subfolder
stringclasses 367
values | filename
stringlengths 13
25
| abstract
stringlengths 1
39.9k
| introduction
stringlengths 0
316k
| conclusions
stringlengths 0
229k
| year
int64 0
99
| month
int64 1
12
| arxiv_id
stringlengths 8
25
|
|---|---|---|---|---|---|---|---|
1609 | 1609.05589_arXiv.txt | In a Mirror Twin World with a maximally symmetric Higgs sector the little hierarchy of the Standard Model can be significantly mitigated, perhaps displacing the cutoff scale above the LHC reach. We show that consistency with observations requires that the $Z_2$ parity exchanging the Standard Model with its mirror be broken in the Yukawa couplings. A minimal such effective field theory, with this sole $Z_2$ breaking, can generate the $Z_2$ breaking in the Higgs sector necessary for the Twin Higgs mechanism. The theory has constrained and correlated signals in Higgs decays, direct Dark Matter Detection and Dark Radiation, all within reach of foreseen experiments, over a region of parameter space where the fine-tuning for the electroweak scale is 10-50\%. For dark matter, both mirror neutrons and a variety of self-interacting mirror atoms are considered. Neutrino mass signals and the effects of a possible additional $Z_2$ breaking from the vacuum expectation values of $B-L$ breaking fields are also discussed. | An intriguing idea, that has persisted over several decades, is that of the Mirror World: the Standard Model, with quarks and leptons $(q,l)$ and gauge interactions $SU(3) \times SU(2) \times U(1)$, is supplemented by an identical sector where mirror quarks and leptons $(q',l')$ interact via mirror gauge interactions $SU(3)' \times SU(2)' \times U(1)'$. There are two prime motivations for this idea. The discrete symmetry that interchanges the ordinary and mirror worlds can be interpreted as spacetime parity, $P$, allowing a neat restoration of parity \cite{Lee:1956qn,Kobzarev:1966qya}. Secondly, mirror baryons are expected to be produced in the early universe and to be sufficiently stable to yield dark matter, and this may lead to an understanding of why the cosmological energy densities of baryons and dark matter are comparable. A third, more recent motivation for the Mirror World arises from the absence so far of new physics at colliders to explain the origin of the weak scale. If the Higgs doublets of the two sectors $(H,H')$ possess a potential with maximal $SO(8)$ symmetry at leading order, the observed Higgs boson becomes a pseudo-Nambu-Goldstone boson with a mass that is insensitive to the usual Standard Model (SM) quadratic divergences; this is the Twin Higgs idea~\cite{Chacko:2005pe}. Furthermore, if the symmetric quartic coupling of this potential, $\lambda$, is large relative to the SM quartic coupling, $\lambda_{\rm SM}$, the Mirror World reduces the amount of fine-tuning by a factor of $ 2 \lambda / \lambda_{\rm SM}$ to reach any particular UV cutoff of the effective theory. Since we now know that $\lambda_{\rm SM} = 0.13$ is small, this improvement can be very significant, allowing a Little Hierarchy between the weak scale and the UV cutoff, which may be beyond the LHC reach. In this paper we formulate a minimal, experimentally viable, low energy effective theory for this idea, Minimal Mirror Twin Higgs, and study its signals. This is a pressing issue: mirror baryon dark matter, the Twin Higgs mechanism and consistency with cosmological limits on the amount of dark radiation all require a breaking of parity, $P$ \cite{Barbieri:2005ri}. How is this to be accomplished? We do not attempt a UV completion, whether supersymmetric~\cite{Chang:2006ra,Craig:2013fga} or with composite Higgs~\cite{Batra:2008jy,Geller:2014kta,Barbieri:2015lqa,Low:2015nqa}, but note that both account for the approximate $SO(8)$ symmetry of the Higgs potential. The $SO(8)$ invariant quartic interaction contains an interaction $H^\dagger H H'^\dagger H'$ thermally coupling the two sectors at cosmological temperatures above a few GeV, so that the bound on dark radiation provides a very severe constraint on Mirror Twin Higgs. The Twin Higgs mechanism requires a parity breaking contribution to the Higgs mass terms in the potential, $\Delta m_H^2$. We find this term by itself to be insufficient to solve the dark radiation problem, nomatter what other interactions connect the two sectors, at least for fine-tunings above the percent level. This then implies that the Yukawa couplings of the two sectors differ, $y' \neq y$. Hence we introduce an effective theory, Minimal Mirror Twin Higgs, where all $P$ violation arises from a breaking of flavor symmetry, yielding different Yukawa couplings in the two sectors. This single source of $P$ violation leads simultaneously to three key results \begin{itemize} \item The $\Delta m_H^2$ term necessary for the Twin Higgs mechanism is generated via $q'$ loops. \item The mirror QCD phase transition temperature is raised above the decoupling temperature of the two sectors, solving the problem of excessive dark radiation. \item The mirror baryon mass is raised, allowing viable dark matter. \end{itemize} A striking signature would be the discovery at the LHC, or a future collider, of the mirror Higgs itself, decaying to $WW$ or $ZZ$; see \cite{Barbieri:2005ri} and Fig.~10 of \cite{Buttazzo:2015bka}. As the mirror Higgs mass depends on the $SO(8)$ invariant quartic, $\lambda$, it could be beyond the LHC range, and here we focus on other signals. The size of $P$ breaking in the Yukawa couplings to obtain the above three results leads to a preferred range of the lightest $q'$ mass of $(2-20)$ GeV, leading us to compute signals for the following quantities \begin{itemize} \item The signal strength, $\mu$, and the invisible width, $\Gamma(h \rightarrow \rm inv)$, of the Higgs boson. \item The amount of dark radiation, $\Delta N_{\rm eff}$. \item The direct detection rate for mirror baryon dark matter from Higgs exchange. \item The effective sum of neutrino masses affecting large scale structure and the CMB. \end{itemize} These signals are tightly correlated as they all depend on the Higgs portal between the two sectors, the ratio of the weak scales, and on the masses of the light $q'$. For dark matter, both mirror neutrons and a variety of self-interacting mirror atoms are considered. After a brief review of the Twin Higgs mechanism in section~\ref{sec:TH}, we demonstrate that the breaking of $P$ in the Yukawa couplings is necessary in section~\ref{sec:Z2BYuk}. We define the Minimal Mirror Twin Higgs theory in section~\ref{sec:MinEFT}, and discuss the consequences of breaking $P$ in the Yukawa couplings. We constrain the $q'$ Yukawa couplings from $\Gamma(h \rightarrow \rm inv)$ and $\Delta m_H^2$, and study how large the mirror QCD phase transition temperature $T'_c$ can be. In section~\ref{sec:cosmology} we study the cosmological history of the two sectors when the only communication between them arises from the Higgs interaction $H^\dagger H H'^\dagger H'$ and find that the decoupling temperature can be lower than $T'_c$, allowing $\Delta N_{\rm eff}$ to lie inside the observational limit. We predict the amount of dark radiation and the effective sum of neutrino masses. In section~\ref{sec:mixing} we examine an alternative cosmology when communication between the sectors is dominated by kinetic mixing of the hypercharge gauge bosons. We study a variety of candidates for mirror dark matter in section~\ref{sec:DM}, and find that the $H^\dagger H H'^\dagger H'$ interaction, together with the enhanced $q'$ Yukawa couplings, will allow direct detection at planned experiments over a large part of the mass range. In section~\ref{sec:heavynu'} we briefly study $\Delta N_{\rm eff}$ and dark matter candidates when additional $P$ breaking arises from the absence of Majorana masses for right-handed mirror neutrinos. In the Appendix we show that a PQ symmetry common to both sectors allows a solution to the strong CP problem, with the axion mass enhanced by the mirror sector by a factor of order $10^3$, leading to the possibility that $f_a$ is of order 10 TeV. | The Twin Higgs mechanism significantly relaxes fine-tuning of the electroweak scale, and allows for a larger cut off scale. The cut off of the Standard Model is \begin{align} \Lambda_{\rm SM} \simeq 1.4~{\rm TeV}\times \left(\frac{\Delta_{\rm SM}}{10}\right)^{1/2}, \end{align} while that of the Twin Higgs theory is \begin{align} \Lambda_{\rm TH} = 5.7~{\rm TeV} \times \left(\frac{\Delta_{\rm TH}}{10}\right)^{1/2} \lambda^{1/2}. \end{align} The minimal theory has no new colored states to be produced at the LHC. It does offer the possibility of discovery modes at the LHC, such as production of the mirror Higgs via Higgs mixing; but the larger cut off may raise the masses of new particles above the LHC reach. Mirror Twin Higgs models, however, predict the existence of extremely light particles, mirror photons and mirror neutrinos, that contribute to the dark radiation of the universe, leading to constraints on a realistic theory. We have found that, independent of the interactions that couple the two sectors, it is {\it necessary} to break the mirror $Z_2$ symmetry in the Yukawa couplings. \noindent \textbf{Minimal Mirror Twin Higgs} \hspace{0.25in} We have constructed a completely realistic effective field theory of Twin Higgs below the cut off $\Lambda_{\rm TH}$. It contains a complete mirror sector, so that a UV completion, which we did not study, can restore spacetime parity symmetry. In Minimal Mirror Twin Higgs, the only $Z_2$ breaking arises from Yukawa couplings and the only communication between the sectors is from Higgs mixing, required by Twin Higgs, and kinetic mixing of hypercharge fields, allowed by gauge invariance. Furthermore, the $Z_2$ breaking Yukawa couplings not only suppress dark radiation to acceptable levels, but generate the $Z_2$ breaking Higgs mass term necessary for the Twin Higgs mechanism and raise the mirror baryon mass as required for realistic dark matter. The $Z_2$ breaking Yukawa couplings induce a sizable invisible branching ratio of the SM-like Higgs boson through its mixing with the mirror Higgs. This, together with the reduction of the Higgs coupling to Standard Model particles, leads to a universal deviation from unity of the Higgs signal-strengths correlated with the masses of mirror fermions, as shown in Fig.~\ref{fig:mass_mu}. Irrespective of the mirror fermions masses, the high-luminosity running of the LHC and the ILC can probe $v'/v<4$ and $10$ respectively. As illustrated in Fig.~\ref{fig:THiggs}, it is non-trivial that Higgs mixing can lead to a decoupling temperature less than the QCD$'$ phase transition temperature, necessary for a solution of the dark radiation problem. Fig.~\ref{fig:THiggs} shows that, with Higgs mixing alone, there must be light mirror fermions with mass $m_{f'}$ in the range of about $(4-28)$ GeV. The upper bound on $m_{f'}$ gives a lower bound of $\Delta N_{\rm eff}\gtrsim 0.3$. The allowed range for $m_{f'}$ depends on $N_{f'}$, the number of light mirror fermion states, and narrows considerably for larger values of $N_{f'}$. As the upper bound on $m_{f'}$ becomes tighter, $\Delta N_{\rm eff}$ increases, which is shown in Fig.~\ref{fig:Neff Higgs}. The lower bounds on $m_{f'}$ from Figs.~\ref{fig:THiggs} and \ref{fig:Neff Higgs} imply lower bounds on the invisible branching ratio of the SM-like Higgs boson and the universal deviation from unity of the Higgs signal-strengths. For $v'/v=3$, the left panel of Fig.~\ref{fig:Brinv} shows that this branching ratio is typically in the range of 0.05-0.15 that can be probed by high luminosity running of the LHC. For $v'/v=5$, the invisible branching ratio is reduced; the right panel of Fig.~\ref{fig:Brinv} shows that in some cases the signal is as small as 0.002 - 0.01, that could be probed by ILC. Essentially the entire parameter space can be probed. If the kinetic mixing parameter $\epsilon$ is large enough to affect the thermal history of the two sectors near the QCD$'$ phase transition, the allowed range of the light mirror fermion masses is enlarged. In section \ref{sec:mixing} we showed that the upper bound on $m_{f'}$ could be extended as high as 50 GeV, while mirror leptons could be as light as 2 GeV. A mirror quark can be lighter than 2 GeV provided the lightest mirror meson is heavier than 2 GeV. Even including kinetic mixing, the lower bound on $\Delta N_{\rm eff}$ remains about 0.3, but the enlarged range of light mirror fermion masses is important for mirror dark matter. Mirror baryons and leptons are natural candidates for dark matter. Dark matter can be composed of mirror neutrons, mirror atoms, or even a mixture of the two. Such dark matter can be directly detected via Higgs exchange, as illustrated in Fig. \ref{fig:NNp scat}. For low $v'/v$, PandaX and LUX have recently excluded large values for the dark matter mass. In the absence of kinetic mixing, the region of dark matter masses allowed by present limits on dark radiation can be fully explored, up to the $1(2) \sigma$ limit by the XENON1T (LZ) experiment. Adding kinetic mixing to the thermal cosmological history, the limit on $\Delta N_{\rm eff}$ is consistent with lighter dark masses that XENON1T and LZ are unable to probe. For mirror neutron dark matter masses in the (1-10) GeV region, present direct detection limits bound $\epsilon < 10^{-3} - 10^{-4}$ from scattering via a dipole moment. Mirror atomic dark matter with a sizable kinetic mixing such that the thermal history is affected is excluded. In the Minimal Mirror Twin Higgs the seesaw mechanism yields light neutrino masses for both sectors. These neutrinos can be Majorana with those in the mirror sector heavier by a factor $(v'/v)^2$ than the observed neutrinos, leading to important effects in both CMB and LSS. The effective sum on neutrino masses relevant for cosmological data is at least 2.3 times greater than in the Standard Model. Small mixing between these Majorana standard and mirror neutrinos could lead to mirror neutrinos being observed as massive sterile neutrinos. Alternatively the seesaw could lead to Dirac neutrinos, with the mirror states as right-handed neutrinos degenerate with the observed left-handed states, leading to only very small effects on CMB and LSS. The predictions of Eqs.~(\ref{eq:mnuM}, \ref{eq:mnuD}) for the masses of mirror neutrinos, however, rely on the assumption that $Z_2$ breaking does not substantially affect either the neutrino Yukawa couplings or right-handed neutrino masses, and therefore our predictions for the effective sum of neutrino masses, $\left(\sum m_\nu \right)_{\rm eff}$, are less robust than the predictions for $\Delta N_{\rm eff}$, $\Gamma(h \rightarrow \rm inv)$ and the mirror baryon dark matter direct detection cross section. \noindent \textbf{Additional $Z_2$ breaking from $B-L$ vevs} \hspace{0.25in} A large $B-L$ vev can implement the seesaw mechanism for the known neutrinos, while a small or zero $B' - L'$ vev can lead to large Dirac mirror neutrino masses. Without light mirror neutrinos, the minimal $\Delta N_{\rm eff}$ is lowered to 0.08 (0.17) without (with) a massless $B'-L'$ gauge field. Furthermore, there are new possibilities for mirror dark matter. If there is a massless $B'-L'$ gauge boson, dark matter could be $B'_{udd}\nu'$ or $B'_{ddd} \bar{e}'$ atoms, and constraints from self interactions and the ionized fraction are weekend if $\alpha'_{B-L} \gsim \alpha$. Without a massless $B'-L'$ gauge boson, dark matter could be a mixture of $B'_{uud}e'$ atoms and mirror neutrons $B'_{udd}$; constraints from self interactions and the ionized fraction are relaxed. | 16 | 9 | 1609.05589 |
1609 | 1609.09244_arXiv.txt | {} {We present the second ROSAT all-sky survey source catalogue, hereafter referred to as the 2RXS catalogue. This is the second publicly released ROSAT catalogue of point-like sources obtained from the ROSAT all-sky survey (RASS) observations performed with the position-sensitive proportional counter (PSPC) between June 1990 and August 1991, and is an extended and revised version of the bright and faint source catalogues. } {We used the latest version of the RASS processing to produce overlapping X-ray images of $6.4 \degr \times 6.4 \degr$ sky regions. To create a source catalogue, a likelihood-based detection algorithm was applied to these, which accounts for the variable point-spread function (PSF) across the PSPC field of view. Improvements in the background determination compared to 1RXS were also implemented. X-ray control images showing the source and background extraction regions were generated, which were visually inspected. Simulations were performed to assess the spurious source content of the 2RXS catalogue. X-ray spectra and light curves were extracted for the 2RXS sources, with spectral and variability parameters derived from these products. } {We obtained about 135,000 X-ray detections in the $0.1-2.4$ keV energy band down to a likelihood threshold of 6.5, as adopted in the 1RXS faint source catalogue. Our simulations show that the expected spurious content of the catalogue is a strong function of detection likelihood, and the full catalogue is expected to contain about $30$~\% spurious detections. A more conservative likelihood threshold of % 9, on the other hand, yields % about 71,000 detections with a $5$~\% spurious fraction. We recommend thresholds appropriate to the scientific application. X-ray images and overlaid X-ray contour lines provide an additional user product to evaluate the detections visually, and we performed our own visual inspections to flag uncertain detections. Intra-day variability in the X-ray light curves was quantified based on the normalised excess variance and a maximum amplitude variability analysis. X-ray spectral fits were performed using three basic models, a power law, a thermal plasma emission model, and black-body emission. Thirty-two large extended regions with diffuse emission and embedded point sources were identified and excluded from the present analysis. } {The 2RXS catalogue provides the deepest and cleanest X-ray all-sky survey catalogue in advance of eROSITA.} | \label{sec:Introduction} The ROSAT all-sky survey (RASS) was the first to scan the whole sky with a powerful imaging X-ray telescope operating in the $0.1 - 2.4$\,keV band \citep{Truemper1982}. The Wolter type I mirror system \citep{Aschenbach1988} was exceptionally well suited for the sky survey operation because of the very low micro-roughness of the mirrors ($< 0.3$\,nm), which was responsible for the excellent contrast of the X-ray images. The focal plane detector used for the sky survey, the position-sensitive proportional counter (PSPC), had a five-sided anti-coincidence system that reduced the particle background with an efficiency of $99.85\,\%$ \citep{Pfeffermann1986,Pfeffermann2003}. This efficient anti-coincidence veto design resulted in a low, non-X-ray (particle) background. Another reason for the exceptionally low particle background of ROSAT was the low Earth orbit with orbital height of $\sim 580$\,km and inclination $53^\circ$ (orbital period 96 minutes). The ROSAT survey observations were performed in scanning mode, where the field of view (FOV) of the PSPC detector scanned a two-degrees-wide strip along a great circle over the ecliptic poles within 96 minutes. With a shift of about one degree per day, an all-sky survey was completed within half a year. Because of periods of very high background or poor attitude values, some parts of the sky were missed, but re-observed during the February survey in 1991 and the August survey in 1991. Before the main survey between August 1990 and January 1991, the July mini-survey in 1990 was performed for testing and is part of the ROSAT all-sky survey. Still some parts of the sky remained unobserved, which were later (February 1997) covered by pointing observations. The analysis of these will be reported in a separate paper \citep{Freyberg2016a}. The RASS sensitivity \citep{Truemper1993} surpassed that of the Uhuru \citep{1978ApJS...38..357F} and HEAO-1 \citep{1984ApJS...56..507W} surveys by a factor >\,100 in the soft X-ray band. The RASS bright source catalogue (RASS BSC), containing \TotalNumberBscVoges\ sources, was first published in electronic form \citep{Voges1996} and later in a printed version \citep{Voges1999}. This catalogue has served a very large scientific community working in different fields - from solar system objects (Moon, comets, and planets) out to clusters of galaxies and quasars at large cosmological distances. The faint part of the ROSAT all-sky survey (RASS FSC) with \TotalNumberFscVoges\ sources down to a detection likelihood limit of 6.5 was published only in an electronic version \citep{Voges2000}. Both RASS BSC and FSC, which taken together constitute the ROSAT 1RXS catalogue, were based on the so-called RASS-2 processing \citep{Voges1999}. An updated version of the processing (RASS-3) was performed subsequently with event files being made public, but without further documentation. After the all-sky survey, ROSAT performed an extended program of pointed observations that covered a significant part of the sky ($\sim 18\,\%$) with deeper observations. Several other satellite missions with imaging telescopes have gathered data over large areas of the sky, producing large catalogues of X-ray point sources. Of particular note are the ROSAT PSPC pointed catalogue \citep{2RXP}, the XMM-Newton catalogue of pointed observations \citep{3XMM}, the XMM-Newton slew survey \citep{XMMSL1}, and the deep Swift X-ray telescope point source catalogue \citep{2014ApJS..210....8E}. \begin{figure*} \centering \includegraphics[angle=-90,width=180mm,clip=]{rass_fields_psplit.png} \caption{ Structure and numbering scheme of the sky fields in the ROSAT all-sky survey in equatorial coordinates (courtesy K.\,Dennerl, available at the 2RXS web site.) Areas around the North and South ecliptic poles with 1$\degr$ and $15\degr$ radius are indicated at the upper left and lower right corner. } \label{Fig0_field_structure} \end{figure*} The aim of this paper is to present a revised point source data base of the ROSAT all-sky survey (ROSAT 2RXS). The main points of improvement are as follows: \begin{enumerate} \item Use of an improved detection algorithm. \item Reduction of systematic time delays between star tracker and photon arrival time. \item A significantly improved reduction of spurious detections by a careful visual screening of each catalogue entry and the exclusion of large, extended emission regions, in particular from the background-map creation process. \item The provision of X-ray images of 1378 sky fields ($6.4\degr \times 6.4\degr$) covering the whole sky. \item The provision of local maps ($40\arcmin \times 40\arcmin$) for each detected X-ray source. \item The creation of source spectra and light curves and deduction of characteristic parameters. \item The creation of new photon event tables through correcting astrometric errors that are present in the publicly available event files (originating from the RASS-3 processing). \item The delivery of a documented and reproducible point source catalogue. \item Performing extensive simulations to estimate the amount of spurious detections in the 2RXS catalogue as a function of the detection likelihood and other source parameters. \end{enumerate} The total number of entries listed in the 2RXS catalogue is \TotalNumberDetections. Of these, \TotalNumberScreened\ have been flagged as uncertain detections (cf.\ Sect.\ref{sec:ScreeningIndividual}). We have also provided the results of cross correlations of the catalogue with major source catalogues from X-rays and other wavelength bands (see Sect.~\ref{sec:CrossMatches}). \begin{figure*}[thp] \centering \includegraphics[clip=]{i_930304.png} \caption{ Example for a source count image of sky field 930304 in the way it is used for source detection. For the colour representation in this and all following images the ESO/MIDAS colour table {\tt heat} is used with linear intensity scaling. The cyan solid lines confine the $6.4\degr \times 6.4\degr$ field and the scan direction is marked with the yellow line through the centre of the field. 2RXS sources are indicated with green circles. Detections that were manually flagged as being uncertain (see Sect.~\ref{sec:Screening}) are marked with cyan crosses. Coordinates refer to the image centre, the exposure time is the median over the whole image. } \label{Fig_sky_field} \end{figure*} | We have re-analysed the photon event files from the ROSAT all-sky survey. The main goal was to create a catalogue of point-like sources, which is referred to as the 2RXS source catalogue. We improved the reliability of detections by an advanced detection algorithm and a complete screening process. New data products were created to allow timing and spectral analysis. Photon event files with corrected astrometry and Moon rejection (RASS-3.1 processing) were made available in FITS format. The 2RXS catalogue will serve as the basic X-ray all-sky survey catalogue until eROSITA data become available. In this paper we list the most interesting objects in terms of their timing and spectral properties. A discussion of the science highlights is beyond the scope of the paper. With the publication of the 2RXS catalogue and its data products, the detailed science specific exploration is now available for the astrophysical community. The experience gained by the High-Energy Group at MPE in creating the new ROSAT all-sky survey X-ray source catalogue is being and will be fed into the data reduction analysis and scientific exploration of the forthcoming eROSITA all-sky survey. | 16 | 9 | 1609.09244 |
1609 | 1609.07259_arXiv.txt | Detection and classification of transients in data from gravitational wave detectors are crucial for efficient searches for true astrophysical events and identification of noise sources. We present a hybrid method for classification of short duration transients seen in gravitational wave data using both supervised and unsupervised machine learning techniques. To train the classifiers we use the relative wavelet energy and the corresponding entropy obtained by applying one-dimensional wavelet decomposition on the data. The prediction accuracy of the trained classifier on nine simulated classes of gravitational wave transients and also LIGO's sixth science run hardware injections are reported. Targeted searches for a couple of known classes of non-astrophysical signals in the first observational run of Advanced LIGO data are also presented. The ability to accurately identify transient classes using minimal training samples makes the proposed method a useful tool for LIGO detector characterization as well as searches for short duration gravitational wave signals. | } Table~\ref{tab:glitches} lists the transients used in our analysis. Standard searches for compact binary coalescences use matched filtering as the base algorithm~\cite{SathyaSanjeev91}, while the burst searches primarily look for excess power in the data along with time coincidence to trigger a detection \cite{Klimenko:2004qh,Klimenko_CWB_2008}. Both these searches are followed by extensive sanity checks, where the auxiliary channels insensitive to astrophysical signals are inspected to rule out possible terrestrial coupling~\cite{abbott2016characterization}. Auxiliary channels are often in thousands and their coupling with the GW strain sensitive channel is seen to fluctuate in time due to the dynamic nature of the instrument. This often makes the auxiliary channel veto procedure a daunting task. Incorporating a machine learning based veto procedure to identify well known classes of non-astrophysical transients can help discern the trigger right at the strain channel and thus reduce false alarms. \begin{figure*}[t] \includegraphics[width=0.49\textwidth]{Fig1a} \includegraphics[width=0.49\textwidth]{Fig1b} \caption{Left panel depicts typical transient events (SNR set to 50 for better visualisation). Wavelet energy median distribution for simulated data (SNR varied from 8 to 100) shown in the right panel}\label{fig:ip}. \end{figure*} | We have convincingly demonstrated the resourcefulness of machine learning in detector characterization and burst signal analysis in LIGO like complex instruments. We showed that an effective feature extraction technique, in conjunction with an efficient classifier, can be used to classify a variety of transients in practical situation involving real data. We used relative wavelet energy, wavelet entropy and kurtosis as a possible parameter set for classifier input. This, coupled with a difference boosting neural network, was very accurate in discerning between classes with slightly different morphology and possibly different physical origin. The usefulness of the method was shown in our analysis where we could do an accurate targeted search for a specific glitch using minimal training sets. The parameter set used here can be expanded to include other features which can aid the classification even when the corresponding values are unavailable for other classes. The special construction of the classifier makes sure that it does not suffer from the curse of dimensionality unlike most neural network classifiers. Hence the feature set can be expanded in future without causing much computational overhead. Combining class information along with multi-channel coincidence analysis will help to narrow down to the cause for a particular kind of transient present in the data. If there is good enough reason to believe that the trigger is non-astrophysical then glitch based vetoes can be applied to those times. This would lower background triggers in search pipelines thus enhancing confidence in the true detections. We plan to develop such a data quality vector which can be used to directly veto low latency triggers produced by search pipelines looking for astrophysical signals. | 16 | 9 | 1609.07259 |
1609 | 1609.07129_arXiv.txt | We report the detection of coherent pulsations from the ultraluminous X-ray source NGC\,7793~P13. The $\approx$0.42\,s nearly sinusoidal pulsations were initially discovered in broadband X-ray observations using \xmm and \nustar taken in 2016. We subsequently also found pulsations in archival \xmm data taken in 2013 and 2014. The significant ($\gg5\sigma$) detection of coherent pulsations demonstrates that the compact object in P13 is a neutron star, and given the observed peak luminosity of $\approx10^{40}$\,\ergps (assuming isotropy), it is well above the Eddington limit for a 1.4\msun accretor. This makes P13 the second ultraluminous X-ray source known to be powered by an accreting neutron star. The pulse period varies between epochs, with a slow but persistent spin up over the 2013--2016 period. This spin-up indicates a magnetic field of $B\approx1.5\times10^{12}$\,G, typical of many Galactic accreting pulsars. The most likely explanation for the extreme luminosity is a high degree of beaming, however this is difficult to reconcile with the sinusoidal pulse profile. | \label{sec:intro} Due to their high luminosities ($L >10^{39}\,\ergps$), most ultraluminous X-ray sources (ULXs) have been thought to harbor black holes (BHs) with masses ranging from $M\approx10\,\msun$, consistent with a stellar remnant accreting above the Eddington rate (e.g., \citealt{Poutanen07, Middleton15}), to intermediate mass BHs \citep[$M\approx10^{2-5}$\,\msun; e.g.,][]{Miller04} in a sub-Eddington disk accretion regime. The discovery of coherent pulsations in the ULX M82~X-2 showed that the compact object in this system is a neutron star \citep{bachetti14a}. M82~X-2 reaches X-ray luminosities of $2\times10^{40}\,\ergps$, demonstrating that accreting neutron stars can reach luminosities more than 100 times Eddington (assuming $M_{\rm{NS}}\approx1.4\,\msun$). Accreting magnetized neutron stars can reach these apparent super-Eddington luminosities through a number of mechanisms. High magnetic fields collimate the accretion flow, allowing material to accrete onto the polar regions while radiation escapes from the sides of the column \citep{basko76a}. In addition, large magnetic fields reduce the scattering cross section for electrons \citep{herold79a}, reducing the radiation pressure and increasing the effective Eddington luminosity. The combination of these effects with the consequent geometric beaming have been used to explain known super-Eddington local sources such as SMC X-1 \citep[e.g.,][]{coe81a}. A very highly magnetized (magnetar-like) neutron star has been invoked to explain the extreme luminosity of M82~X-2 \citep[e.g.,][]{eksi15a,dallosso15a, mushtukov15a}. It is difficult, however, to explain the near-sinusoidal pulse profile in the context of a highly beamed system. In contrast, some theoretical work suggests the field in M82~X-2 may be relatively low ($10^9$\,G), based on the ratio of the spin-up rate to the luminosity, which is an order of magnitude lower than typical X-ray pulsars \citep{kluzniak15a}. These authors argue that a disk truncated at a large radius, as would occur for a high B-field system, would not provide the required lever arm to power the observed spin-up. The nature of ULX pulsars is very much in question, since no model can explain all the observed characteristics. The ULX NGC\,7793 P13 (hereafter P13; \citealt{Read99}) is one the few ULXs with a dynamical mass constraint of the compact object and a well classified optical companion \citep[spectral type B9Ia,][]{Motch11}. Optical monitoring revealed a $\approx$64\,d photometric period also seen in the radial velocity of the \heii\ emission. Adopting this as the orbital period of the binary system, \cite{motch14a} derive a dynamical mass estimate for the accretor of 3--15\,\msun, assuming a BH. This constraint, together with a peak luminosity of $L_{\rm X} > 6\times10^{39}$ \ergps, makes P13 a prime example for a super-Eddington system. Here we report on new \xmm \citep{xmmref} and \nustar \citep{harrison13a} X-ray observations of P13 in which we detect coherent pulsations, requiring that P13 hosts a highly super-Eddington neutron star accretor. \footnote{During preparation of this manuscript, \citet[submitted]{israel16a} reported on an independent discovery of this period in archival \xmm data. Our study includes newer \xmm and \nustar data which extend the investigation to higher energies and cover a longer time range.} | We detected coherent X-ray pulsations from the ULX \ulx, making this only the second confirmed ULX pulsar after M82~X-2 \citep{bachetti14a}. Its properties seem to be in line with a high luminosity extension of known Galactic neutron star binaries. Between observations in 2013 and 2016 we see a significant spin-up from which we estimate the magnetic field strength to be $\approx1.5\times10^{12}$\,G, typical of Galactic systems. Continued monitoring of the pulse period evolution of this remarkable source will be of particular value and help us understand if high variability is a tell-tale sign of super-Eddington neutron stars. | 16 | 9 | 1609.07129 |
1609 | 1609.07403_arXiv.txt | Despite their use as cosmological distance indicators and their importance in the chemical evolution of galaxies, the unequivocal identification of the progenitor systems and explosion mechanism of normal type Ia supernovae (SNe~Ia) remains elusive. The leading hypothesis is that such a supernova is a thermonuclear explosion of a carbon-oxygen white dwarf, but the exact explosion mechanism is still a matter of debate. Observation of a galactic SN~Ia would be of immense value in answering the many open questions related to these events. One potentially useful source of information about the explosion mechanism and progenitor is the neutrino signal because the neutrinos from the different mechanisms possess distinct spectra as a function of time and energy. In this paper, we compute the expected neutrino signal from a gravitationally confined detonation (GCD) explosion scenario for a SN~Ia and show how the flux at Earth contains features in time and energy unique to this scenario. We then calculate the expected event rates in the Super-K, Hyper-K, JUNO, DUNE, and IceCube detectors and find both Hyper-K and IceCube will see a few events for a GCD supernova at 1 kpc or closer, while Super-K, JUNO, and DUNE will see events if the supernova is closer than ${\sim}0.3$ kpc. The distance and detector criteria needed to resolve the time and spectral features arising from the explosion mechanism, neutrino production, and neutrino oscillation processes are also discussed. The neutrino signal from the GCD is then compared with the signal from a deflagration-to-detonation transition (DDT) explosion model computed previously. We find the overall event rate is the most discriminating feature between the two scenarios followed by the event rate time structure. Using the event rate in the Hyper-K detector alone, the DDT can be distinguished from the GCD at 2$\sigma$ if the distance to the supernova is less than $2.3\;{\rm kpc}$ for a normal mass ordering and $3.6\;{\rm kpc}$ for an inverted ordering. | } The stellar explosions known as thermonuclear - or Type Ia - supernovae are important phenomena in astrophysics. SNe~Ia are the distance indicators \cite{Phillips1993,Phillips1999} which indicate the Universe is accelerating \cite{Riess1998,Schmidt1998,Perlmutter1999} and SNe~Ia are also major contributors to the chemical evolution of galaxies \cite{1997NuPhA.621..467N,2009ApJ...707.1466K} leading to notable changes in the rate at which, for example, oxygen and iron are enriched when the contribution from SNe~Ia becomes significant. But despite their importance, little is known about the progenitors of these supernovae. Only in a few cases for nearby SNe~Ia, such as SN 2011fe or SN 2014J, do pre-explosion images place strong constraints upon the luminosity of the progenitor system \cite{2011Natur.480..348L,2014ApJ...790....3K,2014MNRAS.442.3400N}. Nevertheless, various arguments point towards a paradigm that SNe Ia are the explosion of a carbon-oxygen white dwarf in a binary system that accretes sufficient mass from its companion to begin explosive carbon burning. The source of the accreted mass may be another, smaller, white dwarf (a double degenerate scenario) or a main sequence or giant star (a single degenerate scenario). The reader is referred to Maoz, Mannucci and Nelemans \cite{Maoz2014} and Ruiz-Lapuente \cite{Ruiz-Lapuente2014} for reviews and variations upon these scenarios. In addition to the debate over the identity of the companion, another debate is over the mechanism for the explosion itself. Popular mechanisms for the single degenerate scenario at the present time are the delayed detonation transition model (DDT) model \cite{DDToriginal}, the gravitationally confined detonation (GCD) model \cite{Plewa2004}, and the pulsational reverse detonation (PRD) model \cite{PRDoriginal}, though others also exist. We refer the reader to Hillebrandt \emph{et al.} \cite{Hillebrandt2013} for a recent review. Determining the explosion mechanism from observation will be difficult. The unknown identity of the progenitor means observatories will require close to full sky coverage, while the brevity of the explosion - simulations indicate the star becomes unbound within a few seconds - means sub-second integration times. Thus, it may not be until there is a galactic SN~Ia that the explosion mechanism can be determined observationally. The rate of SNe~Ia in the Galaxy is calculated to be $1.4^{+1.4}_{-0.8}$ per century by Adams \emph{et al.} \cite{2013ApJ...778..164A} and represents ${\sim}30\%$ of the total supernova rate. Adams \emph{et al.} calculate the most probable distance to a Galactic SN Ia to be $d = 9\;{\rm kpc}$. At such close proximity, when the next SN Ia occurs within the Galaxy, not only can we exploit electromagnetic observations, it may also be possible to obtain information about the progenitor and explosion mechanism in the gravitational waves and the neutrino signal. The gravitational wave emission of the explosion (not the inspiral) has been considered by Falta \emph{et al.} \cite{2011PhRvL.106t1103F} and Seitenzahl \emph{et al.} \cite{Seitenzahl2015a}. The neutrino emission was calculated by Kunugise and Iwamoto \cite{2007PASJ...59L..57K}, Odrzywolek and Plewa \cite{Odrzywolek2011a}, and Seitenzahl \emph{et al.} \cite{Seitenzahl2015a}. The study by Odrzywolek and Plewa \cite{Odrzywolek2011a} is particularly worth emphasizing because they showed how the neutrinos are capable of distinguishing between explosion mechanisms even when the electromagnetic output is identical. But as with many things neutrino related, detecting the neutrino signal from SNe Ia is not easy. In comparison with core-collapse supernovae, the flux of neutrinos at Earth from a SN Ia is approximately four orders of magnitude smaller for a source at the same distance and also the spectrum peaks in the range of $3\;{\rm MeV}$ \cite{2007PASJ...59L..57K} rather than the $10- 20\;{\rm MeV}$ for the neutrino spectrum from a core-collapse SN. The shift to lower energies reduces the number of events one expects in a detector because of the energy dependence of neutrino cross sections and, furthermore, detecting the events is harder because they occur much closer to, or below, detector thresholds. However one advantage of the neutrinos from SNe Ia compared to core-collapse SNe is that the neutrino signal is more reliably calculated. The densities in the core of SNe~Ia are not sufficient to trap the neutrinos so there is no neutrino transport to follow and there are also no neutrino self-interaction effects \cite{Duan:2006an,Duan:2006jv} because the neutrino density is not large enough. The advent of several next-generation neutrino detectors recently prompted Wright \emph{et al.} \cite{Wright2016} (hereafter called Paper I) to recompute the expected signal from a DDT SN Ia explosion including many effects which had not been taken into account previously. For their calculation, Wright \emph{et al.} took into account the full time and energy dependence of the emission and calculated the neutrino flavor transformation as a function of time and energy through the supernova and along several rays through the simulation in order to study the line-of-sight variability. They then processed the fluxes at Earth using the SNOwGLoBES detector event rate calculation software so as to compute the total and differential event rates in the Super-K, Hyper-K, JUNO, and DUNE detectors. A separate analysis was undertaken for the signal in the IceCube detector. Their conclusion was that despite the difficulties, neutrino detectors are becoming so large and sensitive that the neutrinos from a DDT SN~Ia at the Galactic center can be detected. As the supernova is placed closer to Earth, features in the spectrum as a function of time and energy may be observed. The goal of this paper is to repeat the calculations of Paper I, but this time for the neutrino signal from the alternative scenario of a gravitationally confined detonation. The strategy is very similar to that found in Paper I and the outline of our paper is as follows. In Sec. \ref{sec:SneModel}, we provide an outline of the gravitationally confined detonation simulation we adopt. We then describe in Sec. \ref{sec:Production} the method by which the simulation is post-processed to compute the emitted neutrino spectra followed by Sec. \ref{sec:NeutrinoOscillation}, where we calculate the neutrino flavor transformation through the simulation as a function of time and energy along different rays. In Sec. \ref{sec:NeutrinoDetection}, we combine the emitted spectra and neutrino flavor transition probabilities to compute the flux at Earth and send these fluxes through the SNOwGLoBES software for representative examples of various neutrino detector technologies. The total and differential event rates we compute are investigated in order to determine how well the GCD supernova can be observed in neutrinos for a range of of distances to the event. Finally we compare the GCD and DDT neutrino signals in Sec. \ref{sec:Compare} in order to determine the distance/detector requirements necessary to distinguish the two scenarios. We conclude with Sec. \ref{sec:Conclusion}. | } In this paper, we have studied neutrinos from a gravitationally confined detonation scenario for a SN Ia in order to determine the features which might allow us to discriminate between this and other explosion mechanisms. \WW{Compared to earlier studies e.g.\~ by Odrzywolek \& Plewa \cite{Odrzywolek2011a}, we have added a number of refinements to the calculation. We adopt a 3D SNe~Ia simulation where previously 2D simulation was used, we calculate neutrino emission including more detailed spectral information, we include the effects of neutrino oscillations through the supernova by numerically solving the three-flavor evolution equations as a function of time, energy and for ten lines of sight, and we use a more complete detector modeling code \textsc{SNOwGLoBES} to compute the expected signal as a function of time and energy in a selection of detectors and for a variety of distances to the SN.} The results show that the GCD SN~Ia produces two neutrino bursts. The first is associated with deflagration burning with a peak luminosity of $2.3\times10^{48}$ erg/s at $t=0.45$ s. The second burst is associated with detonation burning with a peak luminosity of $1.2\times10^{48}$ erg/s at $t=2.82$ s. There is very little neutrino emission in between the two bursts. We also find a 10 MeV $\nu_e$ spectral feature associated with electron capture on copper appears at $\sim 1$ s during the deflagration burst which persists for the entirety of the detonation burst. Neutrino oscillations introduce significant flavor conversion and are very line-of-sight dependent due to the discontinuity-ridden density profile and the asymmetrical explosion. The oscillations also deviate from adiabatic evolution across much of the time and energy parameter space. However, even though the oscillations show large line-of-sight dependence, the effect of line-of-sight variation on the total number of detected events is only a few percent. The calculated interaction event rates show that, for a GCD SN Ia at 1 kpc, Hyper-K and IceCube would only see a few interaction events. While Super-K, JUNO, and DUNE would need a GCD SN Ia at ${\sim}0.3$ kpc to see a few interaction events. Thus the conclusion is that one would either need an improbably close SN or order-of-magnitude larger detectors than current and near future proposals to confidently observe the neutrinos from a SN Ia with a GCD explosion mechanism. A comparison of the neutrino signals from the DDT SN and the GCD SN reveals several similar features between them. The line-of-sight analysis shows that neither model would suffer from a great amount ($<10\%$) of line-of-sight variation of the total detected events even though the GCD explosion is much more aspherical than the DDT. In both cases, inverted mass ordering yields more events than normal ordering but both are less then the case of no neutrino oscillations. The 10 MeV neutrino spectral feature that is found in both DDT and GCD is difficult to detect even for next generation neutrino detectors because it appears at late times when the luminosity is low. However the comparison between the DDT and GCD also reveals two very important, explosion model distinguishing, features. Firstly, the DDT SN has an order of magnitude more overall neutrino emission than the GCD SN. If the distance to the supernova can be determined then the large difference in event rates - assuming events are detected - will allow discrimination between the two models. If the supernova is sufficiently close that a star which explodes as a GCD SN Ia produces events in a detector, the distinctive separation in time between the deflagration and detonation burst of this scenario will also allow discrimination between the two scenarios independent of the overall event rate diagnostic. And finally, SN Ia can also provide information about neutrino properties if the explosion mechanism and distance are known. We find, if the mechanism and distance are known, the overall event rate can be used to determine the neutrino mass ordering. | 16 | 9 | 1609.07403 |
1609 | 1609.06319_arXiv.txt | Here we present three-dimensional high resolution simulations of Type Ia supernova in the presence of a non-degenerate companion. We find that the presence of a nearby companion leaves a long-lived hole in the supernova ejecta. In particular, we aim to study the long term evolution of this hole as the supernova ejecta interacts with the surrounding interstellar medium. Using estimates for the x-ray emission, we find that the hole generated by the companion remains for many centuries after the interaction between the ejecta and the interstellar medium. We also show that the hole is discernible over a wide range of viewing angles and companion masses. | Type Ia supernovae (SNe Ia) are very important tools in cosmology considering their standardizable light curves \citep{Pskovskii1977,Phillips1993,Hamuy1996,Phillips1999} which make them excellent standard candles \citep{Colgate1979,Branch1992}. Type Ia supernovae are thought to be the thermonuclear explosions of a carbon-oxygen white dwarf and are characterized by the lack of hydrogen in their spectra and the formation of large amounts of radioactive {}$^{56}$Ni. However, the precise mechanism for producing the explosion remains uncertain. SNe Ia progenitor scenarios generally fall into two categories. In the double degenerate scenario, two white dwarfs combine and detonate. This can be the result of inspiral in a close binary \citep{Pakmor2010,Pakmor2011,Pakmor2012,Dan2012,Guillochon2010,Sato2016}, or the result of a direct collision \citep{Rosswog2009,Raskin2009,Raskin2010,Loren2010}. Even though this latter arrangement requires a dense stellar environments, \eg\ globular clusters or galactic nuclei, these systems can produce a range of luminosities. The second category of SNe Ia progenitors is the single degenerate model. Here, the degenerate white dwarf shares a binary system with a non-degenerate companion. The white dwarf accretes gas from the companion through Roche-lobe overflow \citep{Whelan1973,Nomoto1982,Hillebrandt2000,Hillebrandt2013}. To date, several non-degenerate candidates have been studied, from the canonical hydrogen-burning companions\citep[\eg][]{Hachisu2007} to helium-burning companions \citep{Iben1987,Ruiter2009,Wang2009,Ruiter2011}. The primary challenge for the canonical hydrogen-burning companion is to achieve an accretion rate of $\sim$10$^{-7}$$\rm M_{\odot}/yr$ which allows for a steady increase in the mass of white dwarf while avoiding mass loss from classical novae \citep{Nomoto1991}. At this rate, the white dwarf undergoes thermonuclear runaway as it grows toward the Chandrasekhar mass limit, M$_{ch}$=1.44\msol. In the single degenerate case, much work has been done trying to identify the progenitors that lead to observable SNe Ia. In this scenario, the collision of the expanding supernova ejecta with the companion star is unavoidable. This led \cite{Kasen2010} to calculate theoretical supernova light curves for a supernova interacting with a 1-2 \msol\ red giant star. They found that at early times ($t<8$ days), the luminosity is dominated by the collision. However, only viewing angles that look directly down on the companion will have prominent collision signatures; limiting detection to $\sim10\%$ of the entire population. \cite{Kutsuna2015} built upon this work and found that the expected UV signal is also dependent on the separation between the white dwarf and the companion; finding that for separations $<2.0\times$10$^{13}$cm the UV flux cannot be detected. \cite{Meng2016} found that many of the binary systems that lead to SNe Ia have separations much less than this cutoff, making the measurement of the UV emission much more difficult. Direct observation of the collision between and the supernova and a companion are notoriously difficult. \cite{Cao2015} found a strong but short lived ultraviolet emission from a young SNe Ia within four days of the explosion, which they suggest is the result of the collision between the supernova material and the companion star. However, this requires a separation distance of $\sim4\times$10$^{14}$cm. Given these constraints on the separation distance and viewing angle, the probability of observing this UV signal is low. One important effect (and possible observable signature) produced by the interaction between a companion and the supernova ejecta is the formation of a hole within the ejecta \cite[\eg][]{Fryxell1981,Marietta2000,GarciaSenz2012}. Naturally, ejecta material that interacts with the companion is slowed relative to the rest of the ejecta. This creates a ``mass shadow'' in the ejecta. The presence of this hole is in direct opposition to observations which find that supernova remnants are remarkably spherical \citep{Badenes2010}. Therefore, the ultimate evolution of the hole has important implications for the likelihood of progenitor scenarios for SNe Ia. \cite{GarciaSenz2012} simulated this interaction with a supernova arising from a 1.38\msol\ white dwarf and a 1\msol\ main sequence companion as well as its interaction with the surrounding interstellar medium using a cylindrical coordinates, axisymmetric smooth particle hydrodynamic (SPH) code. They found that the hole generated by the companion is slowly filled in by ejecta due to hydrodynamical instabilities at the edge of the hole. They also estimated the x-ray emission from the ejecta and showed that signatures of a hole in the ejecta should remain visible for an extended period. The strength of this signal depends greatly on the the viewing angle. Here we aim to build on this work by performing a suite of 3D simulations that vary the companion mass. In this paper, we carry out a comprehensive survey of hydrodynamics simulations designed to accurately measure the impact of close binary companion stars on SNe Ia remnants. Each simulation proceeds in two stages. First, we follow the interaction between the supernova ejecta and the companion, accounting for the formation of the ejecta hole and the matter stripped from the companion. Once the ejecta has reached homologous expansion and the mass stripped from the companion has plateaued, we stop that simulation and expand the ejecta to a point where the interaction with the interstellar medium (ISM) becomes important. We then follow the ejecta as it interacts with the interstellar medium. Estimates of the x-ray emission are used to probe the evolution of the hole during this interaction, finding that the hole persists for many centuries after the interaction with the ISM. The structure of the paper is as follows. In \S 2, we describe our numerical setup for each of the binary systems we model, with companion masses ranging from 2--5\msol. In \S 3, we give the results of our simulations and we carry out a number of analyses with the goal of constraining possible observables that can be replicated for real remnants, and in \S 4, we present results from our resolution study and, finally, in \S 5 we discuss our conclusions. | In the single degenerate SNe Ia scenario, the collision between the supernova ejecta and the companion is unavoidable. This has prompted theoretical studies of this interaction as a method of identifying companion properties. While \cite{Kasen2010} computed a series of light curves that suggested this interaction should be visible in $\sim$10$\%$ of all single degenerate SNe Ia, \cite{Kutsuna2015} showed that this emission was highly dependent on the initial separation between the companion and the white dwarf. They established a strict cutoff at $\sim$2$\times$10$^{13}$cm, below which such emission would be invisible. \cite{Meng2016} showed that many existing binary systems have separations below this cutoff. On the other hand, a hole in the supernova ejecta may be observable at late times depending on the evolution of the ejecta as it interacts with the interstellar medium. \cite{GarciaSenz2012} found in their axisymmetric simulations that the hole in the ejecta persisted for many centuries after the interaction with the interstellar medium. However, they also found that over such long timescales, hydrodynamic instabilities at the edges of the hole may eventually close the hole. We have presented here a set of high resolution 3D SPH simulations that aim to illustrate the effect of a companion star on the evolution of a supernova remnant and the observability of the hole produced from the interaction with the companion. Our simulations have two distinct advantages over previous studies in this area. First, all of models are fully 3D to capture the full extent of the degrees of freedom during the interaction. Second, all of the particles have identical mass, avoiding possible numerical artifacts. Each of the simulated companion stars were early stage red giant branch stars that varied in mass from 2\msol\ to 5\msol\ while the supernova was modeled as a 1.0\msol\ white dwarf with an initial explosion energy of $\approx$10$^{51}$ ergs. For each of the simulations, we find that between 0.11-0.22 \msol\ of material was stripped, depending on the companion mass, matching the estimates from \cite{Wheeler1975}. We have presented estimates of the x-ray emission profiles over a wide range of times, companion masses, and angles. The signal from the hole is found as a secondary set of peaks in the x-ray emission. In the head-on $\theta=0$ case, these peaks are found near the primary peaks from the ejecta shell. At larger angles these peaks are found separate and distinct from the ejecta shell. At $\theta=90$ the emission takes on an unmistakably asymmetric appearance. We find that the hole is present at all viewing angles both early ($t=100$ years) and late ($t=300$ years) after the interaction with the ISM. Finally, we find the hole should remain over the long term evolution of the supernova remnant. In a future work, in addition to increasing the resolution, we plan on improving the models presented here. One important feature to include is to use a more realistic ISM. The ISM was modeled here as a uniform medium with constant density. In reality, the ISM is much more complicated and composed of different phases \cite[\eg][]{Cox2005}. The inclusion of nuclear burn network in order to track the formation of important radioactive elements would also improve the estimated x-ray emission. Furthermore, as we have neglecting any cooling physics in the ISM interaction phase, our density profiles are too smooth. Allowing for radiative cooling after the Sedov phase might produce an even stronger signal of the remnant hole. Finally, the inclusion of more realistic treatment of the emission is warranted to validate the results of the estimated x-ray emission. The results presented here provide an essential tool in the study of supernovae remnants. In particular, we have shown that the interaction between the supernova ejecta and the non-degenerate companion leaves a long-lived imprint on the final supernova remnant. We would also like to thank the anonymous referee for their comments which helped to improve this paper. This work was performed under the auspices of the U.S. Department of Energy by Lawrence Livermore National Laboratory under Contract DEAC52-07NA27344. The authors also acknowledge the Livermore Computing Center at Lawrence Livermore Nation Laboratory for providing HPC resources that contributed to the results reported within this paper. LLNL-JRNL-697001. | 16 | 9 | 1609.06319 |
1609 | 1609.01013_arXiv.txt | {Most of the studies on extragalactic $\gamma$-ray propagation performed up to now only accounted for primary gamma-ray absorption and adiabatic losses (``absorption-only model''). However, there is growing evidence that this model is oversimplified and must be modified in some way. In particular, it was found that the intensity extrapolated from the optically-thin energy range of some blazar spectra is insufficient to explain the optically-thick part of these spectra. This effect was interpreted as an indication for $\gamma$-axion-like particle (ALP) oscillation. On the other hand, there are many hints that a secondary component from electromagnetic cascades initiated by primary $\gamma$-rays or nuclei may be observed in the spectra of some blazars.} {We study the impact of electromagnetic cascades from primary $\gamma$-rays or protons on the physical interpretation of blazar spectra obtained with imaging Cherenkov telescopes.} {We use the publicly-available code ELMAG to compute observable spectra of electromagnetic cascades from primary $\gamma$-rays. For the case of primary proton, we develop a simple, fast and reasonably accurate hybrid method to calculate the observable spectrum. We perform the fitting of the observed spectral energy distributions (SEDs) with various physical models: the absorption-only model, the ``electromagnetic cascade model'' (for the case of primary $\gamma$-rays), and several versions of the hadronic cascade model (for the case of primary proton). We distinguish the following species of hadronic cascade models: 1) ``basic hadronic model'', where it is assumed that the proton beam travels undisturbed by extragalactic magnetic field and that all observable $\gamma$-rays are produced by primary protons through photohadronic processes with subsequent development of electromagnetic cascades 2) ``intermediate hadronic model'', the same as the basic hadronic model, but the primary beam is terminated at some redshift $z_{c}$ 3) ``modified hadronic model'' that includes the contribution from primary (redshifted and partially absorbed) $\gamma$-rays.} {Electromagnetic cascades show at least two very distinct regimes labeled by the energy of the primary $\gamma$-ray ($E_{0}$): the one-generation regime for the case of $E_{0}$<10 $TeV$ and the universal regime for $E_{0}$>100 $TeV$ and redshift to the source $z_{s}$>0.02. Spectral signatures of the observable spectrum for the case of the basic hadronic model, $z_{s}$= 0.186 and low energy ($E$<200 $GeV$) are nearly the same as for purely electromagnetic cascade, but for $E$>200 $GeV$ the spectrum is much harder for the case of the basic hadronic model. In the framework of the intermediate hadronic model, the observable spectrum depends only slightly on the primary proton energy, but it strongly depends on $z_{c}$ at $E$>500 $GeV$. As a rule, both electromagnetic and hadronic cascade models provide acceptable fits to the observed SEDs. We show that the best-fit model intensity in the multi-$TeV$ region of the spectrum in the framework of the electromagnetic cascade model is typically greater than the one for the case of the absorption-only model. Finally, for the case of blazar 1ES 0229+200 we provide strong constraints on the intermediate hadronic model assuming the blazar emission model of Tavecchio (2013) and the model of magnetic field around the source according to Meyer et al. (2013).} {} | The development of ground-based $\gamma$-ray astronomy with imaging Cherenkov telescopes has been very fast during the last two decades (for a review see, e.g., Hillas \cite{hillas}). Indeed, only 7 years after the detection of $TeV$ photons from the active galactic nucleus (AGN) Mkn 421 in Punch et al. \cite{punch}, the first detailed study of a blazar (AGN with the jet presumably pointed towards the observer) spectrum was made by Aharonian et al. \cite{aharonian99}. Almost immediately, these observations were utilized to put some constraints on the intensity of extragalactic background light (EBL) (in Stecker \& de Jager \cite{stecker}, de Jager et al. \cite{dejager} for the former observations, and in Aharonian et al. \cite{aharonian99} itself for the latter). Indeed, primary very high energy (VHE, $E$>100 $GeV$) $\gamma$-rays are absorbed on EBL photons by means of the $\gamma \gamma \rightarrow e^{+}e^{-}$ process (Nikishov \cite{nikishov}, Gould \& Shroeder \cite{gould}), and for the case of primary energy $E\sim100$ $TeV$ and higher --- on the cosmic microwave background (CMB) as well (Jelley \cite{jelley}). More recently, the signatures of this absorption process were observed with the Fermi LAT instrument (Ackermann et al. \cite{ackermann}) and the H.E.S.S. Cherenkov telescope (Abramowski et al. (The H.E.S.S. Collaboration) \cite{abramowski}) with high statistical significance ($\sim6\sigma$ and 8.8 $\sigma$, respectively). However, Horns \& Meyer \cite{horns12} found that the strength of absorption at high optical depth ($\tau_{\gamma\gamma}$>2, hereafter simply $\tau$) appears to be lower than expected. This result was obtained on a sample of blazar spectra measured with imaging Cherenkov telescopes by comparing the distributions of the flux points for the $\tau$ regions 1<$\tau$<2 and $\tau$>2 around the intensity extrapolated from the optically-thin regime $\tau$<1. The statistical significance of this effect, according to Horns \& Meyer \cite{horns12}, is 4.2 $\sigma$. While this result was not confirmed by Biteau \& Williams \cite{biteau}, very recently Horns \cite{horns16} again found an indication for this anomaly with another analysis method. Such an anomaly, in fact, closely resembles the so-called ``$TeV$-IR crisis'' (Protheroe \& Meyer \cite{protheroe}) that was derived from the already mentioned observations of Mkn 501 (Aharonian et al. \cite{aharonian99}). The ``$TeV$-IR crisis'', however, was later found to be less severe after the development of more advanced EBL models. On the contrary, the ``new'' anomaly of Horns \& Meyer \cite{horns12}, Horns \cite{horns16} persists for most of these contemporary models of EBL intensity. The authors of Horns \& Meyer \cite{horns12} interpreted their result as an indication for the existence of some non-conventional physical effect, for instance, the process of oscillations of $\gamma$-rays into axion-like particles (ALPs) and back into photons in magnetic field on the way from the source to the observer ($\gamma \rightarrow ALP$). Indeed, a part of photons reconverted from ALPs near the observer can significantly enhance the observed intensity in the $\tau$>2 region. Moreover, once the anomaly is well established, it is possible to put constraints on the gamma-ALP mixing parameters $m_{a}$ (the mass of ALP) and $g_{a\gamma}$ (the photon-ALP coupling constant). In Meyer et al. \cite{meyer} a lower limit on the $g_{a\gamma}$ was found, depending on the $m_{a}$ value. For any fixed $m_{a}$ in the range considered in Meyer et al. \cite{meyer}, some values of $g_{a\gamma}$ greater than the lower limit $g_{a\gamma-min}(m_{a})$ could explain the observed anomaly. Together with the upper limits from the CAST experiment (Andriamonje et al. (CAST Collaboration) \cite{andriamonje}), this result allowed to construct a confidence interval for $g_{a\gamma}$. Besides the anomaly at $\tau$>2, there exists another signature of $\gamma \rightarrow ALP$ oscillation, namely, a step-like irregularity that is situated at the energy lower than the starting energy of the VHE anomaly (e.g, Sanchez-Conde et al. \cite{sanchez-conde}). The drop of intensity associated with this spectral feature is usually about 1/3 as photons (two polarization states) attain equipartition with ALPs (one polarization state). A very recent analysis of Fermi LAT data (Ajello et al. \cite{ajello}) (observations of the NGC 1275 Seyfert galaxy were used), however, did not find this signature. Moreover, other bounds (Ayala et al. \cite{ayala}, Abramowski et al. (H.E.S.S. Collaboration) \cite{abramowski13}, Payez et al. \cite{payez}, Wouters \& Brun \cite{wouters}) on the parameters of $\gamma$-ALP mixing allowed to strongly constrain the scenario considered in Meyer et al. \cite{meyer}. Therefore, the hypothesis that the VHE anomaly is explained by $\gamma \rightarrow ALP$ oscillation appears to be less attractive than before; one needs to search for some other physical mechanism of the anomaly. In this respect, we note that most of extragalactic $\gamma$-ray propagation studies were performed with account of only two elementary processes: the absorption of primary photons (by means of pair-production) and their adiabatic losses. This model (hereafter ``the absorption-only model'') rests on the assumption that the secondary electrons and positrons (hereafter simply ``electrons'' unless otherwise stated) are deflected and delayed by the extragalactic magnetic field (EGMF), thus the cascade photons produced by these electrons by means of the inverse Compton (IC) scattering do not contribute to the observed spectrum. However, the EGMF strength $B$ in voids of the Large Scale Structure (LSS) may be small enough to violate this assumption. The existing constraints (Blasi et al. \cite{blasi}, Pshirkov et al. \cite{pshirkov}, Dolag et al. \cite{dolag}, Neronov \& Vovk \cite{neronov}, Dermer et al. \cite{dermer}, Taylor et al. \cite{taylor}, Vovk et al. \cite{vovk}, Abramowski et al. (H.E.S.S. Collaboration) \cite{abramowski14}, Takahashi et al. \cite{takahashi12}, Takahashi et al. \cite{takahashi13}) on the EGMF strength on characteristic coherence scale 1 $Mpc$ (Akahori \& Ryu \cite{akahori}) (some of them were summarized, e.g., in Fig. 4 of Dzhatdoev \cite{dzhatdoev15a}) do not exclude the probability that cascade emission may contribute to the observed spectrum. The recent work Finke et al. \cite{finke}, using a more conservative method than the above-mentioned references, excluded $B<10^{-19}$ $G$, again on characteristic coherence scale 1 $Mpc$, with statistical significance $5\sigma$. Arlen et al. \cite{arlen} did not find any reason to reject the null hypothesis of $B=0$ at all. Tashiro \& Vachaspati \cite{tashiro} found $B\sim10^{-14}$ $G$. This study was made using the angular correlation pattern of Fermi LAT (Atwood et al. \cite{atwood}) diffuse $\gamma$-rays; however, it is not clear how much this result may be affected by the existence of comparatively strong magnetic field in galaxy clusters. Moreover, there are some hints that the cascade component does contribute to the observed spectrum of blazars at energies $E<300$ $GeV$. Neronov et al. \cite{neronov12}, using observations of Mkn 501 in a flaring state by the Fermi LAT instrument and the VERITAS Cherenkov telescope (Abdo et al. \cite{abdo}) during the 2009 multiwavelength campaign, obtained the energy spectrum $dN/dE$ ($dN$ is the number of photons per energy bin $dE$) from 300 $MeV$ up to 5 $TeV$. It was found that the Fermi LAT spectrum in the 10--200 $GeV$ energy range had a power-law ($dN/dE=C\cdot E^{-\gamma}$) index $\gamma= 1.1\pm0.2$, while in the 300 $GeV$ -- 5 $TeV$ $\gamma \approx 2$. It is interesting that the Fermi LAT lightcurves in the 0.3-3 $GeV$ and 3-30 $GeV$ energy bins do not show any evidence for strong, fast variability, while in the 30-300 $GeV$ energy bin the flare is readily identified on the lightcurve. Neronov et al. \cite{neronov12} found that such a behaviour of the spectral and timing characteristics is typical for intergalactic cascade, assuming $B\sim10^{-16}-10^{-17}$ $G$ on the maximum spatial scale 1 $Mpc$. Another result supporting the incompleteness of the absorption-only model was also obtained with the Fermi LAT telescope. Namely, Furniss et al. \cite{furniss} found a correlation between 10-500 $GeV$ energy flux of blazars with relatively hard ($\gamma<3$) observed spectra above 10 $GeV$ and the fraction along the line of sight occupied by voids in the LSS. This effect may be explained by the same physical mechanism: the cascade emission that is likely to be angle-broadened by the EGMF at some energy below 200-300 $GeV$. Finally, Chen et al. \cite{chen}, again with Fermi LAT data, found the evidence for the existence of ``pair halos'' (extended emission) around the positions of various blazars, thus giving additional support to the hypothesis that the cascade component may contribute to the observed spectra. These works motivated us to study how the inclusion of the cascade component into the fitting of the VHE spectra would influence the data interpretation, especially in the optically thick ($\tau$>1) region. Our paper is by far not the first study of intergalactic cascade spectra; besides the already mentioned work (Neronov et al. \cite{neronov12}), there are many others that included more or less detailed duscussions of these, namely: Aharonian et al. \cite{aharonian99}, Aharonian et al. \cite{aharonian02}, d'Avesac et al. \cite{davezac}, Murase et al. \cite{murase}, Takami et al. \cite{takami}. However, very recently in a conference paper Dzhatdoev \cite{dzhatdoev15b} it was shown that if the primary spectrum is hard enough ($\gamma \approx 1$) and the cutoff in the spectral energy distribution (SED) of the source is situated at a sufficiently high energy ($E_{c}>>1$ $TeV$), then the intersection of a low-energy cascade component and a high-energy primary (absorbed) component forms a kind of an ``ankle'' that to some extent may mimic the $\gamma-ALP$ mixing effect mentioned above. One of the main aims of the present paper is to discuss this effect in more depth. We will see that the last model is qualitatively different from other ``electromagnetic cascade models'', i.e. the models that involve intergalactic cascades initiated by primary gamma-rays. There exists another class of intergalactic cascade models of blazar spectra dealing with primary protons or nuclei that initiate secondary photons by means of photopion losses and pair production with subsequent development of electromagnetic cascades (``hadronic cascade models'') (e.g. Uryson \cite{uryson}, Essey \& Kusenko \cite{essey10a}, Essey et al. \cite{essey10b}, Essey et al. \cite{essey11}, Murase et al. \cite{murase}, Takami et al. \cite{takami}, Essey \& Kusenko \cite{essey14}, Zheng et al. \cite{zheng}). We will compare hadronic models with the electromagnetic cascade model of Dzhatdoev \cite{dzhatdoev15b} to reveal their advantages and diffuculties. The paper is organized as follows. In Sec.~\ref{sec:casc} we discuss some general properties of cascades from primary photon or proton developing on EBL/CMB and give a detailed description of our calculation method for the case of primary proton. In Section 3 we define the sample of blazar spectra to be used in the analysis (in this paper we focus on the observations made with Cherenkov telescopes). In Section 4 we present the fits to observed SEDs. Section 5 contains discussion, where our main results are recalled and discussed in the broader context; and, finally, in Section 6 the conclusions to our work are drawn. | In this work we reviewed the main possibilities of how cascades from primary $\gamma$-rays or protons may influence the data interpretation when testing extragalactic $\gamma$-ray propagation models, mainly dealing with the source redshift range 0.1--0.3. We reviewed the main regimes of electromagnetic cascade development: 1) one-generation regime, which holds when $E_{\gamma0}<$10 $TeV$, 2) the universal regime, which is satisfied when $E_{\gamma0}>$100 $TeV$ 3) a possibility of the extreme energy cascade regime at $E_{\gamma0}>$1 $EeV$. We developed a fast, simple and suffuciently precise hybrid code that allows to calculate the observable spectrum of $\gamma$-rays for the case of primary proton. We performed calculations of observable spectra for this case and investigated several versions of the hadronic cascade model --- the basic model (all gamma-rays are generated by primary protons), the intermediate model (the same, but the proton beam is dissolved at some $z_c$, the basic model is the sub-set of intermediate models with $z_{c}$= 0), the modified model that includes the primary component. We discussed the signatures of the spectrum for the basic hadronic model for $z_{s}$= 0.186. The signatures at a comparatively low observable energy are almost the same as the ones for the case of a purely electromagnetic cascade --- that is, the $E^{-1.60}$ $dN/dE$ spectrum below $\approx 200$ $MeV$ and $E^{-1.85}$ spectrum from $\approx200$ $MeV$ to $\approx200$ $GeV$. However, at higher energy ($E>$200 $GeV$) the spectrum in the hadronic model is much harder than the universal spectrum for the case of a pure electromagnetic cascade. We also performed the comparison of two hybrid calculations of observable spectra for different values of $z_{c}$ from 0 up to nearly $z_{s}$. The results of these calculations with two different methods are in very good agreement for the case of the proton primary energy $E_{p0}\le$ 30 $EeV$, i.e. when the pair-production losses are dominant. For the case of $E_{p0}$= 100 $EeV$ the agreement between the two methods is also good for the basic hadronic model and the intermediate hadronic model with comparatively low $z_{c}<z_{s}/2$, but not so good if $z_{c}\ge z_{s}/2$. The reason of this disagreement may be connected to the possibility of the extreme energy cascade regime. New calculations and measurements of the universal radio background are necessary in order to decide which spectrum is more realistic. While we confirmed the claim that the shape of the observable spectrum in the basic hadronic model depends only slightly on the primary proton energy, it was found that in context of the intermediate hadronic model the shape of the spectrum strongly depends on the $z_{c}$ parameter in the high energy region. We performed a series of fits to the spectra of extreme $TeV$ blazars with various extragalactic $\gamma$-ray propagation models, namely: the absorption-only model, the electromagnetic cascade model, the basic hadronic model, the intermediate hadronic model, the modified hadronic model. Most of the fits presented in this paper are formal best fits. It was found that, as a rule, both the electromagnetic cascade model and the hadronic cascade model are able to provide a reasonable fit to the observed spectral shape. For the case of blazar 1ES 0347-121 and the electromagnetic cascade model we performed the calculation of modification factor $K_{B}$ with respect to the absorption-only model. We found that $K_{B}$>1 in the energy region of 1-10 $TeV$. Therefore, the version of the electromagnetic cascade model that was considered in this work for this source predicts that it is possible that the observable intensity is higher than in the absorption-only model for the same level of EBL without any new physics. For the case of the source 1ES 0229+200 we also performed a fit in the framework of an exotic model with $\gamma\rightarrow ALP$ oscillations and a calculation of the modification factor $K_{B}$ vs. voidiness $K_{V}$ dependence in the framework of the electromagnetic cascade model. We found that when $K_{V}$ is lower than 1, the value of $K_{B}$ does not fall at high energies, but, on the contrary, even grows in the energy range of 1-15 $TeV$, thus making the effects induced by electromagnetic cascades even more pronounced. Finally, using the fact that many blazars are situated in galaxy groups or clusters with comparatively strong magnetic fields (Muriel \cite{muriel}, Oikonomou et al. \cite{oikonomou}), we performed the testing of the emission model of Tavecchio (2014). We found that for $B_{0}$>100 $nG$ this emission model is excluded with significance $Z>$7 $\sigma$. Our results show that the hadronic cascade model experiences significant difficulties. This conclusion is in line with results obtained by Razzaque et al. \cite{razzaque}. Further testing of the hadronic and electromagnetic cascade models may be performed on an event-by-event level with advanced analysis tools such as GammaLib and ctools package (Knoedlseder et al. \cite{knoedlseder}). Such a work is underway in our group. We conclude that cascades from primary $\gamma$-rays or nuclei may induce an appreciable background for astrophysical searches for $\gamma-ALP$ oscillation that needs to be taken into account. The electromagnetic and hadronic cascade models deserve further study, also in context of searches for other exotic phenomena, such as Lorentz invariance violation (Tavecchio \& Bonnoli \cite{tavecchio16}). The knowledge of EGMF would provide a significant insight into the intergalactic cascade phenomena in the Universe. Further measurements with the CTA array of Cherenkov telescopes (Acharya et al. \cite{acharya}) will likely clarify these issues. Together with existing instruments Fermi LAT (Atwood et al. \cite{atwood}) and AGILE (Rappoldi et al. \cite{rappoldi}), future experiments that have either high sensitivity or good angular resolution, such as GAMMA-400 (Galper et al. \cite{galper}) or the novel balloon-borne emulsion $\gamma$-ray telescope GRAINE (Takahashi et al. \cite{takahashi}), may also prove to be helpful in this task. | 16 | 9 | 1609.01013 |
