tauri stars. As they are higher mass, they are shorter lived, but retain their circumstellar disc for long enough to accrete material from them onto their surfaces. Similar scalings of accretion rate to disc mass exist between T Tauri and Herbig Ae \citep{2012A&A...543A..59M}. They also show that the NIR colour excess trend is the same across the T Tauri to Herbig Ae mass range, which can be explained by the reprocessing of both the stellar and accretion luminosities by the inner disc. However there is inconclusive evidence whether Herbig Ae stars are host to strong magnetic fields. Under our current understanding of magnetospheric accretion these strong, stable fields are essential for maintaining a quasi-stable accretion flow. Within this work a large mass range is covered (up to $\sim$ 5 M$_{\odot}$) and similar variations are seen in all objects. One of the larger mass targets, MWC 480 shows an accretion rate spread of 0.012\,-\,0.061 over single observation blocks. Comparing to one of the smallest mass targets in the sample, DR Tau with 0.010\,-\,0.052, shows there is little difference between the two. Indeed, others have found similar accretion rate variations for Herbig Ae stars as is found in this work, and the LAMP sample. For example \citet{2012AN....333..594P} observed accretion variations of amplitude 0.4\,dex over the time-scales of 10 days for one target, while multiple targets showed variations of 0.3\,dex between consecutive days observations. From 38 Herbig Ae/Be stars, \citep{2011A&A...535A..99M} found a typical upper limit of accretion variations of 0.5\,dex over time-scales from days to months. Compiling the samples from both studies (this and the LAMP sample), of 10 low-mass T Tauri and 14 intermediate-mass T Tauri/Herbig Ae stars, suggests that it is the same process that produces the H$\alpha$ variations in T Tauri and Herbig Ae stars, across the entire mass range up to $\sim$ 5 M$_{\odot}$. This variability result is entirely consistent with earlier spectro-polarimetry surveys \citep{2005MNRAS.359.1049V,2003A&A...406..703V,2002MNRAS.337..356V} | 14 | 3 | 1403.4088 |
1403 | 1403.3727_arXiv.txt | One of the main goals of modern observational cosmology is to map the large scale structure of the Universe. A potentially powerful approach for doing this would be to exploit three-dimensional spectral maps, i.e.~the specific intensity of extragalactic light as a function of wavelength and direction on the sky, to measure spatial variations in the total extragalactic light emission and use these as a tracer of the clustering of matter. A main challenge is that the observed intensity as a function of wavelength is a convolution of the source luminosity density with the rest-frame spectral energy distribution. In this paper, we introduce the method of spectral deconvolution as a way to invert this convolution and extract the clustering information. We show how one can use observations of the mean and angular fluctuations of extragalactic light as a function of wavelength, assuming statistical isotropy, to reconstruct jointly the rest-frame spectral energy distribution of the sources and the source spatial density fluctuations. This method is more general than the well known line mapping technique as it does not rely on spectral lines in the emitted spectra. After introducing the general formalism, we discuss its implementation and limitations. This formal paper sets the stage for future more practical studies. | As spectral and imaging capabilities evolve, it is becoming increasingly common in astronomy to think of large three-dimensional data cubes, the specific intensity distribution of sources along many lines of sight that span a range of wavelengths. For studies of internal dynamics of galaxies or molecular clouds, the frequency dimension generally corresponds to elements at different line-of-sight velocities, while for cosmological studies the frequency dimension (ignoring peculiar velocities) corresponds to elements at different cosmological redshifts, and therefore at different distances. Cosmological large scale structure surveys routinely use the spatial distribution of extragalactic light emission or absorption as a tracer of the matter density to constrain cosmological models. Such surveys typically exploit a subset of the three-dimensional data cube of the aggregate intensity of extragalactic light as a function of direction on the sky and of wavelength\footnote{The caveat to this description is that it ignores polarization and time-domain information (and cosmological information beyond the electromagnetic spectrum).}, for instance by working with a catalog of positions of a set of bright objects. In this article, we instead consider the scenario where a full spectral map is available, such as could be obtained\footnote{In realistic applications, this map would of course still be limited by the sky coverage, wavelength range, and angular and spectral resolutions of the survey(s) providing the data.} from integral field spectroscopy (e.g.~HETDEX\footnote{http://hetdex.org} \citealt{hilletal08}) or narrow band imaging surveys (e.g.~J-PAS\footnote{http://j-pas.org/} \citealt{benitez09,molesetal10}, PAU\footnote{http://www.pausurvey.org/} \citealt{benitez09}, or the Alhambra Survey \citealt{molesetal08}). In other words, the data are treated at the level of a continuous three-dimensional map as opposed to a discrete catalog. The spectral intensity as a function of wavelength along a given line of sight in such a map is in general a superposition of the contributions of multiple sources (and absorbers). The goal is then to reconstruct the distribution of these sources as a function of redshift. The difficulty here is that, generically, sources emit over a wide range of wavelengths, so that the mapping from redshift to wavelength is degenerate. One common strategy to evade this degeneracy is to focus on objects that are bright enough that they stand out from the background and dominate the signal along a given line-of-sight (such as bright galaxies or quasars), i.e.~the ``catalog-based'' approach mentioned above. Depending on how well the redshift direction is sampled, this approach has led to measurements of 2-dimensional (angular projected) clustering (e.g.~\citealt{hauserpeebles73,efstmoody01,scrantonetal02,tegmarketal02,frithetal05,coorayetal10,donosoetal13}), ``2 + 1-dimensional'' clustering using photometric redshifts (e.g.~\citealt{padetal07,rossetal11,hoetal12,hoetal13}) and of full 3-dimensional clustering using spectroscopic redshifts (e.g.~\citealt{FKP,coleetal05,eisetal05,beutleretal11,parkinsonetal12,delatorreetal13,Anderson:2013zyy}). An alternative approach to reconstructing the distribution of sources is line mapping, where the signal from well-known spectral lines, e.g.~the HI 21-cm transition, allows a direct mapping between wavelength and redshift (e.g.~\citealt{changetal08,loebwyithe08,visloeb10,Chang:2010jp,pritloeb12,Switzer:2013ewa}). Here, we study the reconstruction of the source luminosity density in the more general case where we do not restrict ourselves to bright sources nor rely on spectral lines only. Instead, we consider reconstruction of the total luminosity density of all sources contributing to the extragalactic signal, optimally exploiting both the spectral and angular variations in the spectral map. While for simplicity we will assume the signal is described in terms of a single effective rest-frame spectral energy distribution (SED) shape, this method does not rely on prior knowledge of this SED form, but instead measures it from the data itself. Formally, our work stems from the following observation. Generally, the specific intensity, $I$, at a given observed wavelength, $\lambda$, can be written as the weighted sum along the line of sight of the contribution of sources of spectrum, $s$, at a a variety of cosmological redshifts: \begin{equation} \label{eq:int} I(\ln \lambda) = \int d z \, w(z) \, s(\ln \lambda,z), \end{equation} where $w(z)$ is a weight function proportional to the density of emitters. If we assume that the spectral energy distribution is non-evolving with redshift and write it in terms of the emitted SED $s_{rest}$, the specific intensity can be written as the convolution \begin{equation} \label{eq:conviso} I(\ln \lambda) = \int d \ln (1 + z) \, \left[(1 + z) w(z) \right] \, s_{rest}(\ln \lambda - \ln (1 + z)). \end{equation} The observation that the specific intensity is a pure convolution, where the rest-frame SED has been convolved with the redshift distribution, bears important implications that we explore in this work. Since both the SED and the redshift distributions are unknown quantities of cosmological interest, we study in particular how this mathematical structure can be fully exploited to reconstruct these quantities. Throughout this work, we will develop an explicit procedure for the spectral deconvolution of Eq.~(\ref{eq:int}) and will elaborate on its domain of validity. We will introduce a general formalism for single type sources and discuss the case of an emitter of known SED in Sec.~\ref{sec:known_sed}. In Sec.~\ref{sec:unknown_sed}, we discuss the use of angular fluctuations in the intensity map, taking advantage of the assumption of statistical isotropy to express the observables in terms of angular (cross-)power spectra. In Sec.~\ref{sec:rec_sed}, we then show how measuring fluctuations over a sufficiently large range of angular scales and wavelengths can potentially allow us to reconstruct both the a priori unknown SED and the underlying density fluctuations. We conclude in Sec.~\ref{sec:conclusion} and discuss in more detail the range of scales at which spectral deconvolution is feasible in Appendix \ref{sec:range validity}. | \label{sec:conclusion} We have introduced a formalism for measuring the projected two-point function of the luminosity density of extragalactic sources from spectral mapping data, i.e.~from a three-dimensional data cube of specific intensity as a function of wavelength and line of sight direction. This method does not rely on the use of spectral lines and works for an arbitrary rest-frame SED form. In fact, the source SED does not even need to be known in advance, but is reconstructed from the data itself. This spectral deconvolution technique makes use of both the mean intensity as a function of wavelength and its angular variations. The wavelength dependence of the rest-frame SED and the redshift dependence of the luminosity density of sources can be disentangled because the observed specific intensity is a convolution of these two quantities. This allows a straightforward spectral deconvolution in Fourier space. The spatial variations in the reconstructed luminosity density contain valuable information on the large scale clustering of matter. While one might expect this approach to only work if the (mean) rest-frame SED is known a priori, we have shown in Section \ref{sec:rec_sed} how to use the statistical isotropy of the observed signal to simultaneously measure the rest-frame SED and the clustering of the luminosity density when this is not the case. Specifically, by first Fourier transforming the spectral map with respect to $\ln\lambda$ (with Fourier conjugate $r$), and then considering the angular power spectrum for mostly transverse modes (multipoles $\ell >> r$), the term describing the clustering of extragalactic sources to a good approximation only depends on multipole $\ell$, while the rest-frame SED term only depends on the wavelength direction $r$. It is this separability of variables that allows for their independent reconstruction. The purpose of this paper has been to give a rather formal presentation of the method of spectral deconvolution, leaving more concrete explorations of how to apply the technique to realistic data for future work. A particularly strong assumption we have made throughout this work is that the extragalactic signal can be described in terms of a single effective rest-frame SED (up to a free normalization). It will be interesting to generalize the method to scenarios where multiple populations with distinct SED's need to be factored in. One additional motivation for studying the information content of three-dimensional spectral maps, beyond the fact that such data will be available from spectral mapping experiments, is that, as discussed in the Introduction, such maps provide a unifying description for a large range of cosmological surveys, with different probes accessing different subsets of the data cube. It would thus be useful to build a more general understanding of all the information that can in principle be extracted from the full data cube. This can then be a guide towards identifying how to optimally exploit these data in the future. | 14 | 3 | 1403.3727 |