1609 | 1609.04770_arXiv.txt | We study the structure of two-point correlators of the inflationary field fluctuations in order to improve the accuracy and efficiency of the existing methods to calculate primordial spectra. We present a description motivated by the separation of the fast and slow evolving components of the spectrum which is based on Cholesky decomposition of the field correlator matrix. Our purpose is to rewrite all the relevant equations of motion in terms of slowly varying quantities. This is important in order to consider the contribution from high-frequency modes to the spectrum without affecting computational performance. The slow-roll approximation is not required to reproduce the main distinctive features in the power spectrum for each specific model of inflation. | \begingroup The Inflationary paradigm has become an important piece in our understanding of the Early Universe. It has been designed and constantly improved since its very first appearance in Ref.~\cite{Guth:1980zm} wherein it solves the main issues of standard Big Bang Cosmology. In the last three decades, we witnessed the emergence of a plethora of models, which not only aim to provide expansion for a sufficient number of e-folds, but also expect to produce a power spectrum of fluctuations consistent with the actual structure of the universe. Direct computation of the power spectrum from equations of motion can be time consuming, especially when resolving specific features of each model in a wide range of energy scales. These problems will be aggravated for multifield models of inflation such as the proposals in Refs.~\cite{Greene:1997fu, Tsujikawa:2002qx}. The primordial power spectrum is valuable in finding sensible ranges of validity for the parameters of any proposed Inflationary model. In the case of models with multiple fields, it encodes vital information about the power transfer between different components. The viability of a model is usually tested when the spectrum is loaded as an input in any of the existing schemes based in Boltzmann transport equations (such as CAMB in Ref.~\cite{Challinor:2005jy} or CLASS in Ref.~\cite{Tram:2013ima}) and compared with data. Being aware of the existing difficulties to design a spectral code, we suggest an approach intended to isolate all the high frequency terms and only use the slowly evolving quantities relevant for calculating the spectrum. These degrees of freedom remain ``frozen'' outside the horizon, which implies the use of large time steps in any numerical evolution scheme. A crucial part of separating fast and slow degrees of freedom relies on focusing on the spectrum, which contains information about all the field correlations. This spectrum is generated by the symmetric product between two field multiplets, forming a correlation matrix. Field dynamics reveals the approximate time-translational invariance of each multiplet component on subhorizon scales. This symmetry must be exploited to define a well-posed Cauchy problem each time vacuum correlations are defined as initial conditions. When the cases of anticorrelation are excluded, positive definite correlation matrices are suitable for Cholesky decomposition into two unique triangular factors. Anticorrelations break the uniqueness of this factorization. This decomposition has been the preferred tool of statisticians to generate correlated samples from any set of unit variance random vectors. Amplitudes of the modes and power of the cross correlations are carried in these Cholesky factors, which act on a rotating basis of solutions, just as in the Schrodinger picture in Quantum Mechanics. The fast rotation of this basis is separated from the slow evolution of amplitudes. We implement a dynamical Cholesky decomposition motivated by the separation of wave solutions into phases and amplitudes. This separation has been explored in single field solutions, (see Refs.~\cite{Brooker:2015iya, Bardeen:1985tr} for more details) with exact results in the case of massless perturbations. The plan for this paper is as follows: in section \ref{sec:review}, we will review the notions of perturbation theory for inflationary models, along with the equations of motion for the background fields. In section \ref{sec:correlators}, we describe the field decomposition technique used in order to separate the fast oscillating phases from the amplitudes. The latter are required to calculate the spectrum. In section \ref{sec:initial}, we discuss the background dynamics and the use of initial conditions based on asymptotic vacuum solutions. Additionally, we introduce a scheme for injecting modes in the system. To conclude, we present our results and discussions. \endgroup \vfill | In this paper, we present a new method to separate fast and slow scales in the context of multifield models of inflation. We describe a scheme based on the Cholesky factorization of any positive definite correlation matrix. As an outcome, we manage to find equations of motion for the ``square root'' of the correlation matrix in terms of slowly varying quantities. After reviewing the dynamical properties and contributions of the background fields, we specify a set of approximate initial conditions for the evolving system, and check that our results are consistent with straightforward averaging over all the realizations of the evolved random fields. We use the new code to calculate spectra of a few well-known models, and check that the shape of the spectrum depends on the choice of background trajectories. This perspective based on dynamical Cholesky decomposition is significantly different from previous efforts in Refs.~\cite{Seery:2012vj, Price:2014xpa}. More recently, the Hamiltonian evolution of the field correlators $\langle\Phi_A,\Phi_B\rangle$ was considered in Ref.~\cite{Dias:2016rjq}, which is a transport scheme quite similar to ours in spirit. However, our method achieves significant computational gains by separating fast time scale of the sub-horizon mode evolution. Additionally, in our evolution scheme $\langle\hat{\chi}_A,\hat{\chi}_B\rangle=\delta_{AB}$ at every instant in time, which makes it convenient to generate properly-correlated random realizations of the fields for Monte-Carlo simulations. This approach can be easily extendable to a diverse number of systems ruled by (almost) any perturbative manifestation arising from hyperbolic differential equations, which present fairly similar structures. We will leave the non-Gaussian extension of this method in application to bispectrum computation for a future project. \appendix | 16 | 9 | 1609.04770 |
1609 | 1609.01694_arXiv.txt | \normalsize\parindent 0pt\parskip 5pt Axions may make a significant contribution to the dark matter of the universe. It has been suggested that these dark matter axions may condense into localized clumps, called ``axion stars." In this paper we argue that collisions of dilute axion stars with neutron stars, of the type known as ``magnetars," may be the origin of most of the observed fast radio bursts. This idea is a variation of an idea originally proposed by Iwazaki. However, instead of the surface effect of Iwazaki, we propose a perhaps stronger volume effect caused by the induced time dependent electric dipole moment of neutrons. | 16 | 9 | 1609.01694 |
||
1609 | 1609.08256_arXiv.txt | We found a molecular cloud connecting from the outer region to the ``Galactic Center Mini-spiral (GCMS)" which is a bundle of the ionized gas streams adjacent to Sgr A$^\ast$. The molecular cloud has a filamentary appearance which is prominent in the CS $J=2-1$ emission line and is continuously connected with the GCMS. The velocity of the molecular cloud is also continuously connected with that of the ionized gas in the GCMS observed in the H42$\alpha$ recombination line. The morphological and kinematic relations suggest that the molecular cloud is falling from the outer region to the vicinity of Sgr A$^\ast$, being disrupted by the tidal shear of Sgr A$^\ast$ and ionized by UV emission from the Central Cluster. We also found the SiO $J=2-1$ emission in the boundary area between the filamentary molecular cloud and the GCMS. There seems to exist shocked gas in the boundary area. | The Galactic Center is the nuclear region of the nearest spiral galaxy, Milky Way. The environment is unique in the galaxy because the region contains several peculiar objects. First, Sagittarius A$^\ast$ (Sgr A$^\ast$) is a counter part of the Galactic Center Black Hole (GCBH) in the regime from radio to X-ray, which is located very near the dynamical center of the galaxy (e.g. \cite[Reid \etal\ 2003]{Reid}) and has a mass of $\sim4\times10^6 $M$_\odot$ (e.g. \cite[Ghez \etal\ 2008]{Ghez}; \cite[Gillessen \etal\ 2009]{Gillessen}). Second, the ``Circum-Nuclear Disk (CND)" is a torus-like molecular gas rotating around Sgr A$^\ast$, which is identified at the distance to a few pc (e.g. \cite[G\"usten \etal\ 1987]{Guesten}). A bundle of the ionized gas streams is located in the inner cavity of the CND. This is called ``Galactic Center Mini-spiral (GCMS)" (e.g. \cite[Ekers \etal\ 1983]{Ekers1983}; \cite[Lo\&Claussen 1983]{LO1983}). The stretched appearance and kinematics with large velocity gradient suggest that it is a tentative structure surrounding Sgr A$^\ast$. Finally, the Central cluster is a star cluster centered at Sgr A$^\ast$. Though the cluster concentrates within $r < 0.5$ pc, it contains $\sim$100 OB and WR stars (e.g. \cite[Genzel \etal\ 1996]{Genzel1996}; \cite[Paumard \etal\ 2006]{Paumard}). There remains considerable controversy about the cluster origin. It would be difficult to form the cluster in a way by which stars are usually formed in the galactic disk because of the following reasons. The tidal force of Sgr A$^\ast$ must have a serious effect on the star formation because the minimum number H$_2$ density for stabilization toward the tidally shearing is $n($H$_2)>3\times10^8$ cm$^{-3}$ at $r < 0.5$ pc (\cite[Christopher \etal\ 2005]{Christopher}, \cite[Montero-Casta\~{n}o \etal\ 2009]{Montero}, \cite[Tsuboi \etal\ 2011]{Tsuboi2011}). In addition, the strong Lyman continuum radiation from the early type stars in the cluster ionizes rapidly the ISM in the region. Because the star formation in the vicinity of the Sgr A$^\ast$ requires an additional mechanism to overcome the difficulties, it is still an open question how the Central cluster is formed. Two distinct but compatible scenarios for the formation of the Central cluster have been proposed. One is current in-situ star formation in such extreme environment of the vicinity of Sgr A$^\ast$. Previous observations have revealed that there are many molecular clumps in the CND(e.g. \cite[Montero-Casta\~{n}o \etal\ 2009]{Montero}, \cite[Mart\'{\i}n \etal\ 2012]{Martin}). Cloud-cloud collision in the CND is proposed as the star formation mechanism in the region (e.g. \cite[Jalali \etal\ 2014]{Jalali}). Unfortunately, we find no cradle dense molecular clouds around massive stars in the inner cavity of the CND at present. The other is that the molecular cloud is falling from the region considerably far from Sgr A$^\ast$ to the vicinity after star formation have already started in the cloud. Thus we searched the falling molecular cloud with star formation in the Sgr A region. \begin{figure}[b] \begin{center} \includegraphics[width=5.5in]{fig1.eps} \caption{Integrated intensity maps of the Sgr A region: {\bf a}: in the CS $J=2-1$ emission line, {\bf b}: in the SiO $v=0, J=2-1$ emission line, {\bf c}: in the H42$\alpha$ recombination line. The central velocity is $V_{\mathrm{LSR}}=100$ km s$^{-1}$. The angular resolution is shown on the lower-left corner of the each panel as a filed oval. The integration velocity widths are 10 km s$^{-1}$ in {\bf a} and {\bf b} and 20 km s$^{-1}$ in {\bf c}. The contours show the continuum emission of the ``Galactic Center Mini-spiral" at 100 GHz for comparison (\cite[Tsuboi \etal\ 2016a]{Tsuboi2016a}). } \label{fig1} \end{center} \end{figure} | Figures 1a and 1b show the integrated intensity maps (pseudo color) of the CND in the CS $J=2-1$ and SiO $v=0, J=2-1$ emission lines, respectively. They are cut from the large area mosaic data mentioned in the previous section. The central velocity is $V_{\mathrm{LSR}}=100$ km s$^{-1}$ both in the maps. The velocity width of each panel is 10 kms$^{-1}$. The CS emission line is a dense molecular gas tracer with $n$(H$_2$)$\gtrsim 10^4$ cm$^{-1}$, while the SiO emission line is a tracer of shocked molecular gas with shock velocity of $V_{\mathrm{s}}\gtrsim 50$km s$^{-1}$. The angular resolutions are $2.3" \times 1.6"$ for the CS $J=2-1$ emission line and $2.5" \times 1.8"$ for the SiO $v=0, J=2-1$ emission line, respectively. They are shown on the lower-left corners of the two panels as filed ovals. The contours in the figures show the continuum emission of the GCMS at 100 GHz for comparison (\cite[Tsuboi \etal\ 2016a]{Tsuboi2016a}). We found a filamentary molecular cloud extending from $\alpha\sim17^h45^m43^s$, $\delta\sim-29^\circ00'00"$ to $\alpha\sim17^h45^m37^s$, $\delta\sim-28^\circ58'40"$ in figure 1a. The north end of the molecular cloud reaches beyond the CND; $r\gtrsim 6$ pc. This molecular cloud is prominent in the velocity ranges from $V_{\mathrm{LSR}}=+70$ km s$^{-1}$ to $V_{\mathrm{LSR}}=+120$ km s$^{-1}$. Along the velocity increases as approaching the south end. Although the filamentary molecular cloud is prominent in the CS 2-1 emission line, this had never been identified by existing telescopes because it is deeply embedded in the CND. However, this is distinguishable clearly from other clouds belong to the CND with the curved filamentary appearance revealed by the high angular resolution of ALMA. On the other hand, the filamentary molecular cloud is not prominent in the SiO emission line except around the south end, $\alpha\sim17^h45^m42^s$, $\delta\sim-28^\circ59'50"$. This indicates that the cloud suffers from strong shock around the south end although there is no evidence of strong shock in the other parts of the cloud. The south end of the molecular cloud has the very wide velocity width of $\Delta V\gtrsim70$ km s$^{-1}$. While there is another component along the ``Western arc" of the GCMS in figures 1a and 1b. This component probably belongs to the CND. Meanwhile, figure 1c shows the integrated intensity map (pseudo color) of the GCMS in the H42$\alpha$ recombination line (pseudo color). The H42$\alpha$ recombination line is an ionized gas tracer. The south end seems to be connected continuously to the north extension of the ``Eastern arm" and ``Northeastern arm" of the GCMS (\cite[Tsuboi \etal\ 2016b]{Tsuboi2016b}). Although the velocity width of the north extension is as large as $\Delta V\sim100$ km s$^{-1}$, the central velocity is $V_{\mathrm{LSR}}\sim100$ km s$^{-1}$ (\cite[Tsuboi \etal\ 2016b]{Tsuboi2016b}) and overlaps to that of the south end of the molecular cloud. The central velocity of the molecular cloud is also continuous with that of the ionized gas in the arms. The morphological and kinematic properties suggest that the molecular cloud continues to the ionized gas physically. As mentioned above, the shocked gas traced by the SiO emission is located around the boundary of these. The molecular cloud is ionized by the Lyman continuum from the Central cluster as approaching to Sgr A$^\ast$ and expands at several hundreds times because the gas temperature increases from several $10$ K to $1\times10^4$ K. Therefore the interaction between the molecular gas and ionized gas probably produces a shock wave around the interface. The filamentary molecular cloud is probably falling from the outer region to the Sgr A region and being disrupted by the tidal shear of Sgr A$^\ast$. The star formation in the cloud may have already started by an external trigger like cloud-cloud collision when the cloud was in the region far from Sgr A$^\ast$. Shock waves induced by cloud-cloud collision may form massive molecular cloud cores (e.g. \cite[Inoue \& Fukui 2013] {Inoue}). A possible remnant of such cloud-cloud collision is observed in the 50MC (\cite[Tsuboi \etal\ 2015]{Tsuboi2015}). In the case of nearly head-on collision, the molecular clouds lose their angular momentums and begin to fall toward Sgr A$^\ast$. Moreover, the formed cores in the clouds can not be destroyed by the tidal force because of their high densities. The massive cores quickly grow to OB stars. The speculation is consistent with the fact that the OB stars are observed on the dust peaks in the GCMS (\cite[Tsuboi \etal\ 2016a]{Tsuboi2016a}). While less massive cores gradually grows to low mass stars. These may be still on the stage of protostars in the GCMS. Recently, several half-shell-like ionized components are found near Sgr A$^\ast$ by JVLA, which would be low mass protostars (\cite[Yusef-Zadeh \etal\ 2015]{Yusef-Zadeh2015}). They are located exclusively on the approaching tip of the ``Northeastern arm". These suggest that the filamentary cloud and GCMS play a role in transferring material including protostars from the outer region of the Sgr A complex to the vicinity of Sgr A$^\ast$. Some of the falling stars may be captured around Sgr A$^\ast$ and others may be scattered far from Sgr A$^\ast$. | 16 | 9 | 1609.08256 |
1609 | 1609.08917_arXiv.txt | {We present a non-parametric model for inferring the three-dimensional (3D) distribution of dust density in the Milky Way. Our approach uses the extinction measured towards stars at different locations in the Galaxy at approximately known distances. Each extinction measurement is proportional to the integrated dust density along its line of sight (l.o.s). Making simple assumptions about the spatial correlation of the dust density, we can infer the most probable 3D distribution of dust across the entire observed region, including along sight lines which were not observed. This is possible because our model employs a Gaussian process to connect all l.o.s. We demonstrate the capability of our model to capture detailed dust density variations using mock data and simulated data from the Gaia Universe Model Snapshot. We then apply our method to a sample of giant stars observed by APOGEE and Kepler to construct a 3D dust map over a small region of the Galaxy. Owing to our smoothness constraint and its isotropy, we provide one of the first maps which does not show the ``fingers of God'' effect. } | Interstellar dust is an integral part of a galaxy. In the cycle of matter, gas and dust are ejected from evolving stars, and eventually this material will form new stars. Dust also attenuates light through reddening and extinction, thereby complicating our interpretation of stellar photometry and spectroscopy, and it imposes complex selection functions on surveys. Any survey with or of stars must account for the effects of interstellar dust. Numerous studies have been undertaken over the years to improve our knowledge of this component of the interstellar medium (ISM). Early extinction maps were emission-based and in two dimensions (2D). A prominent piece of work was that of \citet{1998ApJ...500..525S} who used far-infrared dust emission measured by the IRAS and COBE satellites to build an all-sky map of the dust column density, assuming a standard reddening law. \citet{2006A&A...454..781L} used 2MASS photometric data and the colour excess technique to study molecular clouds and to map the dust column density and extinction in the Pipe nebula. \citet{2006A&A...453..635M} used a Galaxy model to find the intrinsic colours of stars, then, using the measured near infrared colour excess, estimated the extinction and distances towards stars. They applied their model to more than 64\,000 l.o.s to find the 3D distribution of extinction in the inner Galaxy. \citet{2010ApJ...725.1175S} used the Sloan Digital Sky Survey (SDSS) and the blue tip of the stellar locus to measure the colour shift of the main sequence turnoff and thereby calculate the reddening of stars. \citet{2012MNRAS.427.2119S} introduced a hierarchical Bayesian model to simultaneously estimate a distance--extinction relation and the properties of individual stars from multi-band photometry. This improved the precision and accuracy of the maps compared to previous ones. This method was used by \citet{2014MNRAS.443.2907S} to build a 3D extinction map of the northern Galactic plane from IPHAS photometry. A 3D map of interstellar dust reddening for three-quarters of the sky was presented by \citet{2015ApJ...810...25G} using Pan-STARRS1 and 2MASS photometry, based on the method of \citet{2014ApJ...783..114G} (which uses a Bayesian approach similar to that of \citealt{2012MNRAS.427.2119S}). This method was also used by \citet{2014ApJ...789...15S} to make a map of dust reddening out to 4.5 kpc from Pan-STARRS1 stellar photometry which covers the entire sky north of declination -30$^{\circ}$. Their method was especially designed for modelling extinction in the Galactic plane. A similar approach was taken by \citet{2014MNRAS.438.2938H}, who used SDSS and UKIDSS multi-band photometry to map extinction in 3D. They used a Bayesian model to take into account the degeneracy between extinction and stellar effective temperature. The method was previously introduced by \citet{2011MNRAS.411..435B} to estimate effective temperatures and reddenings towards 40\,000 FGK stars from 2MASS and Hipparcos photometry and Hipparcos parallaxes. \citet{2014A&A...561A..91L} applied an inversion method to measurements of stellar colour excess made at optical wavelengths. Together with parallaxes or photometric distances they constructed a map of the ISM within 2.5 kpc of the Sun. \citet{2014MNRAS.445..256S} presented a mapping method in which extinction is modelled as a Gaussian random field with a covariance function which has a Kolmogorov-like power spectrum. This is motivated by the idea that turbulence is responsible for the spatial structure of the ISM. Many of the aforementioned methods consider each l.o.s independently: they do not propagate information between neighbouring l.o.s, even though the dust causing the extinction is likely to be correlated. This produces discontinuities in many published extinction maps, similar to an artefact known as the ``fingers of god''. In this paper, we present a method which uses extinctions and distances towards multiple stars in a collective manner simultaneously, by taking into account the correlation between neighbouring l.o.s; unlike earlier work by \citet{2014MNRAS.445..256S}, the covariance of our Gaussian process prior is in the dust density space; not in the extinction. This enables us to build a 3D map of the dust density which is free of the fingers-of-god effect, as a result of the isotropy of our smooth prior. We do this by modelling the dust density, which is a local property of the ISM, rather than the extinction, which is the integral over all the dust along the l.o.s. We use a non-parametric model which allows us to avoid adopting an explicit -- and inevitably overly simple -- functional form for the variation of dust density in the Galaxy. The true variation of dust is far too complex to be captured by a parametric model, which we can both define in advance and fit well enough with available data. We instead use a Gaussian process, which constrains the variation of the dust densities without choosing a particular functional form for its spatial variation. The Gaussian process instead defines the form of the covariance function between all points in space. This approach permits a wide range of functional variations in 3D space. To build the dust map, we need the extinctions toward and distances of a large number of stars. Such data are being obtained by Gaia, which will soon provide astrometry and spectrophotometry for more than $10^9$ stars brighter than a G-band magnitude of about 20. The expected end-of-mission parallax standard error is around 25 $\mu$as at $G=15$ mag \citep{2014EAS....67...23D}, from which we can estimate distances \citep[e.g.][]{2015PASP..127..994B}. Gaia is also equipped with two low resolution spectrophotometers which will provide the spectral energy distribution of all observed sources. The Gaia Data Processing and Analysis Consortium (DPAC) will use these to estimate the astrophysical parameters of individual stars, including the l.o.s reddening/extinction \citep{2013A&A...559A..74B}. The end-of-mission Gaia $A_{V}$ precision is expected to be around 0.03 to 0.05 mag for stars with $G \leq 15$, increasing to 0.2 mag for sources down to $G=20$ (Andrae et al.\ 2016, in preparation). Individual stellar extinctions and distances are therefore the inputs for our method. This paper is organized as follows. We introduce our method in section \ref{method}, whilst covering some technical details in the appendix. In section \ref{simulations} we demonstrate the model and its ability to recover the true dust values, using both toy simulations and a a simulated Gaia catalogue. We apply the model to real data over a small region of the sky (the Kepler field) in section \ref{real_data}. We summarize in section \ref{discussion} and discuss some aspects of our method including its current strengths and weaknesses. | \label{discussion} We have introduced a new non-parametric method for building a smooth, three-dimensional map of dust density which avoids l.o.s effects. It uses a Gaussian process prior to constrain the variation of the dust density in 3D space, but without assuming a specific functional form for the spatial dependence. It instead uses a covariance function which varies with the separation of points. This allows the model to infer the dust density in unobserved regions. Our model uses the 3D positions of stars together with their l.o.s extinctions as its input data and infers the posterior probability density function (PDF) over the dust at selected points. This PDF is a Gaussian, and we showed that its mean and standard deviation have analytic solutions. While the l.o.s to the observed stars are divided into discrete cells, predictions are made at arbitrary points without any discretization being necessary. We used a truncated covariance function in the Gaussian process which involves two hyperparameters: a correlation length scale $\lambda$ and a dust amplitude $\theta$. The latter can be set from the properties of the input data; adjusting $\lambda$ gives us flexibility to model dust variations on different length and amplitude scales. The only requirement is that $\lambda$ be larger than the cell sizes. While $\lambda$ is some characteristic length scale, our model can and does probe dust structures of much smaller scales (as in figure \ref{fig:mosaiclambdatheta}). We could try to fix these hyperparameters by calculating the Bayesian evidence. The evidence (or ``marginal likelihood'') is the probability of observing the data, for fixed $\lambda$ and $\theta$, averaged over all possible instantiations of the model. We compute this by drawing one sample from the J-dimensional Gaussian process prior (which gives us J values for the dust density), calculating the likelihood for these model dust densities, repeating it for a large number of times (e.g. $K = 10^5$), and then averaging these likelihoods \begin{equation} P(\extvecN | \lambda, \theta) \,=\, \frac{1}{K} \sum_{k=1}^K {P}_{k}(\extvecN | \{{f}_{i}\}, {V}_{N}) \label{eqn:evidence} \end{equation} where $\{{f}_{i}\}$ are calculated attenuations (equation \ref{eqn:dustsumvec}) using dust densities drawn from the prior. Having done this for various $\lambda$ and $\theta$, we then calculate the Bayes factors, which are the ratio of these evidences (for different $\lambda$ and $\theta$) to the one with the specific values of $\lambda$ and $\theta$ used for our APOKASC data (section \ref{real_data}; $\lambda = 2$ kpc and $\theta$ = $4 {\times} {10}^{-7}$). We report these in table \ref{table:evidence}. We get values for the Bayes factors in the case of APOKASC data (section \ref{real_data}) which agree broadly with what we calculated for $\lambda$ and $\theta$. But in the case of the simulated data (section \ref{simulated_data}), the Bayesian evidence does not give us a useful discrimination between models. Most values are very close to zero for a range of $\theta$ and $\lambda$ because our simulated data have high extinctions, which are not well represented by a Gaussian process prior with zero mean. The APOKASC data, in contrast, have smaller extinctions. This shows that using a non-zero mean in the Gaussian process prior will better construct the dust density in regions with higher extinctions, such as the disk of the galaxy and the spiral arms. A different covariance function could also be used in the Gaussian process. We tested various forms of the covariance function, such as truncated exponential forms, but they did not make a significant difference to our results. The covariance function that we are using has the advantage that we can get different variation slopes by changing $\alpha$ (see fig. \ref{fig:gneiting}). \begin{table} \caption[]{Logarithm of the Bayes factors for APOKASC data for different ranges of $\lambda$ and $\theta$.} \begin{tabular}{ ||c||c||c||c||c|| } \hline \multicolumn{5}{|c|}{$\log_{10}$\,(Bayes factor)} \\ \hline $\theta$ $\backslash$ $\lambda$ (pc) & 500 pc & 1000 pc & 2000 pc & 3000 pc \\ \hline $1{\times}10^{-8}$ &$ -1.41$ & $-0.87$ & $-0.43$ & $-0.21$ \\ \hline $1{\times}10^{-7}$ & $0.05$ & $0.14$ & $0.15$ & $0.13$ \\ \hline $1{\times}10^{-6}$ & $0.05$ & $-0.06$ & $-0.18$ & $-0.26$ \\ \hline $1{\times}10^{-5}$ & $-0.40$ & $-0.48$ & $-0.65$ & $-0.69$ \\ \hline \end{tabular} \label{table:evidence} \end{table} One drawback of our model is that its computation time increases non-linearly with the number of stars, $N$, and the number of cells, $J$. As explained in Appendix \ref{sec:acceleration}, the time-consuming part is the (one-off) inversion of the $J\times J$ covariance matrix $\gpcovJ$, which takes time $\ofo(J^n)$ to compute, where $n$ is typically $\lesssim 3$ but can be reduced to around 2.3 \citep{matinvert}, as well as various matrix inversions and multiplications taking time $\ofo(NJ^2)$, which must be done for every prediction. For a problem with $N=230$ and $J=3203$, inverting $\gpcovJ$ took two minutes (using a single core on a modest AMD Opteron 6380 CPU). Making predictions at multiple points can then be done in parallel: 200 predictions took 4 minutes with 40 cores, or 1.2 seconds per point. The computation time for more points is proportional to the number of points. For a problem with $N=1000$ and $J=8185$, it took 40 minutes to invert $\gpcovJ$, and 34 seconds per point to make new predictions for 1000 new points (again with 40 cores). This is 28 times longer than the previous case, which agrees reasonably well with the $\ofo(NJ^2)$ scaling suggested above (which gives $(1000\times 8185^2)/(200\times 3203^2) = 33$). The limiting factor when scaling this up to larger applications may be the memory rather than the run-time. For the case of $J=12\,000$, we needed 8 GB of RAM per core. This number is determined primarily by the number of cells, $J$, because the largest matrix has size $J \times J$. However, as we use sparse matrix methods and a truncated covariance function, the RAM required will not continue to grow as $J^2$. It is rather the density of cells in space, rather than the number of cells, which will ultimately drive the memory requirements. Using the run-time numbers from above, and ignoring memory limitations, then with $N$=10\,000 and $J$=100\,000, the $\gpcovJ$ inversion takes around 30 days (with just one core; this could be accelerated if $\gpcovJ$ inversion is parallelized too). Then even with 10\,000 cores running for 30 days we could only make predictions at 30\,000 points. This (and $N$) is too small to build up a useful dust density map over a large volume of space. One way to accelerate the computations is to use approximate matrix inversions, which can be done in time $\ofo(N^2)$. Alternatively, instead of trying to model the entire volume in one, we could partition it into partially overlapping regions, solve for each separately, and then join them. Thought is required to combine the overlapping regions without discontinuities, but recall that any two points separated by more than $\lambda$ are not connected by our model anyway, due to the truncated covariance matrix (which itself does not produce discontinuities). Optimal partitioning is an area for future investigation. While our method takes into account the uncertainties in the extinction measurements, it does not yet make use of the distance uncertainties. This will be necessary in some practical applications, as even from Gaia more distant and/or fainter stars will have poor distance estimates from the parallaxes \citep[e.g.][]{2015PASP..127..994B}. Including distances in the likelihood model is one approach, but would make the solution non-analytic. We are currently exploring this approach and will report on it in a future paper. Once these developments have been made and tested, the model will be ready to be applied to the Gaia data. | 16 | 9 | 1609.08917 |
1609 | 1609.03833_arXiv.txt | We present a galactic chemical evolution model which adopts updated prescriptions for all the main processes governing the dust cycle. We follow in detail the evolution of the abundances of several chemical species (C, O, S, Si, Fe and Zn) in the gas and dust of a typical dwarf irregular galaxy. The dwarf irregular galaxy is assumed to evolve with a low but continuous level of star formation and experience galactic winds triggered by supernova explosions. We predict the evolution of the gas to dust ratio in such a galaxy and discuss critically the main processes involving dust, such as dust production by AGB stars and Type II SNe, destruction and accretion (gas condensation in clouds). We then apply our model to Damped Lyman-$\alpha$ systems which are believed to be dwarf irregulars, as witnessed by their abundance patterns. Our main conclusions are: i) we can reproduce the observed gas to dust ratio in dwarf galaxies. ii) We find that the process of dust accretion plays a fundamental role in the evolution of dust and in certain cases it becomes the dominant process in the dust cycle. On the other hand, dust destruction seems to be a negligible process in irregulars. iii) Concerning Damped Lyman-$\alpha$ systems, we show that the observed gas-phase abundances of silicon, normalized to volatile elements (zinc and sulfur), are in agreement with our model. iv) The abundances of iron and silicon in DLA systems suggest that the two elements undergo a different history of dust formation and evolution. Our work casts light on the nature of iron-rich dust: the observed depletion pattern of iron is well reproduced only when an additional source of iron dust is considered. Here we explore the possibility of a contribution from Type Ia SNe as well as an efficient accretion of iron nano-particles. | The origin and the evolution of dust is one of the most important problems in Astrophysics. Cosmic dust plays a central role in the physics of the interstellar medium (ISM): it governs the scattering, absorption, re-emission of stellar light (Des\'ert et al. 1990; Witt $\&$ Gordon 2000) and it affects the spectral energy distribution (SED) of background sources (Silva et al. 1999; Granato et al. 2000). Dust properties have been determined from many kind of observations such as infrared continuum emission, depletion patterns in the ISM (Jenkins et al. 2009), isotopic anomalies in meteorites (Gail et al. 2009), extinction (Aguirre et al. 1999) etc. Refractory elements (e.g., Si, Mg, Fe, Ni) are the ones which are subject to elemental depletion since a fraction of their abundances in the ISM is incorporated into dust grains. The circumstellar environments of evolved stars represents the sites where cosmic dust comes from, producing materials of silicate and carbonaceous type, i.e. the most important populations of dust species in the Universe (Draine $\&$ Li 2007). Stellar winds eject these dust particles in the ISM, and then, dust experiences lots of processes, which can decrease or increase its abundance and affect its size (Ferrarotti $\&$ Gail 2006, Zhukovska $\&$ Gail 2008). Thermal sputtering, evaporation in grain-grain collision, thermal sublimation or desorption are some examples of destruction processes, but the most important mechanism for cycling dust back to the gas phase resides in supernova shocks (Jones et al. 1994, McKee 1989). On the other hand, grain growth by dust coagulation and metals accretion onto preexisting grains increases either the dust mass or size and preferably occur in molecular clouds (Liffman $\&$ Clayton 1989, Hirashita 2000, Asano et al. 2013). These clouds are the sites where stars form and where new production of dust occur. All these processes together give rise to the so called ``dust cycle''. Dwek (1998; hereafter D98) developed a chemical evolution model of the Milky Way, taking into account all the processes participating in the dust cycle. Since D98, significant progress has been made concerning dust properties, both in theory and in observations. Calura et al. (2008; hereafter C08) modeled the evolution of dust in galaxies of different morphological types. New theoretical prescriptions about dust processing have appeared in more recent papers (Inoue 2011; Piovan et al. 2011; Asano et al. 2013; Hirashita 2013; Mattsson et al. 2015). High quality observations carried out using satellites and ground based telescopes have shed light on the nature and composition of the dust in local and high-redshift galaxies (Carilli et al. 2001; Draine 2003; Micha{\l}owski et al. 2010; Wang et al. 2010; Gall et al. 2011). In particular, Damped Lyman Alpha (DLA) systems (Wolfe et al. 1986, 2005) offer a great opportunity for studying the composition of the ISM and constraining dust properties at different cosmic times (Pei et al. 1991; Pettini et al. 1994; Vladilo $\&$ P\'eroux 2005; Vladilo et al. 2011). In this work, we present a galactic chemical evolution model that incorporates updated prescriptions for dust production, accretion and destruction. We compare our results with other models widely employed in literature, and we constrain the origin and properties of cosmic dust by comparing these models with data of dwarf irregular galaxies and DLA systems. One of the specific aims is to find a plausible interpretation, in terms of dust evolution, of the rise of iron depletion with increasing metallicity in DLA systems that has been known for a long time but for which there are no clear explanations (Vladilo 2004). In the first part of the paper we present the new chemical model with dust, which adopts updated prescriptions for all the main processes governing the dust cycle. In section~\ref{chemical_model} we present the chemical evolution model adopted while the explanation for the dust model will be given in section~\ref{dust_model}. In section~\ref{dust_model_comparison} we present our results on the amount, composition and evolution of dust in dwarf irregular galaxies. In the second part we show in section~\ref{DLA_sec} the comparison between our dust model and observational data of DLA systems. Finally, in section~\ref{sec_conclusions} some conclusions are drawn. | In this work we have presented a chemical evolution model which takes into account the presence of dust utilizing new updated prescriptions. Dust formation is treated in the same way as first done by Dwek (1998), but with the inclusion of improved condensation efficiencies of Piovan et al. (2011). With respect to other models such as those of Calura et al. (2008) or Grieco et al. (2014), we have also changed the accretion and destruction prescriptions, which are two very important processes in dust evolution. We have applied our model to dwarf irregular galaxies and DLA systems. Our main results can be summarized as follows: \begin{enumerate} \item We studied the dust production rate and the processes occurring in the ISM during the galactic lifetime of a typical irregular galaxy. We have computed the evolution of dust by considering dust production (Type II SNe, AGB stars), destruction and accretion processes. It is worth noting that we excluded the Type Ia SNe as dust producers since there is no observational evidence for that. We have found that dust accretion plays a fundamental role in dust evolution and in certain phases it becomes the dominant process, governing the evolution of the dust mass in the ISM, as predicted by Asano et al. (2013). Moreover our model reproduces the observed dust-to-gas ratios as function of metallicity in such galaxies. \item We investigated the impact of the cut-off of high mass stars (from $18$ to $80Mo$) on the chemical evolution of a typical irregular galaxy. We fail in reproducing the metallicity values observed in dwarf irregulars when the cut-off mass is assumed to be in the range 18-25$M_{\odot}$. On the other hand, this effect does not deeply affect the predicted range of dust-to-gas ratio. \item We compared the dust formation when both P11 and D98 condensation efficiencies are adopted. We found that the rate production of carbon is almost the same using different prescriptions, while the main differences concern silicates: using D98 condensation efficiencies, Type Ia SNe play a significant role and, in addition, a major contribution is given by Type II and LIMS. \item Dust destruction represents a negligible process in dwarf irregulars, whereas the galactic wind is an important mechanism which can affect dust evolution: we showed that it can be the main responsible for stopping the accretion process in the ISM. \item We compared our model for irregulars with the data of DLA systems and we found that these objects can indeed be irregular galaxies, as already suggested in previous papers. We found a particular combination of parameters which best fit the DLAs. In particular, our comparison shows that the depletion pattern of silicon in these objects is well reproduced by the dust contributions of Type II SNe, AGBs and by the accretion process. \item In the case of iron, at variance with the case of silicon, we find a good agreement with the data only when an extra dust source is considered: in particular, we tested the case of dust production by Type Ia SNe and the case of a more efficient accretion in the ISM. The different behavior of iron and silicon that we find brings new evidence that a significant fraction of iron has to be incorporated into a dust population different from silicates, as suggested by previous works. Furthermore, as part of iron dust should be decoupled from silicates, it is possible that such species could originate in a different way: in particular, our results are consistent with a metallicity-dependent accretion of iron nano-particles. \end{enumerate} | 16 | 9 | 1609.03833 |
1609 | 1609.02232_arXiv.txt | We present the discovery of a variable optical counterpart to the unidentified gamma-ray source \fgl\, and argue this is a new compact binary millisecond pulsar (MSP) candidate. We show \fgl\ hosts a semi-detached binary with a 0.86955$\pm$0.00015~d orbital period and a F6-type companion star at an estimated distance of D=1.1$\pm$0.2~kpc, with a radial velocity curve semi-amplitude K$_2$=214.1$\pm$5.0~km~s$^{-1}$ and a projected rotational velocity of Vsin(i)=73.2$\pm$1.6~km~s$^{-1}$. We find a hard X-ray source at the same location with a 0.5--10~keV luminosity L$_\mathrm{X}$=2.6$\times$10$^{32}$~(D/1.1~kpc)$^2$~erg~s$^{-1}$, which strengthens the MSP identification. Our results imply a mass ratio q=M$_2$/M$_1$=0.26$^{+0.02}_{-0.03}$ if the companion star fills its Roche lobe, and q$\gtrsim$0.26 in any case. This classifies \fgl\ as a ``redback'' binary MSP; if its MSP nature is confirmed, this will be the brightest compact binary MSP in the optical band (r'$\simeq$14.3~mag) and will have the longest orbital period among Galactic field systems (nearly 21~hr). Based on the light curve peak-to-peak amplitude ($\Delta$r=0.19~mag), we further suggest that the orbital inclination is high and the putative pulsar mass is close to canonical (M$_1$$\simeq$1.3--1.6~M$_\odot$). Finally, we discuss the lack of heating signatures and asymmetric optical light curves in the context of other redback MSPs. | \label{sec:intro} Nearly one thousand gamma-ray sources from the {\it Fermi} Large Area Telescope (LAT) catalog remain unindentified, about a third of the total sample \citep{Acero15}. This is often due to the lack of counterparts at longer wavelengths, and offers an appealing discovery space. Among the identified Galactic sources, pulsars are the most numerous class \citep{Nolan12,Abdo13}, and {\it Fermi}-LAT is uncovering a new population of nearby binary millisecond pulsars \citep[MSPs; see, e.g.,][]{Hessels11,Ray12,Roberts13}. Dynamical studies of a few MSPs in compact binaries (``black-widow'' and ``redback'' pulsars) have revealed evidence for massive neutron stars, with masses well above the canonical value of 1.4~M$_\odot$ \citep{Kerkwijk11,Romani12b,Kaplan13}. These and related pulsar discoveries have pushed the maximum neutron star mass to more than two solar masses \citep{Demorest10,Antoniadis13}, placing tighter constraints on the equation of state above nuclear saturation density. Finding more such systems is crucial to establish their properties as a class, and constitutes a promising first step towards identifying the most massive neutron stars. \begin{figure*} \begin{center} \resizebox{1.7\columnwidth}{!}{\rotatebox{-90}{\includegraphics[]{charts.ps}}} \caption{ {\it Left:} Full IAC-80-CAMELOT r' band image of the field of \fgl, showing the (95\%) error ellipses from each of the three {\it Fermi}-LAT point source catalogs, as indicated. {\it Right:} Zoomed finding chart showing the location of our variable optical source and redback candidate J0212 (P$_\mathrm{orb} \simeq$20.9~hr, r'$\simeq$14.3mag; red arrow), the Swift X-ray counterpart 1SXPS J021210.6+532136 (red circle), the nearby W~UMa contact binary discovered in this work (P$_\mathrm{orb} \simeq$7.5~hr, r'$\simeq$17mag, purple arrow; see Appendix~\ref{app:wuma}) and the nearby galaxy ZOAG~G134.92-07.63 (black diamond). The brown ellipse shows the Chandra 3-sigma location. } % \label{fig:chart} \end{center} \end{figure*} Radio timing observations of {\it Fermi}-LAT sources have unveiled a flurry of new pulsars \citep{Hessels11,Ray12}. However, black-widow and redback MSPs are often occulted for a large fraction of the orbit \citep{Archibald13}, making their direct detection as radio pulsars challenging. Blind searches for gamma-ray pulsations have met with some success \citep{Pletsch12b}, yet they are computationally challenging, especially when the signal is smeared out by Doppler shifts in short (but unknown) orbital period binaries. Here we take another approach to identify the {\it Fermi}-LAT source \fgl, similar to that of \citet{Romani12} and \citet{Kong12}: we search for and find a variable optical counterpart (Section~\ref{sec:phot}) that matches a previously unidentified X-ray source (Sec.~\ref{sec:xray}). Our spectroscopic study (Section~\ref{sec:spec}) allows us to measure the orbital period, the amplitude of the radial velocity curve, as well as the companion's spectral type and projected rotational velocity. Together with the multi-wavelength properties of the source, which we present in the rest of Section~\ref{sec:results}, this strongly suggests that the binary hosts a recycled ``redback'' MSP. We discuss the system's orbital parameters, potential and peculiarities in Section~\ref{sec:discussion}. | \label{sec:discussion} \subsection{Masses and orbital parameters} \label{sec:m1} We have discovered a variable optical counterpart to the gamma-ray source \fgl, J0212, which coincides with a previously unclassified X-ray source. The multi-wavelength properties of J0212 are consistent with a binary millisecond pulsar in a compact orbit (P$_\mathrm{orb}$=20.869[4]~h) with a F6$\pm$2 main sequence companion star. From the measured Vsin(i)=73.2$\pm$1.6~km~s$^{-1}$ and K$_2$=214.1$\pm$5.0~km~s$^{-1}$ (Sec.~\ref{sec:spec}), assuming a Roche-lobe filling, tidally locked and spherically symmetric companion star, we find a mass ratio q=M$_\mathrm{2}$/M$_\mathrm{1}$=0.26$^{+0.02}_{-0.03}$ \citep[where M$_\mathrm{2}$ and M$_\mathrm{1}$ are the masses of the secondary/companion and the primary/neutron star; see][]{Wade88}. We note this is strictly a lower limit and thus q$\gtrsim$0.26, as the companion may be smaller than its Roche lobe. \begin{figure} \begin{center} \resizebox{0.95\columnwidth}{!}{\rotatebox{-90}{\includegraphics[]{M1-i.ps}}} \caption{ {\it Bottom:} Mass of the primary or neutron star as a function of inclination (i=90$^\circ$ corresponds to the orbital plane viewed edge on). The curves shown use the indicated values of P$_\mathrm{orb}$, K$_2$ and q (Sections~\ref{sec:spec} and \ref{sec:m1}). {\it Top:} Radius of the secondary or companion star as a function of inclination, in units of the semi-major axis a. Horizontal lines show the Roche lobe radius R$_\mathrm{L2}$ for our estimated values of q, as indicated. Curves show R$_2$/a for the observed light curve amplitude $\Delta$r=0.19~mag \citep{Morris85}. Gray-shaded regions show the preferred ranges of i, M$_1$ and R$_2$/a from imposing that R$_2$$\leq$R$_\mathrm{L2}$ (Sec.~\ref{sec:m1}). } % \label{fig:m1} \end{center} \end{figure} These orbital parameters classify \fgl\ as a ``redback'' MSP \citep[which have M$_\mathrm{2}$$\gtrsim$0.1--0.5~M$_\mathrm{\odot}$; e.g.,][]{Roberts11}. The measured absorbing column density (Sec.~\ref{sec:xray}) corresponds to a pulsar dispersion measure DM$\sim$50~pc~cm$^{-3}$ according to the correlation presented in \citet{He13}. This optical ephemeris will allow targeted searches for radio (preferentially around phase 0.5 to avoid pulsar occultation) and gamma-ray millisecond pulsations from this system, not reported to date. \fgl\ does not appear in the NRT radio pulsar search of {\it Fermi}-LAT sources presented by \citet{Guillemot12}, and we find no radio counterpart in the NVSS 1.4~GHz survey \citep[][the field was not covered by FIRST]{Condon98}. If its redback nature is confirmed, J0212 will have the longest P$_\mathrm{orb}$ among the compact binary millisecond pulsars in the Galactic field (both redbacks and black widows). To our knowledge, only two redbacks with longer orbital period are known, both residing in globular clusters: J1748-2446AD (P$_\mathrm{orb}$=1.09~d) in Terzan 5 and J1740-5340 (P$_\mathrm{orb}$=1.35~d) in NGC~6397 \citep[][respectively]{Hessels06,DAmico01}. From our K$_2$ and q measurements we derive the M$_1$-i relation shown in Fig.~\ref{fig:m1} (bottom; using M$_1$=(1+q)$^2$ P K$_2$$^3$ / (2$\pi$G sin$^3$(i))). If the primary is indeed a neutron star, our results imply that i$>$50$^\circ$ and M$_2$$<$0.8~M$_\odot$ for any plausible M$_1$$<$3~M$_\odot$, and M$_1$$\gtrsim$1.3~M$_\odot$ for any i. Furthermore, since we find that irradiation effects are negligible in J0212 (Secs.~\ref{sec:spec} and \ref{sec:colours}), we can constrain the inclination by ascribing the observed light curve amplitude (peak-to-peak amplitude in r' $\Delta$r=0.19mag) to ellipsoidal modulation of the tidally locked companion. Using the analytical method presented by \citet[][equation (6) in particular]{Morris85} and taking limb- and gravity-darkening coefficients for a Solar metallicity F5 star with log(g)=4.5 \citep{Claret11}, we can constrain the companion radius R$_2$/a (Fig.~\ref{fig:m1}, top), where a is the semi-major axis of the orbit. Imposing that R$_2$ is smaller than the corresponding Roche lobe radius \citep[R$_\mathrm{L2}$;][]{Eggleton83}, since there is no evidence for mass transfer and accretion disk lines are not observed (Sec.~\ref{sec:spec}), we find that i$\gtrsim$76$^\circ$ and therefore M$_1$$\simeq$[1.3--1.6]~M$_\odot$ and M$_2$$\simeq$[0.34--0.42]~M$_\odot$. Thus according to the \citet{Morris85} relation, the inclination should be high (i$\gtrsim$76$^\circ$) and the companion should be close to filling its Roche lobe (R$_2$/R$_\mathrm{L2}$$>$98\%) in order to produce the observed ellipsoidal modulation. The M$_2$ constraints above imply that the companion is significantly larger and hotter than an isolated star of its mass. If we take M$_1$=1.5~M$_\odot$ and M$_2$=0.38~M$_\odot$ (i=80$^\circ$), the corresponding R$_2$$\simeq$R$_\mathrm{L2}$=1.3~R$_\odot$ is roughly consistent with the radius of an isolated F6V star. Similar ``stripped'' or ``bloated'' companion stars, hotter and/or larger than isolated stars of the same mass, are seen in redback \citep{Crawford13} and black-widow MSPs \citep{Kerkwijk11}, as well as neutron star transients in quiescence (\citealt{Bildsten01}; see also \citealt{Orosz03} for further discussion). Finally, we note that if the pulsar is detected, J0212 will be an ideal system for an accurate neutron star mass measurement: it has a bright, non-irradiated companion star in a likely high inclination orbit. Pulsar timing and high-resolution spectroscopy can yield much more precise measurements of q and K$_2$, respectively. Detailed modelling of the optical light curve and spectral lines can give tighter and more robust constraints on the inclination angle. \subsection{Colours and broadband SED} \label{sec:colours} After correcting for interstellar reddening using E(B-V) = N$_\mathrm{H}$ / (3.1$\times$1.8$\times$10$^{21}$~cm$^{-2}$) = 0.251 (with the N$_\mathrm{H}$ measured from the X-ray spectrum; Sec.~\ref{sec:xray} and \citealt{Predehl95}), the corresponding (g'-r') and (r'-i') colours are fully consistent with the F6 spectral type we find from optical spectroscopy \citep[][see Sec.~\ref{sec:spec}]{Pecaut13}. The infrared, optical and UV fluxes are also consistent with an F6 main sequence star with radius $\simeq$1.3~R$_\odot$ at D$\simeq$1.1~kpc, as shown in Fig.~\ref{fig:sed}. We note that the W2 flux (the shortest wavelength UV measurement available, at 2120~\AA) is about 40\% higher than the M2 flux (at 2310~\AA). Comparing with stellar atmosphere models of a F5V star (Fig.~\ref{fig:sed}; \citealt{Kurucz93}), we attribute this to a relative drop in M2 flux due to FeII absorption bands in the $\sim$2300--2400~\AA\ range. The spectral energy distribution (Fig.~\ref{fig:sed}) also shows the energy budget of the companion star ($\sim$10$^{34}$~erg~s$^{-1}$) which dominates in the optical band, the gamma-rays from the putative MSP ($\sim$10$^{33}$~erg~s$^{-1}$) and the shock between the MSP and companion winds (intrabinary shock), which presumably powers the X-ray emission ($\sim$10$^{32}$~erg~s$^{-1}$). The colours remain approximately constant along the orbit (Fig.~\ref{fig:lc}), implying little or no temperature change between the different sides of the companion. This complete lack of irradiation or ``heating'' of the companion by the pulsar wind and radiation is exceptional among compact binary MSPs. From the allowed range of temperatures (6640--6150~K; Sec.~\ref{sec:spec}), a semi-major axis of 4.7~R$_\odot$ (i=80$^\circ$) and assuming that the pulsar spin-down power $\dot{E}$ is emitted isotropically, we estimate an upper limit on $\dot{E}$$<$[1--4]$\times$10$^{35}$~erg~s$^{-1}$ \citep[for an irradiation efficiency 10--30\%, following][]{Breton13}. This limit is consistent with the $\dot{E}$ of most MSPs. Thus we suggest that the lack of heating is simply due to the wide orbit of J0212. To our knowledge only PSR~J1740-5340, which is also in a long P$_\mathrm{orb}$$\simeq$32~hr orbit, has shown a similar lack of irradiation \citep{Orosz03}. \subsection{Light curves and distance} We measure phase-zero magnitudes of 14.96, 14.36 and 14.15 in the g', r' an i' bands, respectively, which correspond to a dereddened V=13.83. This makes J0212 the brightest compact binary MSP known to date \citep[about two magnitudes brighter than PSR J1723-2837 in V,][]{Crawford13}. For an F6 main sequence star with absolute magnitude M$_\mathrm{V}$=3.7 \citep{Pecaut13}, implicitly assuming that the companion radius is unperturbed, we estimate a distance to J0212 of D=1.1$\pm$0.2~kpc (where the error corresponds to the allowed range of spectral types, F4--F8). As discussed in Section~\ref{sec:m1}, the radius we infer from our RVC and LC analysis is consistent with this spectral type (but the mass is not). The optical orbital light curves are clearly asymmetric (Fig.~\ref{fig:lc}): the light maximum at phase 0.25 (companion at ascending node) is 0.03, 0.04 and 0.06 magnitudes brighter than the maximum at phase 0.75 (descending node) in the i, r and g bands, respectively. The minimum at phase 0 (companion at inferior conjunction) is about 0.01--0.03 magnitudes brighter than the minimum at phase 0.5. While the asymmetry in the depth of the minima might be partly explained by limb- and gravity-darkening effects, models for compact binary MSP light curves in general, and their asymmetric maxima in particular, are still under development \citep{Breton13,Li14,Salvetti15}. We leave detailed modelling for future work, and simply point out that optical light curves similar to those of J0212 have been observed in confirmed and candidate redback MSPs (PSR~J1628-32, P$_\mathrm{orb}$$\simeq$5~hr, \citealt{Li14}; PSR~J2129-0429, P$_\mathrm{orb}$$\simeq$15.2~hr, \citealt{Bellm16}; 1FGL~J0523.5-2529, P$_\mathrm{orb}$$\simeq$16.5~hr?, \citealt{Strader14}; 3FGL~J2039.6-5618, P$_\mathrm{orb}$$\simeq$5.4~hr?, \citealt{Salvetti15}). \subsection{X-rays and intrabinary shock} J0212 features the highest X-ray luminosity (2.6$\times$10$^{32}$~erg~s$^{-1}$; Sec.~\ref{sec:xray}) among redback (in the pulsar state) and black-widow MSPs \citep{Linares14c}. This places J0212 in the group of relatively X-ray luminous redbacks. PSR~1740-5340, on the other hand, had $\sim$10 times lower L$_\mathrm{X}$ \citep[][and references therein]{Bogdanov10}. The wide orbit and lack of irradiation signatures strongly suggests that the companion wind in these systems is not driven by MSP heating effects. Thus while irradiation of the companion appears to depend critically on P$_\mathrm{orb}$, the luminosity of the intrabinary shock between the pulsar and companion winds is not simply related to P$_\mathrm{orb}$. This may be due to a hotter companion with a larger mass loss rate in the wind, which would compensate the larger orbital separation. \subsection{Towards a systematic search} To conclude, our results highlight the potential of small-aperture optical telescopes like the IAC80 in identifying and characterizing {\it Fermi}-LAT sources. The exceptionally bright optical counterpart to \fgl\ that we have discovered, with r'$\simeq$14.3~mag, sets the record for the brightest MSP with a low-mass companion. A complete photometric survey of unidentified Fermi-LAT sources should thus target a broad range spanning more than 12 magnitudes, from the faintest g'$\sim$27~mag black-widow counterparts \citep[e.g.,][]{Breton13} to the g'$\sim$15~mag of J0212 or brighter. It should be noted, moreover, that intra-night optical variability searches are biased towards short-P$_\mathrm{orb}$ and strongly irradiated systems, which give rise to drastic magnitude changes in only a few hours \citep[e.g.,][]{Schroeder14}. Our findings show that a systematic search for compact binary MSPs should also target low amplitude ($\lesssim$0.05~mag~hr$^{-1}$) and long period ($\gtrsim$12~hr) optical variability. \textbf{Note:} Soon after our manuscript was submitted, Li et al. (ArXiv:1609.02951) published a similar analysis of \fgl. Our results mostly agree. \textbf{Acknowledgments:} This article is based on observations made with the IAC80, WHT, INT and NOT telescopes operated on the islands of Tenerife and La Palma by the IAC, ING and NOTSA. We thank R. Alonso, R. Ashley, M. Balcells, P. Blay, P. Chinchilla, E. G\'omez, A. Oscoz, C. Protasio and J. Telting for granting, assisting or performing some of the observations presented here. {\sc iraf} is distributed by the NOAO, operated by AURA under cooperative agreement with the NSF. We gratefully acknowledge the use of T. Marsh's {\sc molly}, {\sc pamela} and {\sc ultracam} analysis packages. This research was supported by the Spanish MINECO under grants AYA2013--42627 and AYA2012--38700. PRG is supported by a Ram\'on y Cajal fellowship (RYC2010--05762). JC acknowledges support by the Leverhulme Trust through the Visiting Professorship Grant VP2-2015-046. This publication makes use of data products provided by HEASARC, 2MASS and WISE. \vspace{-0.7cm} | 16 | 9 | 1609.02232 |
1609 | 1609.05624_arXiv.txt | An axion is an attractive candidate for the inflation scenario in addition to other phenomenologically favorable scenarios such as a solution of the strong CP problem and the candidate of dark matter in our Universe. Furthermore, string theory predicts a lot of axion particles in the low-energy effective theory through the compactification of extra-dimensional space. When the axion is associated with the higher-dimensional form fields, the form of the axion potential is protected by the higher-dimensional gauge symmetries apart from other matter fields. The axion potential is then generated by the spontaneous or explicit breaking of axionic shift symmetry originating from the higher-dimensional one. The higher-dimensional operators in the axion potential are still controlled, and the cosmological observables induced by the axion inflation are predictable. To construct the inflationary favorable axion potential, the axion decay constant is required to be enough large to obtain the flat direction in the axion potential, in particular, the trans-Planckian decay constant for the natural inflation~\cite{Freese:1990rb}. However, in string theory, the decay constant of a closed string axion is typically around the string scale or grand unification scale $10^{16}\,{\rm GeV}$~\cite{Choi:1985je,Banks:2003sx,Svrcek:2006yi}. When the axion decay constant is of order the Planck scale, the axion inflation generically predicts ${\cal O}(1)$ tensor-to-scalar ratio as can be seen in the Lyth bound~\cite{Lyth:1996im}, which argues that the tensor-to-scalar ratio $r$ is closely related to the inflaton field range, $\Delta \phi$, during the inflation. Under the assumption that a variation of $r$ is negligible over the period $\Delta \phi$, the approximate relation is obtained as~\cite{Lyth:1996im}, \begin{eqnarray} \frac{\Delta \phi}{M_{\rm Pl}} \simeq {\mathcal O}(1) \times \left(\frac{r}{0.01}\right)^{1/2}. \label{eq:Lyth_bound} \end{eqnarray} This indicates that if $\Delta\phi < M_{\rm Pl}$, $r\lesssim 0.01$ is obtained and we call this class of inflation model the small-field inflation throughout this paper. Although the large-field axion inflations ($r\gtrsim 0.01$) are consistent with the recent Planck data~\cite{Planck:2013jfk,Ade:2015lrj}, the weak gravity conjecture~\cite{ArkaniHamed:2006dz} suggests that the higher-order instanton effects give a sizable effect for the axion potential with a trans-Planckian axion decay constant and these would generically violate the slow-roll axion inflation. In this respect, we consider the axion inflations with the decay constant below the Planck scale or string scale, which are favorable from the aspects of the weak gravity conjecture. The flat direction required in the inflation can be realized by choosing the proper parameters in the axion potential. Since the obtained inflaton potential is categorized into the class of small-field axion inflation, it predicts the small amount of gravitational waves and low inflation scale in comparison with the prediction of large-field axion inflation. Such a low-scale inflation is also influential to the isocurvature perturbation originating from the QCD axion. When all the dark matter is dominated by the QCD axion, the current Planck result constrains the Hubble scale during the inflation $H_{\rm inf}$~\cite{Planck:2013jfk}, \begin{align} H_{\rm inf} <0.87 \times 10^7{\rm GeV}\left(\frac{f_{\rm QCD}}{10^{11}\,{\rm GeV}}\right)^{0.408}, \label{eq:iso} \end{align} where $f_{\rm QCD}$ is the decay constant of the QCD axion. It is then possible to avoid the isocurvature constraint by the low-scale inflation, although the upper bound of $f_{\rm QCD}$ depends on the initial misalignment angle of the axion and dilution mechanism after the inflation~\cite{Kawasaki:2004rx,Hattori:2015xla,Akita:2016usy}. Recently, some of the authors conjectured that, in a certain class of small-field axion inflation derived from type IIB superstring theory~\cite{Kobayashi:2015aaa},\footnote{ The model in Ref.~\cite{Kobayashi:2015aaa} can lead to both small-field and large-field inflations.} the tensor-to-scalar ratio $r$ correlates with the axion decay constant $f$ as follows~\cite{Kadota:2016jlw}, \begin{align} r \sim 10^{-6} f^{2q}, \nonumber \end{align} where the fractional number $q$ depends on the model. In the example of Refs.~\cite{Kobayashi:2015aaa,Kadota:2016jlw}, we obtain $q=2$. This behavior originates from sinusoidal functions in the axion inflation potential. The above relation could also predict the magnitude of the inflation potential and the inflaton mass by $f$. In general, superstring theory leads to the axion potential with one or more sinusoidal terms induced by several non-perturbative terms. Thus, it is important to extend the previous analysis to other axion inflation scenarios. In this paper, we further study such dependence of the axion decay constant for not only cosmological observables, but also the reheating temperature and dark matter abundance for the general class of small-field axion inflations, which are the mixture of polynomial and sinusoidal functions suggested in the axion monodromy inflation~\cite{Silverstein:2008sg,Kobayashi:2014ooa,Higaki:2014sja} and general form of sinusoidal functions suggested in the natural and multi-natural inflations~\cite{Freese:1990rb,Choi:2014rja,Czerny:2014wza,Czerny:2014xja}.\footnote{The scalar potential including modular functions in superstring theory can effectively lead to such a multi-natural inflation \cite{Abe:2014xja}.} We constrain the axion decay constant realizing the small-field axion inflations by the isocurvature perturbation originating from the QCD axion, successful Big-Bang nucleosynthesis (BBN) and dark matter abundance. As will be shown, it is quite interesting that the allowed range of the axion decay constant corresponds to the typical decay constant region realized in superstring theory, when our axion is the closed string axion~\cite{Choi:1985je,Banks:2003sx,Svrcek:2006yi}. In the remainder of this paper, we first discuss the conditions leading to the general class of small-field axion inflations and analytical form of cosmological observables as a function of the decay constant in Sec.~\ref{sec:2}. In Sec.~\ref{sec:3}, we derive the constraints for the axion decay constants from the reheating process and dark matter abundance. We summarize our conclusion in Sec.~\ref{sec:con}. | \label{sec:con} We have discussed the general class of small-field axion inflation, which is the mixture of polynomial and sinusoidal functions with an emphasis on the small axion decay constant compared with the Planck scale. In contrast to the large-field axion inflation such as the natural inflation~\cite{Freese:1990rb} and axion monodromy inflation~\cite{Silverstein:2008sg}, the small-field axion inflation predicts the small amount of primordial gravitational waves and low inflation scale. This class of inflation models is motivated by the weak gravity conjecture, which prohibits the trans-Planckian axion decay constant and the constraint from isocurvature perturbation due to the QCD axion. When the axion decay constants and parameters in the scalar potential satisfy the certain conditions leading to the successful small-field axion inflations as discussed in Sec.~\ref{sec:2}, we find that the cosmological observables are written in terms of the axion decay constants in a systematic way. Furthermore, the axion decay constant is severely constrained within the range $10^{14}\,{\rm GeV}\lesssim f\lesssim 10^{17}\,{\rm GeV}$, where the lower bounds are put by $T_{\rm reh}\gtrsim {\cal O}(5)$ MeV in order not to spoil the successful BBN, whereas the upper bounds are set by the constraint from the isocurvature perturbation due to the QCD axion with $f_{\rm QCD}=10^{12}\,{\rm GeV}$. This constrained axion decay constant naturally appears in the string theory, when our discussed axion corresponds to the closed string axion~\cite{Choi:1985je,Banks:2003sx,Svrcek:2006yi}. Although the parameters in the axion potential should be properly chosen to achieve a flat enough direction in the axion potential, the small-field axion inflation is attractive from the theoretical and phenomenological points of view. | 16 | 9 | 1609.05624 |
|
1609 | 1609.00237_arXiv.txt | In this paper we study the abilities of an atmospherical mesoscale model in forecasting the classical atmospherical parameters relevant for astronomical applications at the surface layer (wind speed, wind direction, temperature, relative humidity) on the Large Binocular Telescope (LBT) site - Mount Graham, Arizona. The study is carried out in the framework of the ALTA project aiming at implementing an automated system for the forecasts of atmospherical parameters (Meso-Nh code) and the optical turbulence (Astro-Meso-Nh code) for the service-mode operation of the LBT. The final goal of such an operational tool is to provide predictions with high time frequency of atmospheric and optical parameters for an optimized planning of the telescope operation (dome thermalization, wind-dependent dome orientation, observation planning based on predicted seeing, adaptive optics optimization, etc...). Numerical simulations are carried out with the Meso-Nh and Astro-Meso-Nh codes, which were proven to give excellent results in previous studies focused on the two ESO sites of Cerro Paranal and Cerro Armazones (MOSE Project). In this paper we will focus our attention on the comparison of atmospherical parameters forescasted by the model close to the ground with measurements taken by the observatory instrumentations and stored in the LBT telemetry in order to validate the numerical predictions. As previously done for Cerro Paranal (Lascaux et al., 2015), we will also present an analysis of the model performances based on the method of the contingency tables, that allows us to provide complementary key information with the respect to the bias and RMSE (systematic and statistical errors), such as the percentage of correct detection and the probability to obtain a correct detection inside a defined interval of values. | \label{sec:intro} This paper is part of a general validation study on the forecasts of meteorological parameters and optical turbulence (OT) at the Large Binocular Telescope (LBT) site of Mount Graham (Arizona), performed in the context of Advanced LBT Turbulence and Atmosphere (ALTA\footnote{\url{http://alta.arcetri.astro.it}}) Center project. The ALTA project aims to implement and automate forecast system for LBT using a mesoscale hydrodynamic meteorological model, either for classical meteorological parameters (wind speed and direction, temperature, relative humidity) which are relevant for ground-based astronomy, and astroclimatic parameters ($C_N^2$ profiles, seeing $\epsilon$, isoplanatic angle $\theta_0$, wavefront coherence time $\tau_0$) which are relevant for adaptive optics applications (AO). The final outcome of ALTA project will be the deployment of an operational tool to provide predictions with high time frequency and spatial resolution in order to support and optimize the telescope operation, such as dome thermalization, wind-dependent dome orientation, observation planning based on seeing and other OT parameters condition and optimization of AO systems operation. The first commissioning for the atmospheric parameters is set in June 2016, while the second commissioning for the OT parameters is on December 2016.\\ In order to provide the above results we use MESO-Nh (Lafore et. al. 1998 [\cite{lafore98}]) model developed by the Laboratoire d'Aerologie, CNRM and Meteo France, together with the Astro-Meso-Nh module (Masciadri et. al. 1999 [\cite{masciadri99a}]) which is used to provide forecast of OT parameters. The forecasts are relative to the night time frame, when the LBT telescope is operative.\\ The MESO-Nh model performances for the OT forecasts have already been tested in previous studies performed on major telescope installations, such as Roque de los Muchachos[\cite{masciadri2001a}], San Pedro Martir[\cite{masciadri2004}], Cerro Paranal[\cite{masciadri99b}] and also at very low latitudes such as Antarctica[\cite{lascaux2009,lascaux2010}]. In terms of performances and reliability the critical milestones have been: the proposition of a technic for the model calibration (Masciadri and Jabouille, 2001[\cite{masciadri2001}]) and the validation of such a technique with a sample of 10 nights (a large sample for that epoch) (Masciadri et al., 2004[\cite{masciadri2004}]). A few years ago the model has been validated at Mount Graham (Hagelin et. al. 2011 [\cite{hagelin2011}]) using measurements of a Generlized SCIDAR related to the more extended site testing campaign performed at Mt. Graham (43 nights) (Masciadri et. al 2010 [\cite{masciadri2010}]). Finally the model forecasts for atmospheric and astroclimatic parameters were part of a large validation campaign conducted within the MOSE project, commissioned by the European Souther Observatory (ESO) in order to prove the feasibility of an automated forecast system for their installations in Cerro Paranal and Armazones (VLT and E-ELT respectively) (Masciadri et. al. 2013 [\cite{masciadri2013}], Lascaux et. al. 2013-2015[\cite{lascaux2013,lascaux2015}]).\\ In this paper we will show the preliminary results of an ongoing validation study, in view of the first commissioning of June 2016. In this context, we will limit ourselves to study the performance of Meso-Nh model in forecasting the standard atmospheric parameters near the ground layer, using as a reference the LBT weather stations placed on the telescope dome. The study is performed on a limited sample of nights, which was used to test the model stability and tune the settings in order to obtain the most efficient configuration for an operational tool.\\ A more detailed study, performed on a larger sample of nights, will be published in a forthcoming paper.\\ In section \ref{sec:obs} we will describe the LBT site and the measurements sets that were use for the present validation study. In section \ref{sec:mod_conf} we will describe the model configuration, including the numerical setup used for this study. In section \ref{sec:res} we will present the results of the validation in terms of statistical operators and contingency tables. Finally in section \ref{sec:concl} we will draw the conclusions.\\ | \label{sec:concl} In this paper we presented the preliminary results of the ongoing validation study on the operational forecast system being developed for LBT, as part of the ALTA project, performed on a limited 22 nights test sample. We performed this test on the atmospherical parameters measured near the ground level, using as a reference the weather instruments of the LBT dome, in the context of the first commissioning for the atmospheric forecasts in June 2016. The results we obtained are extraordinarily good and show excellent performance in reconstructing all the tested parameters. Temperature is predicted with an excellent degree of accuracy, with an RMSE~$=0.98 ^\circ C$ and a PC~=~88.3\% with respect to the winter distribution of measurements. Wind speed is reconstructed with an RMSE~=~2.7~m/s and a PC~=~71.6\%, which raises to a POD~$=87.3\%$ in the case of strong winds ($>10$~m/s). The result we obtained for the wind direction is incredibly positive, with an RMSE~$=17^\circ$, RMSE$_{rel}$$\simeq$ 9.5\% and a PC~=~87.4\%, showing a major increase of accuracy with respect to the previous studies conducted on ESO telescope sites. The result for the relative humidity is good, with a global RMSE~=~17.3\%, however we still have problems in correctly detecting high values of the parameter, leading to a unsatisfactory prediction when the value is higher than $\sim 70\%$. This problem will be addressed before the forthcoming final validation study that will be performed on a larger statistical sample of nights. | 16 | 9 | 1609.00237 |
1609 | 1609.05412_arXiv.txt | With the high-resolution data from the \emph{Interface Region Imaging Spectrograph}, we detect a light wall above a sunspot light bridge in the NOAA active region (AR) 12403. In the 1330 {\AA} slit-jaw images, the light wall is brighter than the ambient areas while the wall top and base are much brighter than the wall body, and it keeps oscillating above the light bridge. A C8.0 flare caused by a filament activation occurred in this AR with the peak at 02:52 UT on 2015 August 28, and the flare's one ribbon overlapped the light bridge which was the observational base of the light wall. Consequently, the oscillation of the light wall was evidently disturbed. The mean projective oscillation amplitude of the light wall increased from 0.5 Mm to 1.6 Mm before the flare, and decreased to 0.6 Mm after the flare. We suggest that the light wall shares a group of magnetic field lines with the flare loops, which undergo a magnetic reconnection process, and they constitute a coupled system. When the magnetic field lines are pushed upwards at the pre-flare stage, the light wall turns to the vertical direction, resulting in the increase of the light wall's projective oscillation amplitude. After the magnetic reconnection takes place, a group of new field lines with smaller scales are formed underneath the reconnection site and the light wall inclines. Thus, the projective amplitude decreases remarkably at the post-flare stage. | Solar flares are energetic phenomena in the solar atmosphere, releasing dramatic electromagnetic energy spanning the range from X-ray to radio wavelengths. In the standard two-dimensional (2D) flare model (CSHKP models; Carmichael 1964; Sturrock 1966; Hirayama 1974; Kopp \& Pneuman 1976), a filament rises above the neutral line and then initially drives the flare process. The rising filament pushes the overlying magnetc field lines upwards, and the resulting losses of pressure below form an inward magnetic force towards the neutral sheet. This force drives antiparallel magnetic filed lines to converge, leading to the formation of a current sheet, and magnetic reconnection begins to take place. Thus the released energy heats the coronal plasma and also accelerates particles. The accelerated particles flow downwards from the reconnection site along the newly formed magnetic field lines, and in the lower solar atmosphere, the flare ribbons are generated (Priest \& Forbes 2002). The flare ribbons observed in H$\alpha$ and ultraviolet (UV) wavelengths are the conspicuous characteristics of solar flares and are usually located on either sides of the polarity inversion line. The flare ribbons move apart during the reconnection process and the separation generally stops at the edge of the sunspots. However, Li \& Zhang (2009) reported that flare ribbons sometimes sweep across the whole sunspots. Sunspots are concentrations of magnetic fields and the overturning motion of the plasma is hindered by the strong magnetic field in the sunspot umbra (Gough \& Tayler 1966). Bright structures within the umbra are signatures of not completely suppressed convection, and light bridges are the best known representatives of these structures (Sobotka et al. 1993; Borrero \& Ichimoto 2011). The magnetic field of light bridges is generally weaker and more inclined than the local strong and vertical field (Lites et al. 1991; Rueedi et al. 1995; Leka 1997; Jur{\v c}{\'a}k et al. 2006). Recent simulations and observations have shown that a light bridge's magnetic field is twisted and related to emerging magnetic fields (Louis et al. 2015; Toriumi et al. 2015a,b; Yuan \& Walsh 2016). Above the light bridges, some chromospheric activities have been observed in the forms of jets and surges (Asai et al. 2001; Shimizu et al. 2009; Louis et al. 2014). Recently, Yang et al. (2015) reported an oscillating light wall above a sunspot light bridge. The light wall is brighter than the ambient areas while the wall top is much brighter than the wall body in 1330 {\AA}. Hou et al. (2016) revealed that some light walls are multilayer and multithermal structures which occur along magnetic neutral lines in active regions (ARs), not just above the light bridge. However, the work about the magnetic topology of the light wall is rare. In this Letter, we report that a light wall (oscillating above a sunspot light bridge) is disturbed by a C8.0 flare while one ribbon of this flare intrudes into the sunspot and overlaps the light bridge. The mean projective amplitude of light wall's oscillation increases at the pre-flare stage, and decreases after the flare. Using the coordinated observations from the \emph{Interface Region Imaging Spectrograph} (\emph{IRIS}; De Pontieu et al. 2014) and the \emph{Solar Dynamic Observatory} (\emph{SDO}; Pesnell et al. 2012), we investigate this event in detail for understanding the magnetic configuration of the light wall. | With the high tempo-spatial \emph{IRIS} and \emph{SDO} observations, we detect a light wall oscillating with a mean period of 4.0 minutes above a sunspot light bridge in NOAA AR 12403. The light wall is brighter than the ambient regions while the top and base of the light wall are much brighter than the wall body in 1330 {\AA} channel. On 2015 August 28, a C8.0 flare caused by a filament's activation occurred in this AR with the peak at 02:52 UT. We first observe that one of the flare ribbons intruded into the sunspot and then overlapped the light bridge which was the observational base of the light wall. As a result, the oscillation of light wall was obviously disturbed by the ribbon. The mean projective oscillation amplitude of the light wall increased from 0.5 Mm to 1.6 Mm before the flare, and decreased to 0.6 Mm after the flare. In addition, the images of the LASCO C2 on board the \emph{SOHO} are adopted to study the CME related to the C8.0 flare. This flare is an eruptive flare which results in a CME with an average speed of 253 km s$^{-1}$ and a width angle of 21{\degr}. The light wall has been reported in several works. Yang et al. (2015) reported an oscillating light wall above a sunspot light bridge and interpreted the oscillations of the light wall as the leakage of p-modes from below the photosphere. Hou et al. (2016) revealed that some light walls are multilayer and multithermal structures which occur along magnetic neutral lines in active regions. As a newfound structure, the light wall's driving mechanism and magnetic topology have not been well understood. The present work reports a C8.0 flare disturbing a light wall above a sunspot light bridge, which may contribute to the understanding of the light wall's magnetic configuration. Louis et al. (2014) proposed that the dynamic chromospheric jets above the light bridge seem to be guided by the magnetic field lines. Tian et al.(2014) reported sub-arcsecond bright dots in the transition region above sunspots and suggested that some of these bright dots appear to be located at the bases of magnetic loops. In this event, we suggest that the light wall shares a group of magnetic field lines with the flare loops involved in magnetic reconnection process, and is located at the bases of these magnetic loops. They constitute a coupled system (see Figure 4(a)). When the activated filament begins to rise, the overlying field lines, which connect the positive light bridge fields and the negative plage fields, are pushed upwards before the flare's onset (between T1 and T2 in Figure 3(g)). Then the light wall turns to close to the vertical direction (see Figure 4(b)). Due to the projection effect, the observations along the LOS show the light wall's mean projective maximum height increases from 2.7 Mm to 5.2 Mm. Around the T2 in Figure 3(g), magnetic reconnection takes place (see Figure 4(c)). Underneath the reconnection site, the magnetic filed lines with smaller scales are newly formed and the flare ribbons appear. As a result, the light wall turns to away from the vertical direction (see Figure 4(d)). Thus, the average projective maximum height of light wall decreases remarkably to 3.0 Mm after the flare (after the T2 in Figure 3(g)). Checking the AIA 304 {\AA} observations during the flare's evolution, we detect a quiescent filament lying above the neutral line. When the flare occurred, the dark material flow was observed between the two flare ribbons. Therefore, we consider that this filament was partly activated and subsequently rose, leading to the C8.0 flare. For eruptive flares, the vertical magnetic fields on both sides of the current sheet correspond to the legs of CME-related expanding field lines (Forbes et al. 2006; Aulanier et al. 2010). Here, we detect a CME associated with this flare by LASCO C2 data, which is consistent with the illustration in Figure 4. | 16 | 9 | 1609.05412 |
1609 | 1609.05909_arXiv.txt | We examine the properties of barred disc galaxies in a $\Lambda$CDM cosmological hydrodynamical simulation from the EAGLE project. Our study follows the formation of 269 discs identified at $z=0$ in the stellar mass range $10.6<\log M_{*}/M_{\odot} <11$. These discs show a wide range of bar strengths, from unbarred discs ($\approx 60\%$) to weak bars ($\approx 20\%$) to strongly barred systems ($\approx 20\%$). Bars in these systems develop after redshift $\approx 1.3$, on timescales that depend sensitively on the strength of the pattern. Strong bars develop relatively quickly (in a few Gyr, $\sim 10$ disc rotation periods) in systems that are disc dominated, gas poor, and have declining rotation curves. Weak bars develop more slowly in systems where the disc is less gravitationally important, and are still growing at $z=0$. Unbarred galaxies are comparatively gas-rich discs whose rotation speeds do not exceed the maximum circular velocity of the halos they inhabit. Bar lengths compare favourably with observations, ranging from 0.2 to 0.8 times the radius containing $90\%$ of the stars. Bars slow down remarkably quickly as they grow, causing the inner regions of the surrounding dark halo to expand. At $z=0$ strong bars have corotation radii roughly ten times the bar length. Such slow bars are inconsistent with the few cases where pattern speeds have been measured or inferred observationally, a discrepancy that, if confirmed, might prove a challenge for disc galaxy formation in $\Lambda$CDM. | \label{SecIntro} \begin{figure*} \begin{center} \includegraphics[width=\linewidth,clip]{figs/fig1b.pdf} \end{center} \caption{Disc galaxy sample from EAGLE used in this paper. {\it Left:} Galaxy stellar mass, $M_{*}$, as a function virial mass $M_{200}$. Solid line indicates the prediction of the abundance-matching model of \citet{Guo2010}, for reference. {\it Middle:} Flattening parameter $c/a$, measured as the ratio of the eigenvalues of the principal axes of the inertia tensor of the stars. {\it Right:} Minor axis stellar velocity dispersion, expressed in units of the total. Vertical dashed lines indicate the conditions required to be selected as ``discs'' in our analysis. Discs are shown as coloured circles, spheroidal systems as open triangles, and visually identified ongoing mergers or disturbed systems as crosses. The colour scheme denotes the strength of the bar pattern (see Fig.~\ref{FigBarz0}).} \label{FigGxSample} \end{figure*} The stellar discs of spiral galaxies are dynamically fragile structures prone to morphological and dynamical transformation. These might be triggered by external processes, such as accretion events, mergers, or the tidal effects of satellites and neighbouring galaxies. They may also result from internal processes, which tend to be more subtle and to operate over longer timescales but are nonetheless effective at inducing notable changes in the morphology and structure of the disc. Internal processes invariably redistribute the disc's angular momentum, driving mass inwards while pushing angular momentum outwards. Angular momentum redistribution requires non-axisymmetric features \citep{LBK1972,Tremaine1984}, of which bars---i.e., extended and radially-coherent $m=2$ perturbations to the disc's azimuthal structure---are a particularly clear example. Bars come in many different sizes and shapes; from short inner bars that affect a small fraction of stars to long bars that extend out to the confines of the disc; and from thin rectangular bars that correspond to a single, dominant $m=2$ mode to fat oval structures with sizable contributions from higher even Fourier modes. Taking them all together, bars are an extremely common phenomenon in disc galaxies, and are present in a large fraction of discs \citep[e.g.,][]{Eskridge2000,whyte2002,marinova2007,Sheth2008,Gadotti2011}. The origin of bars has long been an issue of debate. N-body discs quickly turned into bars in early simulations \citep{Miller1968, Hockney1969}, a result that suggested a ``global instability'' that would affect essentially all stellar discs unless stabilized by a suitable mechanism \citep[see][for a review of early work]{Sellwood1993}. One such mechanism was proposed by \citet{Ostriker1973}, who argued, in an influential paper, that cold stellar discs required the presence of a massive non-rotating dark halo in order not to go bar unstable. In this scenario, bars develop quickly in systems where the disc is dominant (perhaps triggered by accretion events or tides), whereas unbarred discs are those whose dynamics is largely dominated by the dark halo \citep[][]{Efstathiou1982}. This idea is still widely in use and criteria for instantaneous ``bar instability'' are a key ingredient of semi-analytic models of galaxy formation that attempt to match the morphological mix of the observed galaxy population \citep[see, e.g.,][and references therein]{Lacey2015}. More recent work, however, has led to a more nuanced view, and it is now recognized that bars, weak and strong, may develop gradually in most stellar discs that are relatively massive and kinematically cold, even when the halo is important. Indeed, in some cases massive halos have even been found to promote bar formation: one clear example is that provided by \citet{Athanassoula2002a}, who shows that bars may develop faster in disc-dominated systems, but they eventually become stronger in halo-dominated ones. Halos apparently do not prevent bars, but, rather, just delay their formation \citep[see][for a recent review]{Athanassoula2013}. Once formed, bars are a conduit for the transfer angular momentum from the disc to other parts of the system. The more angular momentum a bar is able to lose, the longer and thinner (``stronger'') it can become. To grow, then, bars need material to absorb the angular momentum lost by stars that join the bar, be it other stars in the outer disc or particles in the halo that might get trapped in resonances with the bar. A massive halo can therefore aid this process by providing a sink for the angular momentum lost by stars that make up the growing bar \citep{Athanassoula2003}. Bars can therefore develop gradually over many orbital periods, on a timescale that depends mainly on the relative importance of disc vs halo, but likely influenced as well by the velocity dispersion of the disc (hotter discs are less prone to global distortions) and by the potential well depth of the halo (faster moving halo particles are harder to trap into resonances with the bar). The two scenarios---instantaneous bar instability vs gradual bar growth---should in principle yield different predictions for the abundance, size, and pattern speeds of bars, as well as for their evolution with redshift, but detailed predictions in a proper cosmological setting have yet to be worked out. One corollary of gradual bar formation is that bars that grow longer/stronger should slow down \citep{Hernquist1992,Debattista2000}. This is because bars cannot extend beyond corotation, the radius where the angular speed of a circular orbit equals that of the bar pattern \citep{Contopoulos1980}. Angular speeds decrease outwards, so the longer the bar grows the slower its pattern speed must become. The bar cannot grow longer than the disc, of course, but it can continue to slow down, implying that the ratio between corotation radius ($r_{\rm corot}$) and bar length ($l_{\rm bar}$) can provide interesting constraints on the relative importance of the disc and halo, as well as on the time elapsed since the onset of the bar \citep{Debattista2000}. Although the measurements are challenging and often indirect, most observational estimates point to ``fast bars'' where $r_{\rm corot}<1.4\, l_{\rm bar}$ \citep[see, e.g.,][]{Elmegreen1996,Corsini2011}. Interestingly, bar slowdown might also have discernible effects on the dark matter density profile, since it is the halo that absorbs much of the disc angular momentum, especially in the case of strong bars. A number of studies have indeed suggested that the central density cusps expected in cold dark matter halos \citep{Navarro1996,Navarro1997} might be softened and perhaps erased\footnote{Others, however, have argued otherwise, so the issue is still under debate \citep{Sellwood2003,Sellwood2008,Dubinski2009}.} by a bar \citep{Weinberg2002,Holley-Bockelmann2005}. This result has important consequences for models of gamma ray emission by dark matter annihilation in the direction of the Galactic centre \citep{Schaller2016}: the Milky Way is, after all, a barred galaxy \citep[e.g.,][]{Blitz1991}. The discussion above suggests that the abundance of barred galaxies, together with the distribution of bar strengths, lengths, and pattern speeds, may provide interesting constraints on the mass, size, and kinematics of disc galaxies, on the time of their assembly, and on the mass and density profiles of the dark matter halos they inhabit. This is important, because, once a cosmological model has been adopted, the very same properties that govern bar growth are independently specified by other constraints, and cannot be tuned arbitrarily. Success in reproducing the properties of the barred galaxy population in a particular cosmology is thus far from assured. \begin{figure*} \begin{center} \includegraphics[width=\linewidth,clip]{figs/fig3c.pdf} \end{center} \caption{Projected stellar density maps for three examples of an unbarred galaxy (top row), a weak bar (middle row) and a strongly barred disc (bottom row). The leftmost column shows face-on views of the three galaxies. % Dotted, dashed, and solid circles on the images indicate the galaxy radius, $r_{\rm 90}$, the bar length, $l_{\rm bar}$, and the stellar half-mass radius, $r_{50}$. The middle and right panels show face-on and edge-on views, respectively, created with the radiative transfer code SKIRT \citep{Baes2011}. These images show the stellar light based on monochromatic SDSS $u$, $g$ and $r$ band filters and accounting for dust extinction. The rightmost column shows the radial profile of the bar strength parameter, $A_2(R)$, and indicates a few characteristic radii. } \label{FigGxImages} \end{figure*} In $\Lambda$CDM models---the current paradigm of structure formation---the relation between galaxy mass and halo mass may be derived using ``abundance matching'' arguments \citep{Frenk1988,Vale2006,Guo2010,Moster2013,Behroozi2013}. Further, galaxy sizes are also constrained by scaling laws such as the Tully-Fisher relation \citep[see, e.g.,][]{Ferrero2016}. Do galaxies that match those constraints also result in a barred galaxy population whose statistics, bar lengths, and pattern speeds are compatible with observation? We address this question here by examining the properties of disc galaxies in the EAGLE cosmological hydrodynamical simulation of a $\Lambda$CDM universe \citep{Schaye2015,Crain2015}. This paper is organized as follow. In Sec. \ref{SecNumSim} we briefly describe the numerical simulations and the galaxy sample selection. Sec.~\ref{SecRes} presents the results of our analysis, including the frequency of bars (Sec.~\ref{SecBarFreq}); bar lengths (Sec.~\ref{SecBarL}); bar growth (Sec.~\ref{SecBarG}); bar slowdown (Sec.~\ref{SecBarS}); and its effects on the halo mass profile (Sec.~\ref{SecHaloEvol}). We summarize our main conclusions in Sec.~\ref{SecConc}. \begin{figure} \begin{center} \includegraphics[width=\linewidth,clip]{figs/fig3.pdf} \end{center} \caption{Cumulative distribution of the bar strength parameter $A_2^{\rm max}$, compared with observational estimates of the bar fraction of galaxies with comparable stellar mass. The colour scheme assigns different hues of red to strong bars ($A_{2}^{\rm max}>0.4$), of green to weak bars ($0.2<A_{2}^{\rm max}<0.4$), and of blue to unbarred systems ($A_{2}^{\rm max}<0.2$).} \label{FigBarFreq} \end{figure} | \label{SecConc} We have used a $\Lambda$CDM cosmological hydrodynamical simulation from the EAGLE project to study the formation of barred galaxies. The simulation evolves a box $100$ Mpc on a side with $2 \times 1504^3$ particles, half of which are baryonic and half dark matter. Our study focusses on a narrow range of stellar mass, $10.6<\log{M_{*}/M_{\odot}} <11$, which are resolved with at least 22,000 star particles. Of the $495$ galaxies in that mass range identified at $z=0$ we select a sample of $269$ ``discs'', defined as flattened systems with minor-to-major axis ratio $c/a < 0.4$ and relatively low vertical velocity dispersion $\sigma_z/\sigma_{\rm tot}<0.5$. We identify barred galaxies by measuring the amplitude of the normalized $m=2$ Fourier mode of the azimuthal surface density profile as a function of cylindrical radius, and choose as a measure of bar strength the peak amplitude, $A_2^{\rm max}$. We consider ``barred'' all discs with $A_2^{\rm max}>0.2$. We follow the evolution of all these galaxies in order to estimate bar growth timescales, to identify which parameters predict the development of bars best, and to measure the evolution of bar strength, length, and pattern speed. Our main conclusions may be summarized as follows. \begin{itemize} \item About $40\%$ of EAGLE discs in our sample are barred, $20\%$ of them strong bars ($A_2^{\rm max}>0.4$) and another $20\%$ weak bars ($0.2<A_2^{\rm max}<0.4$). This bar frequency seems in reasonable agreement with observational estimates from \citet{Sheth2008,Barazza2008,Nair2010}. \item Bars in our simulated discs span a wide range in terms of length. They typically extend beyond the stellar half-mass radius but rarely exceed the radius containing $90\%$ of the stellar mass, in good agreement with observational estimates. In terms of each of those radii, the median bar length and interquartile range is given by $l_{\rm bar}/r_{50}=1.53^{+0.42}_{-0.41}$ and $l_{\rm bar}/r_{90}=0.35^{+0.18}_{-0.10}$. \item At $z=0$, bar strength correlates strongly with stellar half-mass radius (stronger bars form in smaller discs), hinting that, as expected from earlier work, bars develop preferentially in systems where the disc is gravitationally important. We also find that stronger bars develop in systems that are less gas-rich, and that have formed the bulk of their stars earlier than unbarred discs. \item Strong bars in our sample develop relatively quickly before saturating over a few Gyrs. Weak bars are still growing in strength at $z=0$, and take much longer to develop, with characteristic timescales approaching or even exceeding a Hubble time. Even our strongest, fastest growing bars take roughly $4$-$5$ Gyr ( a few dozen disc rotations) to form fully. \item The gravitational importance of the disc at its half-mass radius may be used to predict which galaxies will develop bars, but its predictive power may be enhanced by considering the overall importance of the disc in the system as a whole. Strong bars form in discs where baryons dominate and whose rotation speeds exceed the maximum circular velocity of the halo. Unbarred galaxies are discs where baryons are less important and whose rotation curves tend to rise in the outskirts. \item Strong bars slow down quickly as they grow and, at $z=0$ are in the ``slow bar'' regime, $r_{\rm corot}/l_{\rm bar}>1.4$. This is in contrast with the few bars whose pattern speeds have been inferred observationally, all of which are ``fast''. This discrepancy may either imply that bar slowdown rates are artificially high in simulations at EAGLE resolution \citep[e.g.,][]{Weinberg2007a}, or, as argued in earlier work, that producing long-lasting ``fast bars'' is a real challenge for $\Lambda$CDM \citep[e.g.,][]{Debattista2000}. \item The bar slowdown induces an expansion of the inner regions of the dark matter halo, as they capture the angular momentum of the forming bar. However, bars form in massive dense discs with heavily contracted halos, so despite the bar-induced expansion barred galaxy halos are still more centrally concentrated than unbarred galaxies of similar stellar mass. Our numerical resolution is not enough to let us ascertain whether this expansion may lead to the formation of constant density ``cores'' in barred galaxy halos. \end{itemize} Our overall conclusion is that current $\Lambda$CDM cosmological hydrodynamical simulations of cosmologically significant volumes such as EAGLE yield a population of simulated discs with bar fractions, lengths, and evolution that are in broad agreement with observational constraints. They also confirm earlier suggestions that ``slow bars'' might pose a severe challenge to this scenario. Although bars form in the manner and frequency expected, they slow down too fast through interaction with the dark halo. Unless the Universe has a population of slow bars that has yet to be recognized, or the bar slowdown we measure is artificially enhanced by limited numerical resolution, accounting for the presence of ``fast bars'' in strongly-barred discs is a clear goal for the next generation of $\Lambda$CDM simulations of galaxy formation. | 16 | 9 | 1609.05909 |
1609 | 1609.09250_arXiv.txt | Circumplanetary disks (CPDs) control the growth of planets, supply material for satellites to form, and provide observational signatures of young forming planets. We have carried out two dimensional hydrodynamical simulations with radiative cooling to study CPDs, and suggested a new mechanism to drive the disk accretion. Two spiral shocks are present in CPDs, excited by the central star. We find that spiral shocks can at least contribute to, if not dominate the angular momentum transport and energy dissipation in CPDs. Meanwhile, dissipation and heating by spiral shocks have a positive feedback on shock-driven accretion itself. As the disk is heated up by spiral shocks, the shocks become more open, leading to more efficient angular momentum transport. This shock driven accretion is, on the other hand, unsteady on a timescale of months/years due to production and destruction of vortices in disks. After being averaged over time, a quasi-steady accretion is reached from the planet's Hill radius all the way to the planet surface, and the disk $\alpha$-coefficient characterizing angular momentum transport due to spiral shocks is $\sim$0.001-0.02. The disk surface density ranges from 10 to 1000 g cm$^{-2}$ in our simulations, which is at least 3 orders of magnitude smaller than the ``minimum mass sub-nebula'' model used to study satellite formation; instead it is more consistent with the ``gas-starved'' satellite formation model. Finally, we calculate the millimeter flux emitted by CPDs at ALMA and EVLA wavelength bands and predict the flux for several recently discovered CPD candidates, which suggests that ALMA is capable of discovering these accreting CPDs. | {Planets form and grow in circumstellar disks.} Initally when a protoplanet forms, the solid core is surrounded by a hydrostatic gaseous envelope that is in contact with the planet's Hill sphere and the rest of the circumstellar disk. As the core's mass increases, the planet envelope contracts due to both the stronger gravity and cooling through radiation. Eventually, when the core mass reaches $\sim 10 M_{\oplus}$, the planet undergoes runaway accretion and contracts significantly. At this stage, the envelop shrinks significantly and detaches from the planet's Hill sphere (Papaloizou \& Nelson 2005). Material that resides beyond the Hill sphere can still flow into the Hill sphere, but forms a circumplanetary disk (CPD) around the protoplanet to conserve angular momentum (Lubow \etal 1999; Ayliffe \& Bate 2009). The accretion of the circumplanetary disk (CPD) onto the planet allows the continuous growth of the planet even after the planet's runaway growth stage. {Circumplanetary disks (CPDs) may provide observational signatures of young forming planets in disks.} Giant planets can be too faint to be detected, but accreting CPDs can be bright and detectable. For example, to form a giant planet of 1$-$10 M$_{J}$ mass within the circumstellar disk's life time ($\sim$ a few million years, Hernandez \etal 2007), the CPD needs to accrete at a rate of $\dot{M}\gtrsim 10^{-9}-10^{-8} \msunyr$. Such an accretion disk will have a luminosity (e.g. Owen 2014, Zhu 2015) of \begin{equation} L_{disk}=\frac{G{ M}_{p}\dot{M}}{2 {\rm R}_{J}}=1.5\times10^{-3}{\rm L}_{\odot} \frac{{M}_{p}}{1 {\rm M}_{J}}\frac{\dot{M}}{10^{-8}\msunyr}\,, \end{equation} which is as bright as a late M-type/early L-type brown dwarf and can be detected by current direct imaging techniques (Zhu 2015). Unlike a planet or brown dwarf which has an almost constant surface temperature, an accreting CPD has a lower temperature at larger disk radii, and those outer disk regions will emit significant infrared flux. Thus, the emission from an accreting CPD is redder than the emission from a planet or a brown dwarf. Direct imaging observations have found several red sources within circumstellar disks (Kraus \& Ireland 2012, Quanz \etal 2013, Biller \etal 2014, Reggiani \etal 2014, Currie \etal 2015) and their photometry at near-IR bands are consistent with accreting CPDs (Zhu \etal 2015). { $H_{\alpha}$ emission lines, which are another observational signature of accretion disks, are also found in some low-mass substellar objects (Zhou \etal 2014, Bowler et al. 2015). } The most direct evidence for accreting CPDs is LkCa 15b, where two accretion tracers, H$_{\alpha}$ line and near-IR thermal emission, have both been detected (Sallum \etal 2015). CPDs are also essential for satellite formation. In our solar system, most satellites around giant planets are in prograde, nearly circular and coplanar orbits, implying that they formed in a shared CPD orbiting within the planet's equatorial plane. The ratio between the total mass of four major Galilean satellites and the Jupiter's mass is $\sim 2\times10^{-4}$. The Saturnian satellite system also has a similar mass fraction with respect to the Saturn's mass (Canup \& Ward 2006). Assuming the gas-to-dust mass ratio is 100, a minimum gas mass of $\sim 0.02$ planet mass is required to produce these satellites (Canup \& Ward 2002, 2006). There are two main scenarios for Galilean satellite formations. One assumes that the satellites form in-situ in a CPD containing this amount of material (e.g. Lunine \& Stevenson 1982; Mosqueira \& Estrada 2003). This disk is referred to as the ``minimum mass sub-nebula'', which is the analog of ``minimum mass nebula'' for the circumstellar disk. Spreading 0.02 Jupiter mass within 30 Jupiter radii (where the four major satellites reside), the minimum mass sub-nebula has a very high gas surface density $\sim 10^{5} {\rm g}\,{\rm cm}^{-2}$. In the other scenario, the satellites form in a CPD with a much lower surface density, but the disk is dynamically evolving and being supplied by the circumstellar disk (Canup \& Ward 2002, 2006). This ``gas-starved'' scenario only requires that the total supplied mass during the whole satellite formation timescale is larger than $0.02$ planet mass. The key question for understanding the structure of CPDs is how CPDs accrete. They may accrete in similar ways as circumstellar disks accrete. Turner \etal (2014) suggest that the surface of CPDs can be ionized by X-rays so that the magnetorotational instability (MRI) can operate at the disk surface leading to accretion, a process similar to the layered accretion proposed in protoplanetary disks (Gammie 1996). However, Fujii \etal (2011, 2014) find that the active layer is so thin ($\Sigma\sim10^{-3}-10^{-2}{\rm g} \,{\rm cm^{-2}}$) in CPDs that the accretion through MRI is negligible. On the other hand, other non-ideal MHD effects (e.g. Hall effects, Kunz \& Lesur 2013, Lesur \etal 2014, Bai 2014) which are important in protoplanetary disks can also be important in CPDs, and CPDs may accrete through magnetic breaking (Keith \& Wardle 2014). Magnetocentrifugal disk wind can also be launched in CPDs, carrying away angular momentum and leading to disk accretion (Quillen \& Trilling 1998, Gressel \etal 2013). Almost all these proposed CPD accretion mechanisms (e.g. layered accretion, non-ideal MHD effects, disk wind) can find their roots in circumstellar disk models. Thus, they are facing the same uncertainties as the accretion mechanisms in circumstellar disks: they sensitively depend on the net magnetic fields assumed and the detailed microphysics in the disk (e.g. dust size distribution). On the other hand, CPDs are different from circumstellar disks in that they are subject to the tidal torque from the central star, and truncated within the Hill sphere of the planet (Martin \& Lubow 2011a). In addition, circumstellar disk material flows through the Hill sphere, continuously replenishing CPDs. These properties make CPDs similar to disks with inflows from companion stars in close binary systems or Cataclysmic Variables (CV). In these binary systems, the tidal torque from the companion star will excite spiral density waves in disks, and when these waves shock in disks, they can transport angular momentum to the disk leading to disk accretion. Previous inviscid isothermal simulations by Rivier \etal (2012) have suggested that accretion due to spiral shocks in CPDs is inefficient with Jupiter's mass doubling time $\sim$ 5 Myrs. Szul{\'a}gyi \etal (2014) have measured a much higher accretion rate in their 3-D isothermal inviscid simulations (10$^{-4}$ M$_{J}$/yr), but this measured accretion rate is not numerically converged with their higher resolution simulations. Instead, by measuring the torque exerted by the star in the simulations, they estimate that the real accretion rate is $\sim 2.5\times 10^{-6}$M$_{J}$/yr. Recent studies on spiral shocks (Ju \etal 2016) have shown that the accretion due to spiral shock dissipation is sensitive to the disk thermodynamics assumed. When the disk is hot and the Mach number, defined as the ratio between the Keplerian speed and the sound speed, is small ($<$10), the equivalent $\alpha$ of shock-driven accretion can reach $\sim$0.01-0.02. The Mach number of CPDs depends on the disk accretion rate and is sensitive to the equation of state applied in simulations (D'Angelo \etal 2003, Ayliffe \& Bate 2009, Machida 2009, Szul{\'a}gyi \etal 2016). Considering the potential importance of spiral shocks driving accretion in CPDs, in this paper we construct two-dimensional inviscid hydrodynamical simulations to study CPDs. These inviscid simulations differ from most previous works that use artificial viscosity (e.g. $\alpha$ viscosity) to sustain the disk accretion. Furthermore, to reduce the numerical viscosity in the CPD region, we adopt a grid structure centered on the planet. Since the thermodynamics which controls the disk Mach number is crucial for the shock-driven accretion, a simple radiative cooling scheme has been included in the simulations, which differs this work from inviscid simulations by Rivier \etal (2012) and Szul{\'a}gyi \etal (2014). We have also measured the mass accretion rate directly from simulations, unlike these two works where they use the torque exerted by the star onto the CPD to estimate the disk accretion rate. As Ju \etal (2016) has pointed out, shock-driven accretion is determined by the shock dissipation (the difference between the tidal torque exerted to the disk and the angular momentum flux carried away by the wave), instead of the total torque alone. We will show that, by allowing the disk being heated up by the shock, the inviscid simulation is numerically converged and the disk can reach a steady state, transferring inflow material from the Hill radius all the way to the central planet quasi-steadily. This shock driven accretion is very efficient with $\alpha\sim 0.001-0.01$. One caveat in our simulations is that our simulations are limited to 2-D and numerous previous 3-D simulations have shown that the infall from circumstellar to circumplanetary disks occurs at high altitudes (Bate \etal 2003, Machida \etal 2008, Tanigawa \etal 2012, Morbidelli \etal 2014, Szul{\'a}gyi \etal 2014, 2016). Furthermore, since the CPD is significantly heated and puffed up, our simple cooling treatment based on the thin disk approximation is inaccurate. 3-D simulations with realistic radiative transfer (Szul{\'a}gyi \etal 2016), thermodynamics, planet evolution (Ward \& Canup 2010) and even planetesimal accretion (D'Angelo \& Podolak 2015) are needed in future to confirm if spiral shocks can lead to efficient accretion in CPDs. In \S 2, our numerical method is introduced. The results are presented in \S 3, including both the disk structure and the accretion rate. After a short discussion in \S 4, the paper will be concluded in \S 5. The detailed energy budget of the accretion process is given in the appendix. | Circumplanetary disks (CPD) control the growth of planets, supply material for satellites to form, and provide observational signatures of young forming planets. In this paper, we provide a new mechanism to explain how CPDs accrete. We have carried out two dimensional hydrodynamical simulations to study CPDs using Athena++. Simple radiative cooling has been considered in the simulation. Different from most previous simulations, we choose a coordinate system centered on the planet, which significantly reduces the numerical error and enables simulations with small numerical viscosity. Our simulation domain extends from the circumstellar disk all the way to the Jupiter's surface. A gap in the circumstellar disk is prescribed, allowing us to control the inflow rate from the circumstellar to the circumplanetary disk. Two spiral shocks are present in CPDs, induced by the tidal force from the central star. We find that these spiral shocks can lead to significant angular momentum transport and energy dissipation in CPDs. { Meanwhile, dissipation and heating by spiral shocks have a positive feedback on shock-driven accretion itself.} As the disk is heated up by spiral shocks, the shocks become more open and propagate further into the inner disk, leading to more efficient angular momentum transport at the inner disk. On the other hand, shock driven accretion cannot guarantee strict steady disk accretion since the angular momentum transport depends on the global wave propagation and local shock dissipation. Mass will pile up in some regions of the disk and sometimes vortices are produced, which produce short-timescale (months to years) variability. The disk is adjusting itself through these variabilities and tries to maintain a quasi-steady state. Eventually, a quasi-steady state of accretion flow is reached in our simulations from the planet's Hill radius all the way to the planet surface. After averaging quantities over a long timescale, angular momentum budget is carefully analyzed. The effective $\alpha$-coefficient characterizing angular momentum transport due to spiral shocks is $\sim$0.001-0.02 even though the disk accretion rates span 4 orders of magnitude. { The $\alpha$ value is higher in a disk with a higher accretion rate due to the shock heating feedback.} With energy budget analysis, we show that radial advection of energy becomes important and the disk generates less infrared radiation than that from the thin disk approximation by a factor of $\sim$2. Thus, if we use the infrared flux calculated from the thin disk approximation to derive the accretion rate of CPDs, we may underestimate the disk accretion rate by a factor of 2. Finally, we calculate the flux from CPDs at ALMA and EVLA wavelength bands and predict the flux to several recent CPD candidates (e.g. HD169142b, HD100546b, LkCa 15b). These CPD candidates should be relatively bright at ALMA wavelength bands. In future, ALMA should be able to discover many accreting CPDs. Unlike near-IR emission which comes from the inner disk region, submm flux comes from the outer disk region which is less subject to vortex production and destruction. Furthermore the dynamical time is much longer at the outer disk. Thus submm flux will be more steady over time than the optical/near-IR flux. We may see that CPDs appear and disappear between two epoch optical/near-IR observations while they are bright all the time at ALMA bands. Although our simulations are limited to 2-D, the possibility that we may have understood how CPDs accrete and its huge implications on observations and satellite formation make it worth being studied in detail in future. | 16 | 9 | 1609.09250 |
1609 | 1609.07767_arXiv.txt | {The Mount Wilson \ion{Ca}{II} index \rhk is the accepted standard metric of calibration for the chromospheric activity versus age relation for FGK stars. Recent results claim its inability to discern activity levels, and thus ages, for stars older than $\sim$2 Gyr, which would severely hamper its application to date disk stars older than the Sun.} {We present a new activity-age calibration of the Mt. Wilson index that explicitly takes mass and \feh biases into account; these biases are implicit in samples of stars selected to have precise ages, which have so far not been appreciated.} {We show that these selection biases tend to blur the activity-age relation for large age ranges. We calibrate the Mt. Wilson index for a sample of field FGK stars with precise ages, covering a wide range of mass and \feh, augmented with data from the Pleiades, Hyades, M67 clusters, and the Ursa Major moving group.} {We further test the calibration with extensive new Gemini/GMOS \rhk data of the old, solar \feh clusters, M67 and NGC 188. The observed NGC 188 activity level is clearly lower than M67. We correctly recover the isochronal age of both clusters and establish the viability of deriving usable chromospheric ages for solar-type stars up to at least $\sim$6 Gyr, where average errors are $\sim$0.14 dex provided that we explicitly account for the mass and \feh dimensions. We test our calibration against asteroseismological ages, finding excellent correlation ($\rho$ = +0.89). We show that our calibration improves the chromospheric age determination for a wide range of ages, masses, and metallicities in comparison to previous age-activity relations.} {} | Stellar ages are fundamental parameters in our understanding of the chemo-dynamical evolution of the Galaxy and other stellar systems as well as exoplanetary systems. Ages, which are very difficult to gauge, are usually only estimated through methods optimized to restricted classes of stars since these are indirect parameters inferred from the time evolution of a range of observational quantities that do not uniquely characterize this range. Strong spectral lines are useful indicators of stellar chromospheric activity (CA) that is physically linked to the efficiency of angular momentum evolution. The stellar rotation and CA in single main-sequence stars decay monotonically with time, under the action of the torque produced by the magnetized stellar wind, as it is a potential indicator of age \citep{soderblometal1991}. The Mount Wilson (MW) project \citep{baliunasetal1995} has been monitoring the widely used S index, which is the ratio of the flux in the line cores of the \ion{Ca}{II} H \& K lines and two nearby continuum regions; this S index can be converted into the \rhk index, which is defined as the absolute line excess flux (line flux - photospheric flux) normalized to the bolometic flux \citep{linsky79,noyesetal1984}. The \rhk is the standard metric in the literature to retrieve stellar ages through CA-age relations \citep[hereafter CAR, e.g.,~][hereafter MH08]{mamajekhillenbrand2008}. Recent claims that the evolution of CA fluxes cannot be traced beyond $\sim$2 Gyr \citep{pace2013} imply that the derivation of CA ages is severely hampered for most of the age dispersion of the Galactic disk, thereby negating its usefulness as a tool to investigate Galactic evolution. Here we analyze the presence, in the CA-age relation, of mass and metallicity ($\mathrm{[Fe/H]}$) biases that are implicit in conventional methods of selecting solar-type stars with precise isochronal ages. We show that these biases have masked structural complexity in the CAR and present a new calibration explicitly relating age, activity, mass, and $\mathrm{[Fe/H]}$. We test this calibration against asteroseismological ages and new Gemini data on the \rhk indexes of the M67 and NGC 188 clusters (ages 4.0 and 6.0 Gyr, respectively). As an extension of MH08 CAR, we provide the activity distribution expected for 6 Gyr solar metallicity stars anchored on 49 NGC 188 members (16$\times$ the M08 sample for NGC 188), constraining the activity average and dispersion beyond the solar age. | We show that mass and metallicity biases are inevitably present in samples of stars selected to have small errors in age, severely distorting the intrinsic distribution of the chromospheric activity $versus$ age plane. The result is a dilution of the decay of chromospheric activity (CA) with time, apparently hampering the derivation of ages through CA beyond $\sim$2 Gyr. These biases can be corrected by means of an age-mass-metallicity-CA relation, which successfully reproduces stellar asteroseismological ages up to 10 Gyr. We further test this calibration by measuring the \rhk index of the 6.0 Gyr-old NGC 188 cluster from new and extensive Gemini/GMOS data. The CA level of NGC 188 is clearly lower than in the well-studied, 4.0 Gyr-old cluster M67. Our calibration successfully recovers correct ages for both clusters: chromospheric ages can be derived within $\sim$0.15 dex. We show that a more complete approach of the CA-age relation, including more variables, appears promising and may imply the ability of the age-mass-metallicity-activity relation to recover reliable stellar ages well beyond the solar age. Present calibrating samples, however, are far from representative of the whole relevant domain of mass, age, metallicity, and stellar activity. Particularly, further data on open clusters with a wide range of age and metallicity are essential to test this approach and possibly push the feasibility of chromospheric age determinations to the full range of age and metallicity of the Galactic disk. | 16 | 9 | 1609.07767 |
1609 | 1609.02142_arXiv.txt | Suppression of ${\rm H_2}$-–cooling in early protogalaxies has important implications for the formation of supermassive black holes seeds, the first generation of stars, and the epoch of reionization. This suppression can occur via photodissociation of ${\rm H_2}$ (by ultraviolet Lyman--Werner [LW] photons) or by photodetachment of ${\rm H^-}$, a precursor in ${\rm H_2}$ formation (by infrared [IR] photons). Previous studies have typically adopted idealised spectra, with a blackbody or a power--law shape, in modeling the chemistry of metal-–free protogalaxies, and utilised a single parameter, the critical UV flux, or $J_{\rm crit}$, to determine whether ${\rm H_2}$-–cooling is prevented. This can be misleading, and that independent of the spectral shape, there is a “critical curve” in the ($k_{\rm LW},k_{\rm H^-}$) plane, where $k_{\rm LW}$ and $k_{\rm H^-}$ are the ${\rm H_2}$–-dissocation rates by LW and IR photons, which determines whether a protogalaxy can cool below $\sim 1000$ Kelvin. We use a one–-zone model to follow the chemical and thermal evolution of gravitationally collapsing protogalactic gas, to compute this critical curve, and provide an accurate analytical fit for it. We improve on previous works by considering a variety of more realistic Pop III or Pop II-–type spectra from population synthesis models and perform fully frequency–-dependent calculations of the ${\rm H_2}$-–photodissociation rates for each spectrum. We compute the ratio $k_{\rm LW}/k_{\rm H^-}$ for each spectrum, as well as the minimum stellar mass $M_\ast$, for various IMFs and metallicities, required to prevent cooling in a neighboring halo a distance d away. We provide critical $M_\ast/d^2$ values for suppression of ${\rm H_2}$-–cooling, with analytic fits, which can be used in future studies. | It has long been known that the cooling of metal--free, primordial gas, from which the first generation of stars form in the first protogalaxies, is dominated by ${\rm H_2}$ molecules \citep{Saslaw+Zipoy}. Furthermore, the abundance of ${\rm H_2}$ molecules in an early protogalaxy is sensitive to radiation impinging on the galaxy, and can be regulated by the early ultraviolet (UV) background in the Lyman-Werner bands~\citep{HRL97}. For unusually soft spectra, infrared (IR) radiation can also play a role, through the photo-detachment of the ${\rm H^-}$ ions, the main catalyst for ${\rm H_2}$ formation in primordial gas (e.g. \citealt{O01}). Recent simulations have suggested that the first stars may not be as massive as had previously been thought (e.g. \citealt{Greif+11}). In the extreme case that the typical first-generation stars had masses as low as a few ${\rm M_\odot}$, the IR radiation from these low-mass stars would dominate radiative feedback and ${\rm H_2}$ chemistry in protogalaxies in the early Universe \citep{WGH12}.\footnote{We do not consider X-rays in this paper, which can also be important for early ${\rm H_2}$ chemistry \citep{HAR00,IT15}.} In recent years, the suppression of ${\rm H_2}$ cooling via radiative feedback has attracted a lot of attention, in the context of forming massive black hole seeds in the early universe. Observations of high-redshift quasars reveal that supermassive black holes (SMBHs) with masses of $\sim 10^9~\msun$ have already formed as early as redshift $z\sim 7$ \citep{Mortlock+11}. A promising way to form such early massive SMBHs is to begin with a massive (say, $\sim 10^5\msun$) seed BH at redshift $z\gsim 10$. This massive seed can then grow further by accretion, and reach $10^9~\msun$ by redshift $z\sim7$. In particular, in contrast to starting with a stellar-mass BH, the accretion rate then need not exceed the value implied by the Eddington limit (see recent reviews by \citealt{VBreview12,SMBHreview13,NatarajanReview14} and \citealt{IHO15}). A promising site for forming such massive BH seeds are in the nuclei of so-called atomic-cooling halos -- i.e. dark matter halos with virial temperatures $T_{\rm vir} \gsim 10^4{\rm K}$, or masses of $\gsim {\rm few}\times 10^7 {\rm M_\odot}$. Assuming that the gas at the center of such a halo is metal-free (or with metalicity at most $\sim 10^{-4}$ of the solar value; \citealt{OSH08}), and also that ${\rm H_2}$--cooling is disabled, the gas remains at temperatures near $10^4{\rm K}$. This leads to accretion rates in the core of the halo as high as $\sim 0.1-1 {\rm M_\odot yr^{-1}}$. Several works have argued that under these conditions, fragmentation and Population III star-formation may be avoided, and a massive seed BH is produced instead, either by direct collapse, or via an intermediate state of a supermassive star \citep{OH02,BL03,KBD04,LN06,SS06,BVR06,VLN08,WA08,RH09b,SSG10,SBH10,WGH11,WHB11,Latif+14}.\footnote{Recent simulations \citep{Regan+14} have suggested that fragmentation might occur at spatial resolution higher than currently numerically feasible. On the other hand, even if fragmentation occurs, a supermassive star may still be the natural outcome, as long as ${\rm H_2}$ cooling is disabled, owing to the rapid migration and coalescence of the central fragments \citep{IH14,Hosokawa+15}.} A crucial assumption in these scenarios is the lack of ${\rm H_2}$--cooling. In the presence of ${\rm H_2}$, the gas in the atomic cooling halos would likely fragment and form Population III stars, similar to the case of lower-mass ``minihalos'' \citep{ABN02,BCL02,Yoshida+03}. It has been well-established in both semi--analytic studies and three--dimensional simulations that a sufficiently strong dissociating LW flux can suppress ${\rm H_2}$-- cooling entirely, keeping the gas close to the virial temperature of the halo throughout the initial stages of collapse \citep{O01,BL03,RH09a,SSG10,SBH10}. As demonstrated in these papers, the spectral shape of the incident radiation is important in determining whether this scenario is plausible. In nearly all previous studies, the incident radiation field is modeled either as a power--law, or as a blackbody spectrum with temperature $10^4 \leq {\rm T_*/K} \leq 10^5$, and a critical flux $J_{\rm crit}$ is defined as the minimum intensity required to prevent cooling. Quoted at the Lyman--limit, $J_{\rm crit}$ for a $T_* = 10^5$K blackbody (hereafter referred to as T5) is usually found to be in the range $10^3- 10^4$ in the customary $J_{21}$ units $J{\rm (13.6eV)} = J_{21} \times 10^{-21} {\rm erg s^{-1} Hz^{-1} cm^{-2} cm^{-2} sr^{-1}}$. For the softest blackbody spectra considered, with $T_*=10^4$K (hereafter T4), the nominal critical flux is much lower than for the T5 type, typically $J_{\rm crit}\sim30$ (e.g. \citealt{O01}). In the T4 case, photodetachment of ${\rm H^-}$, a precursor in the primary formation reaction for ${\rm H_2}$, causes suppression of ${\rm H_2}$--cooling, rather than photodissociation. This is due to the large IR flux near the photodetachment threshold (${\rm h \nu \approx 1-2}$eV) in a T4 spectrum compared to the T5 spectrum {\it for a fixed $J_{21}$}. However, as shown by \citet{WGH12}, this can be misleading: the mass in stars must, in fact, be higher for a T4--type population to produce the same $J_{21}$ as a T5--type. Thus, the lower value of $J_{\rm crit}$ for the soft T4 spectrum makes it more difficult, rather than easier, to suppress ${\rm H_2}$--cooling. The main goal of this paper is to investigate the ${\rm H_2}$--cooling in the case of a ``realistic'' Pop II spectrum for the incident radiation field, rather than an idealised spectrum such as a blackbody or power--law. A single flux, $J_{\rm crit}$, is sufficient for determining the cooling history for a blackbody (power--law), because for a given blackbody temperature (exponent), there is a fixed ratio of flux in the LW bands to the flux at the photodetachment threshold ($0.76$eV). However, for more realistic spectra from population synthesis modeling, it is necessary to consider separately the ${\rm H_2}$--photodissociation ($k_{\rm LW}$) and ${\rm H^-}$--photodetachment rates ($k_{\rm H^-}$), and their evolution over time. We show in \S~\ref{Sec:Results} that {\it there is, in general, a ``critical curve,'' rather than a single $J_{\rm crit}$}, for determining whether ${\rm H_2}$--cooling is suppressed and thus whether the halo is a candidate for a DCBH. This critical curve is a line in the $k_{\rm LW}$ vs $k_{\rm H^-}$ plane, and is independent of the spectral shape. We consider a range of population synthesis model spectra from {\sc starburst99} and show how the results compare, with respect to the critical curve, with results from blackbodies ($T_* = 10^4-10^5$K). We also provide an updated criterion for determining whether ${\rm H_2}$--cooling is suppressed, as a fitting formula for the critical curve. The non--existence of a single $J_{\rm crit}$ and its importance for time--evolving spectra from population synthesis models has also been pointed out recently by \citet{AK15} and \citet{Agarwal+16}. These studies, as well as \citet{Sugimura+14} have investigated the relative importance of $k_{\rm LW}$ and $k_{\rm H^-}$ from {\sc starburst99} and from blackbody spectra in suppressing ${\rm H_2}$--cooling. We here improve on these studies by computing the ${\rm H_2}$--photodissociation rate, summing over all relevant LW lines, rather than assuming that the spectrum is flat in the LW bands, or taking an average of the LW flux. We show that both previous approximations, though computationally efficient, introduce significant errors in the photodissociation rate compared to our full frequency--dependent calculation. This paper is organised as follows: In \S~\ref{Sec:Model} we describe the details of our numerical modeling, We present our results using a variety of {\sc starburst99} models in \S~\ref{Sec:Results}, along with an updated fitting formulae for both the critical curve and critical mass in stars. We briefly discuss the results and implications of our work for future studies, and offer our conclusions in \S~\ref{Sec:Conclusions}. | \label{Sec:Conclusions} In most studies of the direct collapse black hole scenario, a single parameter $J_{\rm crit}$ has been used to delineate the threshold flux above which ${\rm H_2}$--cooling is prevented in an atomic cooling--halo. However, this single-parameter assumption can be misleading, as the value of $J_{\rm crit}$ depends on the spectral shape, and therefore needs to be re-computed for each incident spectrum. As long as the spectrum is assumed to have a fixed shape (in time), such as a blackbody or power--law, specifying the flux at a single frequency (e.g. the Lyman--limit) fixes all relevant photo--chemical rates. In a more general treatment, the ratio of the flux in the Lyman--Werner bands to that at the energies relevant for ${\rm H^-}$--photodetachment ($\sim 1-2$eV), is dependent on the details of the spectrum, which depend on the mass function, metallicity, and age of the irradiating galaxy. Although a new $J_{\rm crit}$ could be computed for each case, we here emphasize that this is unnecessary, and advocate the use, instead of a two-dimension critical curve (Figure~\ref{Fig:CriticalCurve}). Using the population synthesis package {\sc starburst99} to generate realistic spectra, we investigated the relative importance of LW photodissociation and ${\rm H^-}$--photodetachment in suppressing the ${\rm H_2}$ abundance and thereby preventing cooling below the virial temperature of a halo. Very few studies have used spectra from population synthesis modeling, and those that have \citep{Sugimura+14,AK15,Agarwal+16} relied on simplifications to derive a LW photodissociation rate, rather than performing a full, frequency--dependent calculation. We show in \S~\ref{Sec:Results} that these simplifications overestimate the true photodissociation rate by a factor a few in the first $\sim 20$ Myr of a burst. We provide a fitting formula to the critical curve from our one--zone modeling (see Fig.~\ref{Fig:FittingFormula} and Eq.~\ref{Eq:FittingFormula}), which includes important updates to the photochemical model and is accurate at the percent level. The resulting critical stellar masses (Table~\ref{tbl:CriticalHaloProperties} and Figure~\ref{Fig:FittingFormula2}) can be used directly in semi-analytic models and simulations of early structure formation. | 16 | 9 | 1609.02142 |
1609 | 1609.01282_arXiv.txt | Only in the Milky Way is it possible to conduct an experiment which uses stellar streams to detect low-mass dark matter subhaloes. In smooth and static host potentials, tidal tails of disrupting satellites appear highly symmetric. However, perturbations from dark subhaloes, as well as from GMCs and the Milky Way bar, can induce density fluctuations that destroy this symmetry. Motivated by the recent release of unprecedentedly deep and wide imaging data around the Pal~5 stellar stream, we develop a new probabilistic, adaptive and non-parametric technique which allows us to bring the cluster's tidal tails into clear focus. Strikingly, we uncover a stream whose density exhibits visible changes on a variety of angular scales. We detect significant bumps and dips, both narrow and broad: two peaks on either side of the progenitor, each only a fraction of a degree across, and two gaps, $\sim2^{\circ}$ and $\sim9^{\circ}$ wide, the latter accompanied by a gargantuan lump of debris. This largest density feature results in a pronounced inter-tail asymmetry which cannot be made consistent with an unperturbed stream according to a suite of simulations we have produced. We conjecture that the sharp peaks around Pal 5 are epicyclic overdensities, while the two dips are consistent with impacts by subhaloes. Assuming an age of 3.4 Gyr for Pal 5, these two gaps would correspond to the characteristic size of gaps created by subhaloes in the mass range of $10^6-10^7 M_\odot$ and $10^7-10^8 M_\odot$ respectively. In addition to dark substructure, we find that the bar of the Milky Way can plausibly produce the asymmetric density seen in Pal 5 and that GMCs could cause the smaller gap. | Palomar 5 is so diffuse that it was once mistaken for a low surface brightness galaxy by \citet{wilson1955} who ``rediscovered" the globular and called it the Serpens Dwarf, the name that appears surprisingly fitting today, after the detection of the conspicuous S-shaped tails attached to the cluster \citep[][]{pal5disc}. Naturally, both the low stellar density of the satellite and the prominence of its associated stellar stream are tell-tale signs of the ongoing disruption by the Galactic tides. Over the years, Pal 5's tidal tails grew and are currently traced across several tens of degrees on the sky \citep[see e.g][]{rockosi2002,odenkirchen2003,gd_pal5}. Thus, Pal 5 has quickly become a poster child for Milky Way accretion. To date, the cluster's role as a possible powerful and precise Galactic accelerometer has been emphasized by many \citep[see e.g.][]{ibata_et_al_pal5}, but it remains modelled by few \citep[e.g.][]{dehnen2004,kuepper_et_al_pal5}. \citet{dehnen2004} who presented the very first - but nonetheless impressively comprehensive - study of the Pal 5 disruption, established with certainty several key aspects of the satellite's accretion: the cluster's orbit, its mass and size, and the importance of disk shocks in driving the mass loss. However, while getting many observables right, such as the shape of the stream track and the overall behavior of the debris density along the tails, no model in the \citet{dehnen2004} suite could match the level of asymmetry between the star counts in the leading and the trailing tail of the cluster, as displayed in e.g. their Figure 16 and Figure 4 of \citet{odenkirchen2003}. It was then concluded that the observed asymmetry ought to be due to the processes not captured by the simulations. The authors point out that the most likely phenomenon - not included in their numerical setup - which could produce such a small-scale density enhancement in one of the tails is the interaction of the stream with a low mass substructure. They offer three examples of such perturbers: giant molecular clouds, spiral arms and dark matter subhaloes. The irrefutable detection of small-scale density perturbations in the Pal 5 tails as presented by \citet{odenkirchen2003} and emphasized by \citet{dehnen2004} called for an explanation. This inspired \citet{capuzzo2005} to revisit the numerical experiments of \citet{combes1999} who had predicted that globular cluster tails ought to contain low-level stellar clumps. In their simulations, \citet{capuzzo2005} not only confirmed the presence of ubiquitous small-scale substructure in tidal tails, but also provided an intuitive justification of their existence: the stars in the clumps move slightly slower compared to the rest of the surrounding debris in the tail. The deceleration of the stars in the clumps was the clue which helped \citet{kuepper_et_al_2008} establish the genesis of the overdensities: the orbits of the stripped stars in the reference frame of the progenitor are oscillatory, composed of the guiding center circular orbit and the epicyclic ellipse. The two vertices of the ellipse are crossed at the lowest speed, thus leading to episodic bunching of stars along the tidal tail. Although the presence of these epicyclic overdensities was initially shown for progenitors on circular orbits, further efforts showed that it is also present on orbits with a wide range of eccentricities \citep[e.g.][]{kuepper_et_al_2010}. Now, could the clump in the Pal 5's trailing tail simply be an example of an epicyclic overdensity as described above? This seems unlikely as the simulations of \citet{dehnen2004} actually produce epicyclic ``feathering'' of the cluster's tails as is obvious from e.g. Figure 6: from panel to panel, the bunching appears either enhanced or reduced depending on the orbital phase of the progenitor considered. Remarkably, however, the epicyclic clumping is almost undetectable in the snapshot corresponding to Pal 5's current location as shown in their Figure 18, thus begging for a conclusion that the apocentre - near which Pal 5 is currently situated - is not the optimal location for the detection of epicyclic overdensities. Most importantly, however, the overdensities produced are always at a very similar level in the leading and the trailing tails, especially so close to the progenitor. This is exactly the point highlighted by \citet{dehnen2004}: whether the epicycles are strong or not, the density profile of the two tails should be symmetric (see their Figure 16). Note also, that such a (approximate) symmetry of the tidal tails' density profiles might simply be a direct consequence of the symmetry of the Hill's surface \citep[see e.g.][]{bt2008}. Thus, given that the epicyclic clumps - a generic feature of the globular cluster tidal streams - cannot seem to explain the sizeable overdensity in the trailing tail of Pal 5, it is prudent to attempt to establish the actual mechanism behind the apparent asymmetry. Our work is motivated not only by the fact that this conundrum has remained unsolved for more than 10 years, but also by the two recent advances in the tidal tail studies. First, the Pal 5 stream has been mapped by \citet{ibata_et_al_pal5} to an unprecedented depth allowing for a robust determination of the minute details of the debris distribution. Armed with this remarkable dataset, we use a powerful new and robust non-parametric algorithm to determine the stream track and the associated density variation. We confirm with very high confidence, both the detection of a non-monotonic star count evolution along the trailing tail and the asymmetry between the leading and trailing portions of the stream. In addition, we also find a gap-like feature in the leading tail and evidence of two epicyclic overdensities near the progenitor. At first glance, these findings appear to be in tension with \citet{ibata_et_al_pal5} who analyzed the same dataset and found that there are no statistically significant gaps in Pal 5. However, their search was performed on scales smaller than $1$ degree while the search for gaps in this work is focused on larger scales, guided by predictions of the expected distribution of gap sizes from subhaloes in \citep{number_of_gaps} which brings us to the second advance in studies of tidal streams. There is now a much better understanding of the impact of the massive perturber fly-bys on the structure of the tidal tails \citep[see e.g.][]{carlberg_2009,yoon_etal_2011,carlberg_2012,carlberg_2013,three_phases,subhalo_properties,sanders_bovy_erkal_2015,bovy_erkal_sanders_2016,number_of_gaps} - the primary mechanism put forward by \citet{dehnen2004} to explain the unruly star counts in the Pal 5 tails. While it had been known that a fly-by leads to a density depletion around the projected impact point \citep[see e.g.][]{ibata_et_al_2002,johnston_et_al_2002}, it has now been established that the induced stream gap is always accompanied by density hikes on either side \citep[see e.g.][]{carlberg_2012, three_phases}. Moreover, stream gaps go hand in hand with stream wobbles: small-scale perturbations visible in {\it all} phase-space projections of the debris track \citep[see][]{three_phases,subhalo_properties,bovy_erkal_sanders_2016}. Finally, notwithstanding the degeneracy between the age of the gap, the mass of the perturber and its speed, there exists a distinct characteristic gap size for subhaloes with different masses as shown in \citet{number_of_gaps}, with smaller subhaloes tearing smaller holes. However, there also exists a lower bound to the size of the gap. This minimum size emerges because lighter subhaloes on average impart a smaller velocity kick onto the stream stars. Therefore, it takes longer for the density in the gap to drop to detectable levels hence widening it to sizes comparable to those of the gaps induced by more massive subhaloes. For example, according to \citet{number_of_gaps}, it is not feasible to expect DM subhaloes with masses of $10^7 M_{\odot}$ to produce deep gaps less than 5$^{\circ}$ wide in a Pal 5-like stream, with most detectable gaps produced by these subhaloes opening to $\sim10^{\circ}$. Interestingly, the flyby of dark subhaloes is not the only conceivable mechanism that can produce small-scale perturbations in the stream. Naturally, exactly the same generic features described above are also expected from the gaps torn by giant molecular clouds \citep{amorisco_et_al_2016_gmcs}. In addition, in \cite{kohei_rotating_bar} it was shown that the rotating bar of the Milky Way can reshape the stream drastically since different portions of the debris approach their pericenter at different times and hence experience a different force from the bar. The influence of the bar was also studied in \cite{price_whelan_et_al_chaotic_fanning} in terms of the chaos it can create. Sending some of the tidal debris on chaotic orbits can dramatically affect the stream appearance, leading to substantial perturbations of the stream track \citep[e.g.][]{pearson_et_al_pal5}, as well as stream fanning \citep{price_whelan_et_al_chaos}. In this work, through a series of numerical experiments involving N-body simulations of the Pal 5-like cluster disruption as well as the approximate stream models based on modified Lagrange Cloud Stripping \citep[mLCS,][]{gibbons_et_al_2014}, we will demonstrate that the observed small-scale disturbances of the Pal 5 tails are consistent with an impact by two low-mass substructures. If the stream features are indeed caused by the passage of dark subhaloes, we argue that the features in the leading and trailing tails are most likely caused by subhaloes in the mass range $10^6-10^7 M_\odot$ and $10^7-10^8 M_\odot$ respectively. Such subhaloes have long been predicted in $\Lambda$CDM and their detection would represent a stunning confirmation of the theory. However, unfortunately, with the data in hand, we cannot distinguish the higher mass flyby from the influence of the Milky Way bar. In addition, we cannot distinguish the lower mass flyby from a GMC flyby. The complications introduced by the bar and GMCs also suggest that searches for gaps which focus on streams at larger radii will yield detections that can be interpreted more straightforwardly and hence may be more fruitful. This Paper is organized as follows. In \Secref{sec:data} we discuss the deep CFHT photometry published by \citet{ibata_et_al_pal5}. In \Secref{sec:stream_model} we present a novel non-parametric model which we use to extract the stream track and to measure the stellar density variation and the evolution of its width. In \Secref{sec:unperturbed_pal5} we give a review of how tidal streams form in static and smooth potentials to highlight the discrepancy with the observed features. Next, in \Secref{sec:asymmetric_mechanisms}, we study the effect of an impact by substructure and that of the Milky Way's rotating bar and show that they can both contribute to the features seen in Pal 5. In \Secref{sec:other} we discuss several other mechanisms such as the internal Pal 5 rotation, chaos, and perturbations by other globular clusters, as well as how these can be distinguished. The results are compared against expectations in \Secref{sec:discussion}. Finally, we conclude in \Secref{sec:conclusion}. | \label{sec:conclusion} In this paper we have analyzed new, high quality photometry of the Palomar~5 stellar stream published recently by \citet{ibata_et_al_pal5}. In order to fully take advantage of this superb data we have developed a novel non-parametric method to extract the stream properties such as the density along the tails, the centroid track of the debris distribution on the sky, as well as the stream width. Our probabilistic method is adaptive in the sense that the model complexity is not fixed {\it a priori}, but rather is driven by the data in hand. The combination of the quality of the data and the power of the modelling technique yields an exquisite determination of the stream properties. For the first time, we measure significant changes in the stream width and show that the debris cross-section varies differently along each tail. We also detect dramatic stream density fluctuations on a variety of angular scales. First, on the scale of a fraction of a degree, density spikes are detected very close to the progenitor. Second, further away from the Pal~5 cluster, at $\phi_1=-3^{\circ}$ in the leading tail, a density depletion approximately 2 degrees across is measured. It is accompanied by two low-level bumps on either side. Finally, the trailing tail exhibits a prominent density enhancement at a distance of $\sim$ 3 degrees from the progenitor, followed by a smooth drop in star counts, observable for some 8 degrees along the tail, and then a mild bump at $\phi_1 \sim 12^\circ$. As we demonstrate with utmost clarity, the remarkable rise and fall of the trailing debris density do not have counterparts in the leading tail. We interpret the small-scale debris pile-ups in the vicinity of the progenitor as epicyclic overdensities, and conjecture that the two larger scale density perturbations are induced by interactions with small substructure. If dark matter subhaloes were the cause of these stream wrinkles, then their masses are of order of $10^6-10^7$ and $10^7-10^8 M_{\odot}$. Impressively, the size of the larger gap discovered here agrees well with the characteristic gap scale expected in the presence of $\Lambda$CDM sub-structure with masses between $10^5$ and $10^9 M_{\odot}$ as predicted in \citet{number_of_gaps}. It is not easy to over-emphasize the importance of this discovery if the sub-structure that wrought havoc in the stream was non-baryonic. In fact, a subhalo in the $10^6-10^7 M_\odot$ range would increase the lower bound on the warm dark matter particle mass to $>9$-$18$ keV \citep{viel_et_al_2013}. Note, however, that currently we cannot rule out other plausible explanations, such as the effect of the rotating bar, impacts by GMCs, or other complexities in the gravitational potential of the Galaxy, that would lead to mild orbital chaos or would induce substantial variations in the stripping rate of the cluster. In order to distinguish between these mechanisms, we must predict the features each of these can create in the Pal~5 tails. For example, as was shown in \cite{subhalo_properties}, the flyby of a subhalo produces an almost unique signature which can be used to infer the subhalo properties. The precise signatures of the bar, chaos and the pericentre wobble are likely very different, so with additional data, especially the improved radial velocity measurements expected from WEAVE, 4MOST and DESI, it should be possible to determine the culprit. | 16 | 9 | 1609.01282 |