1403 | 1403.4936_arXiv.txt | We conduct the first microlensing simulation in the context of planet formation model. The planet population is taken from the Ida \& Lin core accretion model for $0.3M_\odot$ stars. With $6690$ microlensing events, we find for a simplified Korea Microlensing Telescopes Network (KMTNet) the fraction of planetary events is $2.9\%$ , out of which $5.5\%$ show multiple-planet signatures. The number of super-Earths, super-Neptunes and super-Jupiters detected are expected to be almost equal. Our simulation shows that high-magnification events and massive planets are favored by planet detections, which is consistent with previous expectation. However, we notice that extremely high-magnification events are less sensitive to planets, which is possibly because the 10 min sampling of KMTNet is not intensive enough to capture the subtle anomalies that occur near the peak. This suggests that while KMTNet observations can be systematically analyzed without reference to any follow-up data, follow-up observations will be essential in extracting the full science potential of very high-magnification events. The uniformly high-cadence observations expected for KMTNet also result in $\sim 55\%$ of all detected planets being non-caustic-crossing, and more low-mass planets even down to Mars-mass being detected via planetary caustics. We also find that the distributions of orbital inclinations and planet mass ratios in multiple-planet events agree with the intrinsic distributions. | Gravitational microlensing has discovered more than 50 extrasolar planets, although only about half (27) of them have appeared in the literature. Though few in number, microlensing planets occupy a unique region in the parameter space that is difficult to probe by other techniques \citep[e.g.,][]{gaudi2012,mao2012}. In fact, microlensing planets have already yielded interesting statistical results concerning the frequency of planets around M dwarf stars \citep{gould2010,cassan2012}, and intriguing possibilities of free-floating planets \citep{sumi2011} that can only be found by microlensing. Ever since their first discovery, extrasolar planet systems have challenged the two fashionable models of planet formation - the core accretion and gravitational instability scenarios. In particular, in the core accretion theory, the planet population synthesis models are becoming increasingly sophisticated \citep[e.g.,][]{ida2010}, which take into account effects such as planetesimal accretion, gas accretion, disk evolution, migration and planet-planet interactions etc. For this reason, in this paper, we shall focus on this theory since its predictions are more quantitative and testable. A detailed comparison between planet formation theory and microlensing observations has now become more imperative because of the emergence of next-generation microlensing experiments. In the past, the discovery of extrasolar planets often rested on a combination of work by survey teams [Optical Gravitational Lensing Experiment \citep[OGLE,][]{udalski2003} and Microlensing Observations in Astrophysics \citep[MOA,][]{bond2001}], and follow-up networks with higher-cadence observations [e.g., the Microlensing Follow-up Network \citep[$\mu$FUN,][]{gould2006,gaudi2008}, the Probing Lensing ANomalies NETwork \citep[PLANET,][]{albrow1998}, and RoboNet \citep{tsapras2009}]. However, these joint operations make the selection function sometimes difficult to quantify, although for some subsamples, such as the small number of very high-magnification events, the sample appears to be complete statistically \citep{gould2010}. Such a situation is likely to be changed significantly with the completion of the Korea Microlensing Telescopes Network (KMTNet) by the end of 2014. KMTNet will have three telescopes sited in Chile, Australia and South Africa \citep{kim2010}. Each telescope will have an aperture of 1.6m with 4 deg$^2$ of field of view, surveying about 4 fields with 10 minute cadence. With such high-cadence observations, KMTNet will be able to analyze the data without reference to any follow-up observations. Therefore the selection function will be much simpler. As a result, statistical results will be easier to obtain, which will provide a more robust measurements of planet abundances and distributions, which in turn will better constrain the planet formation theories. \citet{shvartzvald2012} performed a detailed simulation for a next-generation microlensing network of 4 telescopes with aperture ranging from 1m to 1.8m and cadence from 15min to 45min. They use scaled Solar system analogs as the lens systems, and conclude that such a network can find of order 50 planets in 4 years, of which one in six reveals two planets in a single lensing event. This is a factor of several increase in discovery rate over the original alert/follow-up surveys \citep{gaudi2012}. KMTNet, with its larger field of view and higher cadence, is therefore expected to yield more planet detections and bring down the detection limit to lower-mass planets. The present work is the first one that introduces the planet population synthesis model into microlensing simulations. Unlike previous works that use systems with only one planet, or simplified Solar system analogs as the lens system, our simulation is performed fully in the context of the Ida \& Lin core accretion model. Our results will be presented in two papers. The purpose of this paper is to provide a simple and yet somewhat realistic assessment of the fractions of extrasolar planets expected from KMTNet from the planet populations predicted by the core accretion theory; we will also explore how the KMTNet planet population differs from that of the current survey plus followup mode of discovery. In Paper II, we will focus on the multiple-planet events that are detected in our simulation; we will discuss the detection dependence of one planet on the other, the influence of the undetectable planets on the recovery of the parameters of detected planets, and the double/triple degeneracy \citep{gaudi1998,song2014}. In \S2, we present the Ida \& Lin (2010) core accretion model, basics of microlensing and how we simulate microlensing data. In \S3 we describe our method of selecting events with extrasolar planets. In \S4, we present our main statistical results on the expected detection rates of extrasolar planet populations, and finally in \S5, we discuss further our results and implications for future observations. | \label{sec:discussion} We conducted a simple and yet realistic microlensing simulation for a KMTNet-like microlensing survey. The planet population is taken from the Ida \& Lin core accretion model for $0.3M_\odot$ lenses. Our simulation results in 292 planetary events, including 16 double-planet events, from 6690 microlensing events for which the lens system has at least one planet more massive than $0.1M_\oplus$. With the frequency of such planetary systems considered, we find the fraction of planetary events is $2.9\%$, out of which $5.5\%$ show multiple-planet detections. We address the limitations of our simulation here before discussing the implications of our results. (1) We admit that our simulation is not fully realistic in the sense that the Einstein timescale $t_{\rm E}$, which involves distances $D_{\rm S}$ and $D_{\rm L}$, lens mass $M_{\rm L}$ and relative proper motion $\mu_{\rm rel}$, and the source size $\rho$, are fixed to some typical values. This does prevent us from making precise predictions for KMTNet, but predicting the yields of KMTNet or any other specific microlensing experiment is not the main purpose of our work. Our simulation aims to address more general questions, which are not easily clarified if too many observational factors are considered. Moreover, a microlensing simulation with 12 lenses in each system on average will become extremely complicated if all the parameters that we have held fixed (i.e., $D_{\rm L}$, $D_{\rm S}$, $M_{\rm L}$, $\mu_{\rm rel}$ and $\rho$) are set free. (2) The planet population given by Ida \& Lin's model is produced for stars in the Galactic disk, but in our simulation the lens system is placed in the Bulge. Planets forming around Bulge stars may well have very different distributions from those forming around Disk stars, not only because of the different metallicity but also because of the very dense environment \citep{thompson2013}. However, there is as yet no model available to quantitatively predict the planet population in the Galactic Bulge. Therefore, using planet population predictions for Disk stars is our only choice. One positive outcome of this approach is that comparing the results of our simulation with real observations may tell us how different the planet populations are in the Bulge and in the Disk, which is a question that can only be answered by microlensing. (3) When randomly placing the planets on their orbit, we do not take mean motion resonances into account. This may not be correct in the case of resonant systems. However, we notice that the Ida \& Lin's model does not show strong resonance signatures, as is shown in Figure~\ref{fig:period-ratio}. Therefore, the orbit of each planet is mostly unaffected by others in the same system, so randomly placing them on their orbit is acceptable. Our simulation yields more multiple-planet events than our naively expectation. Given that the total number of planets in $6690$ systems is $74560$, the probability for one planet to be detected is therefore $p=0.0041$ if we naively assume such detection does not depend on the characteristics of the microlensing event or properties of the planet. Then the number of single-planet events we would expect to detect in our simulation should be \begin{equation} N_1 = \sum_j n_j p , \end{equation} and the number of double-planet events is \begin{equation} N_2 = \sum_j \frac{n_j(n_j-1)}{2} p^2 , \end{equation} where $n_j$ is the number of planets in the $j^{\rm th}$ system. Given $N_1=276$ in our simulation, we would expect $N_2$ should be $7$ if our assumption holds, which is significantly lower than what we do detect, $N_2=16$. This is reasonable since the detectability of planets depends not only on the physical properties of the planet, but also on the impact parameter of that microlensing event. High-magnification events and massive planets are more favored in multiple-planet microlensing, as is shown in Figures~\ref{fig:q-s-map} and \ref{fig:u0}. Our simulation therefore predicts that multiple-planet events will be detected more than our naive expectation, but they are strongly biased toward massive planets and higher-magnification events. In Section~\ref{subsec:u0-dependence} we have shown that extremely high-magnification events are less sensitive to planet detections than those moderately high-magnification ones in such a KMTNet-like survey program. This apparently conflicts with previous theoretical predictions \citep{griest1998} as well as ongoing observations \citep{gould2010}. The reason might come from the different observing strategies used in our simulation and in current observations. The survey plus follow-up mode used in current microlensing observations can achieve very high cadence (e.g., more than one observation per minute after accounting for multiple observatories) during the peak of high-magnification events, although the survey teams typically obtain only a few observations per night. These intensive observations during the peak make high-magnification events extremely sensitive to planet perturbations. Our simulation is conducted using a strategy similar to that expected for KMTNet, which uses a constant cadence (10 mins) of observations everywhere. Therefore, for extremely high-magnification events, the planet perturbation is so weak that it may be missed by such a 10-min cadence observing strategy. This argues that even in the era of next generation surveys, there is still a need for follow-up of high-magnification events, which will require the next generation surveys to process their data