1609 | 1609.08689_arXiv.txt | The Gemini Planet Imager is a high-contrast near-infrared instrument specifically designed to image exoplanets and circumstellar disks over a narrow field of view. We use science data and AO telemetry taken during the first 1.5~yr of the GPI Exoplanet Survey to quantify the performance of the AO system. In a typical 60~sec H-band exposure, GPI achieves a $5\sigma$ raw contrast of $10^{-4}$ at 0.4''; typical final $5\sigma$ contrasts for full 1~hr sequences are more than 10 times better than raw contrasts. We find that contrast is limited by bandwidth wavefront error over much of the PSF. Preliminary exploratory factor analysis can explain 60--70\% of the variance in raw contrasts with combinations of seeing and wavefront error metrics. We also examine the effect of higher loop gains on contrast by comparing wavefront error maps reconstructed from AO telemetry to concurrent IFS images. These results point to several ways that GPI performance could be improved in software or hardware. | \label{sec:intro} % Over the past two decades, our understanding of exoplanetary systems has progressed thanks to several complementary observational techniques. Transit and radial velocity surveys have yielded a wealth of new terrestrial planets, and a handful of the most tightly-orbiting ones can be characterized spectroscopically. However, no transit or radial velocity survey has a sufficiently long time baseline to detect, much less characterize, the planets in the outer reaches of extrasolar systems. Furthermore, our understanding of the formation and dynamical histories of these systems can be greatly aided by studying the morphology of the protoplanetary disks from which planets form and the debris disks that remain after the systems' evolution has stabilized. Spatially-resolved spectroscopy and polarimetry are powerful tools for investigating these systems. High-contrast, high-resolution instrumentation is necessary to achieve these goals. To image even young, self-luminous gas giant exoplanets around their bright host stars requires near-infrared flux contrasts of $10^{-4} - 10^{-6}$. For even the nearest host stars, the projected separations of interest are often less than 0.5''. The Gemini Planet Imager (GPI)\cite{Macintosh2008, Macintosh2014d} is a high-resolution, high-contrast near-infrared coronagraphic imager at Gemini South specifically designed to meet these requirements. GPI operates in one of two modes: a broadband polarimeter\cite{Perrin2010a} or an integral field spectrograph (IFS) with a wavelength-dependent spectral resolution of approximately 30--80\cite{Larkin2014}. Both modes have selectable bandpasses from Y to K-band, with pixel scales of $14.166 \pm 0.007$~mas/px and fields of view $2.7''$ square. High-contrast science is enabled by GPI's high-order adaptive optics system\cite{Poyneer2014}. For a thorough description and characterization of the AO system, as well as preliminary performance analysis based on 10~months of on-sky data, the reader may refer to Poyneer et al. (2016)\cite{Poyneer2016}. In these proceedings, we provide an updated description the AO system performance, particularly as related to science images, as well as lessons learned, based on data from the first 1.5~yr of the GPI Exoplanet Survey (GPIES). | The GPIES campaign has observed more than 300 stars in 1.5~yr, and these science and AO data provide a wealth of information about the performance of the AO system. Considerable effort has been invested in automated and semi-automated data reduction pipelines for both science data and AO telemetry, greatly facilitating performance data mining. As a result, we know that GPI AO is consistently achieving $10^{-4}$ $5\sigma$ raw contrast at 0.4'' in single 60~sec H-band images, with a stability that enables 1~hr sequences to reach a final $5\sigma$ contrast 10--20 times better. We find several correlations between IFS images, AO telemetry, and seeing conditions. First, we find that bandwidth errors dominate the final error budget at most locations in the PSF, except for the faintest stars and/or separations beyond $\sim0.8''$. This conclusion is reinforced by the good correspondence between 2D reconstructed bandwidth WFE maps and IFS images. Interestingly, seeing values are not well correlated with either WFE or contrast, although the atmospheric coherence time, $\tau$, is correlated. This reinforces the conclusion that GPI AO is not typically limited by the amplitude of seeing phase errors, but by the time lag relative to the speed of the turbulent layers. In a preliminary study, we used exploratory factor analysis to investigate the underlying factors responsible for raw contrast. We find four main underlying factors that track the lower atmosphere, the upper atmosphere, the AO bandwidth residual speckles, and the AO noise errors. A combination of these factors can explain 60--70\% of the variation in raw contrasts. Future work will refine and expand this analysis. We tested the system with higher gains at 500~Hz loop speeds. Until a reliable centroid gain measurement method is implemented, GPI may consider increasing the gain cap on its gain optimizer to compensate for optical gain variations, particularly in sub-median seeing. In an on-sky test, the contrast improvement in high gain IFS images corresponded well to the expected improvement based on 2D reconstructed bandwidth WFE maps. Additional work is needed to ensure that, when operated with a higher gain cap, the gain optimizer will select stable gains in all conditions. These analyses point to key areas for future potential upgrades. The strong dependence of performance on $\tau$ suggests GPI would benefit from wind predictive control and/or increased loop speed. Some initial studies have been done on wind measurement from GPI telemetry\cite{Srinath2016} and predictive Fourier control is being tested at ShaneAO\cite{Rudy2016}, but neither has been implemented on GPI. Several high contrast systems operate or have planned upgrades to operate at $>1$~kHz\cite{Fusco2016, Jovanovic2016, Males2016, Pinna2016}, to clear advantage with respect to lag errors. Additionally, GPI's WFS detector noise limits the system to guide stars of $I<10$. Upgrading to an EMCCD would both increase sky coverage and allow GPI to run faster on fainter targets. | 16 | 9 | 1609.08689 |
1609 | 1609.06007_arXiv.txt | We have recorded three lunar occultations of Aldebaran ($\alpha$~Tau) at different telescopes and using various band-passes, from the ultraviolet to the far red. The data have been analyzed using both model-dependent and model-independent methods. The derived uniform-disc angular diameter values have been converted to limb-darkened values using model atmosphere relations, and are found in broad agreement among themselves and with previous literature values. The limb-darkened diameter is about 20.3~milliarcseconds on average . However, we have found indications that the photospheric brightness profile of Aldebaran may have not been symmetric, a finding already reported by other authors for this and for similar late-type stars. At the sampling scale of our brightness profile, between one and two milliarcsecond, the uniform and limb-darkened disc models may not be a good description for Aldebaran. The asymmetries appear to differ with wavelength and over the 137 days time span of our measurements. Surface spots appear as a likely explanation for the differences between observations and the models. | Aldebaran ($\alpha$~Tau) is one of the brightest and most distinctive stars in the sky, and as a result it has one of the longest records of observations and publications. Being a K5 giant star and located at just 20\,pc from the Sun, it has also one of the largest angular diameters among all stars and it has therefore been the subject of numerous measurements in this sense using a number of techniques. Since Aldebaran is located on the Zodiac, it is regularly occulted by the Moon. We are at present in the middle of one such series of occultations, which will last until early 2018. \citet[RR05 hereafter]{2005A&A...433..305R} presented accurate lunar occultation (LO) and long-baseline interferometry measurements obtained in the near-infrared, and discussed them in the context of previously available determinations. They concluded that the limb-darkened diameter of $\alpha$~Tau is $20.58\pm0.03$ milliarcseconds (mas), or 44~$R_\odot$. Photometric variability is less than 0.01\,mag and the diameter is assumed to be reliably constant. Differences in the angular diameter values available in the literature are indeed present and sometimes significant when taken at face value, however they can often be justified in terms of uniform disc to limb-darkening corrections or by intrinsic limitations in the accuracy. We have recorded three LO events in the present series, and we present here their detailed analysis. Our aim is not so much to confirm or refine the angular diameter determination, but rather to investigate possible asymmetries or surface structure features in the photosphere of this giant star. Indications in this sense had already been presented by \citetalias{1989A&A...226..366R}. | We have recorded three lunar occultation light curves obtained first at the Russian 6-m telescope in the far red, and then 137 days later at the Devasthal 1.3-m telescope in the red and at the 2.4-m Thai National Telescope in the ultraviolet. The analysis by conventional uniform disc (UD) and limb-darkened disc (LD) models leads to values which are approximately consistent with the expected LD value of $20.58\pm0.03$\,mas derived by \citetalias{1989A&A...226..366R} from the combination of accurate occultation and long-baseline interferometry determinations. However, the measurements do not agree at the level of the formal errors, and a close inspection of the fit residuals showed that the UD (and therefore the LD) approximation may not be accurate for this K5 giant. This is consistent with earlier indications of surface asymmetries for Aldebaran as well as for other late-type giants. Analyis by model-independent methods has revealed that the brightness profile of Aldebaran has significant departures from spherical symmetry, at least at the milliarcsecond level, or few percents of its diameter. These asymmetries would be well consistent with cool spots, and lunar occultations provide the means of detecting such spots directly, if coordinated observations are performed for the same event from several sites. We plan to observe more occultations of Aldebaran in the present series which will last until the end of 2017. | 16 | 9 | 1609.06007 |
1609 | 1609.02138_arXiv.txt | *{This is the centenary year of general relativity, it is therefore natural to reflect on what perspective we have evolved in 100 years. I wish to share here a novel perspective, and the insights and directions that ensue from it.} \abstract{This is the centenary year of general relativity, it is therefore natural to reflect on what perspective we have evolved in 100 years. I wish to share here a novel perspective, and the insights and directions that ensue from it.} | General relativity is in many ways unique and different from all other physical theories. The first and foremost among them is the fact that, unlike all other forces, relativistic gravitational law is not prescribed but instead it is dictated by spacetime geometry itself. It naturally arises from inhomogeneity of spacetime and that is why it is universal -- links to everything that physically exists. Presence of any force makes spacetime inhomogeneous for particles to which the force links but not for others. For instance, presence of electric field makes spacetime inhomogeneous for electrically charged particles while for neutral particles it remains homogeneous. By universal force we mean a force that links to everything that physically exists irrespective of particle parameters like mass, charge and spin. Since relativistic gravity is universal and hence it can only be described by spacetime geometry. Thus unlike Newton, Einstein had no freedom to prescribe a relativistic gravitational law because it is entirely governed by spacetime itself which does not obey anyone's dictate or prescription. Since relativistic gravity encompasses Newtonian gravity, it is remarkable that now Newton's inverse square law simply follows from spacetime geometry without any external prescription. \\ Note that spacetime is a universal entity as it is the same for all and equally shared by all and so is the universal force. Hence the two respond to each-other leaving no room for any external intervention. By simply appealing to inhomogeneity of spacetime curvature, we will derive an equation of motion for universal force which would be nothing other than Einsteinian gravity. It is remarkable that we make no reference to gravity at all yet spacetime curvature yields gravitational equation. This happens because both \textit{spacetime and Einstein gravity are universal} \cite{measure-grav}. A general principle that emerges is that \textit{all universal things respond to each other and they must therefore be related}. \\ The equation of motion that emerges from Riemann curvature is non-linear involving square of first derivative of metric. It indicates that gravity is self interactive. As a matter of fact it is the universal character that demands energy in any form must gravitate. Since gravitational field like any other field has energy, it must hence also gravitate -- self interact. Isn't it wonderful that spacetime curvature automatically incorporates this feature through nonlinearity of Riemann tensor? The important aspect of self interaction is that it gravitates without changing the Newtonian inverse square law. This is rather strange because self interaction would, in the classical framework, have asked for $\nabla^2\Phi = 1/2 \Phi'^2$ which would have disturbed the inverse square law. The situation is exactly as it is for photon (light) to feel gravity without having to change its velocity. Within classical framework it is impossible to accommodate these contradictory demands. \\ The answer could however be in the enlargement of framework in which gravity curves space and photon freely floats on it without having to change its velocity. What should curve space and the obvious answer is gravitational field energy which is not supposed to contribute to acceleration, $\nabla\Phi$. Thus gravity self interacts via space curvature and that also facilitates photon's interaction with gravity \cite{ein-new}. These are the two new aspects of Einstein gravity which wonderfully take care of each-other leaving Newtonian inverse square law intact. Einstein is therefore Newton with space curved \cite{ein-new}. Since space and time are bound together in spacetime through universal light velocity, and hence spacetime must be curved. This is how spacetime curvature enters in description of the universal force -- Einstein gravity. In this way the self interaction gets automatically incorporated in Riemann curvature and is reflected through occurrence of square of first derivative of metric. \\ General relativity (GR) is undoubtedly the most elegant and beautiful theory and it is for nothing that Paul Dirac termed it as the greatest feat of human thought! \\ In what follows we would further explore its elegance and richness of structure and form in relation to what new insights and understanding we have gained in past 100 years, and marvel on new questions and directions that ensue. | 16 | 9 | 1609.02138 |
|
1609 | 1609.02248_arXiv.txt | By performing a global magnetohydrodynamical (MHD) simulation for the Milky Way with an axisymmetric gravitational potential, we propose that spatially dependent amplification of magnetic fields possibly explains the observed noncircular motion of the gas in the Galactic centre (GC) region. The radial distribution of the rotation frequency in the bulge region is not monotonic in general. The amplification of the magnetic field is enhanced in regions with stronger differential rotation, because magnetorotational instability and field-line stretching are more effective. The strength of the amplified magnetic field reaches $\gtrsim$ 0.5 mG, and radial flows of the gas are excited by the inhomogeneous transport of angular momentum through turbulent magnetic field that is amplified in a spatially dependent manner. As a result, the simulated position-velocity diagram exhibits a time-dependent asymmetric parallelogram-shape owing to the intermittency of the magnetic turbulence; the present model provides a viable alternative to the bar-potential-driven model for the parallelogram shape of the central molecular zone. In addition, Parker instability (magnetic buoyancy) creates vertical magnetic structure, which would correspond to observed molecular loops, and frequently excited vertical flows. Furthermore, the time-averaged net gas flow is directed outward, whereas the flows are highly time dependent, which would contribute to the outflow from the bulge. | The magnetic field near the GC is much stronger than in the Galactic disc. \cite[Crocker et al.(2011)]{cro11} gave a stringent lower bound $> 50$ $\mu$G on 400 pc scales from $\gamma$-ray observations. Recently, \cite[Pillai et al. (2015)]{pil15} reported that the field strength of a dark infrared cloud, G0.253+0.016, near the GC is $\sim$ a few mG from polarization observations (see also Pillai et al. 2016 in this volume). These results are consistent with an inferred field strength $\sim$ mG from observed complex structure such as the nonthermal filaments \cite[(Yusef-Zadeh et al. 1984; Tsuboi et al. 1986; Lang et al. 1999; Nishiyama et al. 2010)]{yus84,tsu86,lan99,nis10}. Such strong magnetic fields affect the dynamics of the gas in the GC region. One of the manifestations of the magnetic activity is molecular loops \cite[(Fukui et al. 2006; Fujishita et al. 2009; Torii et al. 2010a,b; Riquelme et al. 2010; Kudo et al. 2011)]{fuk06,fuj09,tor10a,tor10b,riq10,kud11}, which are supposed to % rise upwards as a result of magnetic buoyancy \cite[(Parker instability; Parker 1966)]{pak66}. Japanese groups \cite[(e.g., Shibata \& Matsumoto 1991; Nishikori et al. 2006; Machida et al.2009; 2013)]{sm91,mac09,mac13} have investigated magnetic activity in the Galaxy with MHD simulations. Following these works, we \cite[(Suzuki et al. 2015; S15 hereafter)]{suz15} performed a global MHD simulation, particularly focusing on the magnetic activity in the Galactic bulge region. In this proceedings paper, we introduce the results of our global simulation based on S15. | \begin{figure}[h] \begin{center} \includegraphics[width=3.in]{./galaxy28_B_1Dtave1.eps} \caption{Radial distribution of the time-averaged magnetic field strength. Solid, dashed, and dotted lines correspond to the radial, azimuthal, and vertical components, respectively. } \label{fig:B} \end{center} \end{figure} Figure \ref{fig:B} presents the radial profile of each component of the time-averaged magnetic field strength after the amplification of the magnetic field is saturated. In the almost entire region, $>40$ pc, the toroidal component ($B_{\phi}$) dominates the poloidal component ($B_R$ and $B_z$) because the differential rotation efficiently amplifies the magnetic field by the stretching motion. However, the poloidal component is not so weak and the total field strength, $\sqrt{B_R^2 + B_{\phi}^2 + B_z^2}$, reaches $\sim 0.5$ mG inside $\lesssim 0.5$ kpc. \begin{figure}[h] \vspace*{-0.3 cm} \begin{center} \hspace{-2.cm} \includegraphics[width=3.3in]{./glbdsk_velfld28inv_t27-10pi_1correction.eps} \hspace{-1.5cm}\includegraphics[width=3.3in]{./glbdsk_posvel28inv_t27-10pi_2correction.eps} \vspace*{-0.9 cm} \caption{{\it left:} Face-on views of density in units of $n$ cm$^{-3}$ (colour) and velocity field (arrows) at the Galactic plane. {\it Right:} Simulated $l-v$ diagram. Colours indicate column density integrated along line of sight. } \label{fig:faceon} \end{center} \vspace*{-0.2cm} \end{figure} The left panel of Fig. \ref{fig:faceon} shows density (colours) and velocity (arrows) fields at the Galactic plane \begin{lrbox}\myVerb \scriptsize{\verb|http://ea.c.u-tokyo.ac.jp/astro/Members/stakeru/research/galaxypot/faceon28inv_4.gif| } \end{lrbox} \footnote{Movie is available at \\ \usebox\myVerb}. The velocity field shows radial flows, in addition to the background rotating component. These radial flows are excited by the inhomogeneous transport of the angular momentum. The angular velocity of the rotation is not smoothly distributed with radius because the contribution from the radial pressure gradient force is spatially dependent. MHD turbulence, which transports the angular momentum outward, is developed more effectively in regions with stronger differential rotation. Therefore, angular momentum is transported more efficiently for stronger differential rotation. As a result, the radial force balance of gravity, centrifugal force, and pressure gradient force breaks down, and the fluid element moves radially inward or outward. % In addition to this mechanism, the inhomogeneous gradient of magnetic pressure also drives radial flows. Radial (noncircular) flows are also observed as a characteristic feature in a position -- velocity diagram. The right panel of Fig. \ref{fig:faceon} shows column density in a plane of Galactic longitude, $l$, vs. line-of-sight velocity, $v$. One can see a ``parallelogram'' shape in the central part of the diagram, which can be directly compared to the observed ''asymmetric parallelogram'' \cite[(e.g., Liszt \& Burton 1980; Bally et al. 1987; Takeuchi et al. 2010)]{lb80,bal87,tak10}. The shape of the simulated parallelogram changes with time \begin{lrbox}\myVerb \scriptsize{\verb|http://ea.c.u-tokyo.ac.jp/astro/Members/stakeru/research/galaxypot/pv-anim28inv_hr_2.gif| } \end{lrbox} \footnote{Movie is available at \\ \usebox\myVerb}, depending on regions in which fast radial flows are excited. In our simulation vertical flows are also frequently excited by magnetic activity, and such flows contribute to the outflow from the Galactic bulge region (see Kakiuchi et al. 2016 in this volume for the detail). This work was supported by Grants-in-Aid for Scientific Research from the MEXT of Japan, 24224005 (PI: YF). Numerical simulations in this work were carried out at the Cray XC30 (ATERUI) operated in CfCA, National Astrophysical Observatory of Japan, and the Yukawa Institute Computer Facility, SR16000. \nocite{*} | 16 | 9 | 1609.02248 |
1609 | 1609.07976_arXiv.txt | Microscopic liquid brines, especially calcium-perchlorate could emerge by deliquescence on Mars during night time hours. Using climate model computations and orbital humidity observations, the ideal periods and their annual plus daily characteristics at various past, current and future landing sites were compared. Such results provide context for future analysis and targeting the related observations by the next missions for Mars. Based on the analysis, at most (but not all) past missions' landing sites, microscopic brine could emerge during night time for different durations. Analysing the conditions at ExoMars rover's primary landing site at Oxia Planum, the best annual period was found to be between $L_s$ 115 - 225, and in $Local\hspace{0.1cm} Time$ 2 - 5, after midnight. In an ideal case, 4 hours of continuous liquid phase can emerge there. Local conditions might cause values to differ from those estimated by the model. Thermal inertia could especially make such differences (low TI values favour fast cooling and $\textrm{H}_2\textrm{O}$ cold trapping at loose surfaces) and the concentration of calcium-perchlorate salt in the regolith also influences the process (it might occur preferentially at long-term exposed surfaces without recent loose dust coverage). These factors should be taken into account while targeting future liquid water observations on Mars. | The most probable form of liquid water on Mars today and during most of the Amazonian period are the microscopic liquid layers formed by deliquescence or by the interfacial forces at the ice-mineral interface. Confirmation of their existence and the characterization of their consequence could be understood by in-situ analysis. Although remote data provide information on the environmental conditions for the emergence of these liquids, and Earth based laboratory measurements provide information on the micro-physical processes, it is difficult to extrapolate the real microscopic events inside the Martian regolith under the given exact composition, particle size etc. Thus, in-situ analysis could provide the most relevant information here. As near future surface missions aim for low latitude sites, the microscopic liquid between ice and minerals cannot be analysed, while there is a better chance for the observation of deliquescence driven liquefaction. In this work we exploit the potential emergence and observational possibilities of deliquescence driven microscopic liquid water formation at future ExoMars candidate landing sites, in order to give predictions on the targeting of observations there. As the main contribution of this work to the field, we present numerical estimations of a range of parameters regarding which seasonal phase, and local time provide the ideal combinations for the emergence of liquid and possibility for these observations at the candidate landing sites. For this we present the properties of of the climate modelling we used, correlate the results with the threshold values for microscopic liquid formation of Mars-relevant salty solutions, analyse the possibility of the emergence of these liquids at past and future landing sites, and finally predict ideal conditions for the identification of these liquids by the ExoMars rover. \\ Our work and this article could help optimize the laboratory measurements on Earth - what kind of tests and simulations should be run in the future, where and in what season landers should be sent out to Mars for the highest chance of finding and analysing water in its liquid form. Our results could also prove to be useful in the planning and optimization of measurements at the ExoMars rover's landing site, including the earlier candidate sites. \\ \subsection{\textbf{Background information}} This chapter provides an overview of the reasons and background conditions for the possible emergence of liquid water on Mars and the approach we used to predict their occurrence in this study. There is an ongoing debate regarding the potential existence of liquid water on the surface of Mars. Certain theoretical models predict that liquid water may be present \citep{clow1987} \citep{haberle2001} \citep{hecht2002}, while some surface features favour the ephemeral existence of a liquid phase \citep{brass1980} \citep{Mellon2001} \citep{knauth2002} \citep{motazedian2003} \citep{kossacki2004} \citep{kereszturi2009} \citep{szynkiewicz2009} \citep{mcewen2014}. Computations favour the emergence of microscopic liquid water, especially brines \citep{mohlmann2004} \citep{kossacki2008} \citep{martinez2013}. The presence of brines is suggested by the observations at the Phoenix landing site and other locations \citep{kossacki2004} \citep{chevrier2009} \citep{hecht2009} \citep{hecht2009det} \citep{renno2009}, which may act as a possible agent for the formation of some recent flow-like features. \\ Recently, using the meteorological observations from the Curiosity rover, it was demonstrated that the night time conditions are favourable for the emergence of microscopic, thin, liquid film on the surface of hygroscopic mineral grains on the Martian surface and in the shallow subsurface \citep{torres2015}. Taking into account the above mentioned possibilities, it is worth calculating the possible emergence of such microscopic liquids on Mars in order to better orient instrument development, and also target in-situ analysis there. If such microscopic liquids emerge, it happens under restricted physical and chemical conditions, at limited temporal and spatial occurrence, thus suggesting parameters for the targeted analysis is highly important.\\ Under the current usually cold and dry Martian conditions, the \textbf{deliquescence process} could provide microscopic liquid water without ice. Hygroscopic salts could adsorb water vapour from the atmosphere where deliquescence (\textcolor{black}{transition from solid to aqueous phase}), or efflorescence (\textcolor{black}{transition from aqueous to solid phase}) happen \citep{gough2011}. The threshold limits for these processes are the \textcolor{black}{eutonic} relative humidity (\textcolor{black}{$\textrm{RH}_{\textrm{eut}}$}, where deliquescence starts), the deliquescence relative humidity (DRH, where deliquescence is complete) or efflorescence relative humidity (where the thin liquid layer is lost and the material turns to a solid phase). Deliquescence phase transition occurs when the local relative humidity (RH) is equal to or higher than the DRH. \\ We worked with the \textcolor{black}{eutonic} temperature and the water activity of a solution at the \textcolor{black}{eutonic} T values, which are presented in Table \ref{tab:salts}. Data regarding calcium-perchlorate ($\textrm{Ca(ClO}_4)_2$) were acquired from the work of \citep{toner2014}, the magnesium-perchlorate ($\textrm{Mg(ClO}_4)_2$) from \citep{mohlmann2011} and the values for calcium-chloride ($\textrm{CaCl}_2$) from \citep{davila2010}. The characteristics of the salts and brines with the most probable occurrance among the candidate ones are described below. Magnesium-perchlorate ($\textrm{Mg(ClO}_4)_2$), identified by the Phoenix \textcolor{black}{lander}, is one of the most analysed salts \citep{hecht2009det}. The small spheroids observed on the lander's robotic arm appeared to merge during the mission, which some scientists argue could be a sign that the spherules were liquid water \citep{renno2009}. These could have been the liquid brine form of the salt mentioned above, but unfortunately detailed observations are not available. | The Correlation of model- and observation-based annual and daily temperature plus humidity values points to ideal periods for the possible emergence of microscopic liquid brine on Mars. These ideal periods are found at night time, when in cold and vapour oversaturated conditions hygroscopic minerals could produce deliquescence and thus a thin liquid layer on the salts' surfaces. In most cases, because of the slower night time cooling and faster early morning warming, the longest periods of liquid occurrence is before the temperature minimum during the night. The ideal brine for this process among the Mars-relevant candidate ones is $\textrm{Ca(ClO}_4)_2$. Analysis of past landing sites for possible brine emergence allowed insight into the general observational possibility of such liquid - although more focused laboratory tests are required. At most sites (Opportunity, Curiosity, Phoenix, Viking 2) deliquescence could occur in theory (using modelled meteorological values and assuming the existence of hygroscopic salts there). At the landing site of Viking 1 and candidate landing site Hypanis Vallis minimal, at Spirit's landing site little to no deliquescence is expected. \\ Comparing different landing sites, the longest daily duration of liquid was around 6 hours and 15 minutes (at Phoenix lander's landing site), and the longest seasonal phase during which it could appear is 0 - 360 degree in $L_s$ at Curiosity rover's landing site. These findings are consistent with some earlier observations on the possible emergence of microscopic liquid water in the form of candidate droplets at the Phoenix probe \citep{renno2009}, and the ideal conditions identified at Curiosity's night time observations \citep{torres2015}. Analyzing the observational possibility at the planned primary landing site of ESA's ExoMars rover mission (Oxia Planum), for night time brine emergence $L_s$ 115 - 225 and Local Time 1 - 5 am are the ideal values. Around the middle of this seasonal phase, deliquescence is expected to happen for 4 hours continuously. It is also worth making observations beside this ideal range, as those measurements could shed light on the behaviour and possible emergence of microscopic liquid on Mars under different climatic conditions. \\ Calculating for the two backup landing sites of the ExoMars rover (ideal for a possible launch later than 2018) different seasonal periods are ideal for night time deliquescence (see the numerical details in Table \ref{tab:periods}). The longest seasonal period, while deliquescence is expected at night time is probably present at Aram Dorsum, where deliquescence is expected throughout most of the year, while at Oxia Planum and Mawrth Vallis a total duration of three times less is expected between southern spring and summer. The longest total continuous duration of deliquescence (4.7, 3.6 hours) is expected at the Oxia and Mawrth sites, while a much shorter duration (up to 2.2 hours) is expected at Aram Dorsum. \\ Of course, local conditions might differ substantially from the above mentioned general findings, thus these predictions are mere suggestions. Evaluating the possible local conditions, strong slope winds (that could enhance evaporation) are not expected at the Oxia Planum region. For the condensation, low temperature surfaces are favourable by catching the atmospheric $\textrm{H}_2\textrm{O}$. These are mainly low \textbf{thermal inertia} regions, i.e. loose dust covered surfaces, and it is also expected that such poorly consolidated, loose dust or sand covered areas did not spend much time on the Martian surface in a static state, and thus they need not spend a long time there for the possible concentration of hygroscopic salts to accumulate by an evaporation-sublimation process. Searching for the ideal site for deliquescence on Mars, especially at the ExoMars landing site, these factors should be considered. \\ \begin{table}[h!] \centering \small \setlength\tabcolsep{2pt} \begin{tabular}{|c|c|c|} \hline Characteristics & Value & Landing site \\ \hline Longest ideal $L_s$ & & Aeolis Palus \\ period overall & 0-360 [$L_s$] & (Curiosity) \\ \hline Longest ideal continuous & & Green Valley \\ time period overall & 6.26 \textit{[martian hours]} & (Phoenix) \\ \hline & 115-225 [$L_s$] \textit{(total: 110)} & Oxia Planum \\ Longest ideal $L_s$ period & 105-360, 0-75 [$L_s$] \textit{(total: 330)} & Aram Dorsum \\ & 90-195 [$L_s$] \textit{(total: 105)} & Mawrth Vallis \\ \hline Longest ideal continuous & 4.76 \textit{[martian hours]} & Oxia Planum \\ time period & 2.24 \textit{[martian hours]} & Aram Dorsum \\ & 3.66 \textit{[martian hours]} & Mawrth Vallis \\ \hline \end{tabular} \caption{\label{tab:periods}In this table you can see the longest overall time and $L_s$ periods ideal for deliquescence to occur among all former landing sites, and the longest time and $L_s$ ideal periods for the ExoMars primary (Oxia) and two backup (Aram, Mawrth) landing sites.} \end{table} It is an important question how near of far our model prediction is from reality on Mars. The global climate model used was not able to handle regional or local differences and specific effects, but at the same time, this is one of the best available methods to estimate current Martian meteorological conditions, and there is a good chance that it shows real trends, and useful suggestions to target future measurements. | 16 | 9 | 1609.07976 |
1609 | 1609.00551_arXiv.txt | {Since the orbital insertion of the Rosetta spacecraft, comet 67P/Churyumov-Gerasimenko (67P/C-G) has been mapped by OSIRIS camera and VIRTIS spectro-imager, producing a huge quantity of images and spectra of the comet's nucleus.} {The aim of this work is to search for the presence of H$_2$O on the nucleus which, in general, appears very dark and rich in dehydrated organic material. After selecting images of the bright spots which could be good candidates to search for H$_2$O ice, taken at high resolution by OSIRIS, we check for spectral cubes of the selected coordinates to identify these spots observed by VIRTIS.} { The selected OSIRIS images were processed with the OSIRIS standard pipeline and corrected for the illumination conditions for each pixel using the Lommel-Seeliger disk law. The spots with higher I/F were selected and then analysed spectrophotometrically and compared with the surrounding area. We selected 13 spots as good targets to be analysed by VIRTIS to search for the 2 $\mu$m absorption band of water ice in the VIRTIS spectral cubes. } {Out of the 13 selected bright spots, eight of them present positive H$_2$O ice detection on the VIRTIS data. A spectral analysis was performed and the approximate temperature of each spot was computed. The H$_2$O ice content was confirmed by modeling the spectra with mixing (areal and intimate) of H$_2$O ice and dark terrain, using Hapke's radiative transfer modeling. We also present a detailed analysis of the detected spots. } {} | Space exploration has triggered major progress in our understanding of comets beginning in March 1986 with the exploration of comet 1P/Halley by an armada of missions including the ESA Giotto mission (Reinhard and Battrick, 1986). With the arrival of the ESA Rosetta mission at the comet 67P/Churyumov-Gerasimenko (67P/C-G) on July 2014, comets appear more complex and fascinating than ever. All the visited comets show a low visible albedo and heterogeneous surface (Barucci et al. 2011). However 67P/C-G and some other periodic comets reveal the presence of intriguing bright spots on the surface (Sunshine et al. 2006; Sunshine et al. 2012; and Li et al. 2013; Pommerol et al. 2015). To better understand the properties and composition of the comet 67P/C-G is one of the major objectives of the ESA Rosetta mission, and all on-board instruments have so far contributed with high quality and a precious quantity of data. Since the orbital insertion of the Rosetta spacecraft, the comet nucleus has been mapped by both OSIRIS (Optical, Spectroscopic, and Infrared Remote Imaging System), and VIRTIS (Visible InfraRed Thermal Imaging Spectrometer) acquiring a huge quantity of surface images in different wavelength bands and spectra, and producing the most detailed maps at the highest spatial resolution of a cometary nucleus surface. The OSIRIS imaging system (Keller et al. 2007) is composed of the Narrow Angle Camera (NAC) designed to study the nucleus with 11 large band filters at different wavelengths from the ultraviolet (269 nm) to the near-infrared (989 nm), while the Wide Angle Camera (WAC) is devoted to the study of gaseous species in the coma with a set of 14 narrow band filters ranging from the ultraviolet to visible wavelengths. The OSIRIS imaging system was the first instrument capable of mapping a comet surface at high resolution, reaching a maximum resolution of 11cm/px during the closest fly-by that occurred on February 14, 2015, at a distance of $\sim$ 6 km from the nucleus surface. VIRTIS (Coradini et al. 2007) is composed of two channels: VIRTIS-M, a spectro-imager operating both in the visible (0.25$-$1.0 $\mu$m) and infrared (1.0$-$5.0 $\mu$m) ranges at low spectral resolution ($\lambda/\delta\lambda$=70--380), devoted to surface composition, and VIRTIS-H, a single-aperture infrared spectrometer (1.9$-$5.0 $\mu$m) with higher spectral resolution capabilities ($\lambda/\delta\lambda$=1300-3000) devoted to the investigation of activity. \begin{figure*}[t] \includegraphics[width=1\textwidth]{Barucci_Fig1_new.pdf} \caption{Map of comet 67P/Churyumov-Gerasimenko, resulting from merging a more detailed shape model SHAP4S (Preusker et al. 2015) for the northern hemisphere and shape model SHAP5 (Jorda et al. 2016) for the southern hemisphere. In red the selected bright spots are reported, based on OSIRIS images and a spectro-photometric analysis, considered as good targets to be investigated by an analysis of VIRTIS data, plus the two bright spots analysed by Filacchione et al. (2016a). The numbers (1-8) represent the spots with positive detection of H$_2$O ice by VIRTIS analysis discussed in this paper. } \label{spots} \end{figure*} The OSIRIS images of 67P/C-G show a highly shaped, irregular bilobed comet, with a dark, dehydrated, and morphologically complex surface characterized by several terrain types, including numerous diverse geomorphologic features (Sierks et al. 2015). The comet's surface is highly heterogeneous with different geological terrains showing smooth, dust-covered areas, large scale depressions, brittle materials with many pits and circular structures, and exposed consolidated areas (Thomas et al. 2015a; El-Maarry et al. 2015; El-Maarry et al. 2016). Pits have also been connected to activity, possibly accompanied by outbursts (Vincent et al. 2015). The complex surface of comet 67P/C-G shows regions covered by different layers of dust on both lobes, including areas with evidence of transport and redistribution of dust materials (Thomas et al. 2015b). Temporal variations of morphological structures have also been observed on the smooth terrains of the Imhotep region (Groussin et al. 2015), as well as in other regions (Fornasier et al. 2016), in particular when the comet was close to perihelion. The comet shows albedo variation of up to about 25\% and spectrophotometric analysis (Fornasier et al. 2015) identified at least three groups of terrains with different spectral slopes (computed in the 535-882 nm range). These differences have been associated with the local composition variation, but since many different surface characteristics overlap, this makes the interpretation difficult. Oklay et al. (2016a) also studied surface variegation on the comet, detecting local color inhomogeneities connected to active and inactive surface regions. \begin{table*} \begin{center} \caption{Observing conditions for the OSIRIS images as reported in Fig. 2, where the ice spots have been identified. The time (UT) refers to the start time of the first image of each sequence, followed by the number of filters available. The diameter size (d) of the spots along with the location region, phase angle ($\alpha$), distance between Rosetta spacecraft, and comet surface ($\Delta$), spatial resolution (R), latitude (Lat), and longitude (Long) are reported.} \small{ \label{observing} \begin{tabular}{l l l c c c c c c c} \hline N. & Time reference & Filters & d & Region & $\alpha$ & $\Delta$ & R & Lat & Long \\ & & & (m) & &($^{\circ}$) & (Km) & (m/px) & ($^{\circ}$) & ($^{\circ}$) \\ \hline 1 & 2015-06-27T13h26 & F22, F23, F41, F24, F71, F27, F51, & 36 & Imhotep & 89.50 & 191.94 & 3.6 &-5.8 & 189.4 \\ & & F61, F28, F15& & & & & & & \\ 2 & 2015-06-27T17h48 & F22, F23, F41, F24, F71, F27, F51, &45 & Anhur & 89.39 & 188.43 & 3.5 & -41.7 & 63.7 \\ &&F61, F28, F15 & & & & & & & \\ 3 & 2015-04-12T21h42 & F22, F23, F41, F24, F71, F27, F51 &11 & Khonsu & 80.46 & 147.98 & 2.7 & -23.8 & 198.3 \\ &&F61, F28,F16, F15 & & & & & & & \\ 4 & 2014-11-22T04h57& F22, F23, F24, F27, F28, F51, F61 &10 & Atum & 92.70 & 29.50 & 0.54 & -20.7 & 227.4 \\ 5 & 2014-11-22T06h32& F22, F23, F24, F27, F28, F51, F61 &6.5 & Imhotep & 92.78 & 29.50 & 0.54 & -22.0 & 182.8 \\ 6 & 2014-09-19T09h19 & F22, F16, F23, F24, F41 & 2-5 (each) & Khepry & 70.48 & 26.50 & 0.49 & 4.2 & 71.7 \\ 7 & 2014-09-05T05h21 & F22, F23, F27, F16, F28, F41, F71 & 3-5 (each) & Imhotep & 57.23 & 41.44 & 0.77 & -8.1 &188.3 \\ 8 & 2014-09-05T08h00& F22, F23, F27, F16, F28, F41, F71& 6 & Imhotep & 58.43 & 40.76 & 0.75 & -2.4 &174.8\\ \hline \hline \end{tabular} } \end{center} \end{table*} The first results by VIRTIS (Capaccioni et al. 2015) about the spectral analysis showed the presence of a broad absorption feature around 2.9--3.6 $\mu$m present across the entire observed region and compatible with carbon-bearing compounds (opaque minerals associated with organic macromolecular materials) with no evidence of ice-rich patches. Later on, De Sanctis et al. (2015) detected the first evidence for the presence of H$_2$O ice as part of a diurnal cycle on the neck of the comet, while Filacchione et al. (2016a) identified H$_2$O ice on two gravitational debris falls in the Imhotep region exposed on the walls of elevated structures. The latter was interpreted as being possibly extended layering in which the outer dehydrated crust is superimposed over water ice-enriched layers. During the first mapping phase of 67P/C-G nucleus, completed in August-November 2014 (heliocentric distances between 3.6 and 2.7 AU), VIRTIS-M achieved a complete mapping of the illuminated regions in the equatorial and northern hemisphere, which enabled us to retrieve the first compositional maps by using VIS and IR spectral parameters (Filacchione et al. 2016b). During the same period, coma observations performed by VIRTIS-M (Migliorini et al. 2016) and VIRTIS-H (Bockel\'ee-Morvan et al. 2015) channels have traced the H$_2$O vapor emission, which occurs preferentially above the illuminated regions of the northern hemisphere. As limited evidence of exposed H$_2$O ice regions has so far been collected, the aim of this work is to investigate in depth the composition of the 67P/ C-G surface, combining the high spatial resolution images of OSIRIS and the high spectral resolution of VIRTIS for detecting and emphasizing interesting ice spectral signatures. Over 100 meter-sized spots were identified by Pommerol et al (2015), possibly associated with the presence of H$_2$O, on the basis of laboratory experiments, but with no confirmation of the real presence of ice. Deshapriya et al. (in preparation) are collecting a catalogue of large bright spots that are present on the surface of the comet by analyzing the OSIRIS images and spectrophotometry data. For this work we selected the largest spots, good candidates to search for H$_2$O ice that could be detected at the lower angular resolution of VIRTIS. We identify large features with high albedo and low spectro-photometric slope with OSIRIS, we compute accurate coordinates, and we analyze them on the basis of the VIRTIS spectra. Over the large number of spots identified by OSIRIS, 13 of them were checked by VIRTIS. Eight of them show clear evidence of H$_2$O ice in their spectra. In this paper we report on the analysis of the eight bright spots for which we obtained a positive detection of H$_2$O ice in VIRTIS data. In Section 2, the OSIRIS data and the performed analysis are presented and, in Section 3, the VIRTIS data, while in Section 4 the spectral modeling of the selected spots is described. In Section 5, a detailed analysis of the spots and surrounding area is reported, while in Section 6, a possible evolution of the area is discussed. The main aim of this work is to confirm the unambiguous presence of H$_2$O ice by spectral analysis. \begin{figure*} \centering \includegraphics[width=1\textwidth]{Barucci_Fig2a.pdf} \clearpage \end{figure*} \begin{figure*} \centering \includegraphics[width=1\textwidth]{Barucci_Fig2_bis.pdf} \caption{NAC OSIRIS images (first column) for the eight spots reported in Table 1 with a zoom on the spot (second column). The images have been taken with F22 filter (at 649.2 nm). The arrows indicates the spots that have been analysed using boxes of 3x3 pixels. The measured $I/F(\alpha)$ of the bright spots are reported in red, while the surrounding area is reported in black (third column). The relative reflectance (normalized to F23 at 535 nm) of the indicated bright spot in red and the surrounding area in black are represented in the fourth column. } \label{images} \end{figure*} | Comet 67P/Churyumov-Gerasimenko shows a surface rich in heterogeneous geological structures and surface morphological variations that show color and albedo variations across the surface. The high-resolution images obtained by OSIRIS enable us to identify a large quantity of bright spots of different size and located in areas with different properties and high albedo. In this paper, we present for the first time a complementary study of data acquired by the OSIRIS and VIRTIS instruments. A major objective of this paper is to firmly detect the presence of H$_2$O ice on the comet's surface. We confirm the presence of H$_2$O ice on eight new spots and we model the spectra with H$_2$O ice and dark material. Comparing the coordinates of the detected eight H$_2$O ice spots with those of 67P/C-G dust jets, five spots (4, 5, 6, 7 and 8) have been found to lie in the same approximate position of the jets identified by Vincent et al. (2016a) and one (spot 3) among the outbursts observed in the cometary summer (Vincent et al., 2016b). Observational evidence showed that the majority of dust jets also arose from rough terrains and fractured walls rather than smooth areas (Vincent et al., 2016a). Some of these detected H$_2$O ice spots have also been compared by Oklay et al. (2016b) to those of comets 9P and 103P. The detection of H$_2$O ice signatures by VIRTIS on eight of the 13 locations given by OSIRIS data does not mean that the other spots do not contain ice on their surface and this can be explained by not simultaneous observations, unfavorable instrumental signal-to-noise conditions, spatial resolution on the surface, different illumination/viewing geometry, and by the fact that VIRTIS-M channel was unavailable after 4 May, 2015 owing to the failure of the active cooler. ~ The main results of this work can be summarized as follows: \begin{itemize} \item We presented for the first time a complementary analysis of H$_2$O ice-rich areas using data acquired by the OSIRIS and VIRTIS instruments. Comparing high spatial resolution VIS images with extended IR range spectra enables us to study the morphological, thermal and compositional properties of these areas at the same time. \item The analysis of the spectral properties observed by VIRTIS-M indicates that, on these areas, the H$_2$O ice abundance is between 0.1 and 7.2$\%$, mixed in areal or/and in intimate modalities with the dark terrain. \item The ice is distributed on the two lobes of 67P/C-G in locations which remain in shadow for longer. \item The detected bright spots are mostly on consolidated dust free material surfaces, mostly concentrated in equatorial latitudes. \item {The mass release of H$_2$O at the location of the eight ice-rich spots has been estimated.} \item Some spots are stable for several months and others show temporal changes connected to diurnal and seasonal variations. Stability of the spots is corroborated by the temperature retrieved at the surface. The behavior of ice on these locations is in very good agreement with theoretical expectations. \item Six of the detected H$_2$O ice spots are located in approximately the same position of the previously detected cometary jets. \end{itemize} H$_2$O ice is present on the surface substratum where solar illumination plays an important role with seasonal and diurnal variations. During the perihelion orbit passage of the comet, the Rosetta spacecraft was at a greater distance and the available surface OSIRIS images were at lower resolution. Starting in March 2016, the comet is observed again from close distances. With analysis of other available data (in particular from OSIRIS), we will study the surface changes after the perihelion passage to better understand the surface evolution of the comet.\\ | 16 | 9 | 1609.00551 |
1609 | 1609.05973_arXiv.txt | This automated catalogue combines all the largest published optical, radio and X-ray sky catalogues to find probable radio/X-ray associations to optical objects, plus double radio lobes, using uniform processing against all input data. The total count is 1\,002\,855 optical objects so presented. Each object is displayed with J2000 astrometry, optical and radio/X-ray identifiers, red and blue photometry, and calculated probabilities and optical field solutions of the associations. This is the third and final edition of this method. | Astronomy is a large science with many areas of specialty, so few researchers publish across many of the subdisciplines at once. Perhaps this limitation delayed publication of joined optical, radio and X-ray sky data in the early 2000's when {\it ROSAT} X-ray and NVSS \& FIRST radio source catalogs were already available. In 2004 Martin Hardcastle and I endeavoured to fill the gap by publishing the Quasars.org catalogue (QORG: Flesch \& Hardcastle \shortcite{QORG}); this used a uniform data-driven algorithm to align and overlay the optical, radio and X-ray input data while still accomodating the peculiarities of each input catalogue. This was updated in 2010 by the Atlas of Radio/X-ray Associations (ARXA: Flesch \shortcite{ARXA}) with \textit{Chandra} and \textit{XMM-Newton} X-ray data and larger FIRST \cite{FIRST} radio data. Today the dynamic databases like NED\footnote{NASA/IPAC Extragalactic Database, http://ned.ipac.caltech.edu} and SIMBAD\footnote{SIMBAD database at CDS, http://simbad.u-strasbg.fr/simbad} collect all those data and present them with their originally published astrometry. While this resource is essential, it omits the astrometric refinement made possible by optical field solutions (where X-ray fields are aligned to their optical background) and leaves the determination of causal association (i.e., that a particular optical object is the true source of nearby radio/X-ray detections) to the user's devices. Thus the approach of QORG and ARXA is still needed as an aid to astronomers' quick selection of objects of interest. Since ARXA's publication in 2010 there have been continued releases of \textit{XMM-Newton} and \textit{Chandra} catalogued data, and concluding releases from the FIRST project. Also the first catalogue of \textit{Swift} X-ray sources has been published, and multiple identification releases from the prolific SDSS\footnote{Sloan Digital Sky Survey, http://sdss.org} project, and individual publications. Equally important for ARXA-style processing is the quality of the optical background data, and since 2010 I have added the USNO-B \cite{MONET} optical data, plus much SDSS optical data, to make an optical background of over a billion objects with very few one-colour objects compared with before. Combining all these data using the QORG/ARXA processing now results in over a million optical objects with radio/X-ray association of 40\%-100\% likelihood of being true. Thus, this seems the right time to publish this third edition as the ``Million Optical - Radio/X-ray Associations Catalogue'' (hereafter: MORX), which is the final edition of this method. This catalogue is available in both flat file and FITS formats\footnote{at http://quasars.org/morx.htm}, with a ReadMe. Table 1 shows a few sample lines of the flat file with some explanation of the columns, but the ReadMe gives full details of the layout and contents. Figure 1 shows the sky coverage of MORX. \begin{figure*} \includegraphics[scale=0.5, angle=0]{MORX-sky-coverage.jpg} \\ \tiny{Chart produced with TOPCAT \cite{TAYLOR}.} \caption{Sky coverage of the MORX catalogue, darker is denser. Large dense places represent FIRST coverage. The density boundary at declination -40$^{\circ}$ shows the southern edge of NVSS coverage.} \end{figure*} \begin{table*} \caption{Sample lines from the MORX catalogue (left half placed on top of right half)} \includegraphics[scale=0.325, angle=0]{MORX-sample-lines.jpg} \tiny Notes on columns (see ReadMe for full descriptions): \begin{itemize} \item TYPE: R=core radio, X=X-ray, 2=double radio lobes, Q=QSO, A=AGN, q=photometric quasar, G=galaxy, L=LINER, S=star. \item OPT COMM: comment on photometry: p=POSS-I magnitudes, so blue is POSS-I O, j=SERC Bj, +=optically variable, m=nominal proper motion. \item R/B PSFs: '-'=stellar, 1=fuzzy, n=no PSF available, x=not seen in this colour. \item RADIO/X-RAY CONF: calculated percentage confidence that this radio/X-ray source is truly associated to this optical object. \item QSO/GAL/STR PCT: based on its photometry and the association(s), the calculated percentage confidence that this optical object is a QSO/galaxy/star. \item ERR PCT: calculated percentage chance that this association is wrong, equals 100\% minus the combined radio/X-ray confidence. \item RAD/XR OFFSETS: in arcsec, the astrometric offset from the optical object to the best (i.e., highest confidence associated) radio/X-ray source, after that source has been shifted by its optical field solution. \item ``BEST'' CATALOG AND FIELD / SEQUENCE: identifies the radio/X-ray field used in the optical field solution. \item RA MOVE: in arcsec, the East-West shifts of the optical field solution pertaining to this source, calculated in integer RA arcsecs, i.e., 1/3600$^{th}$ of a RA degree. \item DEC MOVE: in arcsec, the North-South shifts of the optical field solution pertaining to this source, calculated in integer arcseconds. \item \# OBJ: the number of sources which were associated to optical objects (with $>$70\% confidence) by this optical field solution. \item TOT SRC: the total number of sources in this radio/X-ray field. \item REF \& ZREF: citations for name and redshift; citations are indexed in the file "MORX-references.txt". \end{itemize} The full table can be downloaded from http://quasars.org/morx.htm, also available in FITS. \end{table*} The QORG paper presents the full details of the matching techniques in its appendix A, but in the sections below I summarize the method and issues involved in constructing this catalogue, including any updated techniques, and give a brief roundup of all the optical, radio and X-ray source catalogues. Thus this paper gives a comprehensive overview suitable for most users of the catalogue. | This catalogue compiles and optically aligns {\it Chandra}, {\it XMM-Newton}, {\it ROSAT} and {\it Swift} X-ray sources, and NVSS, FIRST and SUMSS radio sources, and presents those calculated to be associated to optical objects. Identifications are included to present an informative map to support pointed investigations. Its total count is 1\,002\,855 optical objects with 40\%-100\% likelihood of radio/X-ray associations including double radio lobes. | 16 | 9 | 1609.05973 |
1609 | 1609.07147_arXiv.txt | We use 12 cosmological $N$-body simulations of Local Group systems (the {\sc Apostle} models) to inspect the relation between the virial mass of the main haloes ($M_{\rm vir,1}$ and $M_{\rm vir,2}$), the mass derived from the relative motion of the halo pair ($M_{\rm tim}$), and that inferred from the local Hubble flow ($M_{\rm lhf}$). We show that within the Spherical Collapse Model (SCM), the correspondence between the three mass estimates is exact, i.e. $M_{\rm lhf}=M_{\rm tim}=M_{\rm vir,1}+M_{\rm vir,2}$. However, comparison with {\sc Apostle} simulations reveals that, contrary to what the SCM states, a relatively large fraction of the mass that perturbs the local Hubble flow and drives the relative trajectory of the main galaxies is not contained within $R_{\rm vir}$, and that the amount of ``extra-virial'' mass tends to increase in galaxies with a slow accretion rate. In contrast, modelling the peculiar velocities around the Local Group returns an unbiased constraint on the virial mass ratio of the main galaxy pair. Adopting the outer halo profile found in $N$-body simulations, which scales as $\rho\sim R^{-4}$ at $R\gtrsim R_{\rm vir}$, indicates that the galaxy masses perturbing the local Hubble flow roughly correspond to the asymptotically-convergent (total) masses of the individual haloes. We show that estimates of $M_{\rm vir}$ based on the dynamics of tracers at $R\gg R_{\rm vir}$ require a priori information on the internal matter distribution and the growth rate of the main galaxies, both of which are typically difficult to quantify. | \label{sec:intro} In an expanding, flat Universe the rather intuitive concept of {\it galaxy mass} is ill-defined and difficult to infer observationally. Given that the mass of a galaxy is thought to be a key parameter for the early collapse and subsequent virialization of cosmological substructures (e.g. Rees \& Ostriker 1977; White \& Rees 1978; Blumenthal et al. 1984) as well as for galaxy formation and evolutionary processes (see e.g. Kravtsov et al. 2004; Conroy \& Wechsler 2009; Moster et al. 2010; Behroozi et al. 2010; Gao et al. 2011; Laporte et al. 2013; Sawala et al. 2016; Rodr{\'{\i}}guez-Puebla et al.2016), it becomes crucial to quantify the correspondence between the different methods that have been proposed in the literature to infer galaxy masses from the dynamics of visible tracers. The main difficulty in measuring the mass of a galaxy resides in the limited kinematic information available to an observer in the outskirts of these systems. Indeed, in our current cosmological paradigm galaxies form in the inner-most regions of virialized dark matter structures called {\it haloes} (e.g. Mo, van den Bosch \& White 2010). Within those regions different sorts of kinematic tracers, such as gas, stars and stellar clusters, can be used to constrain the inner mass distribution of a galaxy. However, the amount of baryonic tracers declines progressively at large distances from the halo centre, introducing severe uncertainties in our understanding of the mass distribution in the outer-most regions of galactic haloes, to the point that it becomes extremely challenging to identify the physical {\it edge} of galaxies with the surrounding Universe. The very notion of galaxy edge is thwarted by numerical simulations of structure formation, which show that dark matter haloes exhibit smooth density profiles without well-defined boundaries. These simulations reveal that over-densities in a close-to-homogeneous background induce perturbations in the expansion of the Universe (the so-called Hubble flow) which in principle extend to infinite distances from the source. The gravitational pull slows down the Hubble flow % around over-dense regions, such that at the {\it turn-around distance} ($R_{\rm ta}$) the (local) expansion of Universe halts. Within this volume we find another scale of interest, that where the number of particles moving away from an over-density becomes roughly equal to that moving towards it. By definition, this region is in {\it dynamical equilibrium} and the spherical radius that contains it is called the {\it virial radius}, $R_{\rm vir}$, which provides a natural boundary between a galaxy halo and its environment. A useful, albeit rough, estimate of the size of $R_{\rm vir}$ corresponds to the radius where the free-fall time ($t_{\rm ff}$) is equal to half of the age of the Universe ($t_0$). Within a sphere of mass $M$ and mean density $\langle \rho\rangle =3 M/(4\pi R_{\rm vir}^3)$ the free fall time is $t_{\rm ff}=[3\pi/(32 G\langle \rho\rangle)]^{1/2}=[\pi^2R_{\rm vir}^3/(8 G M)]^{1/2}$. The condition $2t_{\rm ff}=t_0$ thus implies that $R_{\rm vir}=(2GM t_0^2/\pi^2)^{1/3}$. Note that the size of the virialized region expands as $R_{\rm vir}\sim t_0^{2/3}$ at a fixed mass $M$, thus leading to a halo growth that is not driven by accretion of new material or a change in the halo potential, but solely due to a boundary definition based on a timing argument. This effect is usually known as the {\it pseudo-evolution of dark matter haloes} (e.g. Diemand et al. 2007; Cuesta et al. 2008; Diemer et al. 2013; Zemp 2014). Note also that the monotonic expansion of $R_{\rm vir}$ with time suggests that the natural evolution of the cosmic web is towards a set of increasingly isolated haloes in dynamical equilibrium (Busha et al. 2005). Although in principle the virial radius of a galaxy is an observable quantity (i.e. the spherical radius encompassing the volume where the in- and outwards motion of the galaxy constituents balance out), in practice a direct measurement is rarely possible owing to the scarcity of kinematic tracers in the outskirts of dark matter haloes. As a remedy, a widely-used approach to estimate the virial radii of galaxies and galaxy clusters consists of measuring the mass enclosed within the volume populated by kinematic tracers, $M(< r_\star$), which is the {\it extrapolated} to a virial radius, $R_{\rm vir}\gtrsim r_\star$. The extrapolation is typically made upon the adoption of a theoretically-motivated density profile, the most common one being the Navarro, Frenk \& White (1996; 1997) profile, henceforth NFW. However, several studies show that the NFW profile does not describe well the overall shape of dark matter haloes found in cosmological $N$-body simulations (e.g. Prada et al. 2006; Betancort-Rijo et al. 2006; Cuesta et al. 2008). Deviations from the NFW profile at large radii appear to largely correlate with the rate at which haloes accrete mass (e.g. Ludlow et al. 2013; Diemer \& Kravtsov 2014) and may have an impact on the masses of galaxy clusters derived from X-ray observations (e.g. Avestruz et al. 2014) as well as weak-lensing mass measurements (e.g. Oguri \& Hamana 2011; Becker \& Kravtsov 2011). In addition, Taylor et al. (2015) warns about the dangers of adopting halo profiles found in dark matter-only cosmological models to describe galaxies that have been acted on by baryonic feedback. Using hydrodynamical simulations of Milky Way-like galaxies these authors find that extrapolating $M(<50\kpc)$ measurements out to the virial radius systematically overestimates the halo mass and underestimates the halo concentration. For local galaxies like the Milky Way and Andromeda extrapolations of the NFW profile yield virial masses that are uncertain by a factor $\sim 2$--$3$. Tests against mock data sets reveal that the uncertainty is largely dominated by (i) the unknown impact of baryons on the halo density profile and (ii) the unknown orbital distribution of the tracer particles (e.g. Wang et al. 2015). For example, Smith et al. (2007) use high-velocity stars from the RAVE survey (Steinmetz et al. 2006; Zwitter et al. 2008) to measure the local escape speed of our Galaxy. The range of values found by these authors ($498<v_{\rm esc}/\kms<608$) leads to a NFW halo with virial mass $M_{\rm vir}=0.85_{-0.29}^{+0.55}\times 10^{12}M_\odot$. Xue et al. (2008) extend the mass constraints to $D\lesssim 60\kpc$ by modelling the kinematics of blue horizontal branch stars (BHBs), which yields a similar value $0.8_{-0.2}^{+0.2}\times 10^{12}M_\odot$. Sakamoto et al. (2003) and Battaglia et al. (2005) find $2.5_{-1.0}^{+0.5}\times 10^{12}M_\odot$ and $0.8_{-0.2}^{+1.2}\times 10^{12}M_\odot$, respectively, from a mixed sample of globular cluster, giant stars and satellite galaxies. Watkins et al. (2010) applies the Jeans equations to 26 satellite galaxies of the Milky Way and finds a virial mass $0.7$-$3.4\times 10^{12}M_\odot$ depending on the (unknown) orbital distribution of the satellite population. Barber et al. (2014) deal with this uncertainty by directly matching the distribution of subhaloes found in dark matter-only $N$-body simulations against the observed position and velocity of Milky Way dwarf spheroidals. This comparison yields a virial mass in the range $0.6<M_{\rm vir}/(10^{12} M_\odot)<3.1$. A number of recent studies avoid profile extrapolation by modelling the dynamics of tracers that lie {\it beyond} the nominal virial radius of the halo. For example, by demanding that the total momentum of the Local Group should balance to zero Diaz et al. (2014) estimate the individual masses of the Milky Way and Andromeda to be $M_{\rm MW} = (0.8 \pm 0.5) \times 10^{12} M_\odot$ and $M_{\rm M31} = (1.7 \pm 0.3) \times 10^{12} M_\odot$, respectively. Also, accreted material reaching its first apocentre after halo collapse gives rise to a density caustic (Fillmore \& Goldreich 1984; Bertschinger 1985) which can be used to define a halo boundary that is in principle observable (Rines et al. 2013, Patej \& Loeb 2015; More et al. 2016). The outermost caustic manifests itself as a sharp density drop in the halo outskirts at a location known as `back-splash' radius ($R_{\rm sp}$). Unfortunately, the relation between the caustic location and the virial size is very sensitive to the recent mass evolution of the halo, which adds significant uncertainty to the relation between the mass enclosed within the back-splash radius, $M(<R_{\rm sp})$ and $M_{\rm vir}$. Cosmological collisionless simulations show that the back-splash radius varies between $R_{\rm sp}/R_{\rm vir}\sim 0.8$--$1.6$ in haloes with fast or slow accretion rates, respectively (see More et al. 2015 for details). This paper follows up on the work of Pe\~narrubia et al. (2014, 2016; hereinafter P14 and P16) who construct idealized models of structure growth to describe the (local) cosmic expansion around an isolated, overdense region of the Universe. These models are incorporated into a Bayesian framework to fit orbits to measured distances and velocities of galaxies between 0.8 and 3 Mpc from the Local Group (hereinafter LG). The method returns masses for the Milky Way and M31 which do not rely on {\it a priori} premises on the internal matter distribution in those galaxies nor on equilibrium assumptions (see also Banik \& Zhao 2016 for a similar study). However, the relation between the mass perturbing the local Hubble flow and the individual virial masses of the main haloes remains unclear to date. In this contribution we use self-consistent cosmological simulations of twelve LG-like systems (the so-called {\sc Apostle}\footnote{A Project Of Simulating The Local Environement} project, Sawala et al. 2016; Fattahi et al. 2016) to calibrate the correspondence between the mass derived from modelling the local Hubble flow and the individual virial masses of the Milky Way \& M31 analogues. The arrangement of this paper is as follows. In Section~\ref{sec:mass} we introduce the Spherical Collapse Model, which provides a simple, analytical representation of the non-linear growth of structures in a close-to-homogeneous Universe. Section~\ref{sec:num} outlines the construction of mock data sets using cosmological $N$-body models of the Local Group as well as the method we use to fit peculiar velocities around these systems. Section~\ref{sec:res} inspects the relation between the virial masses of the main Local Group haloes and those inferred from the local Hubble flow. We discuss the validity of the Spherical Collapse Model assumptions in Section~\ref{sec:discussion} and attempt to estimate the virial mass of the Milky Way from the dynamics of nearby ($\lesssim 3\mpc$) galaxies. Our findings are summarized in Section~\ref{sec:sum}. | \label{sec:discussion} \subsection{Model assumptions}\label{sec:dis} Although the SCM described in Section~\ref{sec:mass} bears little resemblance to the mass distribution found in the {\sc Apostle} simulations, it provides an accurate representation of the peculiar velocities of nearby tracer galaxies (see Fig.