in real time and produce high-magnification alerts, as is done for current surveys. With $\sim 20\%$ of high-magnification planet detections yielding multiple planets, such follow-ups are important for measuring the number of multiple-planet systems \citep{gaudi1998}. The advantage of conducting uniformly high-cadence observations everywhere in the light curve like KMTNet, in addition to obtaining a well controlled planets sample for statistical studies, is the ability to detect more low-amplitude planetary perturbations and perturbations due to planetary caustics. Low-amplitude perturbations are usually produced by source-star trajectories that do not cross any caustics. If we define caustic crossings as occurring if the closest distance between the source trajectory and caustics is less than two source radii, we find that $55\%$ of all detected planets in our simulation are not due to this caustic crossing, as is listed in Table~\ref{tab:Amax-number}. In contrast, we searched all published microlensing planets and characterized them according to this definition of caustic crossing. We find only three real microlensing planets are due to such non-caustic-crossing events within $26$ published microlensing planets with very good data coverage under the current observation strategy which are listed in Table~\ref{tab:all-planets}. This implies that in future microlensing programs like KMTNet, WFIRST \citep[the Wide-Field InfraRed Survey Telescope,][]{spergel2013} and possibly Euclid \citep{penny2013}, at least half of the microlensing planets will not be detected by crossing caustics. The non-caustic-crossing character of the event makes it more difficult to determine the physical properties of the lens system, since the unknown but important quantity $\theta_{\rm E}$ cannot be determined from the angular size of the source star that is derived from the source color and brightness \citep{yoo2004}. However, in the case of WFIRST it may be possible to measure $\theta_{\rm E}$ by astrometric microlensing \citep{gould2014} or (in the case that the lens is luminous) by taking high-resolution images several years before or after the event. The number of planets detected via planetary, central and resonant caustics are 107, 128 and 78 respectively. \footnote{Planets detected via both planetary and central caustics are counted twice.} The fraction of that by planetary caustics, $35\%$, is slightly higher than but consistent with $27\%$(=7/26) based on real microlensing planets. More planets being detected via planetary caustics and the high cadence observations around the planetary anomaly lead to the detection of very low-mass planets even down to Mars-mass, as the planetary caustic shrinks slower ($\sim \sqrt{q}$) than the central caustic does ($\sim q$) as the planetary mass ratio $q$ decreases. In Figure~\ref{fig:cdf-q} we compare the cumulative distribution of mass ratios of planets detected in our simulation with that of real microlensing planets. we notice the two curves coincide with each other surprisingly well for $q>10^{-3}$, but that the curve from our simulation has a long tail toward very small mass ratio, which means that future microlensing surveys will be able to explore more very low-mass planets than current observations. This tendency is not changed even when we choose a larger $\Delta \chi^2$ cutoff value. To understand which events contribute to this change, we divide all events into two groups: high-magnification events ($A_{\rm max}>100$) and low-magnification events ($A_{\rm max}<100$). The two panels in Figure~\ref{fig:cdf-q-2sample} tell us that most of these low-mass planets are detected in low-magnification events, which is understandable since they are more often detected via planetary caustics (Figure~\ref{fig:q-s-map}). Within all 26 well-understood microlensing planets, 3 are claimed to be super-Jupiters around M-dwarf hosts [OGLE-2005-BLG-071Lb \citep{udalski2005,dong2009a}; MOA-2009-BLG-387Lb \citep{batista2011}; OGLE-2012-BLG-406Lb \citep{poleski2014,tsapras2014}] . The ratio, 3 out of 26, is much higher than the estimation from either core accretion theory \citep[e.g.,][]{kennedy2008} or other exoplanet detection techniques \citep[e.g.,][]{cumming2008}. In our simulation, we find that the small fraction of super-Jupiter systems given by the Ida \& Lin core accretion model is magnified by a factor of $\sim 8$ if observed via microlensing. Therefore, this observational bias should be taken into account when comparing the frequency of massive planets around M dwarfs from microlensing observations with that from planet formation theory. Our simulation also shows that the inclination of the lens system of multiple-planet events obeys the intrinsic distribution of orbital inclinations, and that the mass ratio between the two detected planets also agrees with the intrinsic mass ratio distribution of the planetary system. | 14 | 3 | 1403.4936 |
1403 | 1403.0809_arXiv.txt | {Classical Cepheids are among the most important variable star types due to their nature as standard candles and have a long history of modeling in terms of stellar evolution. The effects of rotation on Cepheids have not yet been discussed in detail in the literature, although some qualitative trends have already been mentioned.} {We aim to improve the understanding of Cepheids from an evolutionary perspective and establish the role of rotation in the Cepheid paradigm. In particular, we are interested in the contribution of rotation to the problem of Cepheid masses, and explore \emph{testable} predictions of quantities that can be confronted with observations.} {Recently developed evolutionary models including a homogeneous and self-consistent treatment of axial rotation are studied in detail during the crossings of the classical instability strip (IS). The dependence of a suite of parameters on initial rotation is studied. These parameters include mass, luminosity, temperature, lifetimes, equatorial velocity, surface abundances, and rates of period change.} {Several key results are obtained: i)\,mass-luminosity (M-L) relations depend on rotation, particularly during the blue loop phase; ii)\,luminosity increases between crossings of the IS. Hence, Cepheid M-L relations at fixed initial rotation rate depend on crossing number (the faster the rotation, the larger the luminosity difference between crossings); iii)\,\emph{the Cepheid mass discrepancy problem vanishes} when rotation and crossing number are taken into account, without a need for high core overshooting values or enhanced mass loss; iv)\,rotation creates dispersion around average parameters predicted at fixed mass and metallicity. This is of particular importance for the period-luminosity-relation, for which rotation is a source of intrinsic dispersion; v)\,enhanced surface abundances do not unambiguously distinguish Cepheids occupying the Hertzsprung gap from ones on blue loops (after dredge-up), since rotational mixing can lead to significantly enhanced Main Sequence (MS) abundances; vi)\,rotating models predict greater Cepheid ages than non-rotating models due to longer MS lifetimes. } {Rotation has a significant evolutionary impact on classical Cepheids and should no longer be neglected in their study.} | Classical Cepheids are objects of interest for many areas of astrophysics. On the one hand, they are excellent standard candles allowing the determination of distances in the Milky Way and up to Virgo cluster distances. On the other hand, they are excellent objects for constraining stellar evolutionary models. Accordingly, Cepheids have played a special role among pulsating (variable) stars and have a long history of modeling efforts both in terms of their evolution and pulsations \citep[see][and references therein]{2013IAUS..289..116B}. Following the development of pulsation models in the late 1960's \citep{1969MNRAS.144..461S,1969MNRAS.144..485S,1969MNRAS.144..511S}, systematic differences between pulsational and evolutionary masses of Cepheids became apparent. These differences were originally referred to as Cepheid mass anomalies \citep{1980ARA&A..18...15C}. Improved opacities \citep{1991ApJ...371..408I,1994MNRAS.266..805S} could mitigate a good fraction of the disaccord in terms of masses. Nowadays, it is common to speak of \emph{the mass discrepancy} as the systematic offset between evolutionary masses and those derived with other methods \citep[e.g.][]{2006MmSAI..77..207B,2008ApJ...677..483K}, with a typical disagreement at the level of $10-20\%$. Several mechanisms have been put forward to resolve the mass discrepancy problem. Among the most prominent in the recent literature are augmented convective core overshooting \citep{2012ApJ...749..108P} and pulsation-enhanced mass-loss \citep{2008ApJ...684..569N}. However, the effect of rotation on evolutionary Cepheid masses has not been discussed in detail although the impact of rotation on the mass-luminosity and period-luminosity relations have already been mentioned previously in the literature \citep[see][]{2000ARA&A..38..143M,2001A&A...373..555M}. The first large grid of models incorporating a homogeneous and consistent treatment of axial rotation was recently presented by \citet[from hereon: paper\,I]{paperI}. Subsequently, \citet[from hereon: paper\,II]{paperII} extended the grid in terms of rotation for stars between $1.7$ and $15\,M_\odot$. Using these state-of-the-art grids it is now possible, for the first time, to consider the effect of rotation on populations of classical Cepheids. Cepheid progenitors are B-type stars on the Main Sequence (MS). Observationally, B-type stars are known for their fast rotation since the first homogeneous study of rotational velocities by \citet{1949ApJ...110..498S}. Thus, it is empirically known that fast rotation is common for the progenitors of Cepheids. \citet{2010ApJ...722..605H} have recently carried out a very detailed investigation of rotational velocities for B-stars of different masses and evolutionary states (on the MS), providing even an empirical distribution of rotation rates. Their distribution can serve as a guideline for our study in the sense that the typical MS rotation rates of B-stars in the mass range appropriate for Cepheids ($\sim 5 - 9 \,M_\odot$) is approximately $v/v_{\rm{crit}}=0.3 - 0.4$, where $v_{\rm{crit}}$ denotes critical rotation velocity. As shown in papers\,I\,\&\,II, such fast rotation can significantly alter the evolutionary path of stars by introducing additional mixing effects that impact MS lifetime, stellar core size, age, and MS surface abundances. Clearly, these effects will propagate into the advanced stages of evolution, such as the Cepheid stage. The blue loop phase of intermediate-mass evolved stars during core helium burning is very sensitive to the input physics \citep[see the ``magnifying glass'' metaphor in][p. 305]{1994sse..book.....K}. Thus, there is a two-fold interest in studying the models presented in papers\,I \&\,II in terms of their predictions for Cepheids: i) properties of Cepheids inferred using models (e.g. mass) can be updated to account for rotation; ii) certain predictions made by the models can be tested immediately using observational data (e.g. surface abundances). This paper is the first in a series devoted to the endeavor of extending the evolutionary paradigm of Cepheids to include the effects of rotation. We here focus mainly on the detailed exploration of predictions made by the models as well as their interpretation. A detailed comparison of these predictions to observed features is in progress and will be presented in a future publication for the sake of brevity. Further projects will include the investigation of the combined effect of metallicity and rotation as well as a self-consistent determination of pulsation periods and instability strip boundaries. The structure of this paper is as follows. In Sec.\,\ref{sec:Physics}, we briefly state a few key aspects of the models presented in papers\,I\,\&\,II relevant in the context of this work. Sec.\,\ref{sec:Predictions} contains the predictions made by the rotating Cepheid models, focusing on the mass discrepancy and features that are accessible to observations. Sec.\,\ref{sec:Discussion} discusses the reliability and implications of the predictions made, and Sec. \ref{sec:Conclusions} summarizes the key aspects. | \label{sec:Conclusions} This paper presents the first detailed investigation of the effect of rotation on evolutionary models of classical Cepheids. The study is based on the latest state-of-the-art rotating Geneva models (cf. paper\,I) that incorporate a homogeneous and self-consistent treatment of rotation over the entire evolutionary cycle for a wide range of stellar masses. A dense grid of evolutionary tracks of different rotation rates for intermediate-mass stars is available from paper\,II and enables a detailed investigation not only as a function of initial rotation rate, $\omega$, but also as a function of time during the IS crossings. Qualitative predictions are made as a function of initial rotation rate for an array of observable quantities, such as surface gravities, surface abundance enhancement, surface velocities, radii, and rates of period change. These will be quantified and compared to observational constraints in a subsequent publication. The key results of this investigation are: \begin{enumerate} \item M-L relations depend on $\omega$. This is true for all stars during all stages of evolution, although the difference is more obvious during the blue loop phase. \item For Cepheids, an M-L relation at fixed $\omega$ furthermore depends on crossing number. The greater $\omega$, the greater the difference between the crossings, since rotation broadens the blue loops. \item The Cepheid mass discrepancy problem vanishes when rotation and crossing number are taken into account, without a need for high core overshooting values or enhanced mass loss. \item Differences in initial rotation rate and crossing number between Cepheids of identical mass and metallicity create intrinsic scatter around the average values predicted. This is true for most parameters considered here, among them in particular the Cepheid period-luminosity-relation, i.e., rotation is a source of intrinsic dispersion for the PLR. \item Rotational mixing can significantly enhance surface abundances during the MS phase. Consequently, enriched surface abundances do not unambiguously distinguish Cepheids on the first crossing from ones on second or third crossings. \item Rotating models predict older Cepheids than currently assumed due to longer MS lifetimes. \end{enumerate} Arguably the two most important results are result numbers 3 and 4, i.e., the solution of the mass discrepancy problem and the prediction of intrinsic scatter for Cepheid observables at fixed mass and metallicity. Since the mass discrepancy has been claimed to be solved by several other studies in the past, one may be reluctant to immediately adopt rotation as the best explanation. However, result 4 may provide the smoking gun for distinguishing between rotation and other effects such as core overshooting. While it would be difficult to explain why stars of identical mass and metallicity should present different overshooting values (leading to the observed scatter in, say, radius or luminosity), dispersion arises naturally when a dispersion in initial velocity is considered, which is an observational fact. As this paper shows, the effects of rotation on Cepheid populations are highly significant, ranging from issues related to inferred masses to the prediction of systematic effects relevant for the distance scale. Further related research is currently in progress. | 14 | 3 | 1403.0809 |
1403 | 1403.0724_arXiv.txt | { We present new results of CO(1-0) spectroscopic observations of 4 SDSS type 2 quasars (QSO2) at z$\sim$0.3, observed with the 30m IRAM telescope. The QSO2 have infrared luminosities in the ULIRG (UltraLuminous Infrared Galaxies) regime. We confirm the CO(1-0) detection in one of our 4 QSO2, SDSS J1543-00, with $L'_{CO}$ and $M_{H_2}$ (1.2$\pm$0.2) $\times$10$^{10}$ K km s$^{-1}$ pc$^2$ and (9.4$\pm$1.4)$\times$10$^9$ M$_{\odot}$, respectively. The CO(1-0) line has $FWHM=$575$\pm$102 km s$^{-1}$. No CO(1-0) emission is detected in SDSS J0903+02, SDSS J1337-01, SDSS J1520-01 above 3 sigma, yielding upper limits on $M(H_2)\sim$ 9.6, 4.3 and 5.1 $\times$10$^9$ M$_{\odot}$ respectively. Together with CO measurements of 9 QSO2 at $z\sim$0.3-1.0 from the ULIRG sample by Combes et al. (2011, 2013), we expand previous studies of the molecular gas content of intermediate $z$ QSO2 into the ULIRG regime. We discuss the location of the 13 ULIRG QSO2 at $z<$1 with available $L'_{CO}$ measurements in the $L'_{CO}$ vs. $z$ and $L'_{CO}$ vs. $L_{FIR}$ diagrams, in comparison with other QSO1 and ULIRG star forming samples. } | Following the discovery in large numbers of type 2 quasars (QSO2) during the last decade, an intensive follow up has been performed at different wavelengths: X-ray, radio, infrared and optical (e.g., Szokoly et al. 2004, Lacy et al. 2007, Mart\'{\i}nez-Sansigre et al. 2006, Zakamska et al. 2003). In spite of this, the molecular gas content of this class of objects has been very scarcely studied. The study of the molecular gas content in these objects is crucial to better understand the conditions required to trigger both the nuclear and star formation activities in the most luminous active galaxies (AGN), since this gaseous phase can provide large amounts of fuel to form stars and feed the nuclear black hole. It can also provide information about the relation between QSO2 and other systems such as type 1 quasars (QSO1), luminous (LIRG, 10$^{11}$ L$_{\odot} \le L_{IR}< $10$^{12}$ L$_{\odot}$) and ultra-luminous (ULIRG, $L_{IR}\ge$10$^{12}$ L$_{\odot}$) infrared galaxies, where $L_{IR}$ is the total infrared luminosity ($\sim$8-1000 $\mu$m range). However, only few molecular gas studies of QSO2 have been carried out, and frequently focussed on $z>$2 objects (see Villar Mart\'\i n et al. 2013 (VM13 hereafter) and references therein). | \label{sec:con} $\bullet$ We present the results of CO(1-0) spectroscopic observations of 4 SDSS QSO2 at $z$ $\sim$ 0.3 observed with the 30m IRAM telescope. These QSO2 have infrared luminosities in the ULIRG regime, expanding our previous work on less infrared luminous QSO2 at $z\sim$0.3 into this regime. The 4 QSO2 have $L_{FIR}\sim$(1-2.5)$\times$10$^{12}$ L$_{\odot}$. We have also added 9 ULIRG QSO2 at $z\sim$0.3-0.9 from Combes et al. (2011,2013), which have $L_{FIR}\sim$(0.2-2.5)$\times$10$^{13}$ L$_{\odot}$. \\ $\bullet$ CO(1-0) detection is confirmed in one of the 4 objects observed by us: SDSS J1543-01. L$'_{CO}$ and $M_{H_2}$ are (1.2$\pm$0.2) $\times$10$^9$ K km s$^{-1}$ pc$^2$ and (9.6$\pm$1.6)$\times$10$^9$ M$_{\odot}$, respectively. The line has $FWHM=$575$\pm$102 km $^{-1}$. No CO(1-0) emission is detected in SDSS J0903+02, SDSS J1337-01, SDSS J1520-01. 3$\sigma$ upper limits on $M(H_2)$ are 9.6, 4.3 and 5.4 $\times$10$^9$ M$_{\odot}$ respectively.\\ $\bullet$ The $L'_{CO}$ (measurements and upper limits) of all 13 ULIRG QSO2 at $z<$1 fall in the $L'_{CO}$ vs. $L_{FIR}$ and $L'_{CO}$ vs. $z$ correlations defined by QSO1 and QSO2 with different $z$ and $L_{IR}$. They overlap as well with $z<$1 star forming ULIRGs. Several QSO1 and QSO2 in Combes et al. (2011, 2013) mark the lower envelope defined by the scatter of the correlation. They might be gas poor objects. Alternatively, this result might be an effect of a significant contribution of the AGN to the $L_{FIR}$. | 14 | 3 | 1403.0724 |
1403 | 1403.7568_arXiv.txt | {} {The nature of the Thick Disc and its relation with the Thin Disc is presently an important subject of debate. In fact, the structural and chemo-dynamical transition between disc populations can be used as a test of the proposed models of Galactic disc formation and evolution. } {We have used the atmospheric parameters, \aabun \ abundances and radial velocities, determined from the Gaia-ESO Survey GIRAFFE spectra of FGK-type stars (first nine months of observations), to provide a chemo-kinematical characterisation of the disc stellar populations. We focuss on a subsample of 1016 stars with high quality parameters, covering the volume $|Z|<$4.5~kpc and R in the range 2-13~kpc. } {We have identified a thin to thick disc separation in the \aabun \ v.s. \met \ plane, thanks to the presence of a low-density region in the number density distribution. The thick disc stars seem to lie in progressively thinner layers above the Galactic plane, as metallicity increases and \aabun \ decreases. On the contrary, the thin disc population presents a constant value of the mean distance to the Galactic plane at all metallicities. In addition, our data confirm the already known correlations between V$_{\phi}$ and \met \ for the two discs. For the thick disc sequence, a study of the possible contamination by thin disc stars suggests a gradient up to 64$\pm$9~km~s$^{-1}$~dex$^{-1}$. The distributions of azimuthal velocity, vertical velocity, and orbital parameters are also analysed for the chemically separated samples. Concerning the gradients with galactocentric radius, we find, for the thin disc, a flat behaviour of the azimuthal velocity, a metallicity gradient equal to -0.058$\pm$0.008~dex~kpc$^{-1}$ and a very small positive \aabun \ gradient. For the thick disc, flat gradients in \met \ and \aabun \ are derived. } {Our chemo-kinematical analysis suggests a picture in which the thick disc % seems to have experienced a settling process, during which its rotation increased progressively, and, possibly, the azimuthal velocity dispersion decreased. At \met$\approx$-0.25~dex and \aabun$\approx$0.1~dex, the mean characteristics of the thick disc in vertical distance to the Galactic plane, rotation, rotational dispersion and stellar orbits eccentricity are in agreement with that of the thin disc stars of the same metallicity, suggesting a possible connection between these two populations at a certain epoch of the disc evolution. Finally, the results presented here, based only on the first months of the Gaia ESO Survey observations, confirm how crucial are today large high-resolution spectroscopic surveys outside the solar neighbourhood for our understanding of the Milky Way history. } | Understanding the chemodynamical evolution of the Milky Way disc stellar populations is a crucial step to reconstruct the history of our Galaxy. This approach can shed new light on the detailed processes that led the Milky Way to form and evolve as a disc-type galaxy, in the more general context of galaxy evolution. In particular, since the first works highlighting the existence of the Milky Way's thin/thick disc dichotomy \citep{Yoshii,GilmoreReid}, the study of the structural and chemodynamical transition between the two disc populations has opened new pathways to constraint Galactic evolution models. External mechanisms, supported by the hierarchical formation of galaxies in the $\Lambda$CDM paradigm, have been invoked to explain the chemodynamical characteristics of thick disc stars. Among them, the accretion of dwarf galaxies \citep[e.g.][]{Statler, Abadi} and minor mergers of satellites \citep[e.g.][]{Quinn, Villalobos} heating dynamically the disc, have been proposed. Similarly, \citet{Jones83} and \citet{Brook1,Brook2} have proposed the accretion of a gas-rich merger from the collapse of which the thick disc would have formed. On the other hand, purely internal formation mechanisms have also been proposed, like the early turbulent phase of the primordial disc \citep[e.g.][]{Bournaud} and the stars radial migration due to resonances with the spiral structure or the bar of the Milky Way \citep[e.g.][]{Schonrich2009,Minchev10,Loebman}. The relative importance of the different proposed physical processes of evolution can now be tested with increasing robustness, thanks to the Galactic spectroscopic surveys targeting the disc stars. Spectroscopic studies of the disc populations were initially based on kinematically selected samples of solar neighbourhood stars \citep[e.g.][]{Prochaska2000,Fuhrmann04,Bensby05,Reddy06}, with some exceptions like \citet{Edvardsson93} or \citet{Fuhrmann98}. In the classical picture that emerged from those studies, the thick disc population is characterized, compared to the thin disc one, as kinematically hotter \citep{CasettiDinescu2011}, and therefore, composed by stars moving in Galactic orbits with a higher scale height \citep[$\sim$900~pc vs. $\sim$300~pc, e.g.][]{Juric2008}. In addition, the thick disc metallicity distribution peaks at lower values \citep[$\sim$-0.5~dex compared to $\sim$-0.2~dex for the thin disc, e.g][]{WyseGilmore95,Kordopatis11b,Lee2011b}. With respect to the thin disc, the thick disc stars are older and have an enhanced ratio of $\alpha$-element abundances over iron \citep[\aabun, e.g.][]{Bensby05,Bensby07,Fuhrmann08}. Nevertheless, it progressively appeared that the distributions in the above mentioned physical parameters displayed significant overlaps between the thin and the thick disc populations \citep[e.g.][]{BensbyFeltzing10}. This blurred our comprehension of the interplay between the discs and questionned the classical kinematically-based definitions \citep[][]{Bovy12Mono}. On the other hand, as discussed in \citet{Bovy2012a}, defining stellar populations by abundances patterns is a better approach than the traditional kinematical criteria, as chemical abundances can correlate with disc structure, but are formally independent of it. Moreover, the identification of possible chemical evolutionary paths and/or gaps in the abundance ratios distributions can be crucial to disentangle the otherwise overlapping populations in the kinematical space. In the recent years, the number of stars analysed with high enough spectroscopic resolution to provide detailed chemical diagnostics has increased from a few hundreds to several tens of thousands. The advent of Milky Way spectroscopic surveys and of automatised chemical analysis techniques have improved both the statistical robustness and the homogeneity of the data as a consequence. From the theoretical side, new approaches based, at least partially, on the chemical identification of disc sub-populations, as the chemical-tagging \citep{FreemanBH} or the mono-abundance populations \citep{RixBovy} methods, are opening promising pathways for constraining the Milky Way's evolutionary processes. This is also challenging our previous conception of the Milky Way disc and the way in which stellar populations can be defined.