~\ref{fig:rv_12}). Such a remarkable --and somewhat unexpected-- agreement is worth investigating in some depth. To this aim we discuss below the validity of some of the key assumptions on which the SCM rests within the Local Group context. Our conclusion is that observations of the local Hubble flow at $R\gtrsim d$, where $d$ is the current separation between the main galaxies, are scarcely sensitive to the detailed mass distribution of the perturber, $M(<d)$, or to its mass assembly history, $M(t)$. \subsubsection{Galaxy pair}\label{sec:pair} Notice first that while Equation~(\ref{eq:kep}) accounts for perturbations in the Hubble flow induced by a {\it single} spherical over-density with radius $R_{\rm vir}$ and mass $M_{\rm vir}$, the Local Group is dominated by {\it two} similar-size galaxies, the Milky Way and Andromeda. Lynden-Bell (1981) argues that the SCM can be still used to describe the dynamics of nearby galaxies upon the adoption of coordinates centred at the barycentre of the Local Group, such that \begin{eqnarray}\label{eq:angmom} 0&=&M_\g{\bf R}_{\g}+ M_\a{\bf R}_{\a}\\ \nonumber 0&=&M_\g{\bf V}_{\g}+ M_\a{\bf V}_{\a}. \end{eqnarray} In this frame galaxies at distances $R\gg d\equiv |{\bf R}_{\rm A}-{\bf R}_{\rm G}|$ feel a potential that can be approximated as \begin{eqnarray}\label{eq:potexp} \Phi(R,\theta)\approx -\frac{G M }{R} + \frac{1}{2}H^2\Omega_{\Lambda}R^2 +\frac{GM f_{\rm m}}{2(1+f_{\rm m})^2}\frac{(1-3 \cos^2\theta)d^2}{R^3}; \end{eqnarray} where $\cos\theta\equiv \hat{\bf R}\cdot \hat{\bf d}$ and $f_m\equiv M_\g/M_\a$. The Keplerian perturbation~(\ref{eq:kep}) is thus recovered by neglecting the right-hand side of Equation~(\ref{eq:potexp}), which corresponds to a quadrupole that decays as $\sim (d/R)^3$. Tests against controlled $N$-body experiments carried by P14 show that, although the potential quadrupole affects the trajectories of galaxies at $R\lesssim d$, its impact on the local Hubble flow is negligible at $R\gtrsim d$. Table~\ref{tab:nbody} shows that the main galaxies in the {\sc Apostle} models have $d\lesssim 1\mpc$ without exception, while the mock catalogues generated in \S\ref{sec:fits} only include galaxies between $1$--$3\mpc$. By design, our lower distance cut removes most of the sensitivity of Equation~(\ref{eq:vd}) to the potential quadrupole generated by the main galaxy pair, as indicated by P14 tests and confirmed in \S\ref{sec:res}. \subsubsection{Accretion and time-dependence of the potential}\label{sec:time} The SCM also neglects the effects of mass accretion, effectively assuming that the mass $M$ is constant throughout the expansion and subsequent collapse of the initial over-density. This is at odds with observations of the outskirts of the Milky Way and M31, which reveal clear evidence of past merger events (e.g. see Belokurov et al. 2006; McConnachie et al. 2009). P14 use dynamical invariants (see Pe\~narrubia 2013 for details) to derive a first-order correction to the perturbed Hubble flow~(\ref{eq:vd}) around time-evolving haloes. For systems with a slow growth rate, $M(t)\approx M_0[1+\epsilon (t-t_0)/t_0]$, the perturbed Hubble flow can be written as \begin{equation}\label{eq:vdt} V(R)\approx \bigg[(1.2 + 0.16 \Omega_\Lambda)\frac{R}{t_0} - 1.1\bigg(\frac{GM_0}{R}\bigg)^{1/2}\bigg]\bigg[1+\epsilon t_0\bigg(\frac{GM_0}{R^3}\bigg)^{1/2}\bigg], \end{equation} where $\epsilon\equiv (\d M/\d t)(t_0/M_0)$ is the dimension-less growth rate and $M_0=M(t_0)$. If the mass evolution is slow one can related this quantity to the mass accretion rate~(\ref{eq:Gamma}) as $\epsilon\simeq \Gamma \Delta\log_{10}a \ln (10) t_0/\Delta t\approx 1.1 ~\Gamma$ for $\epsilon\ll 1$, where $\Delta t$ and $\Delta\log_{10}a$ are calculated between redshifts $z=0.5$ and $z=0$. Note that the first-order correction term on the right-hand side of~(\ref{eq:vdt}) declines at large distances from the galaxy as $\Delta V\sim R^{-3/2}$. For systems that accrete mass recently ($\epsilon >0$) the mass evolution steepens the isochrone $V(R)$ at small distances ($R\lesssim d$) from the matter source, and leaves the perturbed Hubble flow approximately invariant at $R\gtrsim d$. Hence, a time-dependent potential and the Local Group quadrupole influence the dynamics of galaxy tracers within a similar distance range, as discussed in \S\ref{sec:pair}. The mock data catalogues constructed in Section~\ref{sec:apostle} minimize the impact of both, the hierarchical mass accretion of the main haloes and the potential quadrupole by selecting tracer haloes at distances $R\gtrsim d$ from the Local Group barycentre. Under this particular choice the assumption $M=M_0={\rm const.}$ in Equation~(\ref{eq:kep}) does not bias mass mass inferred from the observed Hubble flow. \subsubsection{Large Scale Structure}\label{sec:lss} At distances $R\gg d$ the Hubble flow around the Local Group exhibits pertubations induced by nearby groups of galaxies (e.g. Mohayaee \& Tully 2005; Karachentsev et al. 2009; Courtois et al. 2012) which are not accounted for by Equation~(\ref{eq:kep}). A simple remedy for suppressing the impact of external systems on the virial mass estimates is to exclude galaxies in the vicinity of major perturbers. For example, P14 impose a distance cut to the data set at $R_{\rm max}=3 \mpc$, which roughly corresponds to the distance to the closest associations -- Centaurus A, M81 and IC 342 -- and only fit galaxies whose distance to any of the three major associations is larger than $1 \mpc$. Such hard cuts are motivated by the decreasing accuracy of Equation~(\ref{eq:kep}) at distances where the contribution of the Local Group to the local gravitational acceleration becomes negligible. Here we fit the dynamics of tracer substructures within $3\mpc$ from the main haloes regardless of the distribution of massive haloes beyond this volume. Recall that by design the {\sc Apostle} groups do not contain massive neighbours within $\sim 2.5\mpc$ (see \S\ref{sec:apostle}). Furthermore, the SCM also neglects the overall motion of the Local Group towards peak-density regions in the local volume, such as the Big Attractor (Lynden-Bell et al. 1988), or the vast `wall' of structures connected to the Virgo Cluster (Tully \& Fisher 1987). Numerical simulations of Local Group-like systems show that Large Scale Structures may alter the trajectories of nearby galaxies and induce anisotropic peculiar velocities in the local Hubble flow (Libeskind et al. 2011; Ben{\'{\i}}tez-Llambay et al. 2013). However, recent attempts to detect kinematic anisotropies in the Hubble flow observed within $1$--$3\mpc$ have not yielded any statistically-meaningful deviation from isotropic models (see P16), which suggests that the effects of Large Scale Structure on Equation~(\ref{eq:vd}) may be subdominant with respect to those of the Local Group. In essence, the results of P16 indicate that the galaxies within a $3\mpc$ volume around the Local Group exhibit a coherent motion with respect to the surrounding cosmic web. Tests with {\sc Apostle} simulations support this conclusion and indicate that large scale structures do not unduly affect the dynamical constraints on the Local Group mass derived from the motion of nearby ($R\lesssim 3 \mpc$) galaxies, which greatly simplifies the theoretical framework of \S\ref{sec:mass}. \subsubsection{Massive substructures}\label{sec:masssub} The presence of massive satellites in our Galaxy and Andromeda adds uncertainty into the determination of $({\bf R}_\g$, ${\bf V}_\g)$, and $({\bf R}_\a,{\bf V}_\a)$, respectively, which in turn propagates to the location and the motion of the Local Group barycentre (e.g. G\'omez et al. 2015). P16 incorporates the mass contribution of the largest satellites of the Local Group (the LMC and M33) by fitting $M$ and the ratios $f_m=M_\g/M_\a$, $f_c=M_\lmc/M_{\rm MW}$ and $f_{M33}=M_{\rm M33}/M_{\rm M31}$ simultaneously, where \begin{eqnarray}\label{eq:angmom2} M_\g{\bf V}_{\g}&=&M_\mw{\bf v}_{\mw}+ M_\lmc{\bf v}_{\lmc}\\ \nonumber M_\a{\bf V}_{\a}&=&M_{\rm M31}{\bf v}_{\rm M31}+ M_{\rm M33}{\bf v}_{\rm M33}, \end{eqnarray} and similarly for ${\bf R}_\g$ and ${\bf R}_\a$. However, the physical interpretation of the satellite masses derived from Equation~(\ref{eq:angmom2}) is complicated by the fact that satellites orbiting a larger system tend to lose their dark matter envelopes to tides after a few pericentric passages (e.g. Pe\~narrubia et al. 2008; 2009). Tidal stripping thus calls for a more sophisticated dynamical modelling of the host-satellite interaction beyond the point-mass approximation on which Equation~(\ref{eq:angmom2}) rests. In this contribution we account for the presence of substructures within the main galaxy haloes by computing the combined barycentre of the main halo and its five most massive subhalos at $z=0$ (see \S\ref{sec:apostle}). In addition, we construct mocks that neglect the effects of substructures, finding similar bounds on the model parameters. Indeed, inspection of the {\sc Apostle} models shows that the most massive subhaloes tend to be located in the outskirts of galactic haloes, where their relative velocity with respect to the parent halo, and thus their contribution to Equation~(\ref{eq:angmom2}), is small. A more detailed statistical analysis of cosmological $N$-body models by Deason et al. (2014) shows that the presence of satellites as large as the LMC within the inner-most regions of Milky Way-like haloes is typically rare, in agreement with our findings. \subsection{Model uncertainties \& sample size}\label{sec:uncer} Although the current census of galaxies in the vicinity of the Local Group ($0.8\lesssim R/\mpc \lesssim 3$) is still relatively small ($N_g\approx 35$ according to P14), on-going observational efforts may soon uncover a large population of faint galaxies predicted by $\Lambda$CDM models (e.g. using HI surveys, Tollerud et al. 2015). In this Section we explore to what degree the sample size affects the constraints on our model parameters. For illustration purposes we use the {\sc Apostle} halo with the largest number of identified dark matter structures (AP11-330892, see Table~\ref{tab:nbody} and Fig.~\ref{fig:apostles}) to generate mock catalogues with a varying sample size of subhaloes within $1$--$3\mpc$ from the barycentre of the AP-07 system. Fig.~\ref{fig:number} shows the dependence of the posterior distributions of the parameters $M_{\rm lhf}$ (upper panel), $f_{m}$ (middle panel) and $\sigma_m$ (lower panel) as a function of the number of galaxies in the sample, $N_g$. Error bars indicate 68\% confidence intervals around the median value of the posterior distribution. As expected, the uncertainties of the model parameters tend to decrease for increasing values of $N_g$. All samples produce consistent bounds on $M_{\rm lhf}$ and $f_{m}$ given the computed uncertainties. Interestingly, having large samples is specially important for measuring the mass ratio of the main galaxies $f_m$ with accuracy, whereas the effective mass $M_{\rm lhf}$ appears less sensitive to the value of $N_g$. The parameter $\sigma_m$, however, does not behave well. In particular, the bottom panel of Fig.~\ref{fig:number} shows a sudden increase in the value of $\sigma_m$ at $N_g\gtrsim 100$ which cannot be accounted by statistical uncertainties of the fits. The reason behind the discontinuous dependence of $\sigma_m$ with $N_g$ can be traced back to the presence of long tails in the distribution of peculiar velocities, which only reveal themselves when the number of galaxies in the sample is large enough. As discussed in \S\ref{sec:fits}, the uncertainties of $M_{\rm lhf}$ and $f_{m}$ are directly related with the hyperparameter $\sigma_m$, which itself provides a proxy for the scatter of the Hubble flow within the distance interval were the tracers are located. As a result, we find that the size of the error bars derived from galaxy samples with $N_g\lesssim 100$ members tends to be underestimated. We have checked that AP-07 is the only {\sc Apostle} model in which $\sigma_m$, and therefore the quoted model uncertainties, have a discontinuous dependence on $N_g$. Yet, given the relatively small number of known galaxies within $3\mpc$ of the Local Group, this result calls for caution when interpreting P14 and P16 statistical bounds on the masses of Local Group galaxies. In particular, the value of $\sigma_m\simeq 50\kms$ recently measured using $N_g\simeq 35$ tracer galaxies within $3.0\mpc$ from the Local Group (P14, Banik \& Zhao 2016) must be taken as a lower limit. \begin{figure} \includegraphics[width=82mm]{number.pdf} \caption{Median and 68\% confidence intervals of the posterior distributions of the parameters $M_{\rm lhf}$ (upper panel), $f_{m}$ (middle panel) and $\sigma_m$ (lower panel) of the model {\sc Apostle} halo AP11-330892 (Table~\ref{tab:nbody}) as a function of the number of halo tracers in the mock catalogues. } \label{fig:number} \end{figure} \subsection{The virial mass of the Milky Way} We now turn to the thorny task of estimating the virial mass of a galaxy, $M_{\rm vir}$, from the effective mass perturbing the local Hubble flow, $M_{\rm lhf}$. Note first that while the parameter $M_{\rm lhf}$ can be measured directly from the heliocentric distances and velocities of nearby galaxies, a derivation of the virial mass of the Milky Way ( $M_{\rm MW,vir}$) requires a priori information on the matter distribution of the dark matter halo. Neglecting the baryonic mass and adopting DK14 profile for simplicity reduces the number of free parameters to three: $M_{\rm MW,vir}$, $c_{\rm vir}$ and $\Gamma$; for one single constraint, i.e. the Milky Way mass inferred from the Hubble flow, $M_{\rm MW,lhf}$. Using the statistical correlation between virial mass and concentration $c_{\rm vir}=c_{\rm vir}(M_{\rm vir},z$) at $z=0$ found in cosmological $N$-body simulations (Mu\~noz-Cuartas et al. 2011) removes one further degree of freedom and reduces the problem to estimating the growth rate of our Galaxy, $\Gamma$. As discussed in the Introduction, the virial mass of field haloes grows via (i) the hierarchical accretion of substructures and, (ii) a monotonic decrease of the density threshold with respect to which virial over-densities are defined, which leads to the so-called {\it pseudo-evolution} of dark matter haloes. For haloes with a relatively small growth rate Equation~(\ref{eq:Gamma}) can be linearized as $$\Gamma\approx \Gamma_{\rm acc}+\Gamma_{\rm pseudo}.$$ Diemer et al. (2013) and Zemp (2014) find that in Milky Way-sized haloes pseudo-evolution accounts for $\sim 30\%$ of the virial mass growth from $z=0.5$ to the present, which yields $\Gamma_{\rm pseudo}\sim 0.6$. On the other hand, in our Galaxy $\Gamma_{\rm acc}$ may be dominated by the (total) mass of the brightest satellite galaxy, the Large Magellanic Cloud (LMC), which to date remains strongly debated in the literature. Let us inspect two leading scenarios, one where the LMC contribution to the Galaxy mass is negligible, and a second one where the total LMC mass is introduced as a free parameter in the Hubble flow model. Using the SCM to fit the dynamics of nearby ($0.8\lesssim R/\mpc\lesssim 3$) galaxies under the assumption $M_{\rm LMC}\approx 0$ yields $M_{\rm MW,lhf}=0.8^{+0.4}_{-0.3}\times 10^{12}M_\odot$ (P14). If we adopt the results of Fig.~\ref{fig:mass_appr} at face value and assume $\Gamma= \Gamma_{\rm pseudo}\simeq 0.6$ we find a Milky Way virial mass that may be as low as $M_{\rm MW,vir}\sim M_{\rm MW,lhf}/1.25\sim 0.64\times 10^{12}M_\odot$. In contrast, P16 find that the LMC may be more massive than previously thought (see also Besla et al. 2012 and Belokurov \& Koposov 2016). Implementing the proper motions of the LMC measured recently from HST data in Equation~(\ref{eq:angmom2}) and fitting $M_{\rm LMC}$ and $M_{\rm MW, lfh}$ simultaneously to the same data set as in P14 returns a similar effective Milky Way mass, $M_{\rm MW,lhf}=1.04^{+0.42}_{-0.38}\times 10^{12}M_\odot$, while strongly favouring a massive LMC, $M_{\rm LMC}=0.25^{+0.09}_{-0.08}\times 10^{12}M_\odot$. These values yield $\Gamma_{\rm acc}\simeq 0.53$, and a total value of $\Gamma\sim 1.17$. Hence, combination of P16 constraints with Fig.~\ref{fig:mass_appr} suggests a virial mass $M_{\rm MW,vir}\sim M_{\rm MW,lhf}/1.20\sim 0.87\times 10^{12}M_\odot$, which is in excellent agreement with some of the values discussed in the introduction (e.g. Battaglia et al. 2005; Smith et al. 2007; Xue et al. 2008), but somewhat in tension with the high-mass bounds derived from the motion of Leo I (Sakamoto et al. 2003; Boylan-Kolchin et al. 2013). The above estimates emphasize the practical challenges that must be faced when trying to infer the virial mass of a galaxy from the effective mass measured from the local Hubble flow. In general, our analysis indicates that extending the dynamical range of observational data sets to objects located beyond the virial radius of the Milky Way does {\it not} lead to more stringent constraints on $M_{\rm MW,vir}$. In fact, estimates of $M_{\rm MW,vir}$ based on the measured value of $M_{\rm MW,lhf}$ require additional information on the mass growth of our Galaxy, which is typically difficult to quantity. Unfortunately, the concept of virial mass based on an evolving overdensity threshold leads to pseudo-evolution, which adds further uncertainty to the relation between the two mass estimates. | 16 | 9 | 1609.07147 |
1609 | 1609.08794_arXiv.txt | We have determined the masses and mass-to-light ratios of 50 Galactic globular clusters by comparing their velocity dispersion and surface brightness profiles against a large grid of 900 $N$-body simulations of star clusters of varying initial concentration, size and central black hole mass fraction. Our models follow the evolution of the clusters under the combined effects of stellar evolution and two-body relaxation allowing us to take the effects of mass segregation and energy equipartition between stars self-consistently into account. For a subset of 16 well observed clusters we also derive their kinematic distances. We find an average mass-to-light ratio of Galactic globular clusters of $<M/L_V>=1.98 \pm 0.03$, which agrees very well with the expected $M/L$ ratio if the initial mass function (IMF) of the clusters was a standard Kroupa or Chabrier mass function. We do not find evidence for a decrease of the average mass-to-light ratio with metallicity. The surface brightness and velocity dispersion profiles of most globular clusters are incompatible with the presence of intermediate-mass black holes (IMBHs) with more than a few thousand $M_\odot$ in them. The only clear exception is $\omega$~Cen, where the velocity dispersion profile provides strong evidence for the presence of a $\sim$40,000 M$_\odot$~IMBH in the centre of the cluster. | \label{sec:intro} Globular clusters are among the oldest structures in the universe, having formed within 1 to 2 Gyr after the Big Bang \citep{kravtsovgnedin2005}. Studying their origin and evolution has therefore important implications for our understanding of star formation and the growth of structure in the early universe. In addition, due to their high central densities and high stellar encounter rates, globular clusters are also unique environments for the creation of exotic stars like blue stragglers \citep{bailyn1995,davies2004}, low-mass X-ray binaries \citep{verbunt1993,pooleyetal2003} and millisecond pulsars \citep{manchesteretal1991}. The high stellar densities in globular cluster could also give rise to the creation of intermediate-mass black holes \citep{pzm2002,pzetal2004,gierszetal2015}, which might be the progenitors of supermassive black holes in Galactic centers. Globular clusters are finally important environments for the creation of tight black hole binaries which merge through the emission of gravitational waves \citep{banerjeeetal2010,downingetal2011, rodriguezetal2016a,rodriguezetal2016b,askaretal2016}. In order to understand the rate of creation of exotic stars, it is important to know the mass density profile of globular clusters and how different types of stars are distributed within a globular cluster. This is possible by a detailed modeling of the internal kinematics of globular clusters. Several methods have been suggested in the literature to derive cluster masses from observed density profiles: One can either using analytic formulas which relate a cluster's mass to its radius and velocity dispersion inside some radius \citep[e.g.][]{mandushevetal1991,straderetal2011}, or fit analytic density profiles like Plummer or King models to the observed velocity and surface density profiles of globular clusters \citep[e.g.][]{mclaughlinvandermarel2005,kimmigetal2015}. Finally it is possible to deproject the observed surface density profile and then derive the cluster mass through Jeans modeling and a fit of the observed velocity dispersion profile \citep[e.g.][]{vandevenetal2006,noyolaetal2008,lutzgendorfetal2012,lutzgendorfetal2013}. Most approaches assume a constant mass-to-light ratio inside globular clusters. However, since the relaxation times of globular clusters are generally much smaller than their ages, high-mass stars like compact remnants and giant stars are concentrated towards the cluster centers while low-mass stars are pushed towards the outer cluster parts \citep{baumgardtmakino2003}. Hence the assumption of a constant mass-to-light ratio is not valid for globular clusters. In addition, due to energy equipartition, massive stars move more slowly at a given radius compared to average cluster stars \citep{trentivandermarel2013,bianchinietal2016}. As a result, the velocity dispersion derived from giant stars will underestimate the true velocity dispersion, which leads to an underestimation of the total cluster mass if mass segregation is not properly taken into account \citep{shanahangieles2015}. It is possible to account for mass segregation by e.g. using multi-mass King-Michie models \citep{michie1963,gg1979} or the more recently suggested {\tt LIMEPY} models \citep{gieleszocchi2015,zocchietal2016b}. Multi-mass models have however additional degrees of freedom since the amount of mass segregation between different mass components can in principle be freely chosen in the models. In the present paper we follow a different approach to derive the absolute masses and mass-to-light ratios of globular clusters from their surface density and velocity dispersion profiles. We perform a large grid of $N$-body simulations and scale each model so that it has the same half-light radius as the observed clusters. Scaling is done in such a way that the relaxation time is kept constant, thereby making sure that mass segregation of stars and (partial) energy equipartition between them are taken into account in a self consistent way in the scaled models, i.e. each model has the exact amount of mass segregation which a real globular cluster would have if it started from the same initial condition. We then determine the model which best fits the observed density and velocity dispersion profile for each globular cluster and determine the total mass, mass-to-light ratio and the possible presence of an intermediate mass black hole in the observed clusters from the best-fitting model. Our paper is organised as follows: In section 2 we describe the grid of $N$-body models that we have performed, and in section 3 we describe the selection of the observational data. Section 4 presents our results and we draw our conclusions in section~5. | 16 | 9 | 1609.08794 |
|
1609 | 1609.01998_arXiv.txt | We study the main cosmological properties of the agegraphic dark energy model at the expansion and perturbation levels. Initially, using the latest cosmological data we implement a joint likelihood analysis in order to constrain the cosmological parameters. Then we test the performance of the agegraphic dark energy model at the perturbation level and we define its difference from the usual $\Lambda$CDM model. Within this context, we verify that the growth index of matter fluctuations depends on the choice of the considered agegraphic dark energy (homogeneous or clustered). In particular, assuming a homogeneous agegraphic dark energy we find, for the first time, that the asymptotic value of the growth index is $\gamma \approx 5/9$, which is close to that of the usual $\Lambda$ cosmology, $\gamma^{(\Lambda)} \approx 6/11$. Finally, if the distribution of dark energy is clustered then we obtain $\gamma \approx 1/2$ which is $\sim 8\%$ smaller than that of the $\Lambda$CDM model. | The analysis of various observational data including those of supernovae type Ia (SNIa) \citep{Riess1998,Perlmutter1999,Kowalski2008}, cosmic microwave background (CMB) \citep{Komatsu2009,Komatsu2011,Jarosik:2010iu,Planck2015_XIII}, large scale structure and baryonic acoustic oscillation (BAO) \citep{Percival2010,Tegmark:2003ud,Cole:2005sx,Eisenstein:2005su,Reid:2012sw,Blake:2011rj}, high $z$ galaxies \citep{Alcaniz:2003qy}, high $z$ galaxy clusters \citep{Allen:2004cd,Wang1998}, weak gravitational lensing \citep{Benjamin:2007ys,Amendola:2007rr,Fu:2007qq} strongly suggest an accelerated expansion of the universe. In the context of General Relativity (GR) the so-called dark energy (hereafter DE), which has a negative pressure, is required in order to interpret the cosmic acceleration. It is interesting to mention that from the overall energy density, only $\sim 30\%$ consists of matter (luminous and dark) while the rest corresponds to DE (see the analysis of Planck2015 \citep{Planck2015_XIII} and references therein). From the theoretical perspective, over the last two decades, a large family of phenomenological models has been proposed to study the cosmological features of DE. The simplest DE model is the concordance $\rm \Lambda$ cosmology for which the equation of state (EoS) $ w_{\rm \Lambda}$ is strictly equal to -1. Despite the fact that the $\Lambda$CDM model fits extremely well the cosmological data it suffers from the fine tuning and the cosmic coincidence problems \citep{Weinberg1989,Sahni:1999gb,Carroll2001,Padmanabhan2003,Copeland:2006wr}. Since the nature of DE has yet unknown, different versions of dynamical DE models have been introduced in order to alleviate the above cosmological issues. Generally speaking, it has been proposed that we cannot entirely understand the nature of dark energy before the establishment of a complete theory of quantum gravity \citep{Witten:2000zk}. Nevertheless, it is promising that holographic dark energy (HDE) models \citep{Horava2000,Thomas2002} inspired by the principles of quantum gravity can be suggested and may hopefully provide an efficient explanation for the dynamical nature of DE. Specifically, the holographic principle \citep{Susskind1995} points out that in a finite-size physical system the number of degrees of freedom should be finite and bounded by the area of its boundary \citep{Cohen1999}. In other words, the total energy of a physical system with size $L$ obeys the following inequality $L^3\rho_{\Lambda}\leq Lm_{\rm p}^2$, where $\rho_{\rm \Lambda}$ is the quantum zero-point energy density and $m_{\rm p}=1/\sqrt{8 \pi G}$ is the Planck mass. Applying the latter arguments to cosmological scales it has been found \citep[see ][]{Li2004} that the density of the HDE is given by \begin{equation}\label{eq:density} \rho_{\rm d}=3n^2m_{\rm p}^2L^{-2}\;, \end{equation} where $n$ is a positive constant and the coefficient $3$ is used for convenience. Obviously, in this case the features of DE strongly depend on the definition of the size $L$ in equation (\ref{eq:density}). If we assume $L$ to be the Hubble radius $H^{-1}$ then we cannot produce an accelerated expansion of the universe \citep{Horava2000,Cataldo2001,Thomas2002,Hsu2004}. Another choice would be to replace $L$ with the particle horizon but again we would not be able to extract cosmic acceleration \citep{Horava2000,Cataldo2001,Thomas2002,Hsu2004}. The final choice for $L$ is to use the event horizon (first introduced by \cite{Li2004}). In this case, the HDE model is able to provide cosmic acceleration and it is consistent with observations \citep{Pavon2005,Zimdahl2007,Sheykhi:2011cn}. Notice, that the HDE model has been widely investigated in the literature \citep{Huang:2004wt,Huang2004b,Gong2004,Gong2005a,Gong2005b,Zhang:2005hs,Zhang:2007sh, Elizalde:2005ju,Guberina:2005fb,Guberina:2006qh,BeltranAlmeida:2006is,Wang:2004nqa,Shen:2004ck}. Since the HDE model is obtained by choosing the event horizon length scale, an obvious drawback concerning causality appears in this scenario. Recently, a new DE model, dubbed agegraphic dark energy (ADE) model, has been suggested by \cite{Cai2007} in order to alleviate the above problem. In particular, combining the uncertainty principle in quantum mechanics and the gravitational effects of GR Karolyhazy and his collaborators \citep{Karolyhazy:1966zz,Karolyhazy1982,Karolyhazy1986} made an interesting observation concerning the distance measurement for the Minkowski spacetime through a light-clock Gedanken experiment \citep[see also][]{Maziashvili:2007zz}. They found that the distance $t$ in Minkowski spacetime cannot be known to a better accuracy than $\delta t=\beta t_{\rm p}^{2/3}t^{1/3}$, where $\beta$ is a dimensionless constant of order ${\cal O}(1)$ \citep[see also][]{Maziashvili:2007zz}. \footnote{Through out this work we use the units $\hbar=c=k_{\rm B}=1$. Hence we have $l_{\rm p}=t_{\rm p}=1/m_{\rm p}$ ,where $l_{\rm p}$, $t_{\rm p}$ and $m_{\rm p}$ are the reduced Planck length, time and mass, respectively.} Based on the Karolyhazy relation, \cite{Maziashvili:2007zz} argued that the energy density of metric fluctuations in the Minkowski spacetime is written as \citep[see also][]{Maziashvili:2007dk} \begin{equation}\label{eq:energy2} \rho_{d}\sim\frac{1}{t_{p}^{2}t^{2}}\sim \frac{m_{p}}{t^{2}}\;, \end{equation} where $ m_{p}$ and $t_{p}$ are the reduced Plank mass and the Plank time, respectively \citep[see also][]{Sasakura:1999xp,Ng:1993jb,Ng:1995km,Krauss:2004fb,Christiansen:2005yg,Arzano:2006wp,Ng:2007bp}. Using equation (\ref{eq:energy2}) \cite{Cai2007} proposed another version of holographic DE the so-called agegraphic dark energy (ADE) in which the time scale $t$ is chosen to be equal with the age of the universe $T= \int^{t}_0dt=\int_{0}^{a}\frac{da}{aH} $, with $a$ the scale factor of the universe and $H$ the Hubble parameter. Therefore, the ADE energy density is given by \citep{Cai2007} \begin{equation} \rho_{d}=\frac{3n^{2}m_{p}^{2}}{T^{2}}\label{eq:energy3} \end{equation} where $n$ is a free parameter and the coefficient $3$ appears for convenience. The present value of the age of universe ($T_{\rm 0}\sim H_{\rm 0}^{-1}$) implies that $n$ is of order ${\cal O}(1)$. It has been shown that the condition $n>1$ is required in order to have cosmic acceleration \citep{Cai2007}. Although, the ADE scenario does not suffer from the causality problem \citep{Cai2007} it faces some problems towards describing the matter-dominated epoch \citep{Wei:2007ty,Neupane:2007fw,Wei:2007xu}. To overcome this issue \cite{Wei:2007ty} proposed a new agegraphic dark energy (NADE) model, in which the cosmic time $t$ is replaced by the conformal time $ \eta= \int^{t}_0{\frac{dt}{a(t)}}=\int^{a}_0{\frac{da}{a^{2}H}}$ and thus the energy density in this case becomes \citep{Wei:2007ty} \begin{equation} \rho_{d}=\frac{3n^{2}m_{p}^{2}}{\eta^{2}}\label{roo} \;. \end{equation} It is interesting to mention that \cite{Kim:2007iv} showed that the NADE model provides the proper matter-dominated and radiation-dominated epochs, in the case of $n>2.68$ and $n>2.51$, respectively. Also \cite{Wei:2007xu} found that the coincidence problem can be alleviated naturally in this model and using the cosmological data (SNIa, CMB etc) they obtained $n=2.716^{+0.111}_{-0.109}$. We would like to point out that the cosmological properties of the ADE and the NADE models can be found in \cite{Kim:2008hz,Setare:2010zy,Karami:2010qe,Sheykhi:2010jn,Sheykhi:2009rk,Sheykhi:2009yn, Sheykhi:2009sz,Sheykhi:2009sn, Lee:2008zzw,Jawad:2014ssa,Liu:2012kha, Zhang:2012pr,Farajollahi:2012zz,Zhai:2011pp, Chen:2011rz,Sun:2011vg,Lemets:2010qz,Zhang:2010im,Liu:2010ci,Karami:2010aq,Malekjani:2010uv,Jamil:2010vr, Karami:2009wd,Sheykhi:2009yn,Sheykhi:2009sz,Sheykhi:2009rk,Wu:2008jt,Zhang:2008mb,Kim:2007iv,Neupane:2007ra,Wei:2007xu}.\\ Furthermore, it is well known that beyond the expansion rate of the universe, DE affects the formation of cosmic structures \citep{Peebles1993,Tegmark:2003ud}. Usually, in dynamical DE models with $w_{\rm d}\neq-1$, one can assume that the DE perturbations behave in a similar fashion to matter \citep{Abramo2007,Abramo:2008ip,Abramo2009a,Batista:2013oca,Batista:2014uoa,Mehrabi:2014ema,Malekjani:2015pza,Mehrabi:2015hva}. In principle, the effective sound speed $c_{\rm eff}^{2}=\delta p_{\rm d}/\delta\rho_{\rm d}$ is introduced in order to describe the DE clustering. In particular, if DE is homogeneous then we have $c^{2}_{\rm eff}=1$, while for clustered DE models we use $c^{2}_{\rm eff}=0$. In the homogeneous case, the sound horizon of DE is close to the Hubble length. This does not hold for $c^{2}_{\rm eff}=0$ (clustered DE) which implies that the perturbations of DE grow, via gravitational instability, in sub-Hubble scales \citep{ArmendarizPicon:1999rj,ArmendarizPicon:2000dh,Garriga:1999vw,Akhoury:2011hr}. The growth of matter perturbations in cosmologies where DE is allowed to have clustering has been widely investigated in the literature \citep{Batista:2013oca,Erickson:2001bq,Bean:2003fb,Hu:2004yd,Ballesteros:2008qk, Basilakos:2009mz,dePutter:2010vy,Sapone:2012nh,Dossett:2013npa,Basse:2013zua, Batista:2014uoa,Pace:2014taa,Steigerwald:2014ava,Mehrabi:2014ema,Mehrabi:2015hva,Malekjani:2015pza,Basilakos:2014yda,Nesseris:2014mfa}. In this article, following the lines of the above studies, we attempt to investigate the NADE model at the background and perturbation levels. Specifically, we organize the manuscript as follows: in section (2) we start with a brief presentation of the NADE model and in section (3) we investigate the growth of matter perturbations. In section (4), using the latest cosmological data and the growth rate data, we perform a joint likelihood analysis in order to constraint the free parameters of the model. In section (5) we discuss the growth index and finally, we summarize our results in section (6). | Combining the basic uncertainty principle in quantum mechanics together with the gravitational effects of general relativity, \cite{Cai2007} proposed a new model of DE the so-called agegraphic dark energy model (ADE). Replacing the cosmic time $t$ with the conformal time $\eta$, \cite{Wei:2007ty,Wei:2007xu} introduced in the literature the new agegraphic dark energy (NADE) model in which the corresponding EoS parameter is a function of redshift. The aim of our paper is to investigate the main properties of the NADE model at background and perturbation levels respectively. The current study was performed by a three-step process. Firstly, solving the system of the main differential equations at the background (Friedmann and continuity) and perturbation levels, we investigated the behavior of the basic cosmological quantities $\{H(z), w(z), \Omega_{\rm d}(z), D(z) \}$ in order to understand the global characteristics of the NADE model (see Fig.\ref{figa}). We verified that the EoS parameter remains in the quintessence regime and it lies in the interval $-1 \leq w_{d} \leq -\frac{2}{3}$. Notice, that for large values of $n$ the current value of the EoS parameter tends to -1. As expected, the cosmic expansion depends on the choice of $n$. In particular, we found that for large values of $n>3$, the parameter $\Omega_{\rm d}(z)$ of the NADE model strongly deviates from that of $\Lambda$CDM. In this context we have that the Hubble parameters obey: $H_{\rm NADE}(z)<H_{\Lambda}(z)$. Also, we found that the growth factor $D_{\rm NADE}(z)$ seems to evolve at high redshifts $z>2$ and the amplitude of $D_{\rm NADE} (z)$ increases as a function of $n$. Note that at the same redshift range $D_{\Lambda}(z)$ reached a plateau. Potentially, the latter behavior of $D(z)$ can be used to distinguish between NADE and $\Lambda$CDM at perturbation level. Secondly, we performed a joint statistical analysis, involving the latest geometrical data (SNe type Ia, CMB shift parameter and BAO etc) and growth data and found that the combined statistical analysis, within the context of flat FRW space, can place tight constrains on the main cosmological parameters giving the reader the opportunity to appreciate the precision of our statistical results. In particular, we found $n=2.775\pm 0.053$ ($\sigma_{8}=0.800\pm 0.017$) and $n=2.795\pm 0.054$ ($\sigma_{8}=0.786\pm 0.017$) for homogeneous and inhomogeneous (DE is allowed to clump) NADE models respectively. Notice, that the present value of $\Omega_m$ is close to 0.29 in all cases. It is interesting to mention that the above constraints are in agreement with those of \cite{Wei:2007xu} \citep[see also][]{Wei:2008rv,Zhang:2012pr}. Using the aforementioned cosmological parameters we found that the age of the universe is $t_{0}^{(\rm NADE)}\simeq 13.91$ Gyr which differs from that of the Planck \citep{Planck2015_XIII} by $\sim 1\%$. Lastly, using the basic information criteria (AIC and BIC), we concluded that both smooth and clustered NADE scenarios as well as the $\Lambda$CDM model fit the observationally data equally well. Thirdly, we studied the performance of NADE model at the perturbation level. Specifically, following the methodology of \citep{Steigerwald:2014ava} we estimated for the NADE model the asymptotic value of the growth index of linear matter fluctuations ($\gamma$). Considering a homogeneous NADE model we found $\gamma \approx 5/9$ which is close to that of the traditional $\Lambda$ cosmology $\gamma^{(\Lambda)} \approx 6/11$. On the other hand, if we allow clustering in NADE then we obtained $\gamma \approx 1/2$, which is $\sim 8\%$ smaller than that of the $\Lambda$CDM model. Finally, we extended the growth analysis in the case where $\gamma$ varies with redshift and found that the $\gamma(z)$ is quite large with respect to that of the clustered NADE scenario. This implies that the clustered DE scenario can be differentiated from the other two models on the basis of the growth index evolution. | 16 | 9 | 1609.01998 |
1609 | 1609.03165_arXiv.txt | In this paper, we report our multiwavelength observations of the C3.1 circular-ribbon flare SOL2015-10-16T10:20 in active region (AR) 12434. The flare consisted of a circular flare ribbon (CFR), an inner flare ribbon (IFR) inside, and a pair of short parallel flare ribbons (PFRs). The PFRs located to the north of IFR were most striking in the \textit{Interface Region Imaging Spectrograph} (\textit{IRIS}) 1400 {\AA} and 2796 {\AA} images. For the first time, we observed the circular-ribbon flare in the Ca {\sc ii} H line of the Solar Optical Telescope (SOT) aboard \textit{Hinode}, which has similar shape as observed in the Atmospheric Imaging Assembly (AIA) 1600 {\AA} aboard the Solar Dynamic Observatory (\textit{SDO}). Photospheric line-of-sight magnetograms from the Helioseismic and Magnetic Imager (HMI) aboard \textit{SDO} show that the flare was associated with positive polarities and a negative polarity inside. The IFR and CFR were cospatial with the negative polarity and positive polarities, implying the existence of a magnetic null point (\emph{\textbf{B}}$=$\textbf{0}) and the dome-like spine-fan topology. During the impulsive phase of the flare, ``two-step'' raster observations of \textit{IRIS} with a cadence of 6 s and an exposure time of 2 s show plasma downflow at the CFR in the Si {\sc iv} $\lambda$1402.77 line ($\log T\approx4.8$), suggesting chromospheric condensation. The downflow speeds first increased rapidly from a few km s$^{-1}$ to the peak values of 45$-$52 km s$^{-1}$, before decreasing gradually to the initial levels. The decay timescales of condensation were 3$-$4 minutes, indicating ongoing magnetic reconnection. Interestingly, the downflow speeds are positively correlated with logarithm of the Si {\sc iv} line intensity and time derivative of the \textit{GOES} soft X-ray (SXR) flux in 1$-$8 {\AA}. The radio dynamic spectra are characterized by a type \Rmnum{3} radio burst associated with the flare, which implies that the chromospheric condensation was most probably driven by nonthermal electrons. Using an analytical expression and the peak Doppler velocity, we derived the lower limit of energy flux of the precipitating electrons, i.e., 0.65$\times$10$^{10}$ erg cm$^{-2}$ s$^{-1}$. The Si {\sc iv} line intensity and SXR derivative show quasi-periodic pulsations with periods of 32$-$42 s, which are likely caused by intermittent null-point magnetic reconnections modulated by the fast wave propagating along the fan surface loops at a phase speed of 950$-$1250 km s$^{-1}$. Periodic accelerations and precipitations of the electrons result in periodic heating observed in the Si {\sc iv} line and SXR. | \label{s-intro} Solar flares are one of the most spectacular activities in the solar system. Up to 10$^{29}$$-$10$^{32}$ ergs magnetic free energies are impulsively released, accompanied by rapid increases of emissions in various wavelengths \citep{shi11}. It is generally believed that magnetic reconnection plays a key role in the reconfiguration of magnetic field lines and conversion of magnetic energy into kinetic and thermal energies of plasma \citep{for96,pri00}. In the context of standard solar flare model \citep{car64,stu66,hir74,kopp76}, the accelerated nonthermal electrons propagate downward and heat the chromosphere, forming the bright flare ribbons in H$\alpha$, Ca {\sc ii} H, ultraviolet (UV), and extreme-ultraviolet (EUV) wavelengths. The flare ribbons have diverse shapes, with two parallel flare ribbons (PFRs) being the most commonplace \citep{li15c}. Sometimes, circular flare ribbons (CFRs) appear accompanied by remote brightenings, which is generally believed to be associated with a magnetic null point (\emph{\textbf{B}}$=$\textbf{0}) and the dome-like spine-fan topology \citep{mas09,zqm12,pon16}. More complex flare ribbons have also been observed \citep{jos15}. Apart from the localized heating, hard X-ray (HXR) emissions are generated via Coulomb collisions \citep{bro71}. If open magnetic field lines are involved, it is likely that the nonthermal electrons at speeds of 0.06$-$0.25$c$ ($c$ is the speed of light) escape the Sun into the interplanetary space and generate type \Rmnum{3} radio bursts with frequency ranging from 0.2 to hundreds of MHz \citep{dul87,asch95,zqm15}. Sometimes, the HXR and radio emissions of a flare show quasi-periodic pulsations (QPPs), with their periods ranging from milliseconds \citep{tan10} through a few seconds \citep{kli00,asai01,ning05,naka10,hay16} to several minutes \citep{of06,sych09,naka09,ning14}. So far, QPPs have been extensively explored using both imaging \citep{su12a,su12b} and spectral observations \citep{mar05}. \citet{li15b} studied the X1.6 flare on 2014 September 10. Four-minute QPPs are evident not only in the HXR, EUV, UV, and radio light curves, but also in the temporal evolutions of the Doppler velocities and line widths of the C {\sc i}, O {\sc iv}, Si {\sc iv}, and Fe {\sc xxi} lines. \citet{bro15} studied the M7.3 flare on 2014 April 18. The chromospheric and transition region line emissions show quasi-periodic intensity and velocity fluctuations with periods of $\sim$3 minutes during the first four peaks and $\sim$1.5 minutes during the last four peaks. The overpressure of the chromosphere drives evaporation of hot and dense plasma upward into the newly reconnected coronal loops at speeds of tens to hundreds of km s$^{-1}$ \citep{fis85a,fis85b,fis85c}. At the same time, downward chromospheric condensation takes place at speeds of a few to tens of km s$^{-1}$ as a result of the balance of momentum \citep{ter03}. In the past 40 years, the investigations of chromospheric evaporation and condensation have greatly benefitted from the spectroscopic observations in H$\alpha$ and EUV wavelengths from the ground-based as well as space-borne telescopes \citep[e.g.,][]{cza99,bro04,chen10,li11,mil11,you13,tian14,tian15,li15a,pol15,pol16,bat15,gra15}. \citet{ich84} studied the red asymmetry of H$\alpha$ line profiles of 4 flares near the disk center. The first two on June 20 and June 21 in 1982 were indeed circular-ribbon flares. The red-shifted emission streaks of H$\alpha$ line are caused by the conspicuous downward motion in the flare chromospheric region with velocities of 40$-$100 km s$^{-1}$. Using the spectral data of an M6.8 two-ribbon flare observed by the Coronal Diagnostic Spectrometer (CDS) aboard \textit{SOHO}, \citet{cza99} found that the Fe {\sc xvi} downward velocity of the regions between the flare ribbons and the magnetic neutral line decreases from $\sim$90 km s$^{-1}$ to $\sim$15 km s$^{-1}$ within 2.5 minutes in the gradual phase, indicating ongoing magnetic reconnection. \citet{bro04} studied an M2.0 flare observed by CDS and found that the downward velocity of chromospheric condensation in O {\sc v} line decreases from $\sim$40 km s$^{-1}$ to zero within 3 minutes. Compared with CDS, the Extreme-ultraviolet Imaging Spectrometer \citep[EIS;][]{cul07} aboard \textit{Hinode} \citep{kos07} covers a much wider temperature range and has much higher resolution and cadence. Using the spectral data of a C1.0 flare observed by EIS, \citet{bro13} found that the upward velocities of chromospheric evaporation in Fe {\sc xxiii} line decreases from its maximum value ($>$200 km s$^{-1}$) to zero within 2 minutes. The successful launch of the \textit{Interface Region Imaging Spectrograph} \citep[\textit{IRIS};][]{dep14} opened a new era of flare research, thanks to its unprecedented resolution and cadence. The timescales of chromospheric evaporation range from 5 to 12 minutes in most cases \citep{bat15,li15a,tian15,pol15,pol16,gra15}. However, it can reach $\sim$20 minutes in an M1.0 flare \citep{sad16}. The timescales of chromospheric condensation range from 1 to 7 minutes \citep{gra15,sad16}. A delay of $>$60 s between the starting times of condensation and evaporation has been observed \citep{gra15,you15}. Both the upward evaporation velocities and downward condensation velocities are positively correlated with the HXR fluxes, which is consistent with the numerical model of evaporation driven by nonthermal electrons \citep{tian15,li15a}. So far, chromospheric condensation in circular-ribbon flares has rarely been investigated. Using the imaging and raster observations of a \textit{GOES} C4.2 circular-ribbon flare on 2015 October 16 by \textit{IRIS}, \citet{zqm16} found explosive chromospheric evaporation during the impulsive phase of the flare, which is characterized by simultaneous plasma upflow (35$-$120 km s$^{-1}$) observed in the high-temperature Fe {\sc xxi} $\lambda$1354.09 line ($\log T\approx7.05$) and downflow (10$-$60 km s$^{-1}$) observed in the low-temperature Si {\sc iv} $\lambda$1393.77 line ($\log T\approx4.8$). Based on the quantitative estimation of nonthermal energy flux under the thick-target model and the fact that the inner flare ribbon (IFR) where chromospheric evaporation occurred was cospatial with the single HXR source at 12$-$25 keV, the authors concluded that the explosive chromospheric evaporation was most likely driven by nonthermal electrons accelerated by magnetic reconnection. However, the \textit{IRIS} observation was in the ``sparse synoptic raster'' mode with a cadence of 9.4$\times$36$=$338.4 s, which is difficult to study the temporal evolution of the same position. There were a couple of homologous circular-ribbon flares in the same active region (AR 12434) on the same day as the C4.2 flare. In this paper, we study another one, the C3.1 flare observed by \textit{IRIS}. The raster observation was in the ``two-step raster'' mode with a cadence of 6 s, which is quite suitable for investigating the temporal evolution of the same position of flare ribbon. In addition, we observed QPPs during the impulsive phase of the flare. In Section~\ref{s-data}, we describe the instruments and data analysis using observations from various telescopes. Results and discussions are presented in Section~\ref{s-res} and Section~\ref{s-disc}, respectively. Finally, we draw a conclusion in Section~\ref{s-sum}. | \label{s-disc} \subsection{What Is the Cause of Chromospheric Condensation?} \label{s-cause1} The observational studies of chromospheric evaporation and condensation have a long history. \citet{ich84} studied the temporal evolution of the H$\alpha$ line profiles of 4 flares near the disk center. The timescales of the total evolutions are $\sim$1 minute. The authors proposed that the red asymmetry of H$\alpha$ line results from the downward motion of the compressed chromospheric region produced by the impulsive heating by energetic electron beam or thermal conduction. \citet{li15a} found positive correlations between the HXR emissions and upward Doppler shifts of Fe {\sc xxi} line and downward Doppler shifts of C {\sc i} line ($\log T\approx4.0$) in two X1.6 flares. In the C4.2 circular-ribbon flare, \citet{zqm16} found that the IFR where explosive chromospheric evaporation occurred was cospatial with the single HXR source. A quantitative estimation of the energy flux of the nonthermal electron supports the electron-driven evaporation/condensation. Like in \citet{li15a}, we explored the relationship between the downward $V_D$ of the CFR and SXR derivative for the C3.1 flare. The scatter plots of the two parameters are drawn in the right panels of Figure~\ref{fig10} for H2 and H2*, indicating that they are positively correlated with correlation coefficients $\geq$0.7. In Figure~\ref{fig13}, the radio dynamic spectra from \textit{KRIM} and \textit{BLENSW} are displayed in the top two panels. The most striking feature is the type \Rmnum{3} burst with enhanced emissions and rapid frequency drift during 10:17:00$-$10:18:00 UT. The dynamic spectra from RAD1 and RAD2 aboard \textit{WIND}/WAVES are displayed in the bottom two panels. The frequency of burst drifted rapidly from 13 MHz at 10:18 UT to 1 MHz at 10:20 UT and then drifted slowly to $\sim$0.2 MHz until $\sim$10:40 UT. The radio flux at 42.5 MHz is plotted in Figure~\ref{fig6}(b) with a green line. The peaks of the radio flux were roughly coincident with the peaks of SXR derivative. Considering that the HXR flux and type \Rmnum{3} radio burst origin from the nonthermal electrons propagating downward into the chromosphere along the newly reconnected magnetic field lines and upward into the interplanetary space along the open field lines, the coincidence of the peaks during the impulsive phase of the flare implies the existence of nonthermal electrons accelerated by magnetic reconnection. Combining the positive correlation between $V_D$ of the CFR and SXR derivative, we conclude that the chromospheric condensation of the C3.1 flare was most probably driven by nonthermal electrons. The investigations of chromospheric condensation in theory and numerical simulations have also made huge progress. \citet{fis86} derived an analytical expression of the condensation lifetime: \begin{equation} \label{eqn4} \tau=\pi\sqrt{H/g}, \end{equation} where $H$ and $g$ stand for the gravitational scale height and acceleration of gravity in the chromosphere. The typical value of $\tau$ is $\sim$1 minute. However, superposition of continuous condensations energized at different times within an unresolved observational element naturally result in a longer condensation lifetime \citep{fis89}. In the combined modeling of acceleration, transport, and hydrodynamic response in solar flares, \citet{rub15} found that the condensation timescale can reach 70$-$90 s. In the 1D radiative hydrodynamic numerical simulations, \citet{reep15} found that the timescales of the electron-driven chromospheric evaporation are 3.5$-$4 minutes for different cases. In Table~\ref{tbl-2}, we list the timescales of chromospheric condensation and evaporation in the previous literatures. It is revealed that the timescales of evaporation upflows (2$-$20 minutes) are generally 2$-$3 times larger than the timescales of condensation downflows (1$-$7 minutes). For the C3.1 flare in this study, the condensation timescale is 3$-$4 minutes, which is consistent with previous findings. To our knowledge, the cadence (6 s) of the \textit{IRIS} raster observation is the highest ever for the investigation of chromospheric condensation. Based on the radiative hydrodynamic flare model, \citet{fis89} derived an analytical expression of the peak downflow speed for explosive evaporation: \begin{equation} \label{eqn5} v_{peak}\approx0.6(F_{evap}/\rho_{ch})^{1/3}, \end{equation} where $F_{evap}$ and $\rho_{ch}$ represent the energy flux and chromospheric density. For the C3.1 flare, $v_{peak}=52$ km s$^{-1}$. Taking 10$^{-11}$ g cm$^{-3}$ as the typical value of $\rho_{ch}$ in the chromosphere, $F_{evap}$ is estimated to be 0.65$\times$10$^{10}$ erg cm$^{-2}$ s$^{-1}$, which is close to the threshold energy flux for explosive evaporation ($\sim$10$^{10}$ erg cm$^{-2}$ s$^{-1}$). The estimated $F_{evap}$ should be a lower limit since the slit positions (S1 and S2) are located at the CFR whose maximum intensity is far less than the IFR and PFRs. According to the correlation between the intensity and $V_D$ of Si {\sc iv} line in the left panels of Figure~\ref{fig10}, the peak downward velocities at the IFR and PFRs should be larger and the nonthermal electron energy flux should be higher. \subsection{What Is the Cause of QPPs?} \label{s-cause2} QPPs of solar flares have been extensively observed and reported in various wavelengths. However, the mechanism of QPPs still remains unclear. On one hand, there are diverse MHD waves (kink mode and sausage mode) in the solar atmosphere, such as the slow wave, fast wave, and Alfv\'{e}n wave. The existence and propagation of waves may modulate the emissions in certain wavelengths \citep[e.g.][]{naka04,mar05,tan10,su12a}. On the other hand, the magnetic reconnection rate may be modulated by the $p-$mode wave \citep{chen06}, slow-mode wave \citep{naka11}, Alfv\'{e}n wave \citep{asai01}, fast sausage-mode wave \citep{zai82}, and fast kink-mode wave \citep{naka06}. Therefore, the acceleration and propagation of electrons in the downward and upward directions would be modulated by the periodic magnetic reconnections \citep{asch94,ning05}. Besides, the tearing-mode instability and coalescence of magnetic islands occur iteratively, resulting in intermittent magnetic reconnection and particle acceleration \citep{kli00,kar04}. In a single loop flare on 2005 January 1, \citet{naka10} found QPPs of $\gamma$-ray emission with a period of $\sim$40 s, which are also present in the HXR and microwave emissions. They proposed that the QPPs are created by periodic magnetic reconnections accompanying particle acceleration triggered by a global kink oscillation in a nearby coronal loop. In an M1.8 flare on 2002 October 20, \citet{zim10} found QPPs of both thermal and nonthermal HXR emissions with periods of 16 s and 36 s, which are interpreted in terms of MHD oscillations excited in two interacting systems of flaring loops. In our study, the QPPs with periods of 32$-$42 s are observed in UV (Si {\sc iv} $\lambda$1402.77) and SXR derivative. As described at the beginning of Section~\ref{s-ribb}, the length of a loop from the null point to the solar surface is estimated to be $\sim$20 Mm. If the QPPs are caused by the modulation of magnetic reconnection by fast wave propagating along the fan surface, the phase speed ($2L/P_{QPP}$) is estimated to be 950$-$1250 km s$^{-1}$, which is close to the typical value for fast wave in the corona. Hence, the QPPs in the C3.1 flare are most probably caused by the modulation of magnetic reconnection, particle acceleration, and subsequent precipitation by fast wave propagating along the fan surface loops. The nonthermal electrons propagating downward heat the chromosphere and generate evaporation as well as condensation. Therefore, QPPs are generated in both UV and SXR during the impulsive phase of the flare. In the X1.6 two-ribbon flare on 2014 September 10, QPPs are observed in multiwavelengths from radio to HXR \citep{li15b}. Our results indicate that QPPs exist not only in the typical two-ribbon flares, but also in circular-ribbon flares. The same mechanism may at work in both types. Finally, we briefly discuss the influence of the rest wavelength on the results. As mentioned in Section~\ref{s-iris}, the spectral profiles of Si {\sc iv} $\lambda$1402.77 line could not been fitted by a single-Gaussian function or double-Gaussian function. Using the same method as described in \citet{bro15}, we derived the same rest wavelength of Si {\sc iv} line, i.e., 1402.86$\pm$0.0145 {\AA}. Therefore, our method is reasonable and results are convincing. The blueshifted Doppler velocities of H1 (see Figure~\ref{fig8}(d)) may be a signature of gentle chromospheric evaporation due to a very low energy flux of electrons. However, the blueshifts are too small to draw a solid conclusion. | 16 | 9 | 1609.03165 |
1609 | 1609.01026_arXiv.txt | In 2016 May, the intermediate polar FO~Aqr was detected in a low state for the first time in its observational history. We report time-resolved photometry of the system during its initial recovery from this faint state. Our data, which includes high-speed photometry with cadences of just 2 sec, shows the existence of very strong periodicities at 22.5 min and 11.26 min, equivalent to the spin-orbit beat frequency and twice its value, respectively. A pulse at the spin frequency is also present but at a much lower amplitude than is normally observed in the bright state. By comparing our power spectra with theoretical models, we infer that a substantial amount of accretion was stream-fed during our observations, in contrast to the disk-fed accretion that dominates the bright state. In addition, we find that FO~Aqr's rate of recovery has been unusually slow in comparison to rates of recovery seen in other magnetic cataclysmic variables, with an $e$-folding time of 115$\pm7$ days. The recovery also shows irregular variations in the median brightness of as much as 0.2~mag over a 10-day span. Finally, we show that the arrival times of the spin pulses are dependent upon the system's overall brightness. | An intermediate polar (IP) is an interacting binary system featuring a low-mass donor star which overfills its Roche lobe, transferring mass to a magnetic white dwarf (WD) \citep[for a review, see ][]{patterson94}. As such, it is a subset of the cataclysmic variable stars (CVs), but the magnetism of the WD results in a number of characteristics which differentiate IPs from other CVs. After the accretion flow leaves the inner Lagrangian (L1) point, it follows a ballistic trajectory until it either (1) circularizes into an accretion disk whose inner region is truncated by the WD's magnetic field or (2) directly impacts the WD's magnetosphere \citep{hameury86}. In the former scenario (disk-fed accretion), the WD's magnetic field captures plasma from the inner accretion disk, while in the latter (stream-fed accretion), the plasma is captured at some point along its ballistic trajectory. In some cases, disk-fed and stream-fed accretion can occur simultaneously if part of the accretion stream overflows the disk following the stream-disk interaction. Regardless of the mode of accretion, the plasma will begin to travel along the WD's magnetic field lines when the local magnetic pressure exerted by the WD exceeds the ram pressure of the accretion flow. During its journey along the field lines, the plasma travels out of the binary orbital plane, creating a three-dimensional structure known as an accretion curtain \citep{rosen88}. The material within the curtain finally impacts the WD near one of its magnetic poles and is shocked to X-ray-emitting temperatures. \begin{figure} \epsscale{1.2} \plotone{longterm.pdf} \caption{The long-term light curve for FO Aqr, including the linear fit to the recovery (see Sec.~\ref{LCsec}). We also plot a line showing the slowest-possible single-sloped decline consistent with the three ASAS-SN observations between JD 189-204. The large gap results from solar conjunction. The `DKS' and `HMB' data are from co-authors Dvorak and Hambsch, respectively, and the other data are described in the text. In an effort to reduce contamination of the long-term light curve by short-term variations, the recovery data mostly show the median magnitude of the system during an extended time series; thus, ASAS-SN observations of the recovery are not plotted. The bottom panel shows residuals from the linear recovery model and uses Gaussian smoothing (blue line) to emphasize a bump in the light curve near JD 290 (shaded region). \label{longterm}} \end{figure} IPs show a complex range of periodicities in their optical light curves because the WD's spin frequency ($\omega$) is not synchronized to the orbital frequency ($\Omega$) of the system \citep{warner86}. If accretion is predominantly stream-fed, \citet{fw99} predict that the dominant frequency at optical wavelengths will be either the spin-orbit beat frequency ($\omega-\Omega$) or its first harmonic ($2\omega - 2\Omega$) if both poles are accreting and contributing equally to the light curve. The beat frequency is the rate at which the WD completes a full rotation within the binary rest frame. It is therefore the frequency at which the WD's magnetic field lines (rotating at $\omega$) will interact with stationary structures in the rest frame (rotating at $\Omega$), such as the accretion stream. If, however, the accretion is disk-fed, the energy released by accretion is independent of the orbital phase of the secondary, and optical variations will be seen at $\omega$ \citep{fw99}. \begin{figure*}[ht!] \centering \epsscale{1.1} \plotone{lightcurves2.pdf} \caption{The three MDM light curves, obtained on three consecutive nights between JD 2457578-80 with a cadence of 2 sec. The red line in each panel shows our model of the light curve, which is the sum of three components: a beat pulse, a spin pulse, and an orbital variation. For illustration, the top panel plots the beat and spin components of the model in blue and green, respectively, with a constant magnitude offset added to each for clarity. At $\phi_{orb}\sim 0.9$, the frequency of the modulation switches from $2\omega-2\Omega$ to $\omega-\Omega$, a transition explained by the superposition of the beat and spin pulses. Our model predicts that the spin pulse is in phase with one of the beat maxima at orbital phase 0.22 and the other at orbital phase 0.72. \label{lightcurve}} \end{figure*} FO Aquarii (hereinafter, FO Aqr) is a well-studied IP that has been dubbed the `King of the Intermediate Polars' \citep{patterson83} due to its bright apparent magnitude and large spin pulse amplitude. Over its history, the optical light of FO~Aqr has been dominated by a 20.9-minute spin period, although $\omega -\Omega$, $2\omega-2\Omega$, and various harmonics are also detected, albeit with much less power \citep{kennedy}. Prior to 2016 May, it had been observed exclusively in a high state ($V < 14$). \citet{gs88} found no faint states of FO~Aqr in the Harvard Plate Collection, which contains observations of FO Aqr from 1923-1953, and none are present in long-term data from both the AAVSO and the Catalina Real-Time Sky Survey \citep[CRTS; ][]{drake}. Indeed, there are no published reports of a low state prior to 2016, implying a relatively stable and elevated rate of mass transfer for a long period of time. However, after FO Aqr emerged from solar conjunction in 2016 May, the system was unexpectedly detected near $V\sim15.7$ (see Fig. ~\ref{longterm}), suggesting a major decrease in the mass-transfer rate \citep{atel1}. Thereafter, it gradually rebrightened at an average rate of 0.01 mag day$^{-1}$ between 2016 May and 2016 July and showed a strong 11.26-minute periodicity in its optical light curve \citep{atel2}. Here, we present time-resolved photometry of this archetypal IP's recovery from its unprecedented low state. \begin{figure*}[ht!] \epsscale{1.2} \plotone{power_spectra2.pdf} \caption{ Lomb-Scargle power spectra of FO Aqr during three stages of its recovery (top three rows) compared to the resampled \textit{K2} data from the high state (bottom row). The data are shown with a linear scale in the left column and with a logarithmic scale in the right column. JD refers to JD - 2457000. The top panel shows that for the vast majority of our observations, $2\omega-2\Omega$ was stronger than $\omega$, while $\omega-\Omega$ was clearly detected as well. In comparison to the high state power spectrum (bottom panel), $\omega-\Omega$ and especially $2\omega-2\Omega$ were dramatically stronger during the low state. However, during two segments of the recovery, the power spectrum changed, emphasizing that the recovery is not static. The second row shows an eight-day stretch characterized by a near-total disappearance of $\omega-\Omega$, and the power spectra in the third row feature an increased signal at $\Omega$ along with a diminution of $2\omega-2\Omega$ relative to $\omega$. \label{powerspectra}} \end{figure*} | We show that FO~Aqr was nearly 2.5~mag fainter than normal and at its lowest recorded brightness when it emerged from solar conjunction in 2016 May. The light curve in its low state was significantly different from the spin-dominated variability observed when the system is bright. While photometric variations at the spin frequency are still present in the low state, the dominant signal is at twice the beat frequency. This suggests that FO~Aqr was accreting substantially from its accretion stream during the system's low state. In addition, the $e$-folding time of the recovery ($115\pm7$ days) is unusually slow. The observed recovery is longer than the viscous timescale of the disk, implying that the mass-transfer rate from the secondary has not yet returned to normal. The decreased depth and width of the disk-grazing eclipse suggest that the disk has been at least partially depleted. Finally, we find that the arrival times of the spin pulse are dependent upon the system's overall brightness. It will be important for future studies to determine whether the pulse at $\omega$ is indeed attributable to disk-fed accretion curtains or is instead from other processes that can produce a signal at $\omega$ \citep{fw99}. Nevertheless, the presence of the grazing eclipse during the low state shows that there was at least one prerequisite for disk-fed accretion---namely, a disk. | 16 | 9 | 1609.01026 |
1609 | 1609.03509_arXiv.txt | Dwarf spheroidals are low-luminosity satellite galaxies of the Milky Way highly dominated by dark matter (DM). Therefore, they are prime targets to search for signals from dark matter annihilation using gamma-ray observations. While the typical assumption is that the dark matter density profile of these satellite galaxies can be described by a spherical symmetric Navarro-Frenk-White (NFW) profile, recent observational data of stellar kinematics suggest that the DM halos around these galaxies are better described by axisymmetric profiles. Motivated by such evidence, we analyse about seven years of \texttt{PASS8} \emph{Fermi} data for seven classical dwarf galaxies, including Draco, adopting both the widely used NFW profile and observationally-motivated axisymmetric density profiles. For four of the selected dwarfs (Sextans, Carina, Sculptor and Fornax) axisymmetric mass models suggest a cored density profile rather than the commonly adopted cusped profile. We found that upper limits on the annihilation cross section for some of these dwarfs are significantly higher than the ones achieved using an NFW profile. Therefore, upper limits in the literature obtained using spherical symmetric cusped profiles, such as the NFW, might be overestimated. Our results show that it is extremely important to use observationally motivated density profiles going beyond the usually adopted NFW in order to obtain accurate constraints on the dark matter annihilation cross section. | \label{sec:intro} Most of the matter in the Universe consists of an unknown component that is commonly considered to be made of non-baryonic cold dark matter \cite{2011ApJS..192...18K,Ade:2015xua}. Finding the particle nature of dark matter (DM) is one of the most pressing goals in modern physics. While many particle physics models have been proposed to solve this puzzle, the most favored and extensively studied candidates fall into the category of weakly interacting massive particles (WIMPs) \cite{2005PhR...405..279B}. These are characterised by a relic density matching the observed DM density, and naturally arise in many theories beyond the standard model of particle physics such as supersymmetry or universal extra-dimension models. The self-annihilation of WIMPs can result in the production of standard model particles. The goal of so-called indirect DM searches is to look for these particles in regions of the Universe where we know DM is abundant \cite{Gaskins:2016cha}. High-energy gamma rays are one example of those particles expected as a result of WIMP annihilation. The search for these gamma rays is a very active field of research fueled in the last decade by many gamma-ray observations of Milky Way (MW) satellite galaxies \cite{2009ApJ...697.1299A,2010ApJ...720.1174A,2011JCAP...06..035A,2011PhRvL.107x1302A,2012PhRvD..85f2001A,2014JCAP...02..008A,Abdo:2010ex,Ackermann:2013yva,2014PhRvD..90k2012A,Geringer-Sameth:2014qqa,Ackermann:2015zua,Geringer-Sameth:2015lua,2015ApJ...809L...4D,2016JCAP...02..039M,PhysRevD.93.043518} and other promising sites such as the Galactic center \cite{2011PhRvL.106p1301A,2011PhRvD..84l3005H,2015PhRvL.114h1301A,2015JCAP...03..038C,2016PDU....12....1D,2015PhRvD..91f3003C,2016ApJ...819...44A,Zhou:2014lva} or clusters of galaxies \cite{2010ApJ...710..634A,2010JCAP...05..025A,2012ApJ...750..123A,2012JCAP...07..017A,2012ApJ...757..123A,2015ApJ...812..159A,2016JCAP...02..026A,PhysRevD.93.103525}, both from the ground with imaging Cherenkov telescopes and from space with the \emph{Fermi} Large Area Telescope (LAT). More recently, novel and competitive constraints have been obtained also from the \emph{Fermi} measurements of the extragalactic gamma-ray background \cite{2012PhRvD..85h3007A,2013MNRAS.429.1529F,2013PhRvD..87l3539A,2014PhRvD..90b3514A,2014NIMPA.742..149G,2015PhR...598....1F, 2015ApJS..221...29C,2015JCAP...06..029C,2015ApJ...802L...1F,Regis:2015zka,Ando:2016ang,Fornasa:2016ohl}. In this paper, we focus on dwarf spheroidal galaxies (dSphs) that are low-luminosity satellite galaxies which are known to be highly DM dominated \cite{1998ARA&A..36..435M,2007PhRvD..75h3526S,2008ApJ...678..614S,2011JCAP...12..011S,Chiappo:2016xfs}. Their high mass-to-light ratio, proximity, and very low expected gamma-ray background from other astrophysical sources make them ideal candidates to search for gamma rays from DM annihilation. The main astrophysical uncertainty when dealing with indirect DM searches in dSphs is their DM density profile, which is the most crucial ingredient needed to estimate the rate of DM annihilation we expect from a given object. The common assumption often adopted in the literature is that dSphs are characterised by a spherically symmetric, so-called Navarro-Frenk-White (NFW) profile \cite{Navarro:1996gj}. This cusped profile originally predicted by $N$-body simulations of cold dark matter might not be the best choice for all cases, and other profiles have been extensively discussed in the literature, including the Einasto profile \cite{1965TrAlm...5...87E}. Additional complications come from going beyond simple spherical symmetric mass models. We know, in fact, that the observed stellar components of all MW dSphs have an axisymmetric shape on the sky-plane with typical axial ratios of 0.6--0.8 \cite{2012AJ....144....4M}. Additionally, recent high-resolution $N$-body simulations showed that DM subhalos tend to have axisymmetric shapes rather than triaxial \cite{2014MNRAS.439.2863V}. These considerations prove the need to relax the assumption of spherical symmetry in the mass modeling of dSphs, which is also one of the major systematic uncertainties for the $J$-factor (i.e., the line-of-sight integral of DM density squared) estimations that most of previous studies have not considered. In this paper we investigate the impact of observationally motivated axisymmetric mass models on indirect DM searches with dSphs using gamma-ray observations by \emph{Fermi}. Uncertainties on the J-factor estimates were addressed in Ref.~\cite{Ullio:2016kvy}, where they explore the impact of the observationally unknown star orbital anisotropy. Triaxial density profiles have been investigated in detail in Ref.~\cite{Bonnivard:2014kza}, where they determine the bias on the $J$-factor that arises when using a spherical Jeans analysis for halos that are likely to be triaxial in shape. In our work, we go beyond the $J$-factor estimates and study the impact on the upper limits obtained for the DM cross section when adopting the axisymmetric models of Ref.~\cite{Hayashi:2015yfa} with respect to those obtained using the commonly adopted NFW profile. We analyse about seven years of \texttt{PASS8} \emph{Fermi}-LAT data for seven classical dSphs, namely Draco, Leo I and II, Sextans, Carina, Sculptor, and Fornax. These dSphs are selected as the overlapping part of the samples considered by Ref.~\cite{Hayashi:2015yfa} and Ref.~\cite{Ackermann:2015zua}. We fit each dSph both with NFW and axisymmetric profiles, and compare their cross section upper limits. We underline, in particular, that Sextans, Carina, Sculptor and Fornax are characterised by cored axisymmetric profiles rather than cusped, and their results can differ significantly from those of the NFW profiles. This paper is organised as follows. In Sec.~\ref{sec:theomodel}, we discuss the expected flux from DM annihilation from dSphs in the case of a NFW profile. The axisymmetric mass model is introduced in Sec.~\ref{sec:axisym}, where we also discuss a qualitative comparison with the NFW profile. In Sec.~\ref{sec:datasel}, we discuss the \emph{Fermi}-LAT data analysis for the seven selected dSphs and present our results in Sec.~\ref{sec:results}. We discuss our conclusions in Section~\ref{sec:concl}. | \label{sec:concl} Dwarf spheroidal galaxies are important and well established targets for indirect DM searches. The most common choice for the DM density profile in the analysis of these dSphs is an NFW profile. Recent observational data of stellar kinematics, however, imply that DM halos around these galaxies are better described by an axisymmetric profile, with an axis ratio of 0.6--0.8, either cored or cusped. For this reason, we investigated the impact of adopting observationally-motivated axisymmetric models instead of the commonly adopted NFW profile on the limits obtained for the DM annihilation cross section for seven classical dSphs with \emph{Fermi} gamma-ray data. Draco is the most promising dwarf galaxy among the seven analysed. Although its DM distribution is well described by a cusped oblate profile in the axisymmetric modeling, the total amount of gamma rays yielding from the overall region will be similar to that of an NFW profile (i.e., similar $J$-factors). As a result, we obtained very similar upper limits on the annihilation cross section for Draco using an NFW and axisymmetric model. The same is true for Leo~II, while Leo~I shows some mild differences, even if both feature an inner cusp. By testing ten axisymmetric profiles randomly chosen from a Monte Carlo sample of the analyses of stellar kinematics data of Draco, we find that the current uncertainty on the density profile of Draco will give a systematic uncertainty on the cross section upper limits of about 10\%. This proves that our conclusions are robust. The analyses of the dSphs best described by a cored profile (Sextans, Sculptor, Carina and Fornax) result in a more substantial difference between the two adopted profiles. In particular, for Sextans, the best-fit model of its stellar kinematics data yields a much more extended $J$-factor map. We found that the cross section upper limits were weaker by a factor of a few to several compared with those obtained with an NFW profile. This demonstrates the importance of properly assessing DM density profiles from observational data, and also that upper limits in the literature obtained assuming a cusped spherical model (such as an NFW) might be overestimated. | 16 | 9 | 1609.03509 |