\\ In this context, the low resolution (R $\sim$ 2\,000) SEGUE spectroscopic survey addressed the question of the thin to thick disc transition with a robust statistical approach, inside and outside the solar neighbourhood. \cite{Lee2011b} analysed $\sim$ 17\,300 G-type dwarfs with errors of about 0.23 dex in [Fe/H] \citep{Smolinski2011} and 0.1 dex for \aabun \ \citep{Lee2011a}. This was complemented by \cite{Schlesinger} who analysed 24\,270 G and 16\,847 K dwarfs at distances from 0.2 to 2.3~kpc from the Galactic plane. The \citet{Lee2011b} \aabun ~versus [Fe/H] distribution of number densities (see their Fig. 2) already motivated the authors to chemically separate the thin and the thick disc populations, although no clear gap is observed between the two. The SEGUE data interpretations differ from authors proposing no thin-thick disc distinction \citep{Bovy2012a} to those allowing the existence of a distinct thick-disc component formed through an external mechanism \citep{Liu}. Nevertheless, selection effects have been invoked to explain some of the distinctions between the \citet{Lee2011a} and the \citet{Bovy2012a} chemical distributions. On the other hand, the data interpretation in the sense of no thin-thick disc distinction agrees with the results of the \citet{Shonrich09NoThick} theoretical study. More recently, \citet{Boeche} analysed the chemo-kinematical characteristics of 9\,131 giants included in the last RAVE survey data release, with available \met \ and \aabun. The RAVE survey data have a higher resolution (R$=$7\,500) than SEGUE spectra, but with a wavelength range limited to the infrared Ca~II triplet. The \citet{Boeche} analysis, based on the eccentricity-Zmax plane combined with additional orbital parameters, allowed them to identify three stellar populations that could be associated with the Galactic thin disc, a dissipative component composed mostly of thick-disc stars, and the halo. Nevertheless, their thin and thick disc populations, defined with the above mentioned dynamical criteria, do not clearly separate in their \aabun \ v.s. \met \ plane (c.f. their Fig.~10). On the other hand, recent studies of solar neighbourhood samples, analysed with high resolution spectroscopic data, have revealed the existence of a gap in the \aabun ~versus [Fe/H] plane \citep[e.g][]{Fuhrmann04,Reddy06,Bensby07}. Recently, \citet{Adibekyan2013} used the stellar sample of 1\,111 long-lived FGK dwarf stars from \citet{Adibekyan2012} to separate and characterise the different Galactic stellar subsystems, taking into account the existence of the above mentioned gap. Their \aabun ~measurements are based on Mg, Si and Ti averaged abundances and the typical relative uncertainties in the metallicity and \aabun ~are of about 0.03 dex. \citet{Haywood2013} derived ages of the \citet{Adibekyan2012} sample, concluding that two regimes appear in the age-\aabun ~plane, that they identify as the epochs of the thick and thin disc formation. In the \citet{Haywood2013} scenario, the thick disc ``formed from a well mixed interstellar medium, probably first in starburst, then in a more quiescent mode, over a time scale of 4-5 Gyr''. They also suggest that ``the youngest thick disk set the initial conditions from which the inner thin disk started to form 8 Gyr ago, at [Fe/H] in the range of (-0.1,+0.1) dex and \aabun=0.1 dex''. Finally, they interpret the low-metallicity tail of the thin disc as having an outer disc origin, somewhat disconnected from the inner disc evolution. In addition, similar separations in the \aabun ~versus [Fe/H] plane are distinguishable in the \citet{Ramirez} oxygen-abundance versus [Fe/H] measurements. These authors also report similar age-\aabun ~regimes as in \citet{Haywood2013}. In summary, the increasing number of Galactic disc stars with known kinematical and chemical characteristics has emphasized the question of the continuity or not between the thick and the thin disc components, in structural properties, kinematics, abundances and ages. Different authors like \citet{Norris99} and \citet{Bovy2012a} have even raised the possibility of considering the Galactic disc as a single component, without a thick/thin disc distinction, analysing for instance the mass-weighted scale-height distribution of disc stars. In this context, the GIRAFFE observations of the Gaia-ESO Survey \citep[GES,][]{GESMessenger} offer a unique opportunity to extend the previously high-resolution studies of the solar neighbourhood to larger and more radially extended samples. In its first data release, GES already includes $\sim$10\,000 spectra (R$\sim$20\,000 and R$\sim$16\,000) of stars in the Milky Way field, from the halo, the thick disc, the thin disc and the bulge. The expected final number of targets, at the end of the five years of observations, is $\sim$100\,000. The present paper addresses the thick-thin disc transition as seen in the GES kinematical and chemical (\aabun ~and global metallicity) first release data of FGK-type stars. Our point of view, exploiting the high resolution of the GES data, is that of the chemical characterisation and definition of the disc populations (essentially identified as thin and thick discs). As a consequence, the subsequent analysis of the structural properties, kinematics and orbital parameters will be done in the framework of the chemically defined disc populations. Section~2 describes the data sample, including the stellar parameters and \aabun~ measurements. Section~3 presents the derivation of the stellar distances and kinematics. Section~4 focuses on the analysis of the number density distribution in the \aabun~ versus \met \ plane. Then, the distributions of the distances to the Galactic plane (Section~5), of the azimuthal velocities (Section~6), the azimuthal and vertical velocity dispersions (Section~7) and the stellar orbital parameters (Section~8) are discussed, in the context of a proposed thick-thin disc separation based on the chemical criteria. Section~9 presents the rotational and abundance gradients with galactocentric radius and, finally, we discuss our results in Section~10. | We have used the atmospheric parameters, \aabun \ abundance ratios and radial velocities, determined from the Gaia-ESO Survey GIRAFFE spectra of FGK-type stars (first nine months of observations), to present the kinematical and chemical characterisation of the observed stellar populations. First, we have chosen to identify the thin to thick disc separation in the \aabun \ v.s. \met \ plane, thanks to the presence of a low-density region in the number density distribution. This separation is in agreement with the recently proposed one by \citet{Adibekyan2013} from a sample of solar neighbourhood stars, and the ones observed before by \citet{Fuhrmann04}, \citet{Reddy06} and \citet{Bensby07}, among others. Furthermore, we have studied the distribution of the distances to the Galactic plane, in the \aabun \ v.s. \met \ plane, with particular attention to the chemically defined thin and thick disc populations. Regarding the thick disc, our results show that the stars lay in progressively thinner layers along the sequence (as metallicity increases and \aabun \ decreases). This finding could be in agreement with the \citet{Haywood2013} proposal of a thick disc forming stars in thinner and thinner layers for 4-5 Gyrs, also predicted by the dissipational settling formation scenario of \citet{Burkert92}. Similarly, \citet{Bovy12Mono} already pointed out the scenario of an evolution in gradually thinner layers (as suggested by the dependence of their scale height distribution on \aabun), although no clear separation between the thin and the thick disc sequences appears in their data. In addition, the recent hydrodynamic simulations of \citet{Bird13} and \citet{Stinson13} also show discs in which the younger stars have progressively shorter vertical scale heights. On the other hand, the thin disc sequence presents a constant value of the mean distance to the Galactic plane at all metallicities. As a consequence, the thick to thin disc transition in Z values appears more abrupt in the low metallicity regime. Concerning the kinematics, the mean rotational velocities found for the thick and the thin disc sequences are in agreement with the canonical ones (V$_{\Phi}=$176$\pm$16~km/s and V$_{\Phi}=$208$\pm$6~km/s). In addition, the mean Galactic rotation also seems to increase progressively along the thick disc sequence, as the metallicity increases. Our data also confirm the already observed correlations between V$_{\Phi}$ and \met \ for the two discs. In the case of the thick disc sequence, the mean value of the $\Delta$V$_{\Phi}$/$\Delta$\met \ is 43$\pm$13~km~s$^{-1}$~dex$^{-1}$, is in agreement with the previous findings of \citet{Lee2011b}, \citet{Kordopatis11b} and \citet{Spagna2010}. This would confirm the disagreement with the SDSS result of \citet{Ivezic}, based on photometric metallicities. Moreover, we have explored the possible influence of contaminating stars from the thin disc in the analysis of the correlation. If only the stars with the higher \aabun \ values at each metallicity bin are considered, the derived $\Delta$V$_{\Phi}$/$\Delta$\met \ is 64$\pm$9~km~s$^{-1}$~dex$^{-1}$. As a consequence, this steeper value of the gradient, cleaned from the influence of possible thin disc contaminants could be more characteristic of the real correlation for the thick disc. For the thin disc, a negative gradient of $\Delta$V$_{\Phi}$/$\Delta$\met$=$-17$\pm$6~km~$s^{-1}$~dex$^{-1}$ is found. Consequently, the difference in the characteristic rotational velocities between the thin and the thick disc are again accentuated in the low metallicity regime, as for the mean distances to the Galactic plane. Furthermore, the analysis of the velocity dispersions for the chemically separated thin and thick disc stars confirms a higher dispersion for the thick disc population than for the thin disc one, as expected, for both V$_{\Phi}$ and V$_{z}$. Further, the stars in the thick and the thin disc sequences present also different distributions of their stellar orbit eccentricities, with the thick disc stars presenting a broader distribution peaked at higher eccentricities than the thin disc one. Finally, we have analysed the properties of the thin and thick discs as a function of the Galactocentric radius. For the thin disc, we estimate a flat behaviour of the thin disc rotational velocity, a metallicity gradient equal to -0.058$\pm$0.008~dex~kpc$^{-1}$ and, a very small positive \aabun \ gradient. For the thick disc, no gradients in \met \ and \aabun \ are found. The general picture emerging from our analysis is that of a thick disc that lays in progressively thinner and thinner layers as the metallicity increases and the \aabun \ decreases with time (until about \met~$=-0.25~$dex and \aabun~$=0.1~$dex). During this settling process, the thick disc rotation increases progressively and, possibly, the azimuthal velocity dispersion decreases. No measurable metallicity or \aabun \ gradients with galactocentric distance seem to remain for the thick disc. At about \met~$=-0.25~$dex, the mean characteristics of the thick disc in \aabun, vertical distance to the plane, rotation and rotational dispersion are in agreement with (or just slightly different than) that of the thin disc stars of the same metallicity, suggesting a possible connection between the two populations at a certain moment of the disc evolution. This is in agreement with the \citet{Haywood2013} scenario, in which the inner thin disk (R$<$10kpc) inherited from the chemical conditions left at the end of the thick disk phase (the observed metal-rich end of the thick disc sequence in the \aabun \ vs. \met \ plane), some 8 Gyr ago. In their view, there may be an age gap in the star formation possibly causing the gap in \aabun \ enrichment, after which, star formation proceeds in a thin disk. For the thin disc, no clear gradient of the mean vertical distance to the plane with the stellar metallicity is observed. On the contrary, the stars in the metal-poor thin disc sequence tail seem to have a mean Galactic rotation velocity higher than those in the metal-rich part. Moreover, if confirmed, the existence of a small positive gradient in \aabun \ with R leads to a thin disc that would be slightly more metal-poor and \aabun-rich in the outer parts. These two characteristics, similar to those of the metal-poor end of the thin disc sequence in the \aabun \ v.s. \met \ plane, would favour the picture, suggested by \citet{Haywood2013} and \citet{Haywood2006}, of this metal-poor thin disc being formed in the outer parts. Contrary to the similarities between the thin and the thick discs around \met$\sim-$0.3dex, the stars in the metal-poor overlapping regime are more clearly separated in their mean physical properties (chemistry, kinematics, possibly orbital eccentricities). This separation seems to be confirmed by the gap (or low-density region) in the \aabun \ vs. \met. On the other hand, the tight bimodal distribution of \aabun \ at every value of \met \ poses severe constraints to models predictions. Recent attempts to define an age-metallicity relation for stars currently in the solar neighbourhood \citep[e.g.][]{Haywood2013, Mishenina13} hint at an evolutionary path in which a very wide range of ages corresponds to a single \met \ value. \citet{Haywood2013} results (c.f. their Fig.~8) show that at a value of about \met$=$-0.5~dex stars have an age range of around 10~Gyr. On the other hand, the present work and other studies focussed in the solar neighbourhood \citep[e.g.][]{Adibekyan2013} show that the distribution of \aabun \ at \met$=$-0.5~dex appears narrowly bimodal, with two essentially unresolved but quite discrete values of \aabun \ being found. As a consequence, although stars with \met$\approx$-0.5~dex formed over 10Gyr, they seem to be formed with only two values of \aabun \. This is a quite remarkably severe constraint that can be analysed in the light of Galactic models in which radial migration \citep[e.g.][]{Sellwood02, Schonrich2009, Schonrich2009b, Roskar08}, and churning, is a major factor. \citet{Schonrich2009b} claim that a bimodal, although continuous, distribution of \aabun \ can naturally come out as a consequence of the standard assumptions about star formation rates and metal enrichment, with no need to invoke breaks in the Galaxy's star formation history or accretion events. However, the tightness of the well populated observed thick and thin disc sequences in the \aabun \ vs. \met \ plane, with possibly only two values of \aabun \ for a given metallicity, poses severe constraints. If radial migration occurs, one can suppose an evolutionary path in which enrichment in the high-$\alpha$ branch happened rapidly, up to metallicities close to \met$=-$0.25~dex. This should occur across a large enough radial range of the proto-galactic disk. In fact, to preserve the tightness of the high-$\alpha$ sequence, the stars should had formed in a narrow time interval from radially well-mixed interstellar medium (ISM) gas, through all the galactocentric radii which generated stars that are today at a given R (ex. in the solar neighbourhod). Then, to explain the low-$\alpha$ sequence, one can postulate a significant gas accretion, reducing the disc ISM metallicity to, for instance, -0.6~dex. This should be followed by disc evolution in which large scale gas flows \citep[as in fountains, c.f.][and references therein]{Fraternali13} retain a tight \aabun-\met \ correlation, over the radial region which today populates the solar neighbourhood. In addition, this scenario should, somehow, also conciliate the higher rotation values observed in the metal-poor thin disc regime. Othewise, other possibility to explain the observed \aabun \ vs. \met \ distributions, is to suppose that strong radial migration in the sense of churning did not occure \citep[see also the discussion in][]{Haywood2013}. Finally, this study also shows the importance of precise chemical abundance measurements to disentangle the stellar population puzzle of the Galactic disc. The stellar chemical patterns can guide the definition of useful stellar sub-samples to unveil the evolutionary paths in the disc formation history and to give clearer constraints to the models. Moreover, in the near future, the Gaia mission of the European Space Agency will allow precise distances and ages estimations for all the analysed stars. The results presented here, based only on the first months of the GES observations, confirm how crucial are today large high-resolution surveys outside the solar neighbourhood for our understanding of the Milky Way history. | 14 | 3 | 1403.7568 |
1403 | 1403.0454_arXiv.txt | { The probability distribution function (PDF) of the mass surface density is an essential characteristic of the structure of molecular clouds or the interstellar medium in general. Observations of the PDF of molecular clouds indicate a composition of a broad distribution around the maximum and a decreasing tail at high mass surface densities. The first component is attributed to the random distribution of gas which is modeled using a log-normal function while the second component is attributed to condensed structures modeled using a simple power-law. The aim of this paper is to provide an analytical model of the PDF of condensed structures which can be used by observers to extract information about the condensations. The condensed structures are considered to be either spheres or cylinders with a truncated radial density profile at cloud radius $r_{\rm cl}$. The assumed profile is of the form $\rho(r)=\rho_{\rm c}/(1+(r/r_0)^2)^{n/2}$ for arbitrary power $n$ where $\rho_{\rm c}$ and $r_0$ are the central density and the inner radius, respectively. An implicit function is obtained which either truncates (sphere) or has a pole (cylinder) at maximal mass surface density. The PDF of spherical condensations and the asymptotic PDF of cylinders in the limit of infinite overdensity $\rho_{\rm c}/\rho(r_{\rm cl})$ flattens for steeper density profiles and has a power law asymptote at low and high mass surface densities and a well defined maximum. The power index of the asymptote $\Sigma^{-\gamma}$ of the logarithmic PDF ($\Sigma P(\Sigma)$) in the limit of high mass surface densities is given by $\gamma = (n+1)/(n-1)-1$ (spheres) or by $\gamma=n/(n-1)-1$ (cylinders in the limit of infinite overdensity). } | {Observations of the structure of molecular clouds provide insights about the physical processes in the cold dense phase of the interstellar medium and will give us a better understanding how they evolve and eventually form stars. They are furthermore essential to test theoretical models of the origin of the stellar mass function \citep{Padoan1997a,Elmegreen2001, Padoan2002, Hennebelle2008, Elmegreen2011, Hopkins2013a} and the star formation rate \citep{Krumholz2005,Padoan2011,Hennebelle2011,Federrath2012} which are both thought to be related to the density structure of a turbulent molecular gas. The high resolution and sensitivity of modern telescopes allows a detailed analysis of the 1-point statistic or probability distribution function (PDF) of the mass surface density of molecular cloud gas. They are obtained using either the reddening of stars \citep{Kainulainen2009,Froebrich2010,Lombardi2010, Kainulainen2013, Alves2014} or more recently the infrared emission of dust grains \citep{Hill2011,Hill2012,Schneider2012,Schneider2013a,Schneider2013b,Russeil2013}. Despite the complexity of the molecular clouds the observed PDFs of the mass surface density show very similar properties. They all are characterized by a broad peak and a tail at high mass surface densities approximately given by a power law. The PDF at low mass surface densities around the broad peak is attributed to randomly moving gas commonly referred to as 'turbulence' while the tail is attributed to self-gravitating cloud structures. The relative amount of the two different components seems to be related to the star formation activity in the cloud as discussed by \citet{Kainulainen2009}. While non-star forming clouds as the 'Coalsack' or the 'Lupus V' region show only a very low or no evidence of a tail the PDFs of star forming clouds as 'Taurus' or 'Orion' are characterized by a strong tail with no clear separation between the two components. The observations seem to be broadly consistent with current simulations of turbulent molecular clouds. Turbulence would naturally create a wide range of densities and simulations of driven isothermal turbulence have shown that the corresponding PDF has a log-normal form \citep{Vazquez1994,Padoan1997b,Passot1998}, a result which has been confirmed analytically \citep{Nordlund1999}. The projection of the density of those simulations has also been found to be closely log-normal in shape \citep{Ostriker2001,Vazquez2001,Federrath2010,Brunt2010c}. Deviations are expected for non isothermal turbulence which produces higher probabilities at low or high densities \citep{Scalo1998,Passot1998,Li2003}. More recent simulations of forced turbulence also show depending on the assumed forcing for the PDF of the volume density a deviation from the log-normal function with enhanced probabilities at low densities \citep{Federrath2008b,Konstandin2012,Federrath2013b}. The functional form is as shown by \citet{Federrath2013b} approximately described by a statistical function proposed by \citet{Hopkins2013b}. {Simulations of the time evolution of molecular clouds have shown that at late stage the PDF of the volume density would develop a tail-like structure at high density values \citep{Klessen2000,Dib2005,Vazquez2008}. The same behavior is also seen in the PDF of the mass surface density \citep{Ballesteros2011,Kritsuk2011,Federrath2013a}. Currently, observed PDFs are analyzed using a log-normal function for the peak and a simple power law for the tail, respectively. The log-normal function allows a first estimate of the density contrast of the volume density in a turbulent medium based on the fundamental relation of the statistical properties of the mass surface density and the ones of the volume density as provided by \citet{Fischera2004a} and also by \citet{Brunt2010b,Brunt2010c}. The interpretation of the tail is frequently based upon a simple power law density profile $\rho(r)\propto r^{-n}$ of spheres where the PDF of the mass surface density is also a power law. In case of the logarithmic PDF ($\Sigma P(\Sigma)=P(\ln \Sigma)$)\footnote{The PDF of the logarithmic values of the mass surface density is referred to as logarithmic PDF while the PDF of the absolute values of the mass surface density as linear PDF.