1609 | 1609.09101_arXiv.txt | As part of our laboratory investigation of the theoretical line profiles used in white dwarf atmosphere models, we extend the electron-density ($n_{\rm e}$) range measured by our experiments to higher densities (up to $n_{\rm e}\sim80\times10^{16}$~cm$^{-3}$). Whereas inferred parameters using the hydrogen-$\beta$ spectral line agree among different line-shape models for $n_{\rm e}\lesssim30\times10^{16}$~cm$^{-3}$, we now see divergence between models. These are densities beyond the range previously benchmarked in the laboratory, meaning theoretical profiles in this regime have not been fully validated. Experimentally exploring these higher densities enables us to test and constrain different line-profile models, as the differences in their relative H-Balmer line {\it shapes} are more pronounced at such conditions. These experiments also aid in our study of occupation probabilities because we can measure these from relative line {\it strengths}. | Theoretical line profiles are a critical ingredient of white dwarf (WD) atmosphere models \citep[e.g.,][]{Koester79,Bergeron92b,Koester10}. A modification to the hydrogen line profiles by \citet{Tremblay09} resulted in significant systematic changes to the inferred WD atmospheric parameters (i.e., effective temperature, $T_{\rm e}$, and surface gravity, log\,$g$) from \citet{Liebert05}. These H line profiles have since become the standard in the community and in the comprehensive analysis of thousands of WDs \citep[e.g.,][]{Tremblay11,Girven11,Gianninas11,Kleinman13,Limoges15,Guo15}. Though this {\it spectroscopic} method is powerful, precise, and widely used, its results do not agree with mass determinations using gravitational redshifts \citep{Barstow05,Falcon10} nor inferred atmospheric parameters using photometry \citep{Genest-Beaulieu14}. For this latter example, H line profiles are specifically a suspect for the disagreement. We thus experimentally investigate the spectroscopic method by targeting the theoretical line profiles used in WD atmosphere models. We have performed laboratory experiments at the {\it Z} Pulsed Power Facility \citep[e.g.,][]{McDaniel02,Matzen05,Rose10,Savage11} at Sandia National Laboratories to measure the spectral line profiles present in the high-density ($n_{\rm e}$) plasmas of WD photospheres \citep{Falcon10b,Falcon13,Montgomery15,Schaeuble16}. Having achieved higher densities in the laboratory than previously explored in this way---while measuring multiple spectral lines simultaneously---we now extend our measurements to plasmas at even higher $n_{\rm e}$. This allows us to better discriminate amongst theoretical line profiles, since relative line {\it shapes} (i.e., among Balmer lines) differ between calculations with increasing principal quantum number \citep{Tremblay09} and with increasing $n_{\rm e}$. We can also uniquely investigate occupation probabilities \citep{Hummer88} by measuring relative line {\it strengths}. | The systematic disagreement between our $n_{\rm e}$ inferences using different line-profile calculations is small at the lower values of experiment z2553, but it is apparent and greater than the measurement uncertainties at the higher values of experiment z2832. This is troubling because we fit the measured H$\beta$ line to diagnose our plasma conditions. We chose it for two reasons: (1) because its theoretical line profiles agree with one another at these lower densities, and (2) because the H$\beta$ spectral line has been validated by benchmark experiments \citep{Kelleher93}. While a few benchmark H-line-profile experiments have reached electron densities greater than $n_{\rm e}=10\times10^{16}$~cm$^{-3}$ \citep[e.g.,][who achieve $n_{\rm e}\sim28$, $\sim16$, and $\sim14\times10^{16}$~cm$^{-3}$, respectively]{McLean65,Baessler80,Helbig81}, the highest density achieved by one that measures multiple lines is that by \citet{Wiese72}, who reach $n_{\rm e}\sim9\times10^{16}$~cm$^{-3}$; measuring multiple lines to test relative line shapes and strengths is a critical requirement for our laboratory investigation of theoretical line profiles \citep{Falcon15c}. Our experiment now has the capability of verifying line-profile calculations at these high electron densities. | 16 | 9 | 1609.09101 |
1609 | 1609.05743_arXiv.txt | { A simulation study of the energy released by extensive air showers in the form of MHz radiation is performed using the CoREAS simulation code. We develop an efficient method to extract this radiation energy from air-shower simulations. We determine the longitudinal profile of the radiation energy release and compare it to the longitudinal profile of the energy deposit by the electromagnetic component of the air shower. We find that the radiation energy corrected for the geometric dependence of the geomagnetic emission scales quadratically with the energy in the electromagnetic component of the air shower with a second order dependency on the atmospheric density at the position of the maximum of the shower development $X_\mathrm{max}$. In a measurement where $X_\mathrm{max}$ is not accessible, this second order dependence can be approximated using the zenith angle of the incoming direction of the air shower with only a minor deterioration in accuracy. This method results in an intrinsic uncertainty of 4\% with respect to the electromagnetic shower energy which is well below current experimental uncertainties. } | \label{intro} The measurement of high-energy cosmic rays using short radio pulses emitted by air showers is a quickly evolving field of research \cite{Huege2016}. Recently, a new method to measure the cosmic-ray energy using the radiation energy, i.e., the energy that is emitted by the air shower within the frequency band of the detector, was presented \cite{ICRC2015CGlaser, AERAEnergyPRD, AERAEnergyPRL, ARENA2016GlaserAERA}. In this work, we study the emission of the radiation energy from the theoretical side using Monte Carlo simulations of air showers and the calculation of the radiation from first-principles based on classical electrodynamics. More details of this analysis can be found in \cite{GlaserErad2016a}. | In this work, we presented a prediction of the radiation energy emitted by air showers using first-principles calculations based on classical electrodynamics. We determined the longitudinal profile of the radiation energy release. Furthermore, we studied the dependence on the shower energy and found that the radiation energy scales quadratically with the energy in the electromagnetic cascade of the air shower after correcting for the dependence of the geomagnetic emission on the geomagnetic field. In addition, we found that the radiation energy shows a second-order dependence on the atmospheric density in which the shower develops. This dependence can be parametrized using the atmospheric density at the shower maximum resulting in an intrinsic uncertainty of the method of 3\%. In a more practical parametrization using only the zenith angle, the method shows an intrinsic uncertainty of 4\%. Hence, the results presented here can be used by cosmic-ray radio experiments to improve the precision in the energy reconstruction and to calibrate the energy scale from first-principles calculations. | 16 | 9 | 1609.05743 |
1609 | 1609.01740_arXiv.txt | We present the stellar surface mass density {\it vs.} gas metallicity ($\Sigma_*-Z$) relation for more than 500,000 spatially-resolved star-forming resolution elements (spaxels) from a sample of 653 disk galaxies included in the SDSS IV MaNGA survey. We find a tight relation between these local properties, with higher metallicities as the surface density increases. This relation extends over three orders of magnitude in the surface mass density and a factor of four in metallicity. We show that this local relationship can simultaneously reproduce two well-known properties of disk galaxies: their global mass-metallicity relationship {\it and} their radial metallicity gradients. We also find that the $\Sigma_* - Z$ relation is largely independent of the galaxy's total stellar mass and specific star-formation rate (sSFR), except at low stellar mass and high sSFR. These results suggest that in the present-day universe local properties play a key role in determining the gas-phase metallicity in typical disk galaxies. | The observed oxygen abundance in the interstellar medium and the stellar mass of a galaxy are the result of galaxy evolution. The metallicity measured from the ionized gas emission lines reflects the metal content produced in previous generations of stars, while the stellar mass traces the previous gas available for star formation. Therefore, the analysis of how these properties correlate is fundamental to our understanding of the assembly of galaxies across cosmic time. The physical mechanisms explaining the connection between these two fundamental galactic features are a matter of debate. An important clue would be to establish whether these relations are related primarily to the global properties of the galaxy {\it vs.} the local ones. There are two key systematic properties that have been uncovered. The first is the close relation between total stellar mass and metallicity. This has been explored for almost 40 years. In a seminal study, \cite{1979A&A....80..155L} found a positive correlation between the oxygen abundance and the luminosity (a proxy for the stellar mass) in a sample of irregular and compact galaxies. Later studies found similar trends for samples of different Hubble types, including a wide range of luminosities and metallicities \cite[e.g,][]{1989ApJ...347..875S, 1994ApJ...420...87Z}. With the development of large spectroscopic surveys such as the single-fiber Sloan Digital Sky Survey (SDSS), \cite{2004ApJ...613..898T} confirmed the mass-metallicity relation (MZR) for a sample of $\sim$ 53,000 galaxies at $z \sim$ 0.1 spanning over 3 orders of magnitude in stellar mass and a factor of 10 in metallicity \citep[see an update for star forming galaxies in][]{2016MNRAS.457.2929W}. The MZR has been established at different redshifts \citep[e.g.,][]{2005ApJ...635..260S, 2006ApJ...644..813E, 2008A&A...488..463M} and different environments \citep[e.g.,][]{2007MNRAS.382..801M, 2013A&A...550A.115H}. Several scenarios have been invoked to explain the MZR from variations in the initial mass function with mass \citep{2007MNRAS.375..673K}, metal-poor gas infall \citep{2008MNRAS.385.2181F}, outflows or accretion of gas \citep{2004ApJ...613..898T, 2005MNRAS.362...41G, 2007MNRAS.376.1465K} , selective increase of the star-formation efficiency with the stellar mass \citep[also known as \emph{downsizing},][]{2007ApJ...655L..17B, 2008ApJ...672L.107E, 2009A&A...504..373C, 2009MNRAS.396L..71V}, or a combination of these scenarios. The second key property is the existence of radial gradients in the gas-phase metallicity in star-forming galaxies. This too is a long-known and often-studied problem \citep[e.g.,][]{1994ApJ...420...87Z,2010ApJS..190..233M, 2012ApJ...745...66M, 2015MNRAS.454.3664B, 2015MNRAS.451..210C, 2014A&A...563A..49S,2016A&A...587A..70S}. The MZR and radial metallicity gradients are usually analyzed separately, even though the physical and dynamical processes that have established them could be strongly related. In the present paper we are exploring the possibility that both the MZR and the radial metallicity gradients have a common origin that arises from a local empirical relationship between the metallicity and the local stellar surface mass density in the disk. This possibility is suggested by the strong positive correlation between the effective surface mass density and stellar mass \citep[e.g., ][]{2003MNRAS.341...33K} and a strong systematic decrease in the stellar surface mass density with increasing radius \citep[e.g., ][]{2015ApJ...800..120Z}. The relationship between local metallicity and local surface mass density has been studied before. These studies reported that HII regions with larger stellar densities are more metal rich in comparison to those with lower densities \citep{1992MNRAS.259..121V,1984MNRAS.211..507E}. Later, \cite{2012ApJ...745...66M} showed a trend between the surface mass density and the metallicity in their sample of 174 star-forming galaxies using long-slit spectroscopy. This was further considered and verified by \cite{2015MNRAS.451..210C} for a sample of HI-rich galaxies. Recently, the integral-field spectroscopy (IFS) technique has become possible, allowing the analysis of the spatially resolved properties in relatively large samples of disk galaxies. \cite{2012ApJ...756L..31R} demonstrated the local existence of the surface-mass density vs. metallicity relation using IFS data from HII regions in 38 spiral galaxies. \cite{2013A&A...554A..58S} confirmed this local relation using a sample of 150 star forming galaxies included in the CALIFA survey \citep{2012A&A...538A...8S}. The above IFS studies also find a relation between the surface mass density and the specific star formation rate. Despite these efforts, the samples of these IFS surveys are rather small in size and cover a limited range of global galaxy properties. For instance, the stellar mass coverage of the CALIFA survey is complete for nearby galaxies with 9.5 < $\log(\mathrm{M}_*/M_{\odot})$ < 11.0 \citep{2014A&A...569A...1W}. A larger sample of galaxies covering a wider range of physical parameters will allow us to understand whether the relations derived locally are independent of physical parameters that can affect a galaxy as a whole. The MaNGA survey \citep[Mapping Nearby Galaxies at APO,][]{2015ApJ...798....7B} is ideal for probing the impact of global parameters on the local relations. This on-going survey aims to observe in the next 5 years around 10000 galaxies using the IFS technique. Currently more than 1300 galaxies have been observed, allowing us to explore spatially resolved information for a sample of disk galaxies that covers a wide range of stellar masses and star formation rates. By utilizing this large sample and the wide range it covers in galaxy properties, the present paper is aimed at testing a simple hypothesis: the local gas-phase metallicity at a given location in a disk galaxy is determined mainly by the local stellar surface mass density, irrespective of the mass or the star-formation rate of the galaxy in which it resides. If so, this would imply that at least some global scaling relations in galaxies can have a local origin. This article is organized as follows. In Sec.~\ref{sec:SampleCubes} we describe the main properties of the sample of disk galaxies as well as the data used in this study. In Sec.~\ref{sec:Analisis} we present the derivation of the surface mass density, metallicity, and the selection of the star-forming spaxels to be analyzed. We will then present the \emph{local} surface mass density vs. metallicity relation for the MaNGA disk galaxies and compare it to previous IFS surveys. We will examine the residuals in this relation as a function of the total stellar mass and sSFR (Secs.~\ref{sec:muZ_Mtot}) as well as the metallicity profiles derived from this relation ~\ref{sec:dOH_grad}). In Sec.~\ref{sec:Discussion} we discuss the results of this study. Our main results and conclusions are presented in Sec.\ref{sec:Summary}. We adopt $H_0$~=~70~km~s$^{-1}$~Mpc$^{-1}$,~$\Omega_M$~=~0.3 and $\Omega_\Lambda$~=~0.7. | \label{sec:Discussion} We have confirmed that for our sample of disk galaxies the local oxygen abundance has a tight correlation with the local stellar surface mass density (see Fig.~\ref{fig:mu-OH-ssfr}). This is in agreement with previous IFS surveys \citep[see Fig.~\ref{fig:muZ-Surveys}, ][]{2012ApJ...756L..31R,2013A&A...554A..58S}. Our results confirm that the well-studied MZR is a scaled-up version of the local $\Sigma_*$-Z relation. We have shown that the $\Sigma_*$ vs. Z relation can also explain the observed radial metallicity gradients. These results suggest that both the MZR relation and radial gradients are reflections of a more fundamental and intimate relation between local properties within disk galaxies. The much larger body of data provided by the MaNGA survey compared to prior IFS surveys makes it suitable for further analysis of the impact of global properties of the galaxy on the metallicity derived at local scales. Our results indicate that local oxygen abundance can be described by means of the best-fit of the $\Sigma_*$-Z relation over wide ranges of stellar mass and sSFR (covering almost two orders of magnitude in each parameter, see Fig~\ref{fig:difOH_sfr}). These results indicate that to first order the metal content of a disk galaxy depends more strongly on the local properties than on those of the galaxy as a whole. Similarly, in Sec.~\ref{sec:dOH_grad} we analyzed the differences between the observed metallicity gradients and those derived using the radial surface mass density gradients and the $\Sigma_*$-Z relation for galaxies in different stellar mass bins. Except for the lowest-mass galaxies ($\log(\mathrm{M}_*/M_{\odot})$ < 9.2), the derived radial metallicity gradients are in good agreement (see Fig.~\ref{fig:difOH_grad}). An interplay of different processes has been proposed to explain the chemical evolution of galaxies, including both the MZR and radial metallicity gradients \citep[e.g.,][]{2012MNRAS.421...98D,2013ApJ...772..119L}. The simplest case is the closed-box model in which the metal content of disk galaxies depends only on the metal yield ($y_Z$) and the gas mass fraction ($f_{\mathrm{gas}} = M_{\mathrm{gas}}/[M_* + M_{\mathrm{gas}}]$): $Z = y_Z\, \ln (f_{\mathrm{gas}})$. There is a strong observed inverse correlation between the local values for $f_{\mathrm{gas}}$ and $\Sigma_*$, and this then implies an observed strong inverse correlation between $f_{\mathrm{gas}}$ and $Z$ \citep{2015MNRAS.451..210C}. Thus, part of the $\Sigma_*$ vs. Z relation can be explained by the fact that the lower density regions are more gas rich and therefore simply less chemically-evolved. This is consistent with the standard ``inside-out'' model for the formation of disk galaxies in which the inner (denser) regions form earlier than the outer (less dense) regions which in turn are more chemically-evolved and less gas-rich. It is also consistent with the downsizing phenomenon in which the (progenitors of) more (less) massive galaxies form earlier (later). According to integrated observations, the metallicity should also be regulated by the gas accretion of metal-poor gas and/or ejection of metal rich material \citep[e.g., ][]{2004ApJ...613..898T,2005MNRAS.362...41G,2007MNRAS.376.1465K,2010MNRAS.408.2115M,2010A&A...521L..53L}. Indeed, we observed that the residuals in the $\Sigma_*$ vs. Z at the lowest stellar masses (and highest values of sSFR) may reflect the effect of outflows, given that such galaxies have shallow potential wells and may therefore be less able to retain the metals they produce \citep[e.g.,][]{1974MNRAS.169..229L}. These ideas have recently been considered in the context of two related scenarios for the origin of radial metallicity gradients. \cite{2015MNRAS.451..210C} proposed a scaled-down radial version of the reservoir model \citep[e.g., ][]{2013ApJ...772..119L} in which the metal content in each radial bin of a disk galaxy is regulated by local inflows of metal-poor gas from the halo, the metal rich material carried out of the galaxy by feedback, and the metals produced by massive stars (the stellar yield). With this simple model they were able to fit with relatively good agreement the radial metallicity profiles of 50 nearby galaxies with different contents of molecular gas \citep{2014MNRAS.441.2159W}. Similarly, \cite{2015MNRAS.448.2030H} suggest that metallicity gradients are the result of the coevolution of the stellar and gas disk in a virtual closed-box chemical evolution model. In our next paper, we will explore in more detail the inter-dependences between all the basic {\it local} properties of disk galaxies: the metallicity, the stellar surface-mass density, the escape velocity, the specific SFR, and the luminosity-weighted mean age. This in turn will allow us to gain a better understanding of the physical origin of the global mass-metallicity relation and radial metallicity gradients, and of the implications for galaxy evolution. We have used data from the MaNGA survey to study the role of the {\it local} stellar surface mass density $\Sigma_*$ in determining both the {\it local}metallicity relation and radial metallicity gradients in disk galaxies. While the relationship between $\Sigma_*$ and metallicity (Z) has been known for some time \citep[e.g.,][]{2012ApJ...745...66M, 2012ApJ...756L..31R, 2013A&A...554A..58S}, the much larger data set provided by MaNGA allows us to assess whether this relationship is a universal one, or whether instead the residuals in the relationship correlate systematically with the global properties of galaxies. We found: \begin{itemize} \item In agreement with these earlier studies, the metallicity increases steeply with increasing surface mass density over the range $\Sigma_* \sim 10^{0.5}$ to $\sim 10^{2} M_{\odot}$ pc $^{-2}$. The relation then flattens at higher densities ($\Sigma_* \sim 10^2$ to $10^3 M_{\odot}$ pc$^{-2}$), reaching roughly solar metallicity. The {\it rms} residuals around the best fit relation are about $\pm$0.08 dex. \item The loci of the local values of $\Sigma_*$ and Z as measured in galaxies of different global stellar masses ($M_*$) systematically populate different parts of this relationship. As we move from low-mass to high-mass galaxies, the values for $\Sigma_*$ and $Z$ systematically increase: while all the data points lie along the same $\Sigma_*$ - Z relationship, low (high) mass disk galaxies are preferentially composed of regions of low (high) surface mass density and metallicity. \item Similarly, the loci of the local values of $\Sigma_*$ and Z as measured in galaxies of different global specific star-formation rates (sSFR) systematically populate different parts of this relationship. As we move from galaxies with high to low global sSFR, the values for $\Sigma_*$ and $Z$ systematically increase: while all the data points lie along the same $\Sigma_*$ - Z relationship, disk galaxies with high (low) global sSFR are preferentially composed of regions of low (high) surface mass density and metallicity. \item There are only weak systematic relationships between the residuals with respect to the best fit to the $\Sigma_*$ - Z relationship and either the global $M_*$ or sSFR. The largest systematic residuals are for the galaxies with very low $M_*$ ($\sim$ -0.08 dex below $\sim 10^{9.2} M_*$) or very high sSFR ($\sim$ -0.08 dex above sSFR $\sim 10^{-9.2}$ yr$^{-1}$). \item Using the best-fit to the $\Sigma_*$ - Z relation and the measured radial profiles of $\Sigma_{*}(r)$, we can accurately reproduce the observed radial metallicity gradients in disk galaxies spanning the range in global M$_*$ from $10^9$ to $10^{11}$ M$_{\odot}$. The worst disagreement ($\sim -0.1$ dex) was found in the inner regions ($< R_{50}$) in the lowest mass galaxies ($< 10^{9.5}$ M$_{\odot}$). \end{itemize} The critical point is that a single local relationship between stellar surface mass density and gas-phase metallicity can simultaneously reproduce two different systematic properties of disk galaxies: their mass-metallicity relation and their radial metallicity gradients. These reproductions are not perfect over the full range of galaxy properties. In the next paper in this series we will explore the nature of possible 'secondary parameters' that could affect metallicity. Put in the simplest possible terms, the results here show that the local gas-phase metallicity is very well determined by the local stellar surface-mass density: the outer disk of a massive galaxy and the inner regions of a low-mass galaxy (with the same density) have the same metallicity. This is a rather remarkable result in terms of the way in which galaxies assemble. | 16 | 9 | 1609.01740 |
1609 | 1609.01095_arXiv.txt | { \srcs\ is a gravitationally lensed blazar located at a redshift of 0.944. The gravitational lensing splits the emitted radiation into two components, spatially indistinguishable by gamma-ray instruments, but separated by a 10-12 day delay. In July 2014, \srcs\ experienced a violent flare observed by the \fermilat\ and followed by the MAGIC telescopes. } { The spectral energy distribution of \srcs\ can give information on the energetics of $z \sim 1$ very high energy gamma-ray sources. Moreover the gamma-ray emission can also be used as a probe of the extragalactic background light at $z \sim 1$. } { MAGIC performed observations of \srcs\ during the expected arrival time of the delayed component of the emission. The MAGIC and \fermilat\ observations were accompanied by quasi-simultaneous optical data from the KVA telescope and X-ray observations by \swift . We construct a multiwavelength spectral energy distribution of \srcs\ and use it to model the source. The GeV and sub-TeV data, obtained by \fermilat\ and MAGIC, are used to set constraints on the extragalactic background light. } { Very high energy gamma-ray emission was detected from the direction of \srcs\ by the MAGIC telescopes during the expected time of arrival of the trailing component of the flare, making it the farthest very high energy gamma-ray sources detected to date. The observed emission spans the energy range from 65 to 175 GeV. The combined MAGIC and \fermilat\ spectral energy distribution of \srcs\ is consistent with current extragalactic background light models. The broad band emission can be modeled in the framework of a two zone external Compton scenario, where the GeV emission comes from an emission region in the jet, located outside the broad line region. } {} | Even though there are already over 60 blazars detected in the very high energy (VHE, $\gtrsim100\,$GeV) range, most of them are relatively close-by sources with redshift $z\lesssim0.5$. Until mid 2014, the farthest sources observed in this energy range were 3C\,279 ($z=0.536$, \citealp{al08}), KUV\,00311-1938 ($z>0.506$, \citealp{be12}) and PKS1424+240 ($z=0.601$, \citealp{acc10}). In the last two years the MAGIC (Major Atmospheric Gamma Imaging Cherenkov) telescopes discovered VHE gamma-ray emission from \srcs\ at $z=0.944$ \citep{atel6349} and afterwards PKS1441+25 at $z=0.940$ \citep{ah15} almost doubling the boundaries of the known gamma-ray universe. Observations of distant sources in VHE gamma-rays are difficult due to strong absorption in the extragalactic background light (EBL, see e.g. \citealp{gs66}). At a redshift of $\sim 1$ it results in a cut-off at energies\footnote{Unless specified otherwise, the energies are given in the Earth's frame of reference} $\sim 100\,$GeV. Such energies are at the lower edge of the sensitivity range of the current generation of Imaging Atmospheric Cherenkov Telescopes (IACTs), making such observations challenging. To maximize the chance of detection, the observations are often triggered by a high state observed in lower energy ranges. In particular, \fermilat\ (Large Area Telescope) scanning the whole sky every 3 hours provides alerts on sources with high fluxes and information about the spectral shape of the emission in the GeV range. \srcs , also known as S3\,0218+35, is classified as a flat spectrum radio quasar (FSRQ, \citealp{ac11}). The classification is based on the optical spectrum \citep{co03}. It is located at a redshift of $z_s=0.944\pm0.002$ \citep{co03}. % One of the five features from which \citet{co03} derived the redshift was confirmed by \citep{la15}. The object is gravitationally lensed by the face-on spiral galaxy \srclens\ located at a redshift of $z_l=0.68466\pm0.00004$ \citep{cry93}. Strong gravitational lensing forms multiple images of the source \citep[see e.g.][]{ko04}. The flux magnification of an image is the ratio of the number of photons gravitationally deflected into a small solid angle centered on the observer to the number of photons emitted by the source in such a solid angle. The 22.4 GHz VLA radio image shows two distinct components with an angular separation of only 335\,mas and an Einstein ring of a similar size \citep{od92}. Observations of variability of the two radio components led to a measurement of a delay of 10-12 days between the leading and trailing images \citep{co96,bi99, co00, em11}. In the radio image, the leading component (also dubbed `image A' in literature) is located to the west from the trailing component (image B). The delayed component had a 3.57-3.73 times weaker flux \citep{bi99}. However, the observed ratio of magnification varies with the radio frequency \citep{mi06}, presumably due to free-free absorption in the lensing galaxy \citep{mi07}. In the optical range the leading image is strongly absorbed \citep{fa99}. In 2012 \srcs\ went through a series of outbursts registered by the \fermilat\ \citep{ch14}. Even though \fermilat\ does not have the necessary angular resolution to disentangle the two emission components, the statistical analysis of the light curve auto-correlation function led to a measurement of a time delay of $11.46\pm0.16$ days. Interestingly the average magnification factor, contrary to radio measurements, was estimated to be $\sim1$. Changes in the observed GeV magnification ratio were interpreted as microlensing effects on individual stars in the lensing galaxy \citep{vn15}. Microlensing on larger scale structures has been considered as well \citep{sb16}. The radio follow-up observations of \srcs\ after the 2012 gamma-ray flare did not reveal any correlation between the two bands \citep{sp16}. Another flaring state of \srcs\ was observed by \fermilat\ on 2014 July 13 and 14 \citep{atel6316}. Contrary to the results for the 2012 flaring period, in this case the ratio of the leading to delayed GeV emission was at least 4 \citep{bu15}. The 2014 flare triggered follow-up observations by the MAGIC telescopes, which in turn led to the discovery of VHE gamma-ray emission from \srcs\ \citep{atel6349}. In this work we present the results of the observations by the MAGIC telescopes and supporting multiwavelength instruments of the \srcs\ during the flaring state in July 2014. In Section~\ref{sec:inst} we describe the instruments taking part in those observations and the data reduction. The effect of the lensing galaxy on the observed emission is discussed in Section~\ref{sec:lens}. Section~\ref{sec:results} is devoted to the results of the observations. In Section~\ref{sec:sed} we model the spectral energy distribution (SED) of the source. We use the \fermilat\ and MAGIC observations to discuss constraints on the EBL in Section~\ref{sec:ebl}. | MAGIC has detected VHE gamma-ray emission from \srcs\ during the trailing component of a flare in July 2014. It is currently the most distant source detected with a ground-based gamma-ray telescope, and the only gravitationally lensed source detected in VHE gamma-rays. The VHE gamma-ray emission lasted for two nights achieving the observed flux of $\sim 30\%$ of Crab Nebula at 100\,GeV. Using the EBL model from \citet{do11}, the intrinsic spectral index in this energy range was found to be $2.35\pm0.75_{\rm stat} \pm 0.20_{\rm syst}$. The VHE gamma-ray flare was not accompanied by a simultaneous flux increase in the optical or X-ray energy range. We have modeled the X-ray emission as a sum of two components with different magnifications, the weaker one absorbed with column density of $(2.4\pm0.5)\times 10^{22}\ \mathrm{at.\,cm^{-2}}$. The combined \fermilat\ and MAGIC energy spectrum is consistent with the current EBL models. These constraints are however not very strong, with the EBL density scaling parameter being less than 2.1-2.8 of the one predicted by the tested models. The broadband emission of \srcs\ is modeled in a framework of a two-zone external Compton model. According to this scenario, the quasi-stable optical and X-ray emission originates mostly in the inner zone. The enhanced gamma-ray emission during the flare is produced in the second zone, located outside of the BLR. | 16 | 9 | 1609.01095 |
1609 | 1609.04954_arXiv.txt | {Most of the quantitative information about the magnetic field vector in solar prominences comes from the analysis of the Hanle effect acting on lines formed by scattering. As these lines can be of non-negligible optical thickness, it is of interest to study the line formation process further.} {We investigate the multidimensional effects on the interpretation of spectropolarimetric observations, particularly on the inference of the magnetic field vector. We do this by analyzing the differences between multidimensional models, which involve fully self-consistent radiative transfer computations in the presence of spatial inhomogeneities and velocity fields, and those which rely on simple one-dimensional geometry.} {We study the formation of a prototype line in ad hoc inhomogeneous, isothermal 2D prominence models. We solve the NLTE polarized line formation problem in the presence of a large-scale oriented magnetic field. The resulting polarized line profiles are then interpreted (i.e. inverted) assuming a simple 1D slab model. } {We find that differences between input and the inferred magnetic field vector are non-negligible. Namely, we almost universally find that the inferred field is weaker and more horizontal than the input field.} {Spatial inhomogeneities and radiative transfer have a strong effect on scattering line polarization in the optically thick lines. In real-life situations, ignoring these effects could lead to a serious misinterpretation of spectropolarimetric observations of chromospheric objects such as prominences.} | Magnetic fields play a crucial role in the physics of the outer solar atmosphere. Among other phenomena, magnetic effects are the origin of the formation and evolution of various chromospheric structures such as solar prominences (filaments) and spicules. To understand these objects it is essential to compare the magnetic field vector inferred from the observations with the magnetic field resulting from theoretical studies. The main sources of quantitative information on the magnetic field pervading these objects are spectral lines sensitive to the Hanle and Zeeman effects. The 10830$\,\rm{\AA}$ and D3 (5876$\,\rm{\AA}$) lines of helium are among the most frequently used spectral lines for such an analysis. To interpret the observations performed in these lines, different diagnostic tools have been devised, ranging from simple polarization diagrams \citep[e.g.][]{LL04} to modern inversion codes such as HAZEL \citep{HAZEL} and HELIX \citep{HELIXX}. The magnetic field is not an observable quantity; it is indirectly inferred from the observed Stokes vector. The interpretation of any spectropolarimetric observation is thus done assuming a specific generative model. The observed data are values of Stokes parameters for a given set of wavelengths and their errors are usually Gaussian and follow from the signal-to-noise ratio. The generative model is inevitably simple and relies either on what is known as a single-scattering approximation or on a 1D slab model, which can also involve some elementary 1D radiative transfer. Although these generative models are oversimplified when compared to the ``real'' Sun, we are, in a way, condemned to use them because more sophisticated radiative transfer calculations are computationally very demanding and inversion techniques based on self-consistent multidimensional polarized NLTE radiative transfer are still unfeasible. In this paper we investigate the bias that emerges when a simple 1D model is used to interpret the observations that are, in fact, generated in a more complex (and more realistic) multidimensional, dynamic, and inhomogeneous medium. We study the polarized spectral line formation process in a 2D Cartesian slab standing above the solar surface and illuminated by the solar radiation. The slab is inhomogeneous in density and pervaded by different velocity fields and by a uniform vector magnetic field. We consider a prototype line, formed in the two-level approximation (normal Zeeman triplet), under the assumption of complete frequency redistribution. These assumptions simplify the atomic physics involved and allow us to focus on ``macroscopic'' effects on the line formation process. We compute synthetic spectra from these models and then use a 1D scheme to ``invert'' these synthetic observations. We discuss the differences between the input parameters and the inferred ones and point to possible systematic errors and degeneracies in the process of the inference of the magnetic field vector in such objects. | In this paper we have studied the formation of a prototype two-level atom line in a prominence-like structure situated above the solar surface. The structure is finite and inhomogeneous in the $xy$ plane and infinite and homogeneous in the $z$ direction. The line is assumed to have atomic parameters similar to the He\,10830 line. Taking into account the weak field Zeeman effect and the Hanle effect, we have computed a full Stokes vector emerging from the prominence. We have then applied spatial smearing and binning to obtain synthetic spectra. We first conclude that the presence of spatial inhomogeneities induces the presence of linear polarization in Stokes $U$ even in the absence of a magnetic field. The magnitude of this polarization is comparable to the magnitude of Stokes $Q$. In the presence of a magnetic field, the interplay of the Hanle effect and of the radiative transfer effects on the resulting Stokes profiles is qualitatively hard to interpret and therefore we have conducted an inversion of the synthetic spectra using an affine-invariant MCMC scheme. The model used for the inversion is identical to the one used for the forward synthesis except for the fact that it only accounts for 1D radiative transfer. The magnetic field vector retrieved by our inversion procedure is more directed toward the observer (i.e. the polar angle is closer to $\pi/2$ and the azimuth is closer to zero) with respect to the input magnetic field. Also, the magnitude of the inferred magnetic field is significantly lower ($\approx 2$ versus the original $5$\,gauss). We argue that this is because of the multidimensional and radiative transfer effects which cannot be explained by a simple 1D inversion procedure. We draw a parallel with the real Sun where the 1D inversion schemes are used to interpret line profiles formed in multidimensional environments. We also briefly comment on the influence of systematic velocities and photon noise on the spectra. Moderate, random velocities do not seem to influence the spectrum significantly. The presence of noise, on the other hand, has a strong impact on the inference process. Even the relatively low level of noise we have used here causes broad and asymmetric posterior probabilities of the relevant parameters. It is obvious that prominences are more complex than depicted here. They are, undoubtedly, inhomogeneous along the $z$-axis as well. These are not only inhomogeneities in physical parameters, but also in boundary conditions as lower parts of the prominence ``see'' stronger and less anisotropic radiation. It would be of great importance to study the effects of vertical radiative transfer in this case as the vertical geometrical scales of the problem are similar to the horizontal scales. The most thorough approach would be to synthesize a full Stokes spectrum of realistic spectral lines, coming from a 3D model, and then invert it using state-of-the-art inversion codes (HAZEL, HELIX+). This paper is only the first step in the investigations intended to highlight the problems with using the contemporary inversion techniques in the study of spectral lines formed in non-LTE. We conclude this paper, which we hope was instructive, by asking a question: How relevant is this radiative transfer exercise to the interpretation of realistic solar observations? It is tempting to ascribe the dominance of the horizontal fields found in the interpretation of prominence observations to these radiative transfer effects. In the case where the observed line is optically thin, the line formation is nicely described by a single-scattering approximation, and multidimensional effects are negligible (e.g. the D3 line of helium). Inference methods that rely on multi-line observations are also probably immune to these effects. However, for cases where a single, optically thick line is used (e.g. spicules, some prominences), it is possible to encounter line profiles that require detailed radiative transfer modeling. Unfortunately, inversion schemes for the lines formed by scattering which involve multidimensional radiative transfer are (probably) still far away. Detailed, multidimensional NLTE modeling is our best hope to understand the physical conditions in the medium where these lines are formed, both in the chromosphere and in various chromospheric and transition-region structures. | 16 | 9 | 1609.04954 |
1609 | 1609.03329_arXiv.txt | {The past decade has witnessed a large number of Galactic plane surveys at angular resolutions below $20''$. However, no comparable high-resolution survey exists at long radio wavelengths around 21\,cm in line and continuum emission.} {We remedy this situation by studying the northern Galactic plane at $\sim$20$''$ resolution in emission of atomic, molecular, and ionized gas.} {Employing the Karl G.~Jansky Very Large Array (VLA) in the C-array configuration and a large program, we observe the HI 21\,cm line, four OH lines, nineteen H$n\alpha$ radio recombination lines as well as the continuum emission from 1 to 2\,GHz in full polarization over a large part of the first Galactic quadrant.} {Covering Galactic longitudes from 14.5 to 67.4\,deg and latitudes between $\pm 1.25$\,deg, we image all of these lines and the continuum at $\sim$20$''$ resolution. These data allow us to study the various components of the interstellar medium (ISM): from the atomic phase, traced by the HI line, to the molecular phase, observed by the OH transitions, to the ionized medium, revealed by the cm continuum and the H$n\alpha$ radio recombination lines. Furthermore, the polarized continuum emission enables magnetic field studies. In this overview paper, we discuss the survey outline and present the first data release as well as early results from the different datasets. We now release the first half of the survey; the second half will follow later after the ongoing data processing has been completed. The data in fits format (continuum images and line data cubes) can be accessed through the project web-page http://www.mpia.de/thor.} {The HI/OH/Recombination line survey of the Milky Way (THOR){} opens a new window to the different parts of the ISM. It enables detailed studies of molecular cloud formation, conversion of atomic to molecular gas, and feedback from H{\sc ii} regions as well as the magnetic field in the Milky Way. It is highly complementary to other surveys of our Galaxy, and comparing the different datasets will allow us to address many open questions.} | \label{intro} Over the past decade, the Galactic plane was surveyed comprehensively from near-infrared to cm wavelengths. These surveys enable investigations of not only individual local phenomena such as stars, clusters, ionized gas and molecular or atomic clouds, but studies of our Galaxy as a whole, and we can compare the results to extragalactic studies (see, e.g., \citealt{taylor2003,churchwell2009,carey2009,schuller2009,anderson2011,walsh2011,beuther2012b,ragan2014,wang2015,goodman2014,reid2014,abreu2016}). Particularly important for a general understanding of the different physical processes is the multiwavelength approach because different surveys trace different components of the interstellar medium (ISM) and stellar populations, as well as varying temperature regimes and physical processes. Earlier ideas for such a multiwavelength survey approach were promoted by the Canadian Galactic Plane Survey \citep{taylor2003}, for example. The different phases (atomic, molecular, or ionized gas and dust) are not isolated, but interact and, maybe even more importantly, they change from one phase to the other in the natural matter cycle of the ISM. It is therefore important for our understanding of ISM dynamics and star formation to have surveys at comparable angular resolution. While most of the infrared to mm Galactic plane surveys have an angular resolution better than $20''$, the existing HI Very Large Array Galactic Plane Survey (VGPS) survey conducted with the Very Large Array (VLA) in its compact D-configuration has an angular resolution of only $60''$ \citep{stil2006}. For comparison, the most recent single-dish survey of HI with the Effelsberg telescope has an angular resolution of $10'$ \citep{winkel2016}. Even though the VLA D-configuration as well as single-dish HI surveys are appropriate for studying atomic Galactic structure on large scales, they are less useful for the direct comparison with the other existing surveys mentioned above. For example, previous 60$''$ resolution observations of HI and CO emissions lines showed that large-scale atomic gas envelopes and atomic gas flows in the surrounding environments are needed to form denser molecular gas, and subsequently dense core and massive stars (e.g., \citealt{nguyen2011,motte2014}). However, these data could not yet be used to study the interaction between the atomic and the dense molecular gas structures that may occur on significantly smaller scales (see, e.g., the recent 870\,$\mu$m dust continuum emission Galactic plane survey ATLASGAL at $19''$ resolution, \citealt{schuller2009}). For reference, we mention that 0.5\,pc corresponds to $25''$ at a typical molecular cloud distance of 4\,kpc. Furthermore, the new capabilities of the WIDAR correlator at the VLA allow us to observe many spectral lines simultaneously, in particular several molecular OH transitions, a series of H$n\alpha$ radio recombination lines (RRLs, $n=151 \ldots 186$), and the continuum emission. Combining these data with the HI observations probes the transition of matter in the ISM from the diffuse neutral atomic to the dense molecular and the ionized gas components and back. This combined approach is followed in The HI/OH/Recombination line survey of the Milky Way (THOR) we present here. These new THOR C-configuration HI data ($15''-20''$ resolution corresponding to linear scales of 0.2-0.3\,pc at typical distances of 3\,kpc), when combined with the existing D-configuration and GBT (Green Bank Telescope) observations to include the larger-scale emission \citep{stil2006}, enable us to address many questions associated with atomic hydrogen from large-scale Galactic structure and cloud formation processes down to the scales of individual star-forming regions. At the same time, the OH, RRLs, and continuum data provide a more complete picture of the Galactic ISM. This paper presents an overview of THOR and the first data release. The motivation and goals of the survey are described in Sect. 2, and the parameters of the survey are presented in Sect. 3. The observation details and data analysis are given in Sect. 4, while initial results from this survey are presented in Sect. 5. Finally, the potential of this survey and the future possibilities are discussed in Sect. 6, and a summary is presented in Sect. 7. | Based on the early results presented in the previous section, we envision a multitude of future scientific applications. The advantage of THOR is that we do not have to rely on single case studies but that larger statistical approaches are possible. For example, the HI study of the W43 cloud presented in \citet{bihr2015b} will be extended to many clouds within the Milky Way. Similarly, feedback studies as indicated by the M17 data in Sect. \ref{rrl} will be extended to the whole sample of detected H{\sc ii} region in the radio recombination line emission. In a different application, the HI data enable studying the density fluctuation structure function of the CNM down to a few tens of arcsec scales. This structure function is related to the ISM turbulence and can be directly compared to different theoretical models to constrain turbulence and energy dissipation mechanisms (e.g., \citealt{dickey2001,audit2010}). \begin{figure}[htb] \includegraphics[width=0.49\textwidth]{RMSF.png} \caption{RMSF of THOR has a central lobe with FWHM $\sim$70\,rad\,m$^{-2}$ with near side lobes at the 40\% level around $\pm$200\,rad\,m$^{-2}$. The largest detectable Faraday depth scale is $\sim$130\,rad\,m$^{-2}$.} \label{RMSF-fig} \end{figure} \begin{figure}[htb] \includegraphics[width=0.49\textwidth]{Kes75_PWN.png} \caption{Integrated Faraday depth spectrum of the pulsar wind nebula at the center of the SNR Kes75 ($l=29.7104^{\circ}$ \& $b=-0.2402^{\circ}$). The vertical axis represents the degree of polarization as a percentage of the total flux density, after unwrapping Faraday rotation assuming Faraday depth $\phi$ and averaging over the observed frequency range. The green and red lines indicate the maximum over all Faraday depths of 100 and 300 realizations of the analysis, respectively, which replace the $Q$ and $U$ spectra of the target by noise spectra from off-source positions, integrated over the same solid angle as the target. We detect components at Faraday depth $\phi_1 = 60$ rad m$^{-2}$ and Faraday depth $\phi_2 = 258$ rad m$^{-2}$. The formal error in Faraday depth is $\sim$5 rad m$^{-2}$ for a 10$\sigma$ detection.} \label{kes75-fig} \end{figure} For OH masers and absorption studies, THOR provides the perfect dataset. However, for thermal OH emission, the extended structures are filtered out by our C-configuration observations. We are currently exploring whether complementing the THOR OH data with short spacing from the SPLASH survey \citep{dawson2014} or complementary Effelsberg/GBT observations is sufficient, or if the shorter baselines from the VLA in D-configuration are needed for such an aspect of the ISM studies. Similarly, the continuum data allow us to derive spectral indices for compact structures \citep{bihr2016}, but spectral indices for more extended sources such as SNRs are much harder to determine from THOR data alone. Therefore, we are currently examining whether single-dish short spacings are sufficient for the science goals related to the continuum emission in the survey, or if D-configuration data may be needed. We currently merely scratch the surface of the polarization aspect of the THOR survey. Since we have observed the full polarization, Faraday rotation and magnetic field studies of the Milky Way will be feasible. However, the data calibration, imaging, and analysis aspect of this part of the survey have yet to be realized, and therefore, polarization and magnetic field studies will be presented in forthcoming publications (e.g., Stil et al.~in prep.). In addition to the THOR data as a stand-alone survey, it will obviously be important to combine THOR with existing surveys at other wavelengths. Only then we will be able to address all facets of the Milky Way in its appropriate depth. Understanding Galactically important regions such as the bar--spiral arm interface can directly be compared with extragalactic studies (e.g., THINGS, \citealt{walter2008}) and thus be set into a global context. The combination of Galactic and extragalactic systems allows us to derive a concise and complete picture of the ISM and star formation processes. Furthermore, THOR can also be considered as a precursor of Square Kilometer Array (SKA) pathfinder science because the planed GASKAP survey (The Galactic ASKAP survey) with the Australian SKA Pathfinder telescope will achieve comparable sensitivities and angular resolution elements in the southern hemisphere \citep{dickey2013}. | 16 | 9 | 1609.03329 |
1609 | 1609.04509_arXiv.txt | We confirm that the object DDO216-A1 is a substantial globular cluster at the center of Local Group galaxy DDO216 (the Pegasus dwarf irregular), using Hubble Space Telescope ACS imaging. By fitting isochrones, we find the cluster metallicity [M/H] = $-$1.6 $\pm$0.2, for reddening E(B$-$V) = 0.16$\pm$0.02; the best-fit age is 12.3$\pm$0.8~Gyr. There are $\approx$30 RR~Lyrae variables in the cluster; the magnitude of the fundamental mode pulsators gives a distance modulus of 24.77 $\pm$0.08 --- identical to the host galaxy. The ratio of overtone to fundamental mode variables and their mean periods make DDO216-A1 an Oosterhoff Type I cluster. We find a central surface brightness 20.85$\pm$0.17 F814W~mag~arcsec$^{-2}$, a half-light radius of 3$\farcs$1 (13.4~pc), and an absolute magnitude M$_{814}$ = $-$7.90 $\pm$0.16 (M/M$_{\sun}$ $\approx$10$^5$). King models fit to the cluster give the core radius and concentration index, r$_c$ = 2$\farcs$1$\pm$0$\farcs$9 and $c$ = 1.24$\pm$0.39. The cluster is an ``extended'' cluster somewhat typical of some dwarf galaxies and the outer halo of the Milky Way. The cluster is projected $\lesssim$30~pc south of the center of DDO216, unusually central compared to most dwarf galaxy globular clusters. Analytical models of dynamical friction and tidal destruction suggest that it probably formed at a larger distance, up to $\sim$1~kpc, and migrated inward. DDO216 has an unexceptional cluster specific frequency, S$_N$ = 10. DDO216 is the lowest-luminosity Local Group galaxy to host a 10$^5$~M$_{\odot}$ globular cluster, and the only transition-type (dSph/dIrr) in the Local Group with a globular. | \setcounter{footnote}{0} Dwarf galaxies (M$_* \lesssim 10^8 M_{\sun}$) are the most abundant class of galaxies in the Universe. They occupy an important part of parameter space for understanding the feedback processes that seem to govern the relationships between dark halo mass, baryon fraction, and star formation efficiency. Furthermore, their progenitors at high redshift may have played an important role in reionizing the Universe \citep{rob13}. Despite the ubiquity of dwarf galaxies, it is challenging to reliably measure their physical properties and put them in their appropriate cosmological context \citep[e.g.,][]{boy15}. Most dwarf galaxies are undetectable beyond redshift $z$ $\approx$1--2, even in the Hubble Ultra Deep Field or with planned JWST observations \citep{boy16}. Thus, observations of Local Group galaxies have set the benchmark for the accuracy and precision with which ancient star formation rates and chemical evolution histories can be measured \citep[e.g.,][and references therein]{ski14,col14,wei14}. Long-lived main-sequence stars born at lookback times $\geq$10~Gyr, corresponding to $z$ $\gtrsim$ 2, are a direct probe of galaxy evolution in the Universe from the earliest star-forming period, through the epoch of reionization and its aftermath. In this paper we present a photometric analysis of the understudied star cluster at the center of DDO~216 (the Pegasus dwarf irregular, PegDIG, UGC~12613), which we observed serendipitously during our HST program to measure the complete star formation history of this galaxy. PegDIG (M$_V$ = $-$12.5 $\pm$0.2) is roughly a magnitude fainter than the Fornax and Sagittarius dwarfs, which makes it one of the least luminous galaxies known to host a cluster near the peak of the globular cluster luminosity function \citep[See][for examples of similar clusters in dwarfs with M$_V$ $\approx -$11.5]{geo09a,dac09}. First, we review the basic parameters of the galaxy, and then describe our observations and reductions in section~\ref{sec:obs}. Our analysis of the structure and content of the star cluster DDO216-A1, including age and metallicity estimates based on the color-magnitude diagram and an analysis of the variable star population of the cluster, is given in section~\ref{sec:clus}. We place DDO216-A1 in context with the population of massive star clusters in dwarf galaxies and summarise our results in section~\ref{sec:summ}. \subsection{Globular Clusters in Dwarf Galaxies} \begin{figure*}[t!] \plotone{ACole_fig1.pdf} \caption{DDO216 and its globular cluster. {\it Left:} Coadded ACS/WFC image of DDO216 through filter F814W. The total exposure time is 34.3 ksec, spread across 4 days. The cluster DDO216-A1 is clearly visible, 6$\arcsec$ south of the galaxy center. The box around the cluster is 45$\arcsec$ ($\approx$200~pc) across. {\it Right:} Magnified view of the 45$\arcsec$ region around DDO216-A1, constructed from F475W and F814W images. The cluster is high-surface brightness, dominated by red giants and horizontal branch stars, and by eye appears to be $\sim$8-10$\arcsec$ ($\sim$35--44~pc) in diameter. \label{fig-image}} \end{figure*} Given deep enough observations, it is relatively straightforward to derive star formation rates at high lookback times for galaxies within the Local Group. It is more difficult to identify the triggers of star formation. For example, the majority of dwarf-dwarf major mergers are expected to have occured in the first few billion years after galaxy formation began \citep*[e.g.,][]{dea14}, but because of the destructive nature of mergers most of the obvious evidence for merger activity will have long since vanished. Globular clusters are an important window into this process because they require extreme conditions to form, suggestive of vigorous star formation in the mode often associated with galaxy mergers and interactions \citep[e.g.,][and references therein]{bro06}. Globular clusters are tightly bound and will typically survive for a Hubble time unless disrupted in a hostile tidal environment, but the relatively shallow potential wells of dwarf galaxies are not generally conducive to cluster disruption. As a result, many globular and open clusters are known in dwarf galaxies in the Local Group and beyond, in enough numbers to make statistical associations between properties like host galaxy morphology and cluster colours and sizes \citep*[e.g.,][]{sha05,mil07}. These span a range of sizes from extremely luminous and dense nuclear star clusters to low-mass open cluster or association analogues, in galaxies down to some of the least luminous known. In the Local Group, recent work has discovered examples of modest star clusters even among the smallest galaxies \citep[e.g.,][]{crn16}, and more massive, sometimes extended clusters in some of the larger irregular galaxies \citep{sha07,hwa11}. In light of these discoveries and others, \citet*{zar16} have suggested that some of the outer halo globular clusters thought to be Galactic globular clusters may in fact be hosted by undiscovered low surface brightness galaxies. However, the lowest-luminosity Local Group galaxies with catalogued globular clusters similar to the massive and dense globulars of the Milky Way are Fornax and Sagittarius (M$_V$ = $-$13.4 and $\lesssim -$13.5, respectively). \subsection{The Pegasus Dwarf Irregular, DDO 216} \label{sec:peg} The Pegasus dwarf irregular galaxy, PegDIG, was discovered by A.G. Wilson in the early 1950's on Palomar Schmidt plates \citep{hol58}. From early on it was considered to be a candidate member of the Local Group with a distance of $\sim$1~Mpc. The case for membership was supported by the negative heliocentric \ion{H}{1} radial velocity found by \citet{fis75}. PegDIG is considered a distant M31 satellite \citep[d$_{\mathrm M31}$ $\approx$470~kpc,][]{mcc07}; it is not proven that it has ever interacted with M31, although it is more likely than not that PegDIG has previously been within M31's virial radius \citep{sha13,gar14}. PegDIG is fairly isolated at the present time, its nearest neighbor being the M31 satellite And~VI, just over 200~kpc away. It is quite unlikely that PegDIG has had strong tidal interactions with any other known galaxy during the past several Gyr. PegDIG is a fairly typical small irregular galaxy, with roughly 1:1 gas to stellar mass ratio, although it has virtually no current star formation as measured by H$\alpha$ emission \citep{you03}. This leads to its classification as a transition type dwarf, with properties intermediate between the spheroidal (dSph) and irregular (dIrr) types \citep{mcc12}. It has an ordinary metallicity of [Fe/H] $\approx$ $-$1.4 $\pm$0.3 for its stellar mass of $\approx$10$^{7}$~M$_{\odot}$ \citep{kir13}. Unlike the spheroidal galaxies, it appears to be rotating; both HI \citep{kni09} and stellar \citep{kir14} data suggest a rotation speed (not corrected for inclination) of $\approx$15--20~km/s. \citet{mcc07} drew attention to the cometary appearance of the neutral gas, and attributed the asymmetric morphology to ram pressure stripping by diffuse gas in the Local Group, although this conclusion is disputed, based on much deeper HI observations, by \citet{kni09}. The star formation history of PegDIG has been estimated from ground-based data reaching a limiting absolute magnitude of M$_{\mathrm I}$ $\approx$ $-$2.5 \citep*{apa97}, and from HST/WFPC2 observations \citep{gal98} reaching $\approx$2.5 mag deeper. Within large uncertainties \citep{wei14a}, the picture that emerges from these studies is of star formation that has spanned a Hubble time, likely to be declining over time following an early epoch of high star formation rate (SFR). The SFR has certainly declined with time over the past $\approx$1--2~Gyr, despite the large reservoir of neutral gas. In its extended star formation history, PegDIG appears to have more in common with the dIrr galaxies than with a typical dSph, consistent with its retention of neutral gas to the present day and with the assertion of \citet*{ski03} that transition type galaxies represent the low-mass/low-SFR end of the dwarf irregular population \citep[see also][]{wei11}. The precise star formation history over the full lifetime of the galaxy will be determined in a future paper in this series (Cole et al., in preparation). | \label{sec:summ} \subsection{Rediscovery of a Cluster} The photometric and structural parameters of DDO216-A1 are summarised in Table~\ref{tab:props}. By every available measure, DDO216-A1 is a bona fide, ancient globular cluster, indistinguishable in many ways from the globular clusters of the Milky Way and Large Magellanic Cloud. While the cluster was first identified over 35 years ago, it has only sporadically been recorded in the literature \citep{gal98}. With an absolute magnitude M$_{\mathrm V}$ $\approx$ $-$7.1, DDO216-A1 contributes roughly 0.5\% of the V-band light of PegDIG-- this makes its lack of study all the more remarkable, given the long history of observational studies of PegDIG. The reasons for the omission of DDO216-A1 from lists of globular clusters in dwarf galaxies probably stem from a combination of the location of DDO216-A1, seen in projection against the densest part of PegDIG; the unusually extended nature of the cluster, which both gives it much less contrast with its surroundings and makes it somewhat unlike Galactic globular clusters; and the prevalence of background galaxies in the field. At a distance of 900~kpc, it requires diffraction limited imaging to resolve the cluster even to the level of the horizontal branch. The cluster is not prominently visible in Spitzer space telesope IRAC images \citep{jac06,boy09} due to its lack of asymptotic giant branch stars; with hindsight, it is obvious in SDSS images, although unresolved. \subsection{Cluster Mass} In the absence of spectroscopic information, the cluster mass can only be estimated in a model-dependent way or by comparison to better-studied, similarly luminous clusters. \citet{mcl05} compared the dynamical mass-to-light ratios for a large set of Milky Way and Magellanic Cloud globular clusters to predictions of population synthesis models by \citet{bru03}. Using their preferred initial mass function \citep{cha03}, they found a mean model M/L $\approx$1.9 for the old clusters. The median dynamical mass to light ratio for the clusters in their sample was 82$\pm$7\% of this value, with a substantial scatter. Our absolute magnitude for DDO216-A1 is M$_V$ = $-$7.14 $\pm$0.16, which translates to a V-band luminosity of (5.97 $\pm$0.95)$\times10^4$~L$_{\odot}$. Using the M/L estimates from \citet{mcl05} gives either 1.13$\pm$0.18 (population synthesis) or 0.93$\pm$0.15 (dynamical) $\times$10$^5$~M$_{\odot}$. The true range of possible values is even larger, because of potential variations in the initial mass function and the observed variations between clusters. Given the lack of kinematic constraints, an appropriate way to express the probable mass of the cluster is $\log$(M/M$_{\odot}$) = 5.0$\pm$0.1. This is entirely consistent with mass estimates for similarly bright and extended, old clusters in the Milky Way \citep[e.g., IC4499,][]{han11}, and the LMC \citep[e.g., Reticulum,][]{sun92}. \subsection{Formation, Migration, and Survival} The cluster's projected position near the center of PegDIG raises the question of its provenance and survival. If the cluster formed in situ at the center of the galaxy, then it is natural to ask how it has survived tidal evaporation for 12~Gyr. Alternatively, if DDO216-A1 formed at an arbitrary location in the galaxy, then dynamical friction must be acting efficiently enough to bring it nearly to the center within its lifetime. Survivability of clusters in dwarf galaxies can be calculated probabilistically using analytical models for dynamical friction and cluster evolution (Leaman et al., in preparation). Using the dynamical friction formula from \citet{pet16}, a range of galaxy mass profiles and cluster orbits can be tested to see if there are any plausible initial conditions conducive to cluster inspiral and survival. The half-light radius of PegDIG is $\approx$700~pc \citep{kir14} and its stellar mass is log(M$_*$) $\approx$10$^7$~M$_{\odot}$ \citep{mcc12}, but its mass profile is not extremely well known. To a first approximation it could be taken as similar to a scaled-down version of WLM, which is well-fit by a \citet{nav97} profile with virial mass M$_{\rm vir}$ = 10$^{10}$~M$_{\sun}$ and concentration parameter $c$ = 15 \citep{lea12}. To reflect the uncertainties in the parameters, we ran 2000 trials in which the important unconstrained parameters were drawn at random, and the dynamical friction and tidal destruction timescales were calculated analytically. The parameters and their range of sampled values were the initial distance and orbital eccentricity for DDO216-A1 (evenly distributed from 0--2~kpc and from 0--1, respectively), and the virial mass (log-normal distributed around 10$^{10}$~M$_{\odot}$), concentration index (normally distributed around $c$ = 12.5), and Einasto profile slope (evenly distributed between 0--1) for the PegDIG halo. In this set of trials, 27\% of the clusters are found to have survival times longer than 12~Gyr and dynamical friction timescales shorter than this. Within the range of parameters considered, there were no strong trends of survivability in the concentration index, Einasto profile slope, or virial mass, but the best cluster survivability is found for birthplaces from $\approx$300-1000~pc from the galaxy center; nearly half of 10$^5$~M$_{\odot}$ clusters born within this range sink to within $\lesssim$100~pc of the center without tidal destruction over their lifetime. Clusters born interior to this region tend to be tidally disrupted, and clusters born in the galaxy outskirts have dynamical friction timescales longer than a Hubble time. The general feature of the analysis, that in many cases clusters will be disrupted, but that the most massive will sometimes survive to be observed near the center of the host galaxy, is consistent with advanced numerical simulations of star formation in dwarfs (Charlotte Christensen, private communication). \citet*{gui16} performed hydrodynamical simulations of a larger dwarf (M$_*$ = 10$^{9.5}$~M$_{\odot}$), and observed exactly this behavior, producing a 10$^8$~M$_{\odot}$ nuclear star cluster as the result of inspiral, gas accretion, and merging of an initially 10$^4$~M$_{\odot}$ protocluster formed in the outskirts of the dwarf. Their simulated cluster arrives in the central part of the dwarf after $\approx$1~Gyr and is quenched by a final merger with another large cluster. This raises the possibility that DDO216-A1 might show an extended history of star formation as the result of dry or wet mergers, although there is little evidence for this in the current data. These results show that the cluster location is consistent with formation across a large volume of PegDIG, excluding the central region where it is now observed. Given the propensity of clusters to dissolve when located at the center of the galaxy, it seems unlikely that DDO216-A1 formed at its current location. Because dynamical friction tends to stall when the cluster reaches the radius at which the host galaxy density profile flattens into a core, it is not surprising that the cluster is not observed at the precise center of PegDIG. Both of these factors point to the likelihood that the true distance from the PegDIG center to the cluster is larger than the projected separation. \begin{figure*}[t!] \plotone{ACole_fig8.pdf} \caption{Half-light radius and absolute magnitude for a selection of stellar systems. Following \citet{bro11}, the regions of the (r$_h$,M$_V$) plane are labelled ESC-- extended star clusters, UCD-- ultracompact dwarfs and nuclei, and GC--globular clusters. Legend: {\it purple circled cross:} DDO216-A1; {\it grey circles:} clusters and dwarf galaxies \citep{bro11}; {\it black circles:} Milky Way globulars \citep[][2010 edition]{har96}; {\it red circles:} clusters in early-type dwarfs. Filled circles denote galaxy nuclei and other clusters within 150~pc of the host center \citep{cot06,sha05,geo09a}; {\it cyan circles:} clusters in late-type dwarfs, from the same sources. Filled and open symbols as above; {\it orange asterisks:} clusters in Local Group dwarfs fainter than PegDIG; {\it blue stars:} clusters in NGC~6822; {\it green triangle:} WLM cluster; {\it pink filled square:} Scl-dE1-GC1; {\it light blue open square:} Reticulum; {\it gold triangle:} M54 (Sagittarius nucleus). See text for references to individual objects. \label{fig:rhmv}} \end{figure*} \subsection{DDO216-A1 in Context} Following \citet{bro11}, we show the half-light radius and absolute magnitude of a sample of stellar systems in Figure~\ref{fig:rhmv}. The grey circles are points from \citet{bro11}, and include extragalactic globular (GC) and extended (ESC) clusters, and many ultra-compact dwarf galaxies (UCD). The Milky Way globular clusters are plotted in black. Additional clusters drawn from surveys of dwarf galaxies are plotted in red (early-type galaxies) and cyan (late-type galaxies); clusters identifed as nuclear \citep{cot06} or within projected distance 150~pc of their host centers are plotted as filled circles; the symbols for off-center clusters are open. DDO216-A1 clearly sits within the range of extended clusters in this plane, and appears characteristic of the clusters in late-type dwarf galaxies. Some specific clusters mentioned in the text are highlighted in Fig.~\ref{fig:rhmv} with alternate symbols. A census of globular clusters in Local Group dwarf galaxies has been given in \citet{for00}; the count has updated slightly since then as the result of HST and wide-field surveys of dwarfs at large distances, at low surface brightness, or where foreground confusion is high. However, most of the additions are small clusters at least an order of magnitude less luminous than DDO216-A1\footnote{At least one cluster, the one listed in Table~1 of Forbes et al.\ (2000) for DDO~210, has been removed from the list after subsequent imaging.}. Notable exceptions to this trend are the brightest 4 of the 7 new clusters discovered in NGC~6822 by \citet{hwa11} and \citet{hux13}; NGC~6822-SC1 is of similar age, half-light radius, and luminosity to DDO216-A1, although it is significantly more metal-poor \citep{vel15}. Three Local Group galaxies fainter than PegDIG host star clusters: Eri~II \citep[M$_{\mathrm V}$ = $-$7.1,][]{crn16}, And~XXV \citep[M$_{\mathrm V}$ = $-$9.7,][]{cus16}, and And~I \citep[M$_{\mathrm V}$ = $-$11.7,][]{gre00}. Each hosts a single cluster, but all are far fainter than DDO216-A1 (M$_{\mathrm V}$ $\approx$ $-$3.5 to $-$5), and are quite diffuse and extended by comparison. These galaxies are all dSphs; the cluster DDO216-A1 is as luminous as the entire galaxy Eri~II. These clusters are plotted as orange asterisks in Fig.~\ref{fig:rhmv}. The cluster-hosting dIrr galaxies NGC~6822 (M$_{\mathrm V}$ = $-$15.2) and WLM (M$_{\mathrm V}$ = $-$14.2) are both far more luminous than PegDIG. The lowest-luminosity Local Group galaxies with recorded globular clusters as bright as DDO216-A1 are the Fornax and Sagittarius dwarf spheroidals \citep{mcc12}. Both galaxies are $\approx$1~mag brighter than PegDIG and each has multiple clusters. Fornax hosts 5 globulars, and Sagittarius at least 4 \citep{dac95}, possibly as many as 9 \citep{law10}. Normalising each galaxy to an absolute magnitude of M$_{\mathrm V}$ = $-$15, the specific frequency of globular clusters is 29 for Fornax, 5--9 for Sagittarius, 7 for NGC~6822, 2 for WLM, and 10 for PegDIG. PegDIG thus has a rather ordinary specific cluster frequency compared to other Local Group galaxies and its role as cluster host is not surprising. Statistics of clusters in dwarf galaxies beyond the Local Group are necessarily less complete, but summaries of statistical properties of clusters in dwarfs can be found in, e.g., \citet{sha05,geo09a,zar16} and references therein. DDO216-A1 appears to be a typical old and metal-poor cluster of the type common to both giant and dwarf galaxies, although it is unusual to find a cluster as luminous as DDO216-A1 in a galaxy as faint as PegDIG. The contribution of the cluster to the total light of the galaxy is $\approx$0.5\%; \citet{lar15} has shown that globular clusters in dwarf galaxies can contribute up to 20--25\% of the metal-poor stars in a given dwarf. While a detailed analysis awaits the full SFH of PegDIG, it appears as though the fraction is much smaller in this case. Unlike Fornax, WLM, and NGC~6822, the cluster does not have a dramatically lower metallicity than the galaxy as a whole, only about 0.3~dex less. The structural parameters of the cluster resemble those of the ``faint fuzzy'' clusters found in lenticular galaxies \citep{bro02}, but its color is much bluer and it otherwise appears typical of the globular cluster population of dwarfs. DDO216-A1 is unusually close to the center of its host galaxy, and is also unusually extended for a cluster with small projected galacticentric distance \citep{sha05}. However, it falls below the luminosity and surface brightnesss of the great majority of clusters that are most likely to be identified as galactic nuclei \citep{geo09b,bro11}, although exceptions exist \citep{cot06,geo14}. \citet{sha05} find that among dSphs with globular clusters, more than half show clusters seen in projection against the center of the galaxy. While DDO216-A1 fits among these clusters in luminosity, its large radius distinguishes it from the (usually) compact central clusters. However, it is still quite a bit more compact than very extended clusters like Scl-dE1~GC1, with a half-light radius of 22~pc \citep{dac09}. PegDIG is the lowest-mass gas-rich galaxy in the Local Group with a major star cluster. PegDIG is a transition type galaxy, with characteristics of both irregular and spheroidal galaxies; its principal dSph-like quality is the very low rate of current star formation \citep{ski97}. In the (r$_h$, M$_{\mathrm V}$) plane, DDO216-A1 is more typical of clusters in dSph galaxies than in dIrr \citep{sha05}. If PegDIG is considered to be a dIrr, then DDO216-A1 would be one of the largest clusters known in a small dIrr. More extreme examples are very rare. For example, in the list of \citet{geo09a}, there is only one less-luminous dIrr with a comparably bright cluster, the M$_{\mathrm V}$ = $-$11.9 field galaxy D634-03, at a distance of 9.5~Mpc. By comparison, there are numerous cases of globulars in dSphs at the same host galaxy absolute magnitude. The overall trend with galaxy morphology appears to suggest that globular cluster populations are typically much poorer among the gas-rich dwarfs, which also lack nuclear clusters. DDO216-A1 is more extended than $\approx$90\% of Galactic and Magellanic Cloud globular clusters, although given its absolute magnitude and half-light radius it is not an extreme outlier. It would not be out of place among ``outer halo'' globulars such as IC~4499 or NGC~5053, Magellanic Cloud clusters like Reticulum (LMC), or clusters like NGC~6822-SC1. Contrastingly, it is much more luminous than the extended clusters Fornax 1, Arp 2, Ter 8, or Pal 12, similarly extended clusters that are either definite or probable members of spheroidal galaxies. Given the complexity of the field star background, it is impossible to say whether there are multiple ages or metallicities present in the cluster without spectroscopy. These features would be indicative of cluster mergers or in situ star formation from newly accreted gas in the cluster, either of which could contribute to the formation of a nuclear star cluster \citep[e.g.,][]{gui16}. However, circumstantial evidence argues against it being a nuclear cluster: it would be an outlier from the host mass-cluster radius and -cluster mass relations presented in \citet{geo16}, being both too extended for its mass and slightly too massive for its host compared to other nuclear star clusters. \subsection{Summary} We have imaged the central, extended star cluster in the Pegasus transition dwarf with HST/ACS, and obtained a color-magnitude diagram reaching $\approx$0.5~mag above the cluster main-sequence turnoff. DDO216-A1 is in some respects a typical globular cluster but is more extended than most. We find in particular the following major features: \begin{enumerate} \item We have confirmed that DDO216-A1 is a bona fide globular cluster, with absolute I-magnitude M$_{814}$ = $-7.90 \pm0.16$, a mass of $\sim$10$^5$~M$_{\odot}$, and a half-light radius $r_h \approx 13$~pc. While it is larger than $\approx$90\% of Milky Way globular clusters, it is structurally similar to some of the outer halo Milky Way globular clusters, and is not an extreme example of the type of extended cluster that seems to be characteristic of some dwarf galaxies. \item Based on the color-magnitude diagram, the cluster is ancient, 12.3 $\pm$0.8~Gyr old, and is moderately metal-poor, [Fe/H] = $-$1.6 $\pm$0.2. In its variable star properties, DDO216-A1 harbors $\approx$30 RR~Lyrae stars, and is an Oosterhoff Type I cluster with a specific frequency S $\approx$42. \item DDO216-A1 contributes $\approx$0.5\% of the V-band luminosity of PegDIG, but the galaxy does not have an anamolously high specific frequency of clusters compared to other dwarf galaxies. Despite the low mass of the host, the cluster is very close to the peak of the globular cluster luminosity function for all Local Group galaxies. \item The cluster is seen in projection within 30~pc of the galaxy center, but it is much more extended, for its mass, than the typical nuclear star cluster \citep[e.g.][Figure 3]{geo16}. We have not detected evidence for multiple stellar populations in the cluster. Birth at a distance of $\sim$0.3--1~kpc and subsequent infall due to dynamical friction is a possible scenario resulting in the observed cluster position. Because dynamical friction is inefficient within cored mass distributions, and the cluster has not been tidally disrupted, it is likely that the true distance between the cluster and galaxy center is a few times larger than the projected separation. \end{enumerate} The association of globular clusters with intense episodes of star formation involving very large gas masses \citep{bro06} indicates that PegDIG might have had a tumultuous early history. The observed relationship between the size of the largest star cluster and the peak SFR \citep[e.g.,][]{lar02,coo12} suggests that PegDIG should have experienced its highest SFR around the time that DDO216-A1 formed. Considering the large star formation surface densities associated with the formation of a cluster as massive as DDO216-A1 \citep[e.g.,][]{joh16} and the constraints on the total stellar mass of the galaxy, this raises the likelihood that PegDIG may have formed a large fraction of its stars in an intense burst around the time of cluster formation. The resulting peak in SFR at early times would tend to make PegDIG more similar to a prototypical dSph than to a dIrr. Although PegDIG is not currently close to any other known system, its radial velocity and distance from M31 suggest that it is not likely to be on its first infall into the Local Group. \citet{gar14} have found in their cosmological simulations of structure formation that in mock Local Groups, the majority of dwarf galaxy-sized subhaloes found within 1--1.5 virial radii of a large halo at redshift $z$ = 0 have previously spent time inside the virial radius. Timing arguments using numerical action reconstructions \citep[e.g.,][]{sha13}, while subject to significant uncertainty due to unknown initial conditions, also suggest that PegDIG has not always been isolated. Close encounters with M31 or another dwarf could have dramatically increased the SFR and therefore the statistical probability for a large cluster to be formed. \citet{dea14} predict from cosmological simulations that most dwarf-dwarf major mergers tend to occur near the time of first infall into the virial radius of a larger parent galaxy, consistent with the large age of DDO216-A1. The next paper in this series will examine the complete star formation history of PegDIG; the sample CMD for an outer field indicates that photometry well below the oldest MSTO will allow a detailed reconstruction of the SFH back to the oldest ages. | 16 | 9 | 1609.04509 |