} the corresponding power law would be $\Sigma^{-\gamma}$ with $\gamma=2/(n-1)$ \citep{Kritsuk2011,Federrath2013a}. Applying this relation the slope of the tail of the PDF of a number of star forming molecular clouds indicates a radial density profile with a power index $n\sim 2$ \citep{Schneider2013b} as expected for collapsing clouds. A different approach has been chosen by \citet{Kainulainen2013} who also applied a log-normal function to the tail. However, the analytical functions show partly strong deviations to the observed curves. Most of the PDFs published by \citet{Kainulainen2009} reveal a tail at low mass surface densities below the peak which cannot be explained in terms of the simple analytical function. The tail at high mass surface densities has several features which are not expected using simple power law profiles of the radial density. Foremost, the tail is restricted to a certain range of mass surface densities. For a number of PDFs published by \citet{Kainulainen2009} the tail indicates a strong cutoff or a strong change of the slope around $A_V=6-10~{\rm mag}$. The interpretation of the tail is further complicated by the observational facts that condensed clouds are generally located on a certain background level and that they are restricted to small regions within the cloud complex as e.g. in case of the Rosette molecular cloud \citep{Schneider2012}. The tail is therefore not necessarily a simple power law nor directly related to the radial density profile. Furthermore, the analytical functions do not provide a physical explanation for the peak position of the PDF which occurs in case of a number of molecular clouds around $A_V\sim 1~{\rm mag}$. } In this and the following papers an analytical {physical} model of the {global} PDF of molecular clouds is {developed} which resembles the {main} observed features {and is meant to derive basic physical parameters of star forming molecular clouds as the pressure and the density contrast in the turbulent gas.} This paper focuses on the 1-point statistical properties of individual condensed structures, assumed to be spheres and cylinders. In Sect.~2 an analytical solution of the mass surface density and the corresponding PDFs for the considered shapes is presented which is based on a truncated analytical density profile widely used in astrophysical problems. In Sect.~3 the properties of the PDFs {are} discussed and asymptotes for low and high mass surface densities are provided. Also studied is the location of the maximum position of the logarithmic and linear PDF. A summary is given in Sect.~4. The technical details can be found in the appendices. } | The study summarizes a series of properties of the PDF of the mass surface density of spherical and cylindrical structures {having an analytical radial density profile $\rho =\rho_{\rm c}/(1+(r/r_0)^2)^{n/2}$ where $\rho_c$ is the central density and $r_0$ the inner radius.} The profiles are assumed to be truncated at a cloud radius $r_{\rm cl}$ as expected for cold structures embedded in a considerably warmer medium. The results are therefore applicable to individual condensed structures in star forming molecular clouds. {The PDF for given geometry is determined by the power $n$, the density ratio $q=\rho(r_{\rm cl})/\rho_{\rm c}$, the product $\rho_{\rm c}r_{0}$, and, in case of a cylinder, the inclination angle $i$. It is convenient to describe the properties of the PDF in terms of the unit free mass surface density defined by \begin{equation} X_n = \Sigma_n \frac{\cos^\beta i}{2\rho_{\rm c} r_0}q^{-\frac{n-1}{n}}, \end{equation} where $\beta=0$ for spheres and $\beta=1$ for cylinders. The properties are:} \begin{enumerate} \item {For given geometry and power index $n$ the normalized PDF $(1-q^{2/n})P(X_n)$ is a unique curve expressed through a simple implicit function of the parameter $y_n = (1-q^{2/n})(1-x^2)$ where $x = r/r_{\rm cl}$ is the normalized impact parameter.} \item At the central mass surface density $X_n(0)$ the PDF of spheres has a sharp cut-off while the PDF of cylinders has a pole. \item {At high overdensities the PDF has a well defined maximum at fixed $[X_n]_{\rm max}$.} \item At mass surface densities which are small relative to the maximum {position} the asymptotic PDF approaches asymptotically a power law $P(X_{\rm n})\propto X_n$. \item {In the limit of high overdensities the PDFs approach for $n>1$ at mass surface densities above the peak} asymptotically power laws. They are given by $P(X_n)\propto X_n^{-\frac{n+1}{n-1}}$ {in case of spheres and, with the exception of the pole, by $P(X_n)\propto X^{-\frac{n}{n-1}}$ in case of cylinders. For given overdensity the asymptote is a better approximation for steeper density profiles (larger $n$).} \item {For $n<1$ the PDF has a strong cutoff and is limited to a maximum mass surface density $X_n\le 1/(1-n)$.} \end{enumerate} The slope {of the PDF} at high mass surface densities can also be obtained assuming a simple power law profile for spherical clouds {(e.g. \citealt{Kritsuk2011,Federrath2013a}).} But it should be considered that this profile only is an asymptotic behavior {in} the limit of high overdensities and seems more appropriate for collapsing clouds while most condensations might not be in such a state. As shown in the paper in general the {shape of the PDF} is not a power law. Further, the profile would produce a nonphysical high probability at low mass surface densities. The derived properties are related to background subtracted structures within molecular clouds. They are therefore not directly applicable to measurements of the global PDF of molecular clouds which is a statistical mean of different properties not only of the condensed structures but the surrounding medium as well. For instance the tail at high mass surface densities seen in the PDF of star forming molecular clouds may have different physical explanations. It also need to be considered that the functional form of the PDF is affected by an additional background. In case of filaments the situation is furthermore complicated through a possible variation of the inclination angle. Those problems are addressed in a following paper \citep{Fischera2014b} based on isothermal {self-gravitating} pressurized spheres and cylinders. \begin{acknowledgement} This work was supported by grants from the Natural Sciences and Engineering Research Council of Canada and the Canadian Space Agency. The author likes to thank Prof. P. G. Martin and Dr. Richard Tuffs for his support, Quang Nguyen Luong for reading the manuscript and his helpful comments, and the unknown referee for the suggestions. \end{acknowledgement} \appendix | 14 | 3 | 1403.0454 |
1403 | 1403.7873_arXiv.txt | Fast Radio Bursts (FRBs) are new transient radio sources discovered recently. Because of the angular resolution restriction in radio surveys, no optical counter part has been identified yet so it is hard to determine the progenitor of FRBs. In this paper we propose to use radio lensing survey to constrain FRB progenitors. We show that, different types of progenitors lead to different probabilities for a FRB to be gravitationally lensed by dark matter halos in foreground galaxies, since different type progenitors result in different redshift distributions of FRBs. For example, the redshift distribution of FRBs arising from double stars shifts toward lower redshift than of the FRBs arising from single stars, because double stars and single stars have different evolution timescales. With detailed calculations, we predict that the FRB sample size for producing one lensing event varies significantly for different FRB progenitor models. We argue that this fact can be used to distinguish different FRB models and also discuss the practical possibility of using lensing observation in radio surveys to constrain FRB progenitors. \vspace{0.3cm} \noindent {\bf Star formation, Radio sources, Gravitational lenses and luminous arcs, Cosmology} \vspace{0.3cm} \noindent {\bf PACS numbers:} 97.10.Bt, 98.70.Dk, 98.62.Sb, 98.80.-k | Fast Radio Bursts (FRBs) were discovered in recent years \cite{Thorn2013}. A burst rate of about ten thousand bursts per day over the entire sky was deduced \cite{cor2013}. However, it is still not possible to determine the astrophysical origins and the redshift distribution of FRBs, because no optical counterpart has been discovered yet and hence no redshift of FRBs has been measured. In this paper we propose that strong gravitational lensing (SGL) of FRBs by dark matter halos in foreground galaxies may have a critical role in the study of FRB progenitors. Radio surveys for FRBs suffer from the restriction in angular resolutions, which makes it difficult to identify the respective optical counterparts. But for transient sources like FRBs, the SGL properties may be used for identifying the progenitor of FRBs. Different FRB progenitor models result in different FRB redshift distributions, and hence different FRB lensing probabilities, that is, different expected sample size for producing one lensing case. Due to the transient nature of FRBs and the time-delay between images in gravitational lensing, it is possible to identify lensing events of FRBs in a radio survey despite the restriction of angular resolution. The standard way for identifying gravitational lensing images is based on two criteria: (1) their light curves have identical shapes and identical durations; (2) their spectra have the same shape. The second criteria is critical, because the first criteria is relatively easy to check. If a radio telescope has several frequency channels, then the ratio between fluxes in any two channels can be used to judge if the radio spectra of the two FRB images have the same shape. A similar dispersion measure (DM) may be another criteria to justify a lensing case. The statistical lensing probability obtained from a long period FRB survey for multiple lensing images may be used to constrain the associated distribution profile in the FRB redshift space. In addition, if a FRB strong lensing event is confirmed in a well-studied lens system, the redshift of the FRB may be inferred. We calculate the probability for gravitational lensing of FRBs by the singular isothermal sphere (SIS) dark matter halos in foreground galaxies, assuming that the mass function of the halos is given by the Press-Schechter function and the Universe is described by the standard LCDM model. Later, we build four types of FRB models, each corresponding to a FRB progenitor model with different redshift distributions. Lastly, we calculate the expected FRB sample size for producing one lensing case and compare the results for different FRB models. | We have calculated the probability for a FRB gravitationally lensed by dark matter halos in foreground galaxies. For a given population of FRBs, the needed sample size for producing a lensing event depends on the redshift distribution of the FRBs. We have built four different FRB redshift distribution models for four types of FRB progenitors. We found that, for the four different models, the sizes of the FRB samples needed for producing a lensing event are very different. The results are summarized in Table \ref{tab1}. From Table \ref{tab1} we can note that, for typical parameters adopted in the paper, the number of FRBs contained in the sample can differ by an order of magnitude. At least in principle, we can carry out a large FRB lens survey jointly with many radio telescopes to distinguish different FRB progenitors. The observation of FRBs has revealed a burst rate of $\sim 10^4$ per day \cite{cor2013}, which is sufficiently high. However, we note that the sample size calculated here is for a continuous survey in a fixed field of the sky. If the survey is not continuous, a lensing event may be missed out due to the transient nature of FRBs: for example, for a double-image lensing event, an image may not be detected if it occurs during a gap between two time segments of a survey because of the time-delay of gravitational lensing. Then this lensing event will not be identified. For most of the ground based interferometers, the observation time for a fixed field of the sky is discontinuous with gaps. All FRBs occurring during those gaps will miss the detection, including those FRB lensing events. For example, if the observation is carried on for $12$ hours per day, the sampling efficiency for FRB lensing events will be $f = 1/2$ since one of the double images in a lensing event may not be detected. In this case, the expected sample size given in Table \ref{tab1} should be enlarged by a factor $1/f=2$. However, the ratio of the sample size between different models is not affected. If the fixed field is $S (deg^2)$ and the survey program lasts $T (days)$, we have: \begin{eqnarray} S*T = \frac{S_0 N}{f^2 R}, \label{st01} \end{eqnarray} where $S_0= 41,253 deg^2$ is the area of the whole sky, $N$ is the total number of FRBs needed in the sample, and $R$ is the rate of FRBs. Given that $R = 10^4$ per day over the entire sky, $f = 1/2$, and $N= 490 $ for the FRB model without a time delay, we have $S*T = 8,000 deg^2 \cdot day$. This is a typical value of $S*T$ needed for a radio survey for FRB lensing events. From the above results we can predict that if a continuous FRB survey for a fixed field of sky of approximately $10 {\rm deg}^2$ lasts approximately 2 years, it will be possible to constrain FRB progenitors without counterpart observation. | 14 | 3 | 1403.7873 |