1609 | 1609.00600_arXiv.txt | {The supernova remnant (SNR) W49B originated from a core-collapse supernova that occurred between one and four thousand years ago, and subsequently evolved into a mixed-morphology remnant, which is interacting with molecular clouds (MC). $\gamma$-ray observations of SNR/MC associations are a powerful tool to constrain the origin of Galactic cosmic-rays, as they can probe the acceleration of hadrons through their interaction with the surrounding medium and subsequent emission of non-thermal photons. The detection of a $\gamma$-ray source coincident with W49B at very high energies (VHE; $E > 100$~GeV) with the H.E.S.S. Cherenkov telescopes is reported together with a study of the source with 5 years of \fla\ high energy $\gamma$-ray (0.06 -- 300~GeV) data. The smoothly-connected combined source spectrum, measured from 60~MeV to multi-TeV energies, shows two significant spectral breaks at $304\pm20$~MeV and $8.4_{-2.5}^{+2.2}$~GeV, the latter being constrained by the joint fit from the two instruments. The detected spectral features are similar to those observed in several other SNR/MC associations and are found to be indicative of $\gamma$-ray emission produced through neutral-pion decay. } | The strong shocks associated with supernova remnants (SNRs) are very good candidates for accelerating hadronic Galactic cosmic rays to at least $10^{15}$\,eV through the diffusive shock acceleration mechanism \citep[e.g.][]{Drury1983}. The detection of very-high energy \grs\ above $20$\,TeV in objects such as \object{RX J1713.7-3946} \citep{Aharonian2006b} or \object{RCW 86} \citep{Aharonian2009} shows that these objects can accelerate particles up to at least $100$\,TeV. However, whether the bulk of accelerated particles that radiate in the high energy (HE; $0.1 < E < 100$\,GeV) and very-high energy (VHE; $E > 100$\,GeV) \gr\ range is of leptonic or hadronic nature is still under debate \citep{Blasi2013}. A growing number of SNRs interacting with molecular clouds (SNR/MC) are being revealed in the GeV and TeV \gr\ domain. This includes \object{W44} \citep{Abdo2010e}, \object{W28} \citep{Aharonian2008,Abdo2010c}, \object{CTB 37A} \citep{Aharonian2008b,Castro2010} and \object{IC 443} \citep{Albert2007,Acciari2009,Abdo2010d}. Even though isolated SNRs are obviously cleaner laboratories to study cosmic-ray acceleration processes, SNR/MC associations are good candidates to test the presence of hadronic acceleration in SNRs, partly because the HE and VHE \gr\ emissions from the decay of $\pi^0$ mesons are expected to be strongly enhanced. The neutral pions, produced when high energy protons (or nuclei) collide with interstellar material, each decay into two $\gamma$ rays with equal energies in the pion's rest-frame. This translates into a steep rise below $\sim$200\,MeV in the $\gamma$-ray spectral energy distribution (often referred to as the ``pion-decay bump''). This characteristic spectral feature has been recently detected at high energies for three interacting SNRs: IC~443, W44 \citep{Ackermann2013} and \object{W51C} \citep{Jogler2016}. However, whether this feature is the signature of the acceleration of freshly injected protons may be questioned as re-acceleration of diffuse cosmic rays for a limited time period is also possible \citep{Uchiyama2010, Cardillo2016}. In this context, the \object{W49} region, discovered in the 22\,cm survey of \citet{Westerhout1958}, represents one of the most interesting regions in the Galaxy to study cosmic-ray acceleration. This region contains two remarkable objects: a young SNR (\object{W49B}) and a star-forming region (\object{W49A}). The SNR W49B (G43.3--0.2) is another example of a SNR/MC association. It is a mixed-morphology \citep{Rho1998} SNR with a bright shell of diameter $\sim4'$ resolved at radio wavelengths and centrally filled with thermal X rays \citep{Hwang2000}. With a flux density of $38$\,Jy at $1$\,GHz, this source is one of the brightest SNRs of the Galaxy at radio wavelengths. Extensive infrared and X-ray studies revealed that W49B's progenitor was a supermassive star that created a wind-blown bubble in a dense molecular cloud in which the explosion occurred \citep{Keohane2007}. It has recently been shown that W49B's progenitor experienced a jet-driven core collapse explosion \citep{Lopez2013,Gonzalez-Casanova2014}. Observations of mid-infrared lines from shocked molecular hydrogen show that W49B is interacting with molecular clouds. Near-infrared \ion{Fe}{II} observations revealed filamentary structures which is evidence of radiative shocks \citep{Keohane2007}. This interaction is also suggested by \ion{H}{I} line observations \citep{Brogan2001}. The age of the remnant is estimated to be between $1000$ and $4000$ years \citep{Moffett1994,Hwang2000,Zhou2011} and the distance of this object is still not very well constrained. From \ion{H}{I} absorption analyses, \citet{Radhakrishnan1972} derived a distance, which \citet{Moffett1994} later re-calculated to be $\sim 8~\rm kpc$ (using an updated Galactic rotation model). \citet{Brogan2001}, using more recent VLA data, have shown that an association of W49B with the nearby star forming region W49A is also possible, extending the range of possible distances ($8\leq D \leq 12$\,kpc) for this source. More recently, \citet{Zhu2014} obtained a distance of $\sim $10\,kpc. As mentioned above, the other notable component of the W49 region is W49A \citep{Mezger1967}. It is one of the brightest giant radio \ion{H}{II} regions in the Galaxy. This star forming region is located in the densest $\sim$15\,pc of a $10^6\,\rm{M_\odot}$ giant molecular cloud of $\sim$100\,pc in size. It contains numerous compact and ultra-compact (UC) \ion{H}{II} regions and its emission is equivalent to the presence of about $100$ O7 stars \citep{Brogan2001}. The presence of a very massive star with an initial mass of $100-180\,\mathrm{M_{\odot}}$ has been reported by \citet{Wu2014}. W49A is associated with a molecular outflow and strong $\rm H_2O$ masers \citep{Walker1982,Scoville1986}, the proper motion of which was used by \citet{Zhang2013} to determine a distance of $11.11^{+0.79}_{-0.69}$\,kpc. Star forming regions are considered as potential VHE $\gamma$-ray emitters since they generally host massive stars which could accelerate particles to VHEs through interactive or collective wind effects. As an example, the TeV \gr\ source coincident with Westerlund\,1 \citep{Abramowski2012}, detected with the High Energy Stereoscopic System (H.E.S.S.), may be the site of such processes. In the HE domain, a bright source coincident with W49B is detected with the \emph{Fermi} Large Area Telescope (LAT) \citep[\object{3FGL J1910.9+0906} in the \fla\ 3FGL catalog,][]{3FGL}). It was one of the 205 most significant sources after the first three months of observation\footnote{It was designated as \object{0FGL J1911.0+0905} in this so-called {\it Bright Sources List} \citep{Abdo2009}.} and is one of the 360 sources detected above 50~GeV in 80 months of data \citep{2FHLPaper}. This source is the subject of a detailed analysis presented in \citet{Abdo2010b}. In that paper, the detection of HE \gr\ emission was reported towards W49B at a significance of $38\sigma$ with 17 months of data. The authors disfavored a possible pulsar origin for the emission observed towards W49B in the GeV range. Furthermore, from X-ray measurements, the presence of a neutron star as the result of the progenitor star's collapse appears unlikely \citep{Lopez2013}. In the present paper, the detection of a source coincident with the SNR W49B in the VHE \gr\ domain with H.E.S.S. is reported. The analysis of the \fla\ data is applied in this work to 5 years of data using an updated calibration and updated source and background models. The spectral and morphological results in the GeV and TeV regime are discussed in the context of a SNR interacting with a molecular cloud. | A source of very-high energy $\gamma$ rays is detected towards the supernova remnant W49B with H.E.S.S. and a joint study is conducted together with the analysis of 5 years of Fermi-LAT data towards the source. The point spread functions of these instruments are comparable to the physical size of the SNR and do not allow for detailed morphology studies. However, the significantly increased data set compared to previous publications allowed for the derivation of the $\gamma$-ray spectrum of the source between 60 MeV and a few TeV. Two spectral breaks could be identified in the $\gamma$-ray spectrum: one at 304 MeV and the other at 8.4 GeV. The latter, constrained from a joint fit of the \fla\ and H.E.S.S. data, is similar to the spectral breaks observed in other supernova remnants interacting with molecular clouds and further supports the evidence of interaction observed in other wavelengths for this object. The broad band spectrum of W49B could be explained either by leptonic or hadronic models. However, the sharp break detected with the \fla\ data at 304 MeV seems to favor a hadronic origin of the $\gamma$-ray emission since the leptonic models would require an even harder spectrum for the electron population than the ones tested in this work which would fail to explain the radio observations. In the near future, W49B as well as the nearby star forming region W49A will be of particular interest to study at VHE with improved analysis methods and with the next generation of instruments such as the Cherenkov Telescope Array. Observations or analyses of the W49 region with improved angular resolution and higher sensitivity would help to constrain the morphology and the origin of the emission towards W49B and possibly confirm the hint of emission towards W49A. Furthermore, if one assumes that the distances to W49A and W49B are comparable, then the observed difference between the $\gamma$-ray luminosities of the two sources would become especially interesting. It would imply that in the absence of recognizable supernova remnants -- as in W49A -- the other possible energetic particle sources like the shocks expected from interacting or collective stellar winds appear not very effective for HE and VHE $\gamma$-ray emission in this case. Therefore, a truly reliable distance determination for these sources (see Sect.~\ref{intro}) would be of great astrophysical importance. | 16 | 9 | 1609.00600 |
1609 | 1609.07485_arXiv.txt | The neutron star low-mass X-ray binary and intermittent millisecond X-ray pulsar \source\ returned to quiescence in late 2015, after a prolonged accretion outburst of $\simeq$10~yr. Using a \chan\ observation taken $\simeq$180~d into quiescence we detect the source at a luminosity of $\simeq$$4.5\times10^{31}~\dist~\lum$ (0.5--10 keV). The X-ray spectrum can be described by a neutron star atmosphere model with a temperature of $\simeq$54~eV for an observer at infinity. We perform thermal evolution calculations based on the 2016 quiescent data and a $\lesssim$98~eV temperature upper limit inferred from a \swift\ observation taken during an unusually brief ($\lesssim$2~weeks) quiescent episode in 2007. We find no evidence in the present data that the thermal properties of the crust, such as the heating rate and thermal conductivity, are different than those of non-pulsating neutron stars. Finding this neutron star so cold after its long outburst imposes interesting constraints on the heat capacity of the stellar core; these become even stronger if further cooling were to occur. | Neutron stars are one of the possible remnants of once massive stars that ended their life in a supernova explosion. A defining property of neutron stars is that they are very compact; despite having a mass of $\simeq$1.4$~\Msun$, their radius is only $\simeq$10~km. As a result, their interior density rises beyond the density of atomic nuclei. Neutron stars are therefore of prime interest to understand the properties of ultra-dense matter \citep[e.g.][for a review]{lattimer2011}. Our Galaxy harbors $>$100 neutron stars that are part of low-mass X-ray binaries (LMXBs) and accrete gas from a $\lesssim$1~$\Msun$ companion star via an accretion disk. Many of these systems are transient; during X-ray luminous phases matter is rapidly falling toward the neutron star, but during intervening quiescent phases the accretion rate, and hence the X-ray luminosity, is strongly reduced. Two sub-classes of neutron star LMXBs are the quasi-persistent sources and the accreting millisecond X-ray pulsars (AMXPs). Each make up $\simeq$10 per cent of the current population of neutron star LMXBs, and each have distinct outburst and quiescent properties. Quasi-persistent LMXBs stand out by showing prolonged accretion outbursts of years to decades rather than weeks to months. Moreover, several of these sources show strong thermal emission in quiescence that gradually decreases on a timescale of years \citep[e.g.][]{cackett2013_1659,fridriksson2011,degenaar2014_exo3,homan2014,merritt2016}. The crust of a neutron star is heated during accretion outbursts via pycnonuclear fusion reactions that take place at $\simeq$1~km depth, and electron captures occurring at shallower depth. Together, these processes deposit an energy $\simeq1-2$~MeV per accreted nucleon, heating the crust \citep[e.g.][]{haensel1990a,haensel2008,brown1998}. The temperature evolution in quiescence of five quasi-persistent sources can be successfully explained as cooling of the strongly-heated neutron star crust and offers valuable insight into the structure and composition of these neutron stars \citep[e.g.][]{rutledge2002,wijnands2002,wijnands2004,shternin07,brown08,page2013,medin2014,deibel2015,horowitz2015,turlione2013,cumming2016}.\footnote{Recent studies have also revealed crust cooling in three neutron stars with short outbursts \citep[e.g.][]{degenaar2011_terzan5_3,degenaar2015_ter5x3,waterhouse2016}.} AMXPs distinguish themselves by displaying coherent X-ray pulsations when accreting \citep[e.g.][]{wijnands1998}. It is believed that in these objects the stellar magnetic field is strong enough to disrupt the accretion flow and channel plasma to the magnetic poles of the rapidly rotating neutron star. In quiescence, the persistently-pulsating AMXPs show weak or no thermal X-rays but strong power-law emission \citep[e.g.][]{campana2005_amxps,campana2008,jonker2005,wijnands05_amxps,heinke2009,degenaar2012_amxp}. Such a hard emission component is also seen in the quiescent spectra of some non-pulsating neutron star LMXBs, though it is typically much less prominent. The hard quiescent X-rays are often ascribed to residual accretion or non-thermal emission processes related to the neutron star magnetic field \citep[e.g.][]{campana1998,rutledge2001,degenaar2012_amxp,chakrabarty2014_cenx4,wijnands2014}. \source\ shares properties of both the aforementioned sub-classes of neutron star LMXBs; it is the only known AMXP that accreted for a full decade, and the only quasi-persistent source that acted as an AMXP. The source was first seen in outburst in 2005 June when it exhibited a thermonuclear X-ray burst \citep[][]{vanderspek2005}. X-ray pulsations at a frequency of 377.3~Hz were soon found, but were detected only sporadically after $\simeq$2 months \citep[][]{galloway2008_hete,patruno2012_hete}. The source is therefore referred to as an ``intermittent AMXP''. X-ray burst analysis suggests a source distance of $D=4.7$~kpc \citep[][]{galloway06}. The $\simeq$10-year long outburst of \source\ ended in late 2015. We report on a \chan\ ToO observation obtained $\simeq$180~d after it went quiescent. We combine the obtained temperature measurement with an upper limit from \swift\ in 2007 when the source disappeared for $\lesssim$2 weeks, to perform thermal evolution simulations and to probe the thermal properties of this neutron star. | \label{sec:discuss} We report on \chan\ and \swift\ observations obtained within $\simeq$180~d after the $\simeq$10-yr long outburst of the neutron star LMXB and intermittent AMXP \source. Analysis of the \chan\ data reveals that the quiescent spectrum is dominated by soft, thermal emission that can be described by a neutron star atmosphere model with a temperature of $kT^{\mathrm{\infty}}_{\mathrm{eff}} \simeq 54$~eV. Any hard emission tail contributes $\lesssim$9 per cent to the total 0.5--10 keV luminosity of $L_{\mathrm{X}} \simeq 5 \times 10^{31}~\dist~\lum$, for an assumed power-law spectral shape with a photon index of $\Gamma = 1-2$. \source\ is different from the AMXPs by exhibiting a thermally-dominated quiescent X-ray spectrum. Other AMXPs (both persistently and intermittently pulsating) display strong power-law emission, contributing $\gtrsim$50\% to the total unabsorbed quiescent 0.5--10 keV flux \citep[e.g.][]{campana2005_amxps,wijnands05_amxps,heinke2009,degenaar2012_amxp}. With regard to the quasi-persistent neutron star LMXBs, \source\ is much colder after its long outburst than other sources observed at a similar epoch \citep[$kT^{\mathrm{\infty}}_{\mathrm{eff}} \gtrsim 90$~eV; e.g.][]{wijnands2001,wijnands2002,wijnands2003,wijnands2004,degenaar2010_exo2,fridriksson2010,homan2014}. Moreover, our thermal evolution calculations indicate that the crust of the neutron star in \source\ fully cools in $\simeq$0.5--2~yr, which is short compared to the years-long cooling seen for other quasi-persistent sources \citep[e.g.][]{cackett2013_1659,fridriksson2011,degenaar2014_exo3,homan2014,merritt2016}. Our thermal evolution simulations suggest that both the low post-outburst temperature and the relatively fast cooling time scale of \source\ may be due to its relatively low mass-accretion rate during outburst. Indeed, our inferred mass-accretion rate is a factor $\simeq$2 lower than that of the crust-cooling source \exo, and a factor $>$5 lower than that of others \citep[see table~1 in][and references therein]{degenaar2015_ter5x3}. For a lower mass-accretion rate, the crust is heated less strongly and its temperature is lower, resulting in a shorter cooling time-scale \citep[see e.g. figure 1 in][]{page2013}. Nevertheless, it is striking that with only a factor $\simeq$2 difference in mass-accretion rate, \exo\ was detected at $kT^{\mathrm{\infty}}_{\mathrm{eff}} \simeq 115$~eV at $\simeq$180 days after its $\simeq$24-yr long outburst \citep[][]{degenaar2010_exo2}. This likely points to other differences between these two sources, e.g. contrasting core temperatures or different crust properties. Notably, \source\ is the only quasi-persistent LMXB that acted as an AXMP, and the disappearance of its pulsations has been explained as its magnetic field being ``buried'' by accretion \citep[e.g.][]{cumming2008,patruno2012_hete}. It is currently unclear if/how this should affect the thermal properties. If the magnetic field is strongly folded into the crust so that it reaches $\simeq 10^{11}$ G, it may create an insulating layer that prevents the accretion-induced heat to propagate to the surface until the magnetic field re-emerges. If so, the neutron star temperature (and perhaps any hard emission) could possibly increase after the Ohmic diffusion timescale. This could be as short as tens of days for \source\ \citep[e.g.][]{cumming2008}. Nevertheless, our work suggests that the quiescent data obtained so far can be explained with similar physical parameters as inferred for other crust-cooling sources. At present there are thus no indications that its magnetic field has a notable impact its thermal properties. Our simulations show that if the crust of the neutron star is highly conductive with $Q_{\mathrm{imp}} = 1$, no further cooling is expected in \source. Our \chan\ measurement then reflects the temperature of the neutron star core, $\tilde{T}_{0} \simeq 2 \times10^{7}$~K. However, the present data allow for a higher impurity parameter, up to $Q_{\mathrm{imp}} \simeq 8$, for which the core temperature may be as low as $\tilde{T}_{0} \simeq 4\times10^{6}$~K and further cooling of the crust is expected. Future observations of \source\ in quiescence can thus further constrain the impurity impurity of the crust and temperature of the core, as well as possible effects of the magnetic field on the thermal properties. The total energy released during the $\simeq$10 yr accretion outburst, from both the standard nuclear processes and additional shallow heating, amount to $\simeq 2\times 10^{43}$~erg. As the present work was in progress, \citet{cumming2016} found that such an amount of energy can raise the core temperature and provides us with a lower limit on its heat capacity that depends on the tantalizing core composition. Comparing with their figure~7, our $\tilde{T}_0 \simeq 2 \times 10^7$ K provides the strongest constraint to date on the core heat capacity: it must be $C>10^{37} (\tilde{T}_{0}/10^8 \, \mathrm{K})$ erg K$^{-1}$, a value that is close to the minimum provided by the leptons in the core. This minimum can only be reached if all baryons are strongly paired, i.e. superfluid or superconducting, hence have a negligible contribution to the total heat capacity. Observing further cooling in \source\ would push this lower limit further down and may limit the maximum fraction of baryons that are paired, and consequently the minimum fraction of baryons that are not paired. \vspace{-0.2cm} | 16 | 9 | 1609.07485 |
1609 | 1609.09498_arXiv.txt | Mergers and the spin of the dark matter halo are factors traditionally believed to determine the morphology of galaxies within a $\Lambda$CDM cosmology. We study this hypothesis by considering approximately 18,000 central galaxies at $z=0$ with stellar masses $M_{\ast} = 10^{9}$--$10^{12} \, {\rm M}_{\odot}$ selected from the Illustris cosmological hydrodynamic simulation. The fraction of accreted stars -- which measures the importance of massive, recent and dry mergers -- increases steeply with galaxy stellar mass, from less than 5 per cent in dwarfs to 80 per cent in the most massive objects, and the impact of mergers on galaxy morphology increases accordingly. For galaxies with $M_{\ast} \gtrsim 10^{11} \, {\rm M}_{\odot}$, mergers have the expected effect: if gas-poor they promote the formation of spheroidal galaxies, whereas gas-rich mergers favour the formation and survivability of massive discs. This trend, however, breaks at lower masses. For objects with $M_{\ast} \lesssim 10^{11} \, {\rm M}_{\odot}$, mergers do not seem to play any significant role in determining the morphology, with accreted stellar fractions and mean merger gas fractions that are indistinguishable between spheroidal and disc-dominated galaxies. On the other hand, halo spin correlates with morphology primarily in the {\it least} massive objects in the sample ($M_{\ast} \lesssim 10^{10} \, {\rm M}_{\odot}$), but only weakly for galaxies above that mass. Our results support a scenario where (1) mergers play a dominant role in shaping the morphology of massive galaxies, (2) halo spin is important for the morphology of dwarfs, and (3) the morphology of medium-sized galaxies -- including the Milky Way -- shows little dependence on galaxy assembly history or halo spin, at least when these two factors are considered individually. | \label{sec:intro} \renewcommand{\thefootnote}{\fnsymbol{footnote}} \footnotetext[1]{E-mail: [email protected]} \renewcommand{\thefootnote}{\arabic{footnote}} Since \cite{Hubble1926} proposed his galaxy classification scheme, numerous studies have investigated the physical mechanisms that lead to the formation of spiral and elliptical galaxies. Given the nonlinearity of the physical processes involved, many such studies have used numerical simulation as their main tool, with \cite{Toomre1972} and \cite{White1978} being some of the first examples. To this day, one of the main questions in the field of galaxy formation remains to understand which properties of a halo (or its environment) determine the morphology of the galaxy formed at its centre \citep[e.g.][]{Parry2009, Stinson2010, Sales2012, Teklu2015, Zavala2016}. Galactic discs are believed to form through the dissipational collapse of gas in dark matter (DM) haloes, which acquire their angular momentum through tidal torques in the early Universe \citep{Peebles1969, Doroshkevich1970}. Despite the fact that galactic outflows, misaligned gas accretion, and mergers tend to complicate the detailed conservation of specific angular momentum during the formation of a galaxy \citep[e.g.][]{VandenBosch2002, Brook2011, Zjupa2017}, it is nevertheless expected that the largest galactic discs form preferentially in haloes with higher angular momentum \citep[e.g.][]{Fall1980, Fall1983, Mo1998}. Initial attempts to verify this prediction in cosmological simulations suffered from numerical issues leading to a `catastrophic' loss of angular momentum \citep{Navarro1995}, but more recent studies have been able to verify such a trend \citep[e.g.][]{Teklu2015}. On the other hand, elliptical galaxies are believed to be the remnants of repeated galaxy mergers \citep{Toomre1977}. This picture gained support as the hierarchical nature of structure formation \citep{White1978a} started to become recognized, although a number of objections \citep[e.g.][]{Ostriker1980} indicated that the situation was more complicated than this. Subsequent studies using idealized merger simulations have shown that the outcome of a galaxy merger is significantly affected by the masses, gas fractions, and morphologies of the progenitors, as well as by their orbital parameters \citep[e.g.][]{Barnes1992, Hernquist1992, Hernquist1993, Barnes1996, Naab2003, Cox2006}. Because of this, it is necessary to place such merger simulations in a `cosmological context' in order to evaluate the impact of galaxy mergers from a statistical perspective. A relatively inexpensive way of doing this consists in deriving simple prescriptions from such idealized merger simulations and then combining them with semi-analytic models (SAMs) of galaxy formation \citep[e.g.][]{Khochfar2006, Parry2009, DeLucia2011, Zavala2012, Avila-Reese2014, Fontanot2015, Kannan2015} or semi-empirical halo occupation models \citep[e.g.][]{Hopkins2009a, Stewart2009a, Hopkins2010, Moster2014}. A more straightforward approach, although far more computationally expensive, is to run a hydrodynamic cosmological simulation on a large comoving volume and then examine the morphologies of the resulting galaxies directly (see \citealt{Naab2014}, section 2; \citealt{Somerville2015}, section 4.2, for reviews). With this technique, it has been found that both the numerical treatment of hydrodynamics and the feedback implementation can have a dramatic effect on galaxy morphology \citep[e.g.][]{Brook2004, Okamoto2005, Governato2007, Scannapieco2008, Scannapieco2012, Ceverino2009, Sales2010, Agertz2011, Torrey2012, Ubler2014, Snyder2015a, Genel2015, Agertz2016, Dubois2016}. Despite all the inevitable uncertainties in the modelling of feedback, the latest generation of cosmological hydrodynamic simulations \citep[e.g.][]{Vogelsberger2014, Dubois2014, Schaye2015} has been able to produce galaxy populations displaying a `morphological mix' that agrees reasonably well with observations. The latter approach, i.e. using hydrodynamic cosmological simulations to directly investigate the origin of galaxy morphology, has presented some serious puzzles for galaxy formation. For example, by studying a sample of Milky Way-sized simulated galaxies, \cite{Sales2012} found no correlation between galaxy morphology and properties such as halo spin and merging history, challenging the `standard' model of the formation of discs and spheroids. Instead, \cite{Sales2012} proposed that a disc forms when the angular momentum of freshly accreted gas is aligned with that of earlier gas accretion. Furthermore, \cite{Zavala2016} found that galaxy morphology is correlated with the assembly history of the {\it inner} DM halo. In particular, they found that the `angular momentum loss' of a galaxy's stellar component is strongly linked to that of the inner DM halo, whereas the angular momentum history of a galaxy's cold gas is statistically correlated with that of the whole DM halo. In this work we explore the origin of discs and spheroids using the Illustris simulation \citep{Vogelsberger2014, Vogelsberger2014a, Genel2014a, Sijacki2015}, a hydrodynamic cosmological simulation that has been shown to reproduce several galaxy observables reasonably well, including stellar angular momenta \citep{Genel2015} and quantitative optical morphologies \citep{Torrey2015, Snyder2015}, over a wide range of stellar masses. This makes the Illustris simulation a powerful tool to study the physical mechanisms that shape galaxy morphology in a cosmological context. The results presented in this paper complement those of \cite{Genel2015}, who showed that feedback -- both from galactic winds and from active galactic nuclei -- is crucial in order to produce simulated galaxies with stellar angular momenta in reasonable agreement with observations. Here, on the other hand, we adopt the fiducial model from the Illustris simulation and focus on how the morphology of a central galaxy at $z=0$ depends on its merging history and on the spin of its host halo. The current paper is organized as follows. In Section \ref{sec:methodology} we briefly describe the set of cosmological simulations known as the Illustris Project, along with their merger trees and other post-processed catalogues, and present the galaxy sample considered for this study. An overview of galaxy morphology in the Illustris simulation, covering a wide range of stellar masses, is presented in Section \ref{sec:measuring_morphology}. In Section \ref{sec:what_drives_galaxy_morphology} we investigate the dependence of galaxy morphology on merging history (Section \ref{subsec:merging_history}), halo spin (Section \ref{subsec:halo_spin}), and a combination of both (Section \ref{subsec:mergers_and_halo_spin}). Finally, we summarize and discuss our results in Section \ref{sec:discussion_and_conclusions}. | \label{sec:discussion_and_conclusions} We have investigated the connection between galaxy morphology, halo spin, and merging history in the Illustris cosmological simulation, considering approximately 18,000 central galaxies at $z=0$ over a wide range of stellar masses ($M_{\ast} = 10^{9}$--$10^{12} \Msun$). We showed that disc-like and spheroidal galaxies arise naturally in the Illustris simulation over the mass range considered (Figs. \ref{fig:disks} and \ref{fig:spheroids}), and then proceeded to investigate the effects of mergers and halo spin on galaxy morphology as a function of stellar mass. Throughout this work, we quantified galaxy morphology using the $\kappa_{\rm rot}$ parameter \citep{Sales2012}, defined in equation (\ref{eq:kappa}). This quantity is closely related to other kinematic morphological parameters, such as the kinematic disc-to-total fraction (sometimes denoted by D/T) that is often used in other simulation studies (see Section \ref{subsec:morphology_definitions}). Based on the arguments presented in Appendix \ref{app:kappa_expanded}, we employed the $\kappa_{\rm rot}$ parameter to classify galaxies as rotation-dominated ($\kappa_{\rm rot} \geq 0.5$) or dispersion-dominated ($\kappa_{\rm rot} < 0.5$), as illustrated in Fig. \ref{fig:morphology_vs_mass}. In order to test the idea that mergers play an important role in determining galaxy morphology \citep[e.g.][]{Toomre1977, White1978, Barnes1996, Naab2006, Naab2014}, we used the \textit{ex situ} stellar mass fraction $f_{\rm acc}$ to measure the overall impact of galaxy mergers (relative to \textit{in situ} star formation). We found that mergers are indeed an important transformational mechanism in massive galaxies ($M_{\ast} \gtrsim 10^{11} \Msun$), but their importance diminishes at lower masses (Figs. \ref{fig:f_acc_vs_mstar} and \ref{fig:kappa_vs_f_acc}). To the best of our knowledge, this is the first time that a statistical demonstration of the `merger hypothesis' has been carried out with a cosmological hydrodynamic simulation over such a wide range of stellar masses (however, see \citealt{DeLucia2011} for a study using SAMs). Beyond the \textit{ex situ} stellar mass fraction, $f_{\rm acc}$, we also considered three additional merger statistics: the mean merger gas fraction, the mean merger lookback time, and the mean merger mass ratio. By comparing these quantities to $f_{\rm acc}$ we found that -- at any fixed stellar mass -- a higher $f_{\rm acc}$ is associated with a larger number of \textit{massive}, \textit{recent}, and \textit{dry} mergers (Fig. \ref{fig:merging_history}), which happen to be the kind of mergers that are believed to contribute most to the formation of spheroidal galaxies. In fact, in Fig. \ref{fig:kappa_vs_merger_gas_fraction} we showed explicitly that gas-poor mergers contribute more to spheroid formation than gas-rich mergers, although the trend is strong only for massive galaxies ($M_{\ast} \gtrsim 10^{11} \Msun$). Having investigated the role of merging history in shaping galaxy morphology, we proceeded to test the hypothesis that, through conservation of specific angular momentum, the spin of a halo also plays a major role in determining the morphology of the galaxy formed at its centre. According to this picture, DM haloes with high angular momentum content favour the formation of rotationally supported discs. In order to rule out baryonic effects on halo spin, we matched each Illustris halo to its counterpart in Illustris-Dark and used the spin parameter of the latter for all our analyses. We showed in Figs. \ref{fig:spin_parameter_vs_mstar} and \ref{fig:kappa_vs_spin_parameter} that $\kappa_{\rm rot}$ is somewhat determined by halo spin at $M_{\ast} \lesssim 10^{10} \Msun$, but the correlation between these two quantities becomes weaker at higher masses and eventually disappears, presumably due to the increasing impact of galaxy mergers. Finally, we investigated the joint effect of halo spin and merging history on galaxy morphology. We found that, in general, galaxies with high (low) $f_{\rm acc}$ which are also located at the centres of slowly (fast) rotating haloes are more likely to be dispersion-dominated (rotation-dominated) systems, as manifested by the `diagonal' $\kappa_{\rm rot}$ gradients in Fig. \ref{fig:morphology_diagram_manymasses}. However, the relative importance between halo spin and $f_{\rm acc}$ is mass-dependent, as evidenced by the varying \textit{direction} of the $\kappa_{\rm rot}$ gradients: halo spin is the dominant driver of galaxy morphology at low masses, even for galaxies which have undergone major mergers, while dry mergers are more important in determining the morphology of more massive systems. The morphology of Milky Way-sized galaxies, at the transition between these two regimes ($M_{\ast} \approx 10^{10}$--$10^{11} \, \Msun$), depends on a combination of $f_{\rm acc}$ and halo spin, as described above, but shows a weak dependence on any of these two factors alone. We argued that the physical origin of this mass-dependent response of galaxy morphology to mergers is largely due to the nature of the mergers undergone by galaxies in different mass ranges. Indeed, as discussed in Section \ref{subsec:merging_history}, two galaxies of different masses with the same $f_{\rm acc}$ value need not have `equivalent' merging histories. For example, Fig. \ref{fig:merging_history} shows that low-mass galaxies are more likely to have gas-rich mergers, while the merging history of a more massive system usually consists of numerous dry mergers. This has crucial consequences for galaxy morphology. In fact, it has been proposed that dry mergers are important in the formation of spheroidal systems \citep[e.g.][]{Khochfar2003, Naab2006a}, while gas-rich major mergers have been shown to produce disc galaxies \citep[e.g.][]{Springel2005, Robertson2006, Governato2007}. Both of these ideas seem to be manifested statistically in the different panels of Fig. \ref{fig:morphology_diagram_manymasses}, where the `likelihood' of galaxy mergers producing a spheroid increases with stellar mass, along with the `dryness' of the galaxy mergers in each mass range (Fig. \ref{fig:merging_history}). Furthermore, we know that low-mass galaxies undergo major mergers somewhat less frequently than massive galaxies, by about an order of magnitude \citep{Rodriguez-Gomez2015}. This reduced merger frequency, along with the higher gas fractions (and specific star formation rates) present at low stellar masses, further adds to the apparent `resilience' of small galaxies to the destructive effects of major mergers. In other words, we interpret the shifting importance of mergers in determining galaxy morphology as being largely due to the varying nature of galaxy mergers in different mass ranges: mergers between low-mass galaxies tend to be gas-rich and less frequent, while the merging history of a massive galaxy is dominated by repeated gas-poor mergers. In addition, low-mass galaxies have larger gas reservoirs which allow them to regrow their stellar discs. Thus, our work supports a scenario in which galaxy mergers manifest themselves as different kinds of transformational mechanisms, potentially giving rise to both discs and spheroids, depending on the mass range considered. | 16 | 9 | 1609.09498 |
1609 | 1609.07943_arXiv.txt | The TBT project is being developed under ESA's General Studies and Technology Programme (GSTP), and shall implement a test-bed for the validation of an autonomous optical observing system in a realistic scenario within the Space Situational Awareness (SSA) programme of the European Space Agency (ESA). The goal of the project is to provide two fully robotic telescopes, which will serve as prototypes for development of a future network. The system consists of two telescopes, one in Spain and the second one in the Southern Hemisphere. The telescope is a fast astrograph with a large Field of View (FoV) of 2.5 x 2.5 square-degrees and a plate scale of 2.2 arcsec/pixel. The tube is mounted on a fast direct-drive mount moving with speed up to 20 degrees per second. The focal plane hosts a 2-port 4K x 4K back-illuminated CCD with readout speeds up to 1MHz per port. All these characteristics ensure good survey performance for transients and fast moving objects. Detection software and hardware are optimised for the detection of NEOs and objects in high Earth orbits (objects moving from 0.1-40 arcsec/second). Nominal exposures are in the range from 2 to 30 seconds, depending on the observational strategy. Part of the validation scenario involves the scheduling concept integrated in the robotic operations for both sensors. Every night it takes all the input needed and prepares a schedule following predefined rules allocating tasks for the telescopes. Telescopes are managed by RTS2 control software, that performs the real-time scheduling of the observation and manages all the devices at the observatory \cite{kubanek2010rts2}. At the end of the night the observing systems report astrometric positions and photometry of the objects detected. The first telescope was installed in Cebreros Satellite Tracking Station in mid-2015. It is currently in the commissioning phase and we present here the first results of the telescope. We evaluate the site characteristics and the performance of the TBT Cebreros telescope in the different modes and strategies. Average residuals for asteroids are under 0.5 arcsecond, while they are around 1 arcsecond for upper-MEO\footnote{Medium Earth Orbit} and GEO \footnote{geosynchronous satellite } satellites. The survey depth is dimmer than magnitude 18.5 for 30-second exposures with the usual seeing around 4 arcseconds. | \label{sec:intro} % Space Situational Awareness (SSA) program of European Space Agency foresees the deployment of several robotic telescopes to provide surveillance and tracking services for man-made as well as natural Near-Earth Objects (NEOs). These ground-based optical sensors are very efficient systems to detect and track faint objects in high altitude orbital regions. The TBT project is being developed under ESA's General Studies and Technology Programme (GSTP), and shall implement a test-bed for the validation of an autonomous optical observing system in a realistic scenario, consisting of two telescopes located in Spain and a location in the Southern Hemisphere. The TBT consortium consists of companies and institutions from the countries funding the project. The prime contractor is the Spanish company ISDEFE, while other members of the consortium are Ixion Aerospace and Industry and the Fabra-ROA Telescope at Montsec \footnote{TFRM http://www.am.ub.edu/bnc/} are members of the consortium. The other partner is IGUASSU \footnote{Iguassu Software Systems (ISS), Prague, Czech Republic} from Czech Republic. The project itself was driven by geographical return and based on COTS components, both for hardware and software. According to the final selection in the Critical Design Review, the operation modes and strategies design were driven by the elements finally purchased. \begin{figure} [ht] \begin{center} \begin{tabular}{c} % \includegraphics[height=5cm]{view.jpg} \end{tabular} \end{center} \caption[example] { \label{fig:view} View of the Cebreros TBT Observatory at sunset. All the telescope equipment, control and processing computers are hosted inside of the 4.2~m diameter clamshell dome.} \end{figure} In Winter 2016 the Cebreros telescope has entered in the commissioning phase, led mainly by ISDEFE as the integrator of all the subsystems. After this phase, the second telescope will be installed and both observatories will be tested as a complete autonomous and robotic system. For this Site Acceptance Test, the telescopes will be monitored from ESOC using the Human-Machine Interface HMI, implemented by Ixion, and programmed by the scheduler TRANSITO, developed by ISDEFE. The images are processed and the results sent automatically by TOTAS\cite{koschny2015teide}, software developed by M. Busch and tailored by IGUASSU for the TBT Project. \subsection{Telescope design} The telescope is a fast (f/2.5) prime-focus astrograph with a 3-lens Wynne corrector designed by Dr.~V.~Yu~Terebizh (Crimean Astrophyical Observatory). It has a 56-cm diameter mirror and a resulting unvignetted angular FoV of 3.5 degrees in diameter. \begin{figure} [ht] \begin{center} \begin{tabular}{c} % \includegraphics[height=6cm]{layout.jpg} \end{tabular} \end{center} \caption[Optical layout] { \label{fig:layout} Detail optical layout of the VT56s design from V.Yu. Terebizh for the TBT telescopes. From the first lens of the Wynne corrector (right) to the focal plane (left). It includes the shutter, filter and the camera window.} \end{figure} \begin{figure} [ht] \begin{center} \begin{tabular}{c} % \includegraphics[height=6cm]{spot.jpg} \end{tabular} \end{center} \caption[Spot diagram] { \label{fig:spot} Spot diagram in the integral light between 380~nm and 900~nm. Star images correspond to the following field angles (in degrees): 0, 0.5, 1.0, 1.25, 1.5 and 1.75. Box size is 15 microns (2.2\"). Diffraction Airy spot is shown by the circle.} \end{figure} Mirror was polished by LOMO \footnote{Leningrad Optical Mechanical Association, Saint Petersburg, Russia} and Wynne corrector manufactured by TEC \footnote{Telescope Engineering Company, Colorado, USA}. The tube, mount and integration of the four of them were done by APM Telescopes \footnote{Amateur- und Pr\"azisionsoptik Mechanik, Germany}. \begin{table}[ht] \caption{Optical parameters and performance of the VT56s design from V.Yu. Terebizh for the TBT telescopes.} \label{tab:optical} \begin{center} \begin{tabular}{|l|c|} \hline \rule[-1ex]{0pt}{3.5ex} Entrance pupil diameter & 560 mm \\ \hline \rule[-1ex]{0pt}{3.5ex} Focal lenght & 1415.5 mm \\ \hline \rule[-1ex]{0pt}{3.5ex} Image space f\# & 2.53 \\ \hline \rule[-1ex]{0pt}{3.5ex} Plate scale & 6.86 $\mu$m/arcsec \\ \hline \rule[-1ex]{0pt}{3.5ex} Angular field of view & 3.5 deg \\ \hline \rule[-1ex]{0pt}{3.5ex} Linear obscuration & 0.41 \\ \hline \rule[-1ex]{0pt}{3.5ex} Effective aperture diameter & 511 mm \\ \hline \rule[-1ex]{0pt}{3.5ex} Back focal lenght & 121 mm \\ \hline \rule[-1ex]{0pt}{3.5ex} $D_{80}$ (from center to edge) & 9.1 – 11.4 $\mu$m (~1.3\" ~- 1.7\"~) \\ \hline \end{tabular} \end{center} \end{table} \subsection{Mount and pier} The mount is an APM GE-300 manufactured by Michael Knopf. It is able to load up to 300~kg, although our tube and accessories are significantly lighter. As the telescope is hosted in a clamshell dome, this mount was intentionally selected to perform better while tracking with moderated winds. The mount can move up to 20 degrees per second, but is limited at 12 degrees per second for TBT to increase slewing performance. This speed allows re-positioning in 2 seconds for neighbouring fields and maximum 20 seconds between any two locations in the sky. TBT telescopes are controlled by Sitech \footnote {Sidereal Technology, Oregon, USA} Telescope Control Systems. The mount has Renishaw absolute encoders with a resolution better than a tenth of arsecond. The servo controller system actuates over two direct driver motors in a close-loop at 2000~Hz. It has proportional-integral-derivative (PID) controller which has different slewing and tracking modes. These PID parameters were tuned for the requirements of the project. The mount is set on a pier that is tilted according to the latitude of the site. This design allows the movement of the tube under the mount, removing the collision risk. This position is only used to point to the flatscreen hanging on the wall of the dome. \subsection{CCD Camera and shutter} TBT telescopes are equipped with 800S Spectral Instruments cameras. This model has two 16-bit outputs, with readout values between 500~KHz and 1~MHz. Each camera hosts an E2V 230-84 chip with a midband astronomical coating, with a quantum efficiency QE peak \textgreater 95\%. It is a 4K~x~4K chip with 15-micron pixels, covering up to 61mm x 61mm. The resulting plate scale is 2.18 arcseconds per pixel, and the FoV is 2.5 degrees by 2.5 degrees. The camera reads the image from these chips in 12 seconds in the fastest mode. The read-out noise is less than 16 electrons for this mode, while the dark noise is always under 0.2 electrons per second and pixel, at -20 degrees Celsius. For the 500~KHz readout speed, the readout time is 22~s, but the readout noise is under 8 electrons. Linearity and CTE are better than 99\% and 0.999999 respectively. For survey the use of binning could reduce the readout time to 2 seconds and thus matching it with the repositioning dead time of the telescope, leading to more than 350 square degrees surveyed per hour (for 30 second exposure,magnitude $\sim$ 19 depth). For follow-up the use of windowing could also reduce the read-out time under 1 second, that is interesting for fast objects. The optical beam has an 80-mm shutter manufactured by BS \footnote{Bonn Shutter, Bonn, Germany}, to ensure a 1~ms timing accuracy, with a homogeneity of 0.5\% across the FoV. The shutter needs 150~ms to cover/uncover the whole beam, but the time across the chip can be mapped and corrected thanks to a high level of reproducibility. \subsection{Autonomous Emergency System} In parallel to the nominal systems, the TBT project has the Autonomous Emergency System\cite{ocana2015test,ocana_AES}. It is an Arduino-based system that monitors the environment with redundant sensors (rain, wind, light, power) and closes the dome in case of risk for the operations. It is designed to override the nominal system, though its safety parameters are more relaxed and it only takes control when the nominal system fails to react at the change of conditions. The system relies on two 12~V-batteries that are able to close the dome even at 25\% of charge when there is a shutdown. \subsection{Ancillary hardware} Apart from the main systems described here, the TBT project consists of many other elements that we describe briefly in this section. The control system takes the weather input from a Davis Vantage Pro 2 weather station. The outdoor sensors are located on a pole next to a nearby building, 4 meters away from the dome. The station is equipped with instruments to measure humidity/temperature (inside and outside the dome), rain and wind speed (gust and averages). The calibration of the images includes the use of flat images. The use of dusk/dawn sky images was discarded due to the large gradient present in the FoV. Therefore we use an electroluminiscent flatscreen with a smooth flat spectra covering from 400~nm to 1100~nm. The results are good, but the use of night images is promising. When the telescope archive have enough images, the use of night-sky flats is foreseen. In order to avoid hardware failures due to hard shutdowns, all the computers and sensitive electronics rely on a 3000 VA Uninterrupted Power Supply. This UPS is monitored by RTS2 in order to shut down properly all the equipment in case of a prolonged power outage. However the observatory relies on a short-break line, that is never down for more than a few seconds during the automatic start of the station backup diesel generators. The camera has a two-stage refrigeration system, where the hot plate of the peltier is cooled down by a 15 degrees Celsius water flux. We use a Thermocube chiller to circulate 2 liters of water per minute. The chiller is controlled and monitored by RTS2. Alarms for low/high temperature and low level of water are active, though in case of temperature increase in the peltier, the camera is automatically shutdown for protection. | \label{sec:conclusions} % TBT project will deploy two robotic telescopes to perform SST and NEO observations for the European Space Agency. The first telescope was installed in Cebreros ESTRACK Station and has been successfully commissioned. The second telescope will be installed and commissioned during 2017, after the confirmation of the location at the Southern Hemisphere. | 16 | 9 | 1609.07943 |
1609 | 1609.02561_arXiv.txt | We study the origin of the cold molecular clumps in quasar outflows, recently detected in CO and HCN emission. We first describe the physical properties of such radiation-driven outflows and show that a transition from a momentum- to an energy-driven flow must occur at a radial distance of $R \approx 0.25$ kpc. During this transition, the shell of swept up material fragments due to Rayleigh-Taylor instabilities, but these clumps contain little mass and are likely to be rapidly ablated by the hot gas in which they are immersed. We then explore an alternative scenario in which clumps form from thermal instabilities at $R \simgt 1$ kpc, possibly containing enough dust to catalyze molecule formation. We investigate this processes with 3D two-fluid (gas+dust) numerical simulations of a kpc$^3$ patch of the outflow, including atomic and dust cooling, thermal conduction, dust sputtering, and photoionization from the QSO radiation field. In all cases, dust grains are rapidly destroyed in $\approx 10^4$ years; and while some cold clumps form at later times, they are present only as transient features, which disappear as cooling becomes more widespread. In fact, we only find a stable two-phase medium with dense clumps if we artificially enhance the QSO radiation field by a factor 100. This result, together with the complete destruction of dust grains, renders the interpretation of molecular outflows a very challenging problem. | \label{Intro} Cold, fast-moving molecular clumps have been detected at distances up to a few kpc from the central engines of a few dozen active galactic nucleii (AGN), mostly highly-obscured AGN in dusty star forming galaxies \citep{Feruglio15, Cicone2014, Veilleux13}. In the nearby QSO Markarian 231, for example, a $1000 M_\sun$ yr$^{-1}$ kpc-scale outflow with a radial velocity of $750-1000 \kms$ has been well detected by several groups \citep{Feruglio10,Rupke11,Sturm2011}. Similarly, surveys of ultra-luminous infrared galaxies (ULIRGs) and quasar-hosts suggest that the presence of an AGN can boost molecular outflow rates by large factors \citep{Sturm2011,Cicone2014}. On the other hand, the origin of these molecular clumps is unclear. One possibility, in analogy with the prevalent picture for starburst driven winds, is that the clouds are driven out of the host galaxy by ram pressure acceleration \citep[e.g.][]{Veilleux2005}. However this hypothesis runs into serious difficulties, both because: (i) shocks and conduction from the exterior medium tend to compress the clouds perpendicular to the direction of the flow, greatly reducing the momentum flux they receive; and (ii) instabilities and evaporation lead to rapid cloud disruption \citep{Klein94,1994ApJ...434L..33M,2006A&A...457..545O,2008ApJ...678..274O,Scannapieco15,2016ApJ...822...31B}. Together these effects imply that the lifetimes of the clumps are likely to be much shorter that the times required to accelerate them to the observed speeds. A second possibility is that the clumps are formed from the cooling of very high temperature ($\approx 10^7$ K) shock-heated gas, already moving at high radial velocities. In, fact a similar picture has been suggested in the context of starburst-driven galaxy outflows \citep{1995ApJ...444..590W, 2000MNRAS.317..697E, 2003ApJ...590..791S, 2004ApJ...610..226S, 2007ApJ...658.1196T, 2011ApJ...740...75W, 2016MNRAS.455.1830T}, and ultra-luminous infrared galaxies hosting starbursts and AGN \citep{2015ApJ...803....6M}. However, in this case, one has not only to reconcile the clump formation process with the prevailing outflow scenario, but also explain how molecules like CO and HCN can form. These species require solid surfaces to trigger their production as gas phase production is likely to be highly inefficient in the presence of a strong UV background, and thus dust grains must be preserved throughout the entire heating/cooling cycle. This is not an easy requirement because grains are sputtered by nuclear and grain-grain collisions and thus rapidly destroyed in such hot environments. The main motivation of this paper is to explore the possibility that the observed molecular clumps condense out of the quasar outflow material as a result of thermal instabilities occurring in the cooling gas. Carrying out a detailed examination of the dynamical structure of QSO outflows, we conclude that the optimal location for the formation of dusty clumps is within the shocked gas during the energy-driven stage of the evolution. We then carry out a set of detailed numerical simulations of clump formation at this stage, which include dust cooling and destruction, electron thermal conduction, and atomic/ionic cooling and heating process that account for the QSO radiation field. The inclusion of this extended number of physical processes allows us to draw a series of robust conclusions as to the feasibility of the formation molecular clumps though condensation in QSO driven outflows. The structure of this paper is as follows. In \S2 we describe the physics of QSO driven outflows and the cooling of shocked material. In \S3 we describe our initial conductions and simulation methods, and in \S4 we present our results. We conclude in \S5 with a discussion of the implications of these results for clump formation in QSO outflows, and we include appendices that discuss the formation of multiphase media and self-shielding in further detail. | We have studied the origin of the cold molecular clumps recently detected in CO and HCN emission in QSO outflows. We first described the physical properties of the radiation-driven outflow and show that a transition from a momentum- to a energy-driven case must occur at radial distances of $R \approx 0.25$ kpc. During this transition, the dense shell fragments due to Rayleigh-Taylor instabilities, but we find that these clumps are likely to contain very little mass, and they are rapidly ablated and destroyed by the hot gas in which they are immersed. We have thus also explored an alternative scenario in which clumps form from thermal instabilities at $R \geq 1$ kpc, possibly containing enough dust to catalyze molecule formation. We investigate this processes with 3D two-fluid (gas + dust) numerical simulations of a kpc$^3$ patch of the outflow, including atomic and dust cooling, thermal conduction, dust sputtering, photoionization from the QSO radiation field, and self-shielding (see Appendix B). We find that in all cases dust grains are rapidly destroyed during $\approx 10^4$ years; however, and while cool clumps are present in the fiducial run, they appear only as a transient feature that is washed away as cooling is completed. In fact, a stable two-phase medium with dense clumps is found only if we artificially enhance the QSO radiation field by a factor 100. This result, together with the complete destruction of the dust grains, renders the interpretation of molecular outflows a very challenging problem. We pause for a caveat. We have shown that the cooling post-shock gas cannot be piled-up in a well-formed shell due to the fact that QSO outflows, expanding in the steep halo density profile, are Rayleigh-Taylor unstable at all times. In fact, the arguments given for the momentum-driven shell in Sec. 2.2 apply exactly to the energy-driven phase we are discussing. The hot shocked wind gas wraps clumps around and rapidly shreds them. However, if the gas is thermally unstable, fragments form by thermal instability on a timescale $< 1$ Myr (from our simulations) that is comparable or shorter than the RT growth timescale (see eq. 17, but with $R_C$ now at least 10 times larger). This might lead to a different scenario, that is the one we have explored here. Cold clumps are embedded in a hot interclump medium; the two components are in pressure equilibrium. Hence, clumps are protected from ablation by the hot gas surrounding them, and might survive longer. A more robust conclusion on the clump survival in a accelerated two-phase medium requires additional study. Nevertheless, the above qualitative arguments suggest that this situation is more promising than the RT unstable case. How the physical conditions leading to clump formation might be realized in the environment of quasars is unclear. Contrary to intuition, cold clumps do not form in the fiducial, standard conditions because the heating is insufficient to support a stable hot phase until the medium as a whole has cooled. Instead, a stable two phase medium would require a much stronger energy input from the quasar that is inconsistent with the radiated power. One might speculate that other heating mechanisms might be at work. The most likely among these is the contribution of relativistic particles. Although this solution appears unlikely, it is certainly worth further scrutiny, given the thorny questions open by our investigation. Our simulations include all the most relevant physics apart from magnetic fields, self-gravity, and the impact of the bulk motion of the medium as clump formation occurs. Although anisotropic conduction due to the presence of magnetic fields could slow electron thermal conduction, conduction not appear to be a key process in determining the final state of the gas. Furthermore, while gravity in principle could produce a collapse of the cold clumps, we show in Appendix A, self-gravity can be safely neglected as a result of the outflow expansion, as the gas is like torn apart by expansion on a much shorter time-scale that it can gravitationally collapse. This point adds yet another puzzling problem. If self-gravity cannot confine the clumps (in case they form), then they must be pressure confined. However, as the outflow expands to larger radii, the internal density drops and the clumps are progressively less shielded agains the UV quasar radiation acting to dissociate CO molecules. Moreover, if densities become low enough, the cold gas can also become ionized. This might be broadly consistent with the evidence of ionized gas recently found at large distances from the QSO but still clearly associated with it \citep{Nesvadba08, Cresci15, Carniani15}. We conclude that the presence of cold molecular clumps in QSO outflows represents a difficult theoretical challenge. Our study shows that neither a scenario in which these components are galactic clouds entrained and accelerated by the outflow, nor the one in which they condense out of the fast, outward moving gas appear to be viable under ``normal'' conditions. Thus, the solution to the problem is in demand of further investigations. | 16 | 9 | 1609.02561 |
1609 | 1609.07172_arXiv.txt | The emission of the white dwarf-M dwarf binary AR Sco is driven by the rapid synchronization of its white dwarf, rather than by accretion. Synchronization requires a magnetic field $\sim 100$ gauss at the M dwarf and $\sim 10^8$ gauss on the white dwarf, larger than the fields of most intermediate polars but within the range of fields of known magnetic white dwarfs. The spindown power is dissipated in the atmosphere of the M dwarf, within the near zone of the rotating white dwarf's field, by magnetic reconnection, accelerating particles that produce the observed synchrotron radiation. The displacement of the optical maximum from conjunction may be explained either by dissipation in a bow wave as the white dwarf's magnetic field sweeps past the M dwarf or by a misaligned white dwarf rotation axis and oblique magnetic moment. In the latter case the rotation axis precesses with a period of decades, predicting a drift in the orbital phase of maximum. Binaries whose emission is powered by synchronization may be termed synchronars, in analogy to magnetars. | \cite{M16,B16} recently discovered that AR Sco is a M dwarf/white dwarf binary with an orbital period of 3.56 h and a white dwarf spin period of 1.95 m. The white dwarf's spin is slowing with a characteristic time $P/{\dot P} = 0.9 \times 10^7$ y. Most of the radiated power is emitted from the M dwarf near the L$_1$ inner Lagrange point, but the phase of maximum is displaced from conjunction by about 0.15 of the orbit. AR Sco emits X-rays but they are only a small fraction of the system's luminosity, implying that it is not powered by accretion. Spindown is the remaining plausible source of energy, with a power $- I \omega {\dot \omega}$, where $I$ is the white dwarf's moment of inertia and $\omega$ its rotation rate\footnote{Properly, $\omega$ is the WD rotation rate in a rotating frame in which the total angular momentum of the binary is zero. To an excellent approximation, this is the frame rotating at the orbital rate, in which the WD spin rate is about 1\% less than in an inertial frame. The observed optical modulation is at this lower frequency.}. The spindown power is converted to radiation with efficiency several times the 2--3\% of $4\pi$ sterad subtended at the white dwarf by its companion. This excludes reprocessing of roughly isotropic radiation from the white dwarf as the source of the observed radiation. I follow \cite{M16} in assuming the white dwarf's rotational energy is dissipated by interaction of its magnetic field with the M dwarf's atmosphere. The spectral energy distribution and polarization \citep{B16} indicate a substantial contribution of synchrotron radiation, implying that the dissipation is by nonthermal processes, such as magnetic reconnection, that accelerate energetic particles. The rapid spindown suggests that AR Sco may be a missing link between synchronously rotating polars (AM Her stars) whose synchronism is maintained by magnetostatic interaction \citep{JKR79} to extreme accuracy \citep{K89} and the asynchronous intermediate polars (IP). Because AR Sco is still far from synchronism the rate of dissipation of its white dwarf's rotational energy is much greater than that in almost-synchronous systems like V1500 Cyg \citep{K91a,K91b}. The visible brightness of AR Sco does not peak at conjunction (orbital phase $\phi = 0.5$), when we might expect to see the heated side of the M dwarf most fully, but rather around phase $\phi \approx 0.35$. The purpose of this paper is to investigate the origin of this surprising observation. I consider two hypotheses: \begin{enumerate} \item The power dissipated by interaction between the white dwarf's magnetic field and the M dwarf is greater on the latter's leading face where a bow wave may form (Fig.~\ref{ARScoF1}) than on its trailing face. The dissipation rate does not depend on orbital phase but we view the hottest portion of the M star more fully before conjunction. \item The white dwarf's spin axis is not aligned with the orbital axis and its magnetic moment is not aligned with its spin axis. The dissipation rate in the M dwarf's atmosphere depends on the orbital phase (Fig.~\ref{ARScoF2}). Precession of the spin axis makes the orbital phase of maximum drift. \end{enumerate} In either hypothesis the optical maximum is displaced from conjunction. In the first hypothesis it precedes conjunction if the white dwarf spin is prograde (and follows it if retrograde). In the second hypothesis any phase is possible, but precession makes it recede (if the spin is prograde). \begin{figure} \centering \includegraphics[width=3.3in]{ARScoF1.pdf} \caption{\label{ARScoF1} The white dwarf's field is moving past the M dwarf at a speed $a\omega$, where $a$ is the separation of the stars. This is fast enough (15--20\% of $c$) to induce an asymmetry in the magnetic interaction between the ``upstream'' and ``downstream'' faces of the M dwarf, even though there is not expected to be a particle wind from the white dwarf. The white dwarf's magnetic field is an electromagnetic wind, even inside the radiation zone (though not deep in the near zone, close to the white dwarf, where a magnetostatic treatment should be valid). More energy is released in a bow wave on the upstream side. The Figure shows a magnetic moment misaligned with the white dwarf's spin axis, but this is not necessary; an upstream/downstream asymmetry and a torque would occur even if the orbital, spin and magnetic axes were parallel, provided $a \omega {\hskip 0.4mm \not \hskip -0.4mm \ll} c$.} \end{figure} \begin{figure} \centering \includegraphics[width=3.3in]{ARScoF2.pdf} \caption{\label{ARScoF2} The white dwarf's spin $\vec \omega$ makes an angle $\epsilon$ with the orbital axis, and the spin axis precesses. If the magnetic moment is oblique (shown here as perpendicular) to the spin axis, as indicated by the large amplitude modulation at the sideband of the spin frequency, the magnetic field is maximum in a broad fan, indicated by dashed lines, around the rotational equator, and dissipation is maximum at orbital phases when the M dwarf (shown at three orbital phases) is in that fan. In the Figure that occurs at $\phi \approx 0.35$ and $\phi \approx 0.85$ (at $\phi \approx 0.85$ the heated surface is not visible to the observer).} \end{figure} \cite{GZH16} and \cite{B16} have suggested that the white dwarf acts as a pulsar. But the M dwarf is well within the near zone of the white dwarf's dipole radiation field, where the non-radiative Poynting vector exceeds the radiative Poynting vector by ${\cal O} (c/a \omega)^4 \sim 10^3$--$10^4$. \cite{M16} also pointed out the difficulty of coupling an estimated 9\% of the spin-down power, if radiated, to a M dwarf that occupies only 2--3\% of the white dwarf's sky. This again argues against mechanisms that radiate the energy in a broad beam, such as that of a rotating dipole, and indicates non-radiative coupling. I therefore attribute the dissipation to the interaction of the white dwarf's quasi-static vacuum magnetic field with the M dwarf's atmosphere. This interaction and the dissipation rate vary with orbital phase, shifting maximum light from conjunction, even if the white dwarf's rotation and magnetic axes are aligned (that alignment is excluded in AR Sco because of the large amplitude modulation at the sideband of the spin period). The irregularly fluctuating brightness near the orbital phase of maximum may be attributed to the magnetic storms characteristic of magnetic reconnection. | The remarkable properties of AR Sco support the suggestion of \cite{M16} that it differs from most intermediate polars in that it is powered by synchronization of the white dwarf spin with the orbital motion (a ``synchronar'') rather than by accretion. The displacement of the optical maximum from conjunction implies that energy is deposited asymmetrically on the surface of the M dwarf. This may be explained either as a consequence of the subrelativistic phase speed of the rotating white dwarf magnetosphere at the M dwarf or of a white dwarf spin axis that is oblique to the orbital axis. In the former hypothesis the orbital phase of maximum is likely to be fixed; in the latter hypothesis the spin axis precesses with a period of decades, producing a drift in the orbital phase of maximum. In at least one other IP, AE Aqr \citep{IB08,M09}, synchronization is also the chief source of energy and radio emission indicates particle acceleration. The photometry of AE Aqr is complicated by its flaring, but the amplitude of the periodic modulation is low and is explicable as an ellipsoidal variation without evidence of a reflection effect or other interaction with the white dwarf \citep{vP89,B91}. The secondary star of AE Aqr is of earlier spectral type (K4V) than that of AR Sco (M5V) and likely intrinsically more luminous, reducing the effect of magnetic heating on the light curve. AR Sco and AE Aqr may be members of a class of synchronars, magnetic binaries whose emission is powered by the synchronization of their compact (white dwarf) asynchronously rotating components. AR Sco is an extreme example in which the contribution of intrinsic stellar and accretional luminosity is particularly small. Accretion contributes only a small fraction, perhaps none, of the luminosity of AR Sco. This is likely attributable to its strongly magnetized white dwarf. AR Sco also differs from non-accreting binary (neutron star) pulsars with nondegenerate companions because the binary pulsars' companions are in the radiation zones of the fast-spinning pulsars' magnetic fields, so any interaction between the companion and the field does not couple back to the pulsars' rotation. In such a system pulsar spin energy may be deposited in the companion, but the pulsar radiates in a broad dipole pattern. This limits the coupling efficiency to no more than twice the companion's subtended solid angle at the pulsar, inconsistent with observations of AR Sco, and the pulsar despins as would an isolated object. In contrast, in AR Sco the coupling is near-zone, though not quite quasi-static, rather than radiative. Its efficiency is not limited by the subtended solid angle and, as in a well-designed transformer, can approach 100\%. | 16 | 9 | 1609.07172 |