1403 | 1403.0381_arXiv.txt | On 20 August 2010 an energetic disturbance triggered large-amplitude longitudinal oscillations in a nearby filament. The triggering mechanism appears to be episodic jets connecting the energetic event with the filament threads. In the present work we analyze this periodic motion in a large fraction of the filament to characterize the underlying physics of the oscillation as well as the filament properties. The results support our previous theoretical conclusions that the restoring force of large-amplitude longitudinal oscillations is solar gravity, and the damping mechanism is the ongoing accumulation of mass onto the oscillating threads. Based on our previous work, we used the fitted parameters to determine the magnitude and radius of curvature of the dipped magnetic field along the filament, as well as the mass accretion rate onto the filament threads. These derived properties are nearly uniform along the filament, indicating a remarkable degree of cohesiveness throughout the filament channel. Moreover, the estimated mass accretion rate implies that the footpoint heating responsible for the thread formation, according to the thermal nonequilibrium model, agrees with previous coronal heating estimates. We estimate the magnitude of the energy released in the nearby event by studying the dynamic response of the filament threads, and discuss the implications of our study for filament structure and heating. | \label{sec:intro} High-cadence H$\alpha$ observations of large-amplitude longitudinal (LAL) oscillations in solar filaments were first reported by \citet{jing2003}. These oscillations consist of rapid motions of the plasma along the filament, with displacements comparable to the filament length. Since then, a few more events have been identified \citep{jing2006,vrsnak2007,li2012,zhang2012}. In all of the observed oscillations the period ranges from 0.7 to 2.7 hours, with velocity amplitudes from 30 to $100~\mathrm{km~s^{-1}}$. The accelerations are considerable, in many cases comparable to the solar gravitational acceleration. In addition, the oscillations are always triggered by a small energetic event close to the filament. Explaining the LAL oscillations is very challenging because the restoring force responsible for the huge acceleration must be very strong. The energy of the oscillation is also enormous, because the filament is massive and large velocities are generated. However, the motions damp quickly in a few periods, implying a very effective damping mechanism. Several models have been proposed to explain the restoring force and damping mechanism of the LAL oscillations \citep[see review by][]{tripathi2009}, but most do not successfully describe the thread motions. Recently, we studied the oscillations of threads forming the basic components of a filament \citep{luna2012a} in a 3D sheared arcade \citep{devore2005}. In this model, the threads reside in large-scale dips on highly sheared field lines within the overall magnetic structure. We found that the restoring force is mainly gravity, and the pressure forces are small \citep{luna2012b}. This result has been confirmed with numerical simulations by \citet{li2012} and \citet{zhang2013a}. We also studied analytically the normal modes of plasma on a dipped field line, and found that the oscillation generally involves the superposition of two components: gravity-driven modes and slow modes associated with pressure gradients \citep{luna2012c}. However, for typical filament plasma properties the gravity-driven modes dominate. This type of oscillation resembles the motion of a gravity-driven pendulum, where the frequency depends only on the solar gravity and the field-line dip curvature. We estimated the minimum value of the magnetic field at the dips and found agreement with previous estimates and observed values. Additionally, this study revealed a new method for measuring the radius of curvature of the filament dips. These studies demonstrated that the LAL oscillations are strongly related to the filament-channel geometry. We identified the damping of the LAL oscillations as a natural consequence of the thermal nonequilibrium process most likely responsible for the formation and maintenance of the cool filament threads \citep{luna2012b}. In this process, the localized footpoint heating produces chromospheric evaporation and subsequent collapse of the evaporated mass into cool condensations in the corona \citep[e.g.,][]{antiochos1999, karpen2006}. As long as the heating remains steady, the threads continuously accrete mass and grow in length. Non-adiabatic effects (i.e., thermal conduction and optically thin radiation) also contribute weakly to the damping \citep{luna2012b, zhang2013a}. The damping times obtained with this model generally agree with the observed values. Because LAL oscillations in filaments are observed to occur after small energetic events nearby, we speculated that local heating and/or flows from the energetic event triggers and drives oscillations in filament threads magnetically linked to the event site \citep{luna2012b}. \citet{zhang2013a} modeled the effects of two types of perturbations on a filament thread -- impulsive heating at one leg of the loop and impulsive momentum deposition -- and determined that both can cause oscillations. In this work we study a LAL filament oscillation, its trigger, and subsequent damping observed on 20 August 2010 by the Atmospheric Imaging Assembly (AIA) instrument on the Solar Dynamics Observatory (SDO) \citep{Lemen2012a}. We also utilize co-temporal magnetograms from the Helioseismic and Magnetic Imager (HMI) instrument on SDO \citep{scherrer2012}, to provide the overall magnetic context and local connectivity between the filament and the triggering event. The oscillation properties yield the geometry of the magnetic field supporting the cool plasma and the mass accretion rate. This information leads to fundamental conclusions about the filament structure, as well as the dynamics and energetics of the triggering event. Some preliminary results of this work were shown in \citet{Knizhnik2014}. In \S\ref{sec:observation} we describe the data and sequence of events, while \S\ref{sec:trigger} discusses the event that triggered the oscillations. In \S\ref{sec:analysis} our methodology for measuring the filament oscillation and deriving key parameters is described. \S\ref{sec:results} presents the results and examines their implications for the oscillation mechanism, mass accretion, and energization. Based on these results, we describe the most likely filament structure in \S\ref{sec:structure}. Our findings and conclusions are summarized in \S\ref{sec:conclusion}. | \label{sec:conclusion} In this work we have studied the LAL oscillations of a filament and the associated trigger, using observations and theory to determine key properties of the filament and trigger. The LAL oscillations are displacements of low $\beta$ plasma supported by the magnetic field, so the direction of motion reflects the direction of the local magnetic field. Our results agree with the few direct measurements of the orientation of the magnetic field in filament threads with respect to the associated PIL. Thus, the oscillation analysis described in this work is a novel tool to determine the orientation of the filament magnetic field. We determined fundamental characteristics of the LAL oscillations and the triggering flows by fitting curves to the SDO/AIA time-distance diagrams. We used a exponentially decaying sinusoid and a modified Bessel function to fit the oscillation data, respectively representing constant-mass and mass-accreting solutions. Both fits generally match the data well, but the Bessel function fits the initial stage of the oscillation significantly better. Therefore mass accretion is likely to play a major role in damping the oscillations rapidly. We conclude that our earlier model for the damped LAL oscillations \citep{luna2012b, luna2012c} accurately explains the behavior of this filament. Using our earlier analytic approximations, we determined the radius of curvature and minimum strength of the magnetic field lines that support the filament plasma, and inferred the magnetic structure of the oscillating portions of the filament (Figures \ref{fig:sketch} and \ref{fig:cartoon}). We found that the geometry varies little along the filament, demonstrating that the different parts of the filament form a quasi-coherent structure whose origin and subsequent evolution remain linked. The resulting structure, as well as the presence of LAL oscillations, are compatible with the two leading magnetic-structure models --- the flux rope and the sheared arcade --- which predict that the bulk of the filament plasma resides in the dips \citep{mackay2010}. Although both models predict dipped field lines that can host oscillating threads, the distribution of cool mass condensed in these dips through thermal nonequilibrium has been studied in depth only for the sheared arcade model. The thermal nonequilibrium model for filament mass formation predicts that the existing threads accrete material at the same rate as the chromospheric evaporation rate, as long as the standard coronal heating is localized at the footpoints. We established previously that this continuous accretion of mass is responsible for the strong damping of the LAL oscillations \citep{luna2012b}. The mass accretion rate of the filament threads computed in the present study agrees with the predictions of our thermal nonequilibrium model \citep{karpen2006, luna2012c}, and hence with the well-established quiet-Sun coronal heating rate of \cite{withbroe1977}. Based on our results, we suggest that LAL oscillations provide a new opportunity for constraining coronal heating models beyond the usual analyses of AR coronal loops. We propose the following general picture of the event and the filament structure. A reconnection process takes place close to PIL at the northern side. The resulting jet plasma flows along the filament channel field lines at a projected speed of $\sim95~\mathrm{km~s^{-1}}$. These field lines only connect with some parts of the filament, such that the flow reaches the SE part and the NW barb. Threads in these regions are pushed by the hot flows, then oscillate in the dips with a motion resembling a pendulum. Other complex physical phenomena could take place when the hot flows reach the cool plasma, but a detailed study of this interaction is beyond the scope of this work. The restoring force of the oscillations is the projected gravity along the dips, as our model predicts. Continuous, localized coronal heating produces evaporation of chromospheric plasma that accretes onto the already formed filament threads. This mass accretion is responsible for the initial strong damping of the oscillations. More observations of LAL oscillations in filaments and the associated triggering events, together with detailed simulations of the response of filament threads to hot flows, are needed to improve our understanding of this intriguing phenomenon. Additional theoretical modeling of LAL oscillations also will advance the use of seismology to probe the ambient physical conditions in filaments, as demonstrated here. Further analyses of LAL oscillations would benefit as well from greater understanding of the coronal heating mechanism, as the likely driver of mass accretion onto filament threads. We anticipate significant progress on these questions to be made in the near future by the combined capabilities of SDO, IRIS, and the upcoming Solar Orbiter mission. | 14 | 3 | 1403.0381 |
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