1609 | 1609.07491_arXiv.txt | The observed scale heights of extraplanar diffuse ionized gas (eDIG) layers exceed their thermal scale heights by a factor of a few in the Milky Way and other nearby edge-on disk galaxies. Here, we test a dynamical equilibrium model of the extraplanar diffuse ionized gas layer in NGC 891, where we ask whether the thermal, turbulent, magnetic field, and cosmic ray pressure gradients are sufficient to support the layer. In optical emission line spectroscopy from the SparsePak integral field unit on the WIYN 3.5-meter telescope, the H$\alpha$ emission in position-velocity space suggests that the eDIG is found in a ring between galactocentric radii of $R_{min} \le R \le 8$ kpc, where $R_{min} \ge 2$ kpc. We find that the thermal ($\sigma_{th} = 11$ \kms) and turbulent ($\sigma_{turb} = 25$ \kms) velocity dispersions are insufficient to satisfy the hydrostatic equilibrium equation given an exponential electron scale height of $h_{z} = 1.0$ kpc. Using a literature analysis of radio continuum observations from the CHANG-ES survey, we demonstrate that the magnetic field and cosmic ray pressure gradients are sufficient to stably support the gas at $R \geq 8$ kpc if the cosmic rays are sufficiently coupled to the system ($\gamma_{cr} = 1.45$). Thus, a stable dynamical equilibrium model is viable only if the extraplanar diffuse ionized gas is found in a thin ring around $R = 8$ kpc, and non-equilibrium models such as a galactic fountain flow are of interest for further study. | Multi-wavelength observations of nearby edge-on disk galaxies have revealed multi-phase gaseous halos that include molecular, neutral, and warm and hot ionized phases. The extraplanar diffuse ionized gas (eDIG) layers in the Milky Way and other nearby edge-on disk galaxies are remarkable in that their observed scale heights generally exceed their thermal scale heights by a factor of a few \citep[e.g.,][]{Rand1997, Haffner1999, Collins2001, Gaensler2008, Voigtlander2013}. These warm ($T \sim 10^{4}$ K), diffuse ($\langle n_{e,0} \rangle \sim 0.1$ cm$^{-3}$) layers have a range of diffuse, clumpy, and filamentary morphologies, a photoionization power requirement that is satisfied by the O and B stars in the disk, and rotational velocity profiles that suggest a probable disk origin \citep{Lehnert1995, Lehnert1996, Rossa2000, Tullmann2000, Miller2003b, Miller2003a, Rossa2003a, Rossa2003b, Heald2006b, Heald2006a, Heald2007}. Additionally, the detection of eDIG layers is positively correlated with the star formation rate per unit area for starburst, star-forming, and quiescent galaxies \citep{Rossa2003a}. It is also spatially correlated with soft X-ray emission from hot halo gas \citep{Strickland2004, Tullmann2006b, Tullmann2006a}, as well as with radio continuum emission associated with extraplanar magnetic fields and cosmic rays \citep{Dahlem1994, Collins2000, Tullmann2000, Li2016}. The observation that eDIG layers are associated with a minimum star formation rate per unit area is consistent with models of a star-formation driven disk-halo flow. ``Superbubble'' \citep{MacLow1988}, ``galactic chimney'' \citep{Norman1989}, ``galactic fountain'' \citep{Shapiro1976}, and galactic wind \citep[e.g.,][]{Veilleux2005} models all describe the local or global circulation of gas between the disk and the halo due to star formation activity in OB associations. There is observational evidence of bubbles, arcs, and filaments in the halo that are spatially associated with HII regions as well as ultraviolet continuum from young, hot stellar populations in the disk \citep{Dettmar1990, Rand1990, Rand1996, Howk1997, Howk1999, Howk2000, Rossa2000, Rossa2003b, Tullmann2006b}. Thus, a general framework has emerged in which eDIG layers are found in multi-phase gaseous, magnetic field, and cosmic ray halos in galaxies with sufficient star formation rates per unit area. Within this framework, the vertical structure, support, and dynamical state of these layers are not yet fully understood. A range of dynamical models exist to explain the column densities, scale heights, and three-dimensional kinematics of the extraplanar interstellar medium (ISM). One class of models treats extraplanar gas layers as fluid disks that satisfy the hydrostatic equilibrium equation \citep{Boulares1990, Barnabe2006, Henriksen2016}, while another treats the layers as collections of clouds that travel ballistically through the galactic gravitational potential \citep{Collins2002, Fraternali2006}. Some authors suggest that a combination of hydrodynamic and ballistic effects may be closest to reality \citep[e.g.,][]{Benjamin2000}, while others seek to understand the effects of magnetohydrodynamics on the disk-halo interface in a turbulent, star-forming ISM \citep{Hill2012}. Discriminating between dynamical models for each phase of the extraplanar ISM is important for undestanding how each phase is formed, evolves, and participates in the transfer of mass and energy between the disk, halo, and intergalactic environment. Here, we study the dynamical state of the eDIG layer in the nearby edge-on disk galaxy NGC 891. This galaxy is an ideal candidate for this study due to its proximity ($D = 9.9$ Mpc; $1" = 48$ pc; \citealt{Ciardullo1991}), inclination angle ($i > 89^{\circ}$; \citealt{Oosterloo2007}), and well-studied multi-phase gaseous halo. It is classified as an Sb galaxy in the Third Reference Catalogue of Bright Galaxies \citep{deVaucouleurs1991}, but there is evidence at multiple wavelengths for a bar \citep[e.g.,][]{Sofue1993, GarciaBurillo1995, SchechtmanRook2013}. Due to similarities in mass, morphology, and bolometric luminosity, NGC 891 is often considered a Milky Way analog \citep{vanderKruit1984}; however, the far-infrared star formation rate is somewhat higher in the former at $3.8$ \Msun yr$^{-1}$ \citep{Popescu2004}. There is evidence in the HII region number density as well as the far-infrared and radio emission morphology that the star formation rate is highest in the northeast side of the disk \citep[e.g.,][]{Wainscoat1987, Dettmar1990}. There is not evidence of a major disturbance of the stellar disk in deep optical and near-infrared photometry \citep{Morrison1997, SchechtmanRook2013}. However, \citet{Mapelli2008} show that the slight lopsidedness of the disk suggests a mild flyby interaction with the companion UGC 1807. \citet{Oosterloo2007} demonstrate that HI clouds with counter-rotating velocities and an HI filament near systemic velocity reaching over 20 kpc from the disk in projection towards the companion are evidence of interaction and/or accretion. This system also includes extraplanar dust \citep{Howk1997, Howk2000, Seon2014}, diffuse ionized gas \citep{Dettmar1990, Rand1990, Rand1997}, and hot ionized gas \citep{HodgesKluck2013}, as well as extraplanar magnetic fields and cosmic rays \citep{Dahlem1994}. The eDIG in NGC 891 is among the brightest, most spatially extended, and most well-studied eDIG layers known. Discovered in H$\alpha$ narrowband imaging by \citet{Dettmar1990} and \citet*{Rand1990}, the brightest and most vertically extended eDIG is found on the northeast side of the galaxy, where it appears to be spatially associated with the elevated star formation rate \citep{Dettmar1990, Hoopes1999}. The morphology of the layer has both smooth and filamentary components; \citet{Rossa2004} obtained high spectal resolution ($0.1'' = 4.8$ pc) H$\alpha$ narrowband imaging with the WFPC2 camera on the \textit{Hubble Space Telescope} that revealed a diffuse background intersected by filaments, arcs, plumes, bubbles, and supershells. Notably, \citet{Howk2000} and \citet{Rossa2004} detect arcs and filaments that have dimensions of tens of pc wide and several kpc long, are highly collimated to large heights above the disk, and appear to have one or both ends in star-forming regions. Qualitatively, these observations suggest that star formation activity drives the warm ionized gas out of the disk by way of galactic chimneys formed from the bursting of superbubbles associated with spatially correlated supernovae \citep{Shapiro1976, Norman1989}. The photoionization energy requirement of the eDIG layer is met by massive stars in the disk if the ISM is sufficiently porous to UV photons \citep[e.g.,][]{Dettmar1990}. The emission line spectrum is broadly consistent with a photoionized gas in the near-ultraviolet \citep{Otte2001, Otte2002}, optical \citep[e.g.,][]{Rand1997, Rand1998}, and infrared \citep{Rand2008b, Rand2011}. However, in NGC 891, the Milky Way, and other galaxies, the emission line ratios as a function of height above the disk requires a supplemental source of heating and/or ionization; such sources may include shocks \citep{Rand1998}, turbulent mixing layers \citep{Rand1998, Binette2009}, hot low-mass evolved stars \citep{Sokolowski1991, FloresFajardo2011}, and/or cosmic rays \citep{Wiener2013}. A remarkable feature of the eDIG layer in NGC 891 is its considerable spatial extent above and below the midplane. The vertical electron density distribution is well-described by an exponential of the form $\langle n_{e}(z) \rangle = \langle n_{e,0} \rangle e^{-|z|/h_{z}}$, where $\langle n_{e,0} \rangle $ is the mean electron number density in the disk and $h_{z}$ is the electron scale height. The eDIG layer in NGC 891 is well-fit by a scale height of $h_{z} = 1.0$ kpc on the Northeast side of the disk \citep{Dettmar1990, Rand1990, Dettmar1991, Keppel1991}. An improved fit is found if the electron density distribution is expressed as the sum of a thick disk component with $h_{z,disk} = 1.0$ kpc and a halo component with $h_{z,halo} \sim$ a few kpc \citep{Rand1997, Hoopes1999}. The thick disk and halo components may be produced by different processes; for example, the former may be rising out of the disk via galactic chimneys and supershells, while the latter may be condensing onto the disk out of a hot halo phase \citep{Rand1997}. Changing emission line ratios with distance from the midplane suggest that the large scale height is due to true extraplanar emission and not to HII region emission scattered by dust. Additionally, \citet{Ferrara1996} use Monte Carlo radiative transfer simulations of H$\alpha$ photon propagation through the dusty disk of NGC 891 to argue that scattered HII region emission is only 10\% of eDIG emission at $z = 600$ pc. Here, we test a dynamical equilibrium model of the eDIG layer in NGC 891. Although the observed scale height, lack of flaring, and general inhomogeneity of the layer suggest a system out of dynamical equilibrium \citep[e.g.,][]{Dettmar1990}, the various sources of vertical support have yet to be fully quantified for any eDIG layer. Thus, we use optical emission line spectroscopy from the SparsePak integral field unit (IFU; \citealt{Bershady2004, Bershady2005}) on the WIYN 3.5-meter telescope at Kitt Peak National Observatory, as well as radio continuum observations from the CHANG-ES survey from \citet[][in preparation]{Schmidt2016}, to determine the thermal, turbulent, magnetic field, and cosmic ray pressure gradients in the eDIG layer. By comparing the observed and required pressure gradients to support the eDIG layer at its observed scale height, we consider whether the system is best characterized by equilibrium or non-equilibrium (i.e., galactic fountain, galactic wind) models. The paper is laid out as follows. In \S2, we create a mass model to determine the galactic gravitational potential of NGC 891. We give a statement of the problem and the model to be tested in \S3, and we discuss the collection and reduction of optical emission line spectroscopy using the SparsePak IFU in \S4. In \S5.1, we construct a model of the three-dimensional density distribution of the eDIG layer, and we constrain the velocity dispersion along the minor axis from the H$\alpha$, [NII] $\lambda$6583, and [SII] $\lambda$6716, 6731 emission line widths in \S5.2. In \S5.3, we determine the vertical magnetic field and cosmic ray pressure gradients in this system using radio continuum observations from the CHANG-ES survey analyzed by \citet[][in preparation]{Schmidt2016}, and in \S5.4 we assess whether a magnetized eDIG layer in dynamical equilibrium is stable against the Parker instability. We discuss our results in the context of multi-wavelength observations and our knowledge of the Milky Way Galaxy in \S6. In \S7, we conclude that a dynamical equilibrium model dominated by a magnetic pressure gradient is viable for the eDIG layer in NGC 891 only over a limited range of galactocentric radii ($R \sim 8$ kpc). We include an Appendix to illustrate the robustness of this result against variations in the assumed mass-to-light ratio of the stellar disk. | \label{sec_conc} We sought to determine the dynamical state of the eDIG layer in NGC 891. This layer is remarkable due to an observed scale height that exceeds its thermal scale height by a factor of a few. Specifically, we tested a dynamical equilibrium model by quantifying the thermal, turbulent, magnetic field, and cosmic ray pressure gradients in this galaxy. We summarize our results as follows: \begin{enumerate} \item We obtained optical emission line spectroscopy of the eDIG layer using the SparsePak IFU on the WIYN telescope. We probed a wide range in height above and below the disk ($0 \leq |z| \leq 3.2$ kpc) and in projected radius ($-1.65 \leq R' \leq 7.35$ kpc) with moderate spectral resolution ($\sigma = 17$ \kms at H$\alpha$). We found a thick disk exponential electron scale height of $h_{z} = 0.8$ kpc and $h_{z} = 1.2$ kpc on the East and West sides of the galaxy, respectively. This is consistent with past measurements in the literature. \item Several pieces of evidence point to the eDIG being found preferentially at moderate galactocentric radius ($R_{min} \leq R \leq 8$ kpc, where $R_{min} \ge 2$ kpc). These include the comparable $I_{H\alpha}(z)$ on and off of the minor axis, the lack of low- and high-velocity emission line wings, and the location of the observed velocity centroids in position-velocity space. \item We measured the H$\alpha$, [NII] $\lambda$6583, and [SII] $\lambda$6716, 6731 emission line widths along the minor axis, and show that they are consistent with a turbulent medium with a sonic Mach number of $M = 2 - 3$. The thermal ($\sigma_{th} = 11$ \kms) and turbulent ($\sigma_{turb} = 25$ \kms) velocity dispersions are far below that required to support the eDIG layer by thermal and turbulent pressure gradients between $R = 0$ and the observed cutoff at $R = 8$ kpc ($\sigma = 210 - 40$ \kms). The observed turbulent velocity dispersion is supersonic for the warm phase, but is subsonic for the hot halo; this is consistent with the eDIG being a collection of cool clouds embedded in a hot, potentially outflowing phase. \item We referred to an analysis by \citet[][in preparation]{Schmidt2016} of the synchrotron halo in NGC 891 using CHANG-ES radio continuum observations to determine the magnetic field and cosmic ray pressure gradients in this system. The combined thermal, turbulent, magnetic field, and cosmic ray pressure gradients are sufficient to support the eDIG layer at a scale height of $h_{z} = 1$ kpc at galactocentric radii of $R \geq 8$ kpc. The uncertainty on the eDIG filling factor, magnetic field strength, and magnetic scale height yield an uncertainty of a few kpc on this galactocentric radius, and thus it is possible that the eDIG is supported in dynamical equilibrium in a thin ring between $6 \le R \le 8$ kpc. \item Our dynamical equilibrium model of a magnetized eDIG layer is Parker stable if the cosmic rays are sufficiently coupled to the system ($\gamma_{cr} = 1.45$). \item Similarities between the thermal and turbulent properties of the warm ionized gas as well as the synchrotron halos in NGC 891 and the Milky Way Galaxy suggest that extraplanar magnetic fields and cosmic rays may play an important role in the dynamical state of the Reynolds layer. \item In future work, a simultaneous treatment of the extraplanar cold, warm, and hot gas is desired to understand the dynamical state of the multi-phase gaseous halo. \end{enumerate} Studying the eDIG layers of nearby edge-on disk galaxies has both advantages and disadvantages. We are able to directly determine the vertical scale height of the layer and quantify the gas, magnetic field, and cosmic ray pressure as functions of height above the disk. However, we cannot directly measure the vertical velocity dispersion or look for evidence of vertical inflow, outflow, or other non-equilibrium phenomena. We also cannot definitively determine the location of the gas along the line of sight, or search for spatial correlations between the properties of the halo gas and the underlying disk. Thus, fully characterizing the dynamical state of eDIG layers requires observations of disk galaxies with a range of inclination angles. Future work will focus on face-on disk galaxies at high spectral resolution to directly measure the vertial velocity dispersion and any vertical bulk flows, determine the radial distribution of the gas, explore connections between the gaseous halo, the underlying stellar disk, and the intergalactic environment, and assess the generality of the model presented here (Boettcher, Gallagher, Zweibel, \& Benjamin, in preparation). The future is bright for understanding the dynamics of extraplanar gas through careful studies of small samples of nearby galaxies, as well as through ongoing surveys studying the gaseous, magnetic field, and cosmic ray halos in these systems. Large IFU surveys including the Mapping Nearby Galaxies at Apache Point Observatory (MaNGA) project are enabling the statistical study of extraplanar gas properties in hundreds of low-redshift galaxies as part of the fourth-generation Sloan Digital Sky Survey (SDSS-IV) \citep{Bundy2015}. Additionally, the CHANG-ES survey will allow the extraplanar magnetic field and cosmic ray properties of tens of nearby galaxies to be characterized, and thus the relationship between gaseous and synchrotron halos to the studied in a statistical sense. | 16 | 9 | 1609.07491 |
1609 | 1609.09503_arXiv.txt | We present a systematic search for periodically varying quasars and supermassive black hole binary (SMBHB) candidates in the Pan-STARRS1 (PS1) Medium Deep Survey's MD09 field. From a color-selected sample of 670 quasars extracted from a multi-band deep-stack catalog of point sources, we locally select variable quasars and look for coherent periods with the Lomb--Scargle periodogram. Three candidates from our sample demonstrate strong variability for more than $\sim$ 3 cycles, and their PS1 light curves are well fitted to sinusoidal functions. We test the persistence of the candidates' apparent periodic variations detected during the $4.2$ years of the PS1 survey with archival photometric data from the SDSS Stripe 82 survey or new monitoring with the Large Monolithic Imager at the Discovery Channel Telescope. None of the three periodic candidates (including PSO J334.2028+1.4075) remain persistent over the extended baseline of $7 - 14$ years, corresponding to a detection rate of $<$ 1 in 670 quasars in a search area of $\approx$ 5 deg$^2$. Even though SMBHBs should be a common product of the hierarchal growth of galaxies, and periodic variability in SMBHBs has been theoretically predicted, a systematic search for such signatures in a large optical survey is strongly limited by its temporal baseline and the ``red noise'' associated with normal quasar variability. We show that follow-up long-term monitoring ($\gtrsim 5$ cycles) is crucial to our search for these systems. | \label{sec:intro} Supermassive black holes (SMBHs) appear to be at the centers of most, perhaps all, massive galaxies (e.g. \citealt{Kormendy1995}). Thus, when two massive galaxies merge in the $\Lambda$CDM Universe, it is expected that their nuclei will form a supermassive black hole binary (SMBHB; e.g. \citealt{Springel2005}). As the binary coalesces, the early stage of its orbital decay is driven by exchanging angular momentum with the circumbinary gas disk through viscosity; at smaller separations ($a < 1$ pc), its orbital decay becomes more dominated by gravitational wave (GW) radiation (e.g. \citealt{Begelman1980}). However, sub-parsec separation SMBHBs at cosmological distances are too compact to resolve with current, or even future, telescopes. Indirect searches so far, therefore, have been focused on spectroscopy, looking for offset broad lines that suggest two broad line emission regions, each likely associated with each black hole in the binary system \citep{Boroson2009}, or offset or shifted peak of the broad line region (e.g. \citealt{Dotti2009, Eracleous2012}). Another observational aspect of SMBHBs, however, was much under-exploited until recently --- their potential optical variability. One of the first sub-parsec SMBHB candidates identified via its variability was OJ287 \citep{Sillanpaa1988}, which showed quasi-periodic optical outbursts at intervals of 12 years, with the physical interpretation of the burst being the secondary black hole passing through the accretion disk of the primary (e.g. \citealt{Lehto1996, Valtonen2008, Valtonen2011}). More recently, another sub-parsec SMBHB candidate, PG 1302-102 \citep{Graham2015Nat}, was discovered by the Catalina Real-time Transient Survey (CRTS; \citealt{Drake2009}). Its $V$-band light curve can be fitted to a sinusoidal function with period of 1,884 days and amplitude of 0.14 mag. A physical interpretation of PG 1302-102's periodic variability is relativistic Doppler boosting \citep{DOrazio2015Nat}: in this scenario, where the luminosity is dominated by the steadily accreting secondary black hole and the system is viewed at a high inclination angle, emission from the minidisk of the secondary is Doppler-boosted as the black hole orbits at a moderately relativistic speed (along the line of sight). Another possible scenario that could give rise to periodic variability is modulated mass accretion in the system. Simulations of an SMBHB embedded in a circumbinary disk show that although the binary tidal torque clears and maintains a low gas density cavity at radius $< 2a$ (where $a$ is the binary separation), materials can penetrate the cavity through a pair of streams and be accreted onto the binary. These simulations have the similar results that for a mass ratio $0.01 \lesssim q \leqslant 1$ --- as expected in the merger of two massive galaxies --- mass accretion through the circumbinary disk is strongly modulated as a result of the binary's orbital motion within the circumbinary disk, including two-dimensional (2D) hydrodynamical \citep{MacFadyen2008}, 3D Newtonian magnetohydrodynamical (MHD) \citep{Shi2012} and Post-Newtonian MHD \citep{Noble2012} for an equal mass binary, and general relativistic (GR) MHD \citep{Gold2014} and 2D hydrodynamical simulations \citep{DOrazio2013} for various mass ratios. In these simulations, the accretion rate varies on a time scale that is on the order of the binary orbital time scale, which is in turn a function of the total black hole mass and orbital separation by virtue of Kepler's law. Assuming that luminosity tracks mass accretion of the circumbinary disk, the former should then vary as the latter varies. For a typical black hole mass of $10^7 M_\odot$ and typical separation $10^3 R_s$, the orbital period is on the order of $\sim$ year, an observationally feasible time scale for current time-domain surveys: $t_{\rm orb} = 0.88\,\mbox{yr} \left(\frac{M}{10^{7} M_\odot}\right) \left(\frac{a}{10^3 R_{s}}\right)^{3/2}$ (where $R_s$ is the Schwarzschild radius: $R_{\rm s} = 2GM/c^2$). These theoretically explored variability signatures of an SMBHB, as well as encouraging predictions for the detection rates of periodically varying quasars from SMBHBs in a cosmological context \citep{Haiman2009}, motivated several recent systematic searches in large optical time-domain surveys with a temporal baseline of several years --- \cite{Graham2015Nat} and \cite{Graham2015}, with the CRTS; \cite{Charisi2016}, with Palomar Transient Factory (PTF) and additional data from intermediate-PTF and CRTS; \cite{Zheng2015}, with the Sloan Digital Sky Survey (SDSS) and CRTS; and \cite{Liu2015}, with the Pan-STARRS1 Medium Deep Survey (MDS). In our pilot study (\citealt{Liu2015}, hereafter L15), we performed a systematic search for SMBHB candidates in MDS's MD09 field and reported our first significant detection of such a candidate, PSO J334.2028+1.4075. As reported in L15, PSO J334.2028+1.4075 has a coherent period of P = 542 $\pm$ 15 days in $g_{\rm P1}$ $r_{\rm P1}$ $i_{\rm P1}$ $z_{\rm P1}$ filters, corresponding to almost 3 cycles of variation that is well fitted to a sinusoidal function. It also has an archival $V$-band light curve from CRTS \citep{Drake2009}. Even though the photometric precisions are not comparable, the CRTS light curve is consistent (in the residual sense) with the PS1 only (PV1) sinusoidal fit over $\sim$ 9 years, or $\sim$ 6 cycles. It is also a radio loud quasar (R = $\log{(f_{\rm 5 GHz}/f_{2500\AA})}$ = 2.30; \citealt{Becker2001}) from the VLA FIRST catalog (FIRST J221648.6+012427; \citealt{White1997}). Since then, we have repeated our analysis of MD09 with data Processing Version 2 (PV2) which was made available late-2014 and includes extra data from the final phase of the PS1 survey (Fig. \ref{fig:lett_cand}). We find three periodic quasar candidates that satisfy our selection criteria: a coherent period in at least three filters, an S/N for a sinusoidal fit of $>$ 3 in at least one filter, and a variation over at least 1.5 cycles. In addition, we use extended baseline data (from archival and new monitoring observations) to test the persistence of our periodic candidates over $5 - 12$ cycles. Recently, it has been pointed out by \cite{Vaughan2016} that the intrinsic red noise (increasing power at lower frequencies) characteristic of quasar variability can easily mimic periodic variability over a small number of cycles, and they emphasize the importance of demonstrating persistence of periodicity over $\gtrsim$ 5 cycles. \begin{figure}[h] \centering \epsfig{file=ps1_offsetlc_RA334to3345mb3783_pv1,width=0.41\textwidth,clip=} \epsfig{file=ps1_offsetlc_RA334to3345mb7398_pv2,width=0.41\textwidth,clip=} \caption{In L15, we analyzed the periodic quasar candidate PSO J334.2028+1.4075 based on its light curves in PV1 (upper panel), while its analysis in this paper is based on its light curves from PV2 (lower panel). We note the extra data from the last phase of PS1 MDS are included in PV2 (dashed box), while our conclusions from our new analysis on its significance as a periodic quasar candidate did not change. 4.5 $\sigma$ outliers in $g$ and $z$ filters in both versions have been clipped. The dashed lines are a sinusoid of $P=558$ days (see text for details).} \label{fig:lett_cand} \end{figure} This paper thus presents our detailed analysis with MD09 PV2 and is organized as follows. In \S\ref{sec:mds} we introduce the time domain data set used in this study: MD09 from the Pan-STARRS1 MDS. In \S\ref{sec:methods} we describe our methods of variability selection and periodicity search; we also discuss our biases in selecting variable active galactic nuclei (AGNs) in a flux-limited survey like PS1 MDS. In \S\ref{sec:ext}, we test the persistence of the candidates' periodicity with archival light curves and follow-up imaging. In \S\ref{sec:mass}, we measure the black hole mass of binary candidates and calculate their inferred binary parameters. Finally, in \S\ref{sec:conclude}, we conclude with implications for searches for periodic quasars in a large time-domain survey. Throughout this paper, we adopt cosmological parameters for a flat universe: $\Omega_{\rm m}$ = 0.3, $\Omega_{\rm \lambda}$ = 0.7, $H_{\rm 0}$ = 70 km s$^{-1}$ Mpc$^{-1}$. | \label{sec:conclude} Periodic variability in quasars on the timescales of months to years has been theoretically predicted as a signature of an SMBHB. Recent simulations show that in triaxial galaxies (e.g. \citealt{Vasiliev2015}), the ``final parsec problem'' (e.g. review by \citealt{Milosavljevic2003}) is no longer an insurmountable problem that stalls binary evolution at $a>$ 1 pc separations and that binaries can evolve into the GW-dominated regime ($a \lesssim 10^{-3}$ pc) within a few Gyrs. A systematic search for periodic quasars in a large synoptic survey thus provides a novel method to search for SMBHBs in the final phase of their evolution and can potentially yield GW sources in the nano-Hz frequency regime which is accessible to pulsar timing arrays (PTAs) including NANOGrav \citep{McLaughlin2013} and the Parkes Pulsar Timing Array \citep{Hobbs2013}. Our systematic search in the Pan-STARRS1 (PS1) MD09 field resulted in three periodic quasar candidates, from an initial sample of $\sim$ 700 color-selected quasars, that are apparently periodic over the PS1 baseline of 4 years. We further tested the persistence of their periodicity with archival light curves from SDSS Stripe 82 and followed up with imaging with the DCT. Archival \textsl{GALEX} photometry also confirms a larger amplitude of variation at shorter wavelengths, consistent with previous quasar variability studies. These extended-baseline data with photometric precision comparable to that of PS1 disfavor a simple sinusoidal model for the three candidates over an extended baseline of $\sim$ 5 -- 12 cycles. This corresponds to a detection rate of $<$ 1 out of 670 quasars ($\lesssim1.5\times10^{-3}$), which is still compatible with the theoretically predicted sub-parsec binary quasar fraction of $\lesssim$ 10$^{-3}$ out to $z = 1$ from cosmological SMBH merger simulations \citep{Volonteri2009}. The detection rate per area ($<$ 1 in 5 deg$^{2}$) is also in agreement with the theoretical prediction of 100 quasars per 1000 deg$^{2}$ of search area (or 0.5 periodic quasars in 5 deg$^{2}$) from \citet{Haiman2009} for a flux-limited survey of quasars with $m_i < 22.5$ mag. Our ongoing search over all 10 PS1 MDS fields, together with using nightly stacked images in the future which are $\sim$ 1 mag deeper, should increase our sensitivity to true SMBHBs by a factor of 100, and yield tens of promising SMBHB candidates for extended baseline monitoring and multiwavelength studies. In comparison to other SMBHB searches, we note that there are two binary candidates with double broad-line features from a sample of $\sim$ 17,500 SDSS quasars (or a detection rate of $\approx$ 10$^{-4}$) for $z < 0.7$ \citep{Boroson2009}, consistent with the predicted SMBHB rate of $\sim$10$^{-4}$ ($z < 0.7$) by \cite{Volonteri2009}. We also note that \cite{Graham2015} imply a similar detection rate to our study of 68/$\sim$75,000 $\sim$ 0.9$\times$10$^{-3}$ (for quasars $z < 1$), and \cite{Charisi2016} find a detection rate of $\approx$ 1.4$\times$10$^{-3}$ for $z<3$ (or 0.9$\times$10$^{-3}$ for the sub-sample that remained significant after their re-analysis with extended data); however, see our discussions in \S\ref{sec:persistence} and the relevant parts in \S\ref{sec:mass} on the robustness of those claimed candidates. We have demonstrated the power of an extended baseline in testing periodic quasar candidates in surveys whose temporal baselines (covering only $1.5-4$ cycles) are susceptible to false detections from red noise characteristic of normal quasar variability. Fortunately, for most of the periodic quasar candidates discovered in recent optical time domain surveys, continued monitoring over the next few years can robustly test the persistence of the periodicity over a necessary number of cycles ($> 5$) to filter out false alarms, and verify strong SMBHB candidates for direct detection in GWs by PTAs. | 16 | 9 | 1609.09503 |
1609 | 1609.09818_arXiv.txt | We present a new version of the MURaM radiative MHD code that allows for simulations spanning from the upper convection zone into the solar corona. We implemented the relevant coronal physics in terms of optically thin radiative loss, field aligned heat conduction and an equilibrium ionization equation of state. We artificially limit the coronal Alfv{\'e}n and heat conduction speeds to computationally manageable values using an approximation to semi-relativistic MHD with an artificially reduced speed of light (Boris correction). We present example solutions ranging from quiet to active Sun in order to verify the validity of our approach. We quantify the role of numerical diffusivity for the effective coronal heating. We find that the (numerical) magnetic Prandtl number determines the ratio of resistive to viscous heating and that owing to the very large magnetic Prandtl number of the solar corona, heating is expected to happen predominantly through viscous dissipation. We find that reasonable solutions can be obtained with values of the reduced speed of light just marginally larger than the maximum sound speed. Overall this leads to a fully explicit code that can compute the time evolution of the solar corona in response to photospheric driving using numerical time steps not much smaller than $0.1$ seconds. Numerical simulations of the coronal response to flux emergence covering a time span of a few days are well within reach using this approach. | Comprehensive numerical simulations of the solar corona have been developed by a few teams in the past decade \citep{Gudiksen:Nordlund:2002,Gudiksen:Nordlund:2005a, Gudiksen:Nordlund:2005b,Abbett:2007,Gudiksen:etal:2011,Bingert:Peter:2011,Bingert:Peter:2013,Bourdin:etal:2013:obs_driven,Chen:etal:2014}. These simulations are often referred to as ``realistic'' in the sense that they include the relevant macro-physics in terms of MHD, field aligned heat conduction, optically thin radiative loss, and a solar mixture equation of state including partial ionization effects. The heating of the corona is implicitly handled through MHD (upward directed Poynting flux above photosphere) in combination with either explicit or implicit (i.e. numerical) magnetic and viscous diffusivities, which can be considered as a numerical representation of the heating due to braiding of field lines as first suggested by \citet{Parker:1972:topological_heating,Parker:1983:sheets}, see also the review by \citet{Klimchuk:2006:review} for a general overview. It was demonstrated by \citet{Gudiksen:Nordlund:2002,Gudiksen:Nordlund:2005b,Gudiksen:Nordlund:2005a} that this process is sufficient to maintain the corona at MK temperatures. While the treatment of energy dissipation is not realistic (i.e. these models do not include the correct micro-physics) it has been found by \citet{Peter:etal:2004,Peter:etal:2006} that these models do compare remarkably well with observations in a statistical sense. A later analysis by \citet{Bingert:Peter:2013} showed that the statistical properties of the energy deposition agree with the predictions of the nanoflare model \citep{Parker:1988:nanoflare}. While the models of \citet{Bingert:Peter:2011,Bingert:Peter:2013,Bourdin:etal:2013:obs_driven,Chen:etal:2014} do not include the upper convection zone and are driven by a boundary condition either taken from observations or other numerical simulations, the model of \citet{Gudiksen:etal:2011} does self-consistently treat the coupling from the upper convection zone into the solar corona, including also realistic physics for the photosphere (3D radiative transfer) and chromosphere including non-local thermal equilibrium (NLTE) physics. The coupling from the upper convection zone into the solar corona was also addressed by \citet{Abbett:2007}, although their model did not use radiative transfer in the photosphere and relied in addition on empirical heating terms since the braiding of magnetic field lines by photospheric motions turned out to be insufficient at least in the quiet Sun setup they considered. The use of empirical heating descriptions is more common in codes that aim at modeling the larger scale corona \citep[see, e.g.][]{Mok:etal:2005,Mok:etal:2008,vdHolst:etal:2014:AWSoM}, where the resolution is in general insufficient to capture the processes involved in coronal heating directly. In this paper we present a new version of the MURaM radiative MHD code \citep{Voegler:etal:2005,Rempel:2014:SSD}, in which we implement a treatment of the corona along the lines of the above mentioned ``realistic'' simulations. Unlike \citet{Gudiksen:etal:2011} we do not implement at this time a realistic treatment of the chromosphere, i.e. our ``chromosphere'' is treated assuming local thermal equilibrium LTE. Our main emphasis is on introducing efficient ways to deal with the two most stringent numerical time step constraints that are encountered in direct simulations of the solar corona: high Alfv{\'e}n velocities, which might even exceed the speed of light in the classical approximation, and severe time step constraints from field aligned heat conduction. In this paper we explore the potential of the so called Boris correction \citep{Boris:1970:BC} to deal with the high Alfv{\'e}n velocities (semi-relativistic MHD with an artificially reduced speed of light), and apply a conceptually similar approach also to heat conduction. The paper is organized as follows. In Section \ref{sec:num} we summarize the approximations used and present the full set of equations solved. In Section \ref{sec:results} we test the validity of the approximations for simulation of the solar corona by considering four different setups (quiet Sun, open flux, coronal arcade and active region). Section \ref{sec:heating} presents an analysis of the time scales that govern coronal energy transport and release. We briefly discuss numerical efficiency in Section \ref{sec:efficiency} and present our conclusion in Section \ref{sec:concl}. | \label{sec:concl} We presented a new version of MURaM radiative MHD code that has been extended into the solar corona. For the corona we implemented optically thin radiative loss and field aligned heat conduction, while we treat the ``chromosphere'' at this point in LTE. Similar to other so called ``realistic'' MHD simulations of the solar corona \citep[e.g.,][]{Gudiksen:etal:2011,Bingert:Peter:2011,Bingert:Peter:2013,Chen:etal:2014} we do not use any parameterizations of coronal heating. Our setup includes the upper convection zone and photosphere where magneto-convection leads to the generation of a Poynting flux, which self-consistently heats the upper layers of the simulation domain through a combination of (numerical) resistive and viscous energy dissipation. Our implementation uses a fully explicit treatment and circumvents stringent time-step arising from the coronal Alfv{\'e}n velocity and heat conduction through the use of the ``Boris-correction'' \citep{Boris:1970:BC} and a hyperbolic treatment of heat conduction that imposes a maximum characteristic speed for conductive heat transport. Both approaches are inspired by semi-relativistic MHD as they are based on equations with a well defined maximum propagation speed ``speed of light'', which is artificially reduced to a lower values for computational efficiency. We applied our code to four different coronal settings: quiet Sun, open flux, coronal arcade and active region and explored the dependence of the solutions on details of our adopted numerical diffusivity (process that heats the corona in our simulation) and the chosen value for the peak Alfv{\'e}n as well as heat conduction speed ``reduced speed of light''. We tested the latter using the active region setup in which the horizontally averaged Alfv{\'e}n velocity reaches values of $10,000$ km~s$^{-1}$ and the peak value exceeds $100,000$ km~s$^{-1}$. We explored values of $c$ from $400$ km~s$^{-1}$ to about $1600$ km~s$^{-1}$ and found no significant influence on the resulting heating and mean temperature of the corona. It is remarkable that even values of $c$ similar to the speed of sound lead to reasonable results when using the Boris correction as long as the flow velocity is artificially limited to assure the validity of the semi-relativistic approach. From our experiments we concluded that a setting of $c=\mbox{max}(C_S,\,3\,v)$ is a good compromise between computational speed and sufficiently accurate treatment of MHD. This translates to values of $c$ in the $400-800$ km~s$^{-1}$ for the setups considered here. We found that our treatment of hyperbolic heat conduction was sufficiently accurate in all these cases except for features on the scale of the grid. Formally the solution of the hyperbolic heat conduction equation corresponds to a solution computed with a time averaged heat flux. Throughout most of the corona the associated averaging time scale is on the order of seconds, i.e. short compared to typical time scales of interest. With respect to the treatment of numerical diffusivities we compared three setups differing in their effective numerical magnetic Prandtl number. While our reference setup has a high $P_{\rm m}$ (due to low magnetic diffusivity), we considered also a $P_{\rm m}\sim1$ case with comparable diffusivities for velocity and magnetic field as well as a low $P_{\rm m}$ setup with significantly reduced viscosity. The most striking difference is found in the ratio of resistive to viscous heating, which is strongly reduced in a high magnetic Prandtl number setting. This is very similar to the behavior that was found by \citet{Brandenburg:2011:SSD_low_Pm,Brandenburg:2014:Pm} for (high $\beta$) turbulence. While the ratio of resistive and viscous heating depends on numerical details, the sum of both terms (total coronal heating) is found to be very robust. The dependence on $P_{\rm m}$ ultimately illustrates that the microphysics are important if the goal is to determine how dissipation occurs. If the goal is just to quantify the total amount of energy dissipation, details of the dissipation process are less important. Furthermore, the fact that the Corona is essentially a very high $P_{\rm m}$ regime strongly suggests that resistive heating is negligible and dissipation happens through viscosity on scales that are large enough to be captured in current numerical simulations. We investigated four different magnetic field configurations that are representative of quiet Sun, open flux, coronal arcade and active regions. In all four setups we have in the convection zone part of the domain a small-scale dynamo operating that maintains a small-scale mixed polarity field. In the case of the quiet Sun setup the Poynting flux resulting from the small scale dynamo alone is sufficient to maintain an about 1 MK hot corona. Adding a small $3$~G vertical mean field (open flux region) does not change the corona temperature significantly in spite of enhanced heating in the upper half of the simulation domain. The primary reason for that is more efficient cooling through heat conduction, which operates mostly in the vertical direction in this case. It is likely that the temperature would be lower if we would use more appropriate top boundary conditions (fully wave transmitting) in this case, but we did not investigate that further. We find with about $2$~MK a significantly hotter corona in the coronal arcade setup (a $\pm 50$ G vertical mean field in the left and right half of the domain). The active region setup is even hotter in the low corona (5-10 Mm above the photosphere) where we find hot loops with temperatures of up to $5$~MK, but the mean temperature drops significantly towards the top boundary, where it becomes comparable to the quiet Sun setup. The magnetic field reaching into the upper parts of the simulation domain connects mostly to the umbrae of the spot pair in the photosphere where the Poynting flux is strongly suppressed. \citet{Chen:etal:2015NatPh,Cheung:etal:2015:DEM} found the strongest Poynting flux at the outer boundary of the umbra, which explains the hotter loops in the low corona for this setup. We analyzed the time scales of motions that contribute to the coronal energy transport and found that for the closed field QS and CA setups the corresponding time scales are in the 20 to 50 minute range in a height of $8.9$~Mm above the photosphere. At the same time the viscous and resistive heating is highly intermittent with energy releases that are comparable to those expected in the nanoflare picture. Overall this supports the picture of ``DC'' heating by braiding of field lines on long time scales \citep{Parker:1972:topological_heating,Parker:1983:sheets} and intermittent energy release in form of nanoflares \citep{Parker:1988:nanoflare}. Similar to \citet{Ballegooijen:2014:AC-DC} we find that braiding in the photosphere does excite waves with substantial amplitudes of about $30$~km~s$^{-1}$, however, filtering that component out does only moderately impact the Poynting flux. Overall we conclude that numerical ``tricks'' based on semi-relativistic MHD with an artificially reduced speed of light (Boris correction) enable a rather inexpensive modeling of the solar corona by effectively limiting both Alfv{\'e}n and heat conduction speed. We did not find significant drawbacks from this approach in the setups considered here. The effects we found in the mean temperature from both artificial limitation of the Alfv{\'e}n velocity and treatment of numerical diffusivity are comparable to those expected from the intrinsic uncertainty of the input physics (e.g. the assumed coronal element abundance affecting the radiative loss and the numerical value of the heat conduction coefficient in the Spitzer formulation) as well as the intrinsic variability found in these simulations. The computational expense for our active region setup is about $5,000$ core hours for 1 hour of solar time. With the computing resources available today 3D realistic simulation of the solar corona covering the full time span of active region formation to decay are well within reach. \appendix We present here a derivation of the semi-relativistic momentum equation following \citet{Gombosi:etal:2002:SR}. We start from the the MHD momentum equation of the form \begin{equation} \varrho\frac{\partial \vec{v}}{\partial t}=-\varrho(\vec{v}\cdot\nabla)\vec{v}-\nabla p+\varrho\vec{g}+\frac{1}{c}\vec{j}\times\vec{B}\label{eq:app1} \end{equation} With the Maxwell equation \begin{equation} \nabla\times\vec{B}=\frac{4\pi}{c}\vec{j}+\frac{1}{c}\frac{\partial \vec{E}}{\partial t} \end{equation} we can rewrite the Lorentz force as \begin{equation} \frac{1}{c}\vec{j}\times\vec{B}=\frac{1}{4\pi}(\nabla\times\vec{B})\times\vec{B}+\frac{1}{4\pi c}\vec{B}\times\frac{\partial \vec{E}}{\partial t} \end{equation} With the relation $\vec{E}=-{1\over c} \vec{v}\times\vec{B}$ we can relate the time derivative of $\vec{E}$ to already known quantities: \begin{equation} \frac{\partial \vec{E}}{\partial t}=-\frac{1}{c}\left(\frac{\partial \vec{v}}{\partial t}\times\vec{B} +\vec{v}\times\frac{\partial \vec{B}}{\partial t}\right) \end{equation} It is primarily the contribution from the first term that is responsible for limiting the Alfv{\'e}n velocity to values less than $c$ and we keep only that term in the following derivation. The second term leads to additional forces perpendicular to the magnetic field that are important for an exact treatment of semi-relativistic MHD, but not required if we focus only on the reduction of Alfv{\'e}n velocity. In addition it also follows that the second term is of order $v^2\over v_A^2$, i.e. small in the regime where semi-relativistic MHD is valid $v\ll c < v_A$: \begin{equation} \frac{\vert\vec{v}\times\frac{\partial \vec{B}}{\partial t}\vert}{\vert\frac{\partial \vec{v}}{\partial t}\times\vec{B}\vert}\sim\frac{\vert\vec{v}\times(\nabla\times(\vec{v}\times\vec{B}))\vert}{\vert{1\over 4\pi\varrho}((\nabla\times\vec{B})\times\vec{B})\times\vec{B}\vert}\sim\frac{v^2}{v_A^2}<\frac{v^2}{c^2} \end{equation} Here $v_A=|\vec{B}|/\sqrt{4\pi\varrho}$ denotes the (classic) Alfv{\'e}n velocity. With this term the expression for the Lorentz force is given by: \begin{eqnarray} \frac{1}{c}\vec{j}\times\vec{B}&=&\frac{1}{4\pi}(\nabla\times\vec{B})\times\vec{B}+\frac{1}{4\pi c^2}\vec{B}\times(\vec{B}\times\frac{\partial \vec{v}}{\partial t})\\ &=&\frac{1}{4\pi}(\nabla\times\vec{B})\times\vec{B}-\frac{v_A^2}{c^2}[\mathcal{I}-\hat{\vec{b}}\hat{\vec{b}}]\varrho\frac{\partial \vec{v}}{\partial t} \end{eqnarray} where $\hat{\vec{b}}=\vec{B}/|\vec{B}|$ denotes the unit vector in the direction of $\vec{B}$. Substituting this expression back into Eq. (\ref{eq:app1}) yields (using $x_A=v_A/c$): \begin{equation} \left[\mathcal{I}+x_A^2(\mathcal{I}-\hat{\vec{b}}\hat{\vec{b}})\right]\varrho\frac{\partial \vec{v}}{\partial t}=-\varrho(\vec{v}\cdot\nabla)\vec{v}-\nabla p+\varrho\vec{g}+\frac{1}{4\pi}(\nabla\times\vec{B})\times\vec{B} \end{equation} The inverse of the ``enhanced inertia'' matrix on the left hand side is given by \citep{Gombosi:etal:2002:SR}: \begin{equation} \left[\mathcal{I}+x_A^2(\mathcal{I}-\hat{\vec{b}}\hat{\vec{b}})\right]^{-1}=\frac{1}{1+x_A^2}\left[\mathcal{I}+x_A^2\hat{\vec{b}}\hat{\vec{b}}\right]= \mathcal{I}-\frac{x_A^2}{1+x_A^2}\left[\mathcal{I}-\hat{\vec{b}}\hat{\vec{b}}\right] \end{equation} This leads to a momentum equation of the form: \begin{equation} \frac{\partial \varrho\vec{v}}{\partial t}+\nabla\cdot(\varrho\vec{v}\vec{v}+\mathcal{I}p)=\varrho\vec{g}+\frac{1}{4\pi}(\nabla\times\vec{B})\times\vec{B}+\vec{F}_{SR} \end{equation} where the ``semi-relativistic'' correction term is given by \begin{equation} \vec{F}_{SR}=-\frac{x_A^2}{1+x_A^2}\left[\mathcal{I}-\hat{\vec{b}}\hat{\vec{b}}\right]\left(-\varrho(\vec{v}\cdot\nabla)\vec{v}-\nabla p+\varrho\vec{g}+\frac{1}{4\pi}(\nabla\times\vec{B})\times\vec{B}\right)\label{eq:FSR} \end{equation} Since our aim is not to compute an exact solution of semi-relativistic MHD, but rather to use the minimal amount of correction terms needed to limit the Alfv{\'e}n velocity, we can use some freedom in determining the quantity $x_A^2/(1+x_A^2)$ in front of the projection operator. We generalize this expression as $1-f_A$ and use \begin{equation} f_A=\frac{1}{\sqrt{1+({v_A\over c})^4}} \end{equation} which leads to a limitation of the Alfv{\'e}n velocity in the following form: \begin{equation} v_A^2 \longrightarrow \frac{v_A^2}{\sqrt{1+({v_A\over c})^4}}\;.\label{eq:va_lim} \end{equation} While we did not find that the detailed functional form of $f_A$ matters as long as the limited Alfv{\'e}n velocity Eq. (\ref{eq:va_lim}) remains a monotonic function of $v_A$ and $f_A$ asymptotes as $ x_A^{-2}$ for large values of $x_A$, we prefer to use expressions for $f_A$ that have a sharper transition than the $(1+x_A^2)^{-1}$ that follows from semi-relativistic MHD in an attempt to minimize the volume of the simulation domain where the correction term Eq. (\ref{eq:FSR}) contributes. For solving the MHD equations we need to determine a maximum characteristic velocity that will be used for determining the time step as well as for computing numerical diffusivities. As shown by \citet{Gombosi:etal:2002:SR} the wave speeds in semi-relativistic MHD can be quite complicated. We use the following approximate expression: \begin{equation} C_{\rm max}={\rm max}\left(C_S,\sqrt{f_A(C_S^2+v_A^2)}\right)+\vert\vec{v}\vert\label{Eq:Cmax} \end{equation} | 16 | 9 | 1609.09818 |
1609 | 1609.03421_arXiv.txt | { Using CO (4-3) and (2-1) Atacama Large Millimeter Array (ALMA) data, we prove that the molecular gas in the jet-driven winds of the galaxy IC5063 is more highly excited than the rest of the molecular gas in the disk of the same galaxy. On average, the \coft /\coto\ flux ratio is 1 for the disk and 5 for the jet accelerated or impacted gas. Spatially-resolved maps reveal that in regions associated with winds, the \coft /\coto\ flux ratio significantly exceeds the upper limit of 4 for optically thick gas. It frequently takes values between 5 and 11, and it occasionally further approaches the upper limit of 16 for optically thin gas. Excitation temperatures of 30-100\,K are common for the molecules in these regions. If all of the outflowing molecular gas is optically thin, at 30-50\,K, then its mass is 2$\times$10$^6$\msun . This lower mass limit is an order of magnitude below the mass derived from the \coto\ flux in the case of optically thick gas. Molecular winds can thus be less massive, but more easily detectable at high z than they were previously thought to be. } | \label{sec:intro} Jets launched by supermassive black holes are presently being evaluated as important mechanisms for galaxy evolution, affecting multiple phases of the interstellar medium (ISM). By depositing energy in the hot and tenuous gas, jets are proposed to delay the inflow of intergalactic gas into ellipticals and moderate their growth \citep{bower06, croton06}. By depositing energy in the cold and dense gas, jets are proposed to influence the star-formation properties of disks \citep{wagner11, gaibler12}. Observational evidence that the impact of a jet upon dense clouds initiates molecular gas winds has been successfully gathered \citep[e.g.,][]{morganti05, rupke11, sakamoto14, garcia-burillo14, dasyra15, aalto16}. In this letter, we provide observational evidence that the molecular gas accelerated or impacted by a jet is more highly excited and less optically thick than the rest of the ISM. We prove this for the galaxy IC5063, an elliptical with a molecular disk in its center. The jet launched from the black hole of IC5063 propagates through the disk, nearly parallel to its plane. Once the jet encounters clouds, it scatters gas, forming atomic and molecular winds \citep{oosterloo00,morganti07,morganti13,tadhunter14}. Using ultra-deep ESO VLT near infrared (NIR) data, we identified at least four discrete starting points of ionized and warm molecular gas winds near the jet trail \citep{dasyra15}. This constitutes the highest number of known jet-driven winds in a galaxy. The winds sweep the galaxy's inner kpc$^2$, allowing for a spatially-resolved study of the accelerated gas properties. Atacama Large Millimeter Array (ALMA) data have revealed the presence of cold molecular gas in most of these winds \citep{morganti15}. | \label{sec:discussion} Our analysis indicates the presence of at least two distinct molecular gas components: a low-excitation component associated with clouds in the disk, and a high-excitation component associated with gas in the jet-driven winds or near jet-cloud impact points. To be bright in \coft, which has a critical density of 1.7$\times$10$^5$ cm$^{-3}$, both components must contain clouds of high volume density. However, while the low excitation component is optically thick, the high excitation component is (partly) optically thin: \fft\ values $>$4 can only be consistent with emission from optically thin gas at high temperatures. From the antenna temperature definition, assuming that the molecular gas is in local thermodynamic equilibrium (LTE) and that its excitation temperature \tex\ greatly exceeds that of the cosmic microwave background (see below), we obtain analytic solutions for \fft\ (with fluxes in units of Jy\kms ). For optically thick gas, \fft\ = ${ \rm (\nu_{43}/\nu_{21})^3 [e^{(h \nu_{21} /kT_{ex})}-1]/[e^{(h \nu_{43} /kT_{ex})}-1]}$. An upper limit of 4 is reached at the Rayleigh-Jeans limit (as expected from the same line brightness temperature). For optically thin gas, \fft\ = ${\rm (\nu_{43}/\nu_{21})^4 e^{-7h \nu_{10} /kT_{ex}}}$. An asymptotic value of 16 is obtained for \tex $>$$>$7h$\nu_{10}$/k $\sim$ 39 K. From this relation, we find that \tex =32K, if all of the gas kinematically or excitationally affected by the jet (with \fft =5) is optically thin. For the fastest gas in the wind (with \fft =8), \tex\ is 56K. Near the biconical outflow base, where \fft\ is 11, \tex\ is 100K. The temperatures of the gas in the wind can thus be quite high for LTE, and they can be even higher for non-LTE. It is unlikely that this result reflects infrared (IR) pumping of the populations of quantum states with high rotational number J. The molecules should see an IR source of $>$160K with a large filling factor \citep{carroll81} in all pertinent regions of Figs.~\ref{fig:residual_maps} and \ref{fig:ratio_maps}. Infrared pumping has been deemed unimportant for HCN in the wind of Mrk231 \citep{aalto12}. The high \tex\ is instead related to the creation of dense fragments once the jet impinges on molecular clouds. The contribution of the jet to the dissipation and heating of the clouds could thus have a negative impact on the overall star formation of IC5063. Still, it doesn't rule out a local enhancement of star formation in compressed clouds. Given the above bimodality, we determined the gas mass in the wind and in the disk using the appropriate CO-luminosity-to-\htwo -mass conversion factor, $\alpha_{CO}$, for optically thin and thick emission, respectively. For the optically thick gas, we use Eq. 3 of \citet{solomon97} with $\alpha_{CO}$=4.6\msun /(K\kms\,pc$^2$), as the far infrared luminosity of IC5063 is 2$\times$10$^{10}$\lsun . For \hbox{$\rm F_{21}/F_{10}$}=4, a mass of 4.0($\pm$0.5)$\times$10$^8$\msun\ is obtained for the disk. For the optically thin gas, \hbox{$\rm F_{21}/F_{10}$} = $ {\rm (\nu_{21}/\nu_{10})^4 e^{-2h \nu_{10} /kT_{ex}}}$ and $\alpha_{CO}=0.26\,(\tex /30)\,e^{5.53/\tex-0.184}$ \msun / (K\kms\,pc$^2$), for a Galactic CO/\htwo\ abundance \citep{bolatto13}. \hbox{For $\rm F_{21}$=12.3($\pm$2.5) Jy\kms\ and \fft=5, then $\alpha_{CO}$} \hbox{is 0.27\msun/\,(K\kms\,pc$^2$) and the wind mass is 1.7($\pm$0.3)} \hbox{$\times$10$^6$\msun . For \fft=8, $\alpha_{CO}$ is 0.44\msun /\,(K\kms\,pc$^2$)} and the wind mass is 2.4($\pm$0.5)$\times$10$^6$\msun . An order-of-magnitude difference in $\alpha_{CO}$ between the disk and the wind was already assumed by \citet{morganti15}, but the use of an optically-thick \hbox{$\rm F_{21}/F_{10}$} value led to a wind mass of 1.9$\times$10$^7$\msun . Our current knowledge that the wind is partly optically thin leads us to revise its molecular content down: 2$\times$10$^6$\msun\ is its lower mass limit, if all of its gas is optically thin. This is comparable to the atomic wind mass \citep[3.6$\times$10$^6$\msun ; ][]{morganti07}. The actual molecular wind mass will be derived when all data, including new CO (1-0) and (3-2) observations, are analyzed assuming non-LTE conditions: when both high-temperature gas and low-temperature gas (contributing to low J lines) are used to fit the CO spectral line energy distribution (SLED). Still, our simplified LTE analysis proves that molecular wind masses can be overestimated \citep[contrary to what is known from Mrk231;][]{cicone12}. | 16 | 9 | 1609.03421 |
1609 | 1609.03567_arXiv.txt | We present the first results of the High cadence Transient Survey (HiTS), a survey whose objective is to detect and follow up optical transients with characteristic timescales from hours to days, especially the earliest hours of supernova (SN) explosions. HiTS uses the Dark Energy Camera (DECam) and a custom made pipeline for image subtraction, candidate filtering and candidate visualization, which runs in real--time to be able to react rapidly to the new transients. We discuss the survey design, the technical challenges associated with the real--time analysis of these large volumes of data and our first results. In our 2013, 2014 and 2015 campaigns we have detected more than 120 young SN candidates, but we did not find a clear signature from the short--lived SN shock breakouts (SBOs) originating after the core collapse of red supergiant stars, which was the initial science aim of this survey. Using the empirical distribution of limiting--magnitudes from our observational campaigns we measured the expected recovery fraction of randomly injected SN light curves which included SBO optical peaks produced with models from \citet{2011ApJS..193...20T} and \citet{2010ApJ...725..904N}. From this analysis we cannot rule out the models from \citet{2011ApJS..193...20T} under any reasonable distributions of progenitor masses, but we can marginally rule out the brighter and longer--lived SBO models from \citet{2010ApJ...725..904N} under our best--guess distribution of progenitor masses. Finally, we highlight the implications of this work for future massive datasets produced by astronomical observatories such as LSST. | The advent of a new generation of large field--of--view astronomical optical CCD cameras already operational (e.g. iPTF, \citealt{2009PASP..121.1395L}; SkyMapper, \citealt{2007ASPC..364..177K}; Pan--STARRS, \citealt{2004AN....325..636H}; OmegaCam, \citealt{2002Msngr.110...15K}; DECam, \citealt{2015AJ....150..150F}; Hyper Suprime--Cam, \citealt{2010SPIE.7740E..2IF}; KMTNET, \citealt{2011SPIE.8151E..1BK}) and planned (e.g. ZTF, \verb+http://www.ptf.caltech.edu/ztf+; and LSST, \citealt{2009arXiv0912.0201L}) is revolutionizing our understanding of the Universe because of their surveying capabilities. Thanks to these instruments, large regions in the sky are being mapped up to very large distances, and also rare, short--lived optical transients are being found as these large regions of the Universe are observed with a high cadence. The latter presents not only new opportunities for the study of astrophysical phenomena, but also new challenges from the point of view of the data analysis. Large volumes of data will have to be processed in real--time in order to trigger follow up observations that would help disentangle the physical nature of the transients detected \citep[see e.g.][]{2014Natur.509..471G}. The High cadence Transient Survey (HiTS) is a discovery survey that takes advantage of the large \emph{etendue}, the product of collecting area and field--of--view, of the Dark Energy Camera (DECam) mounted on the 4 m Blanco telescope at the Cerro Tololo Interamerican Observatory (CTIO), the fast connectivity available between CTIO and the Center for Mathematical Modelling (CMM~@~U.~Chile), and the computing capabilities of the Chilean National Laboratory for High Performance Computing (NLHPC) that allows us to observe and analyze high cadence DECam data in real--time. Because DECam is the largest \emph{etendue} project in the southern hemisphere until the arrival of the full LSST project, HiTS can be considered a precursor project for some of the challenges regarding the fast analysis of large data volumes, the high cadence observations of deep drilling fields and, depending on the cadence, the exploration of the hour--timescale transient population. HiTS aims to explore the population of transient or periodic objects with characteristic timescales between a few hours and days \citep[c.f.][]{2010ApJ...723L..98K} and apparent magnitudes down to about 24.5 mag. Its main science driver was to discover the elusive shock breakout (SBO) phase of extended red supergiant star (RSG) progenitors undergoing core collapse \citep{1978ApJ...225L.133F, 2008Sci...321..223S, 2008ApJ...683L.131G}, but it also focuses on the study of young SN explosions in general. The latter includes shock--hit companion stars in multiple progenitor systems of Type Ia SNe explosions \citep[see e.g.][]{2000ApJS..128..615M, 2010ApJ...708.1025K, 2011ApJ...741...20B, 2012ApJ...744L..17B, 2014ApJ...784L..12G, 2016ApJ...820...92M, 2015Natur.521..328C}; the shock cooling phase of core collapse supernovae (SNe), which provide constraints on their progenitors size and circumstellar environments \citep{2011MNRAS.415..199M, 2015MNRAS.451.2212G, 2016ApJ...819...35A}; and the early light curves of Type Ia SNe, whose diversity could be driven by different radioactive profiles in their outermost layers \citep{2014ApJ...784...85P}. The structure of this manuscript is the following: in Section~\ref{sec:sne} we will describe some of the relevant physics of red supergiant SN explosions during their earliest observable phases; in Section~\ref{sec:design} we will discuss how the survey was designed and the details of our observation strategy; in Section~\ref{sec:data} we will show how the real--time analysis of the data was performed, including a brief description of a newly developed image subtraction, candidate classification and visualization pipeline; in Section~\ref{sec:results} we will discuss some of the first results, including a detailed discussion on the limiting--magnitude of the survey and its implications for SBO model constraints; in Section~\ref{sec:summary} we summarize the main results from this paper and in Section~\ref{sec:discussion} we discuss the implications from this work. We note that in this manuscript we only present our conclusions about the presence or absence of RSG SBOs, leaving the discussion on other classes of transient events for subsequent publications. | \label{sec:summary} The first results of the High cadence Transient Survey (HiTS) search for supernova (SN) shock breakouts (SBOs) were presented. With the current calibration scheme and data analysis pipeline we see no clear evidence for red supergiant SBO optical early--time optical peaks in light curves resembling SNe II (with a SNR $> 5$). Based on a joint analysis of the three observational campaigns with our empirically derived limiting--magnitudes we conclude that ensembles of explosion models assuming a Salpeter--like initial mass function with an upper mass limit of either 16.5 or 30 $M_\odot$ (M16.5 or M30 distributions, respectively) cannot be excluded for the \citet{2011ApJS..193...20T} models, but we can be marginally excluded for the \citet{2010ApJ...725..904N} models (see Table~\ref{tab:probs}). This result should be taken with caution given all the uncertainties associated to the distribution of SN II progenitor properties. HiTS run in the optical using the Dark Energy Camera (DECam) during the 13A, 14A and 15A survey campaigns. The survey strategy could be described as several contiguous nights (4 nights in 13A, 5 nights in 14A and 6 nights in 15A) of high cadence (2 hr in 13A and 14A and 1.6 hr in 15A), monochromatic ($u$ band in 13A and $g$ band in 14A and 15A), untargeted, varying airmass observations towards a large area of the sky (120 deg$^2$ in 13A and 14A and 150 deg$^2$ in 15A) with single epoch depths between 23 and 24.5 magnitudes. In March of 2013 a pilot phase of the survey was performed, with data being analyzed after the observation run had finished. In March of 2014 we performed the data processing and candidate filtering in real--time, the first real--time analysis of DECam data to our knowledge, although with very simple visualization tools which delayed our reaction capability by a few hours. In February 2015 we achieved the full data analysis, candidate filtering and visualization process in real--time thanks to significant improvements in our visualization tools, which highlights the importance of fast visualization for real--time surveys. We processed more than $10^{12}$ pixels in a stream of 40 Mbps, which after processing resulted in a stream of about 3 new candidates per minute, 5 to 6 minutes after every exposure. As a result, more than 120 SN candidates were detected in total in real--time. We computed empirical 50\% completeness magnitudes analyzing deep stacked DECam images in relation to the individual epochs of the survey. The depth of the survey varied within each night typically by more than one magnitude as the observations had to be performed at varying airmasses in order to achieve the required cadence during the full night. We compared these values with the predictions of the public exposure time calculators (ETCs) versions 5 and 6 (v5 and v6, respectively), which we modified to include the effect of airmass. We validated our modified ETCs studying the relation between these limiting--magnitudes with the observed FWHM and airmasses. The elusive SBOs may have been detected if the actual survey depth had matched the initial limiting--magnitude estimations. However, the survey depth was overestimated due to several factors: an overly optimistic ETC available at the time (ETC v5), a stronger than expected airmass effect, worse than planned observing conditions and errors introduced by the image difference process. During survey design we used two figures of merit to determine the quality of an observational strategy, both defined in Section~\ref{sec:simulations}: 1) the number of SBO optical peak detections and 2) the number of SNe detected at least twice during the first rest--frame day after shock emergence. We have shown that a SBO optical peak detection is much harder to obtain than a double detection of the SN within the first rest--frame day. An obvious consequence is that an example of 2) does not mean that the SBO optical peak should have been seen in the data, which may explain why we have not seen optical SBO peaks in some datasets \citep[see e.g.][]{2016ApJ...818....3K}. Using the empirical limiting--magnitudes with the models described in this analysis \citep[from][]{2011ApJS..193...20T, 2010ApJ...725..904N} we evaluated these two figures of merit and found that the number of predicted events was most sensitive to the upper mass limit of the SN II initial mass function (IMF) distribution. An upper mass limit of 30 $M_\odot$ would lead to about four times more SBO detections than an upper mass limit of 16.5 $M_\odot$, and about twice the number of SNe younger than one day after shock emergence with at least two detections. The number of SNe detected twice during the first day after shock emergence appears to be even more sensitive to the explosion energy, varying by as much as the energy variation factor in the models tested. An important consequence of marginally favouring the relatively dimmer and shorter--lived SBO optical peaks from the \citet{2011ApJS..193...20T} models is that with our typical cadences a real--time detection of a SBO will be unlikely to happen fast enough to react and observe it with other instruments. We expect more than 82\% of our SBO detections to have only one detection before the end of the SBO optical peak with these models, thus we rely on the subsequent early rising SN light for their online identification. In fact, for the shock cooling SN light curve to rise to at least half a magnitude below the SBO optical peak maximum it takes typically more than half a day according to the more realistic \citet{2011ApJS..193...20T} models (see Figure~\ref{fig:day1_LCs}), making the identification of SNe within a few hours of shock emergence incredibly challenging. This highlights the importance of having continuous high cadence observations during several nights followed by low cadence follow up observations in order to aid with the candidate selection and SBO identification via post--processing of the high cadence phase data. Apart from the previously described SBO model constraints, an important contribution from this survey will be the detection of SN candidates younger than one day after shock emergence. A preliminary analysis show that they have a very fast initial rise in their light curve inconsistent with the model light curves used in this analysis. They seem more consistent with shocks breaking into high density circumstellar material (CSM) in RSG stars \citep{2015ApJ...804...28G, 2015MNRAS.451.2212G, 2016ApJ...818....3K, 2016arXiv160103261T, 2016ApJ...820...23G}. In order to compare their light curves to existing models we are compiling host galaxy redshifts and post--SN explosion images for more precise and accurate absolute calibrations. If the shock--CSM interaction in normal RSG stars significantly affects the SBO properties, which could be the case for the high density CSM shock breakouts, the evolution of SN II light curves during the first rest--frame day after shock emergence could be very different to that suggested by the models considered in this work. Thus, the detection of SNe during the first days after shock emergence could be a tool to constrain the properties of RSGs and their CSM. Given the discrepancy in derived mass loss rates between early and late times implied by these recent works, this could be a clue about the wind structure in RSGs \citep{2014Natur.512..282M} or the latest stages of nuclear burning before explosion in these stars \citep{2014MNRAS.445.2492M}. These factors and our non--detection of red supergiant SBOs suggest that HiTS should switch to a lower cadence, multiwavelength survey mode. If SBOs can be detected in a systematic fashion, they could provide an alternative probe for the upper mass limit of the SN II IMF distribution than pre--explosion progenitor detections \citep{2009MNRAS.395.1409S}. This is because the UV radiation during the SBO phase is expected to sublimate most of the CSM dust which could be obscuring the SN II progenitors in pre--explosion images. Some evidence for dust column density changes after SBO exist for SN 2012aw \citep{2012ApJ...759L..13F, 2012ApJ...756..131V}, which appears to be on the high end of the progenitor mass distribution (although see \citealt{2012ApJ...759...20K}). The high cadence strategy will be very important for future SBO surveys, since it is difficult to confirm a SBO detection without a previous non--detection and a subsequent drop in the light curve with a timescale comparable to the SBO optical peak. In a multichromatic survey it may be possible to differentiate SBOs from early SN light curves based on their colors, but this requires either simultaneous multiwavelength observations or high cadence filter changes, which may be expensive for large etendue telescopes. In the case of LSST, intranight multicolor observations will be limited because of the limited number of allowed filter changes during its lifetime, so most SBO detections will most likely rely on high cadence monochromatic observations in the deep--drilling fields to be defined. We have scaled our simulations to LSST's larger FoV, larger mirror area, smaller pixel size and shorter overhead times and found that assuming 30 sec exposure times we could observe 170 LSST fields per night with a cadence of 1.6 hours and would be able to find 13 times more SBO and very young SNe than in our 15A strategy assuming 6 continuous nights of observations. Given that in our simulations 30 sec LSST exposures would produce similar limiting--magnitudes than our 87 sec DECam exposures (a simple scaling suggests that 30 sec $g$ band LSST observations would be 0.1 mag deeper than 87 sec $g$ band DECam observations), this can be approximately explained by a combination of an increased number of fields that can be observed up to a limiting--magnitude during one night (3.4 x larger) and LSST's larger FoV (3.2 x), and is only slightly more than what would be obtained scaling the two telescope's etendues. We did a similar exercise for the Hyper Suprime--Cam instrument and we obtained a similar result, i.e. that the predicted number of SBO and young SNe should approximately scale with etendue. Continuous observations with large etendue telescopes from space or from polar regions of the earth could provide high cadence observations without the large airmass and limiting--magnitude variations experienced by this survey. In fact, the recent detection of a red supergiant SBO candidate with the large etendue, space--based \emph{Kepler} observatory \citep{2016ApJ...820...23G} combines both the continuous high cadence and large etendue required for this purpose. Alternatively, moderately large etendue, space based X--ray or UV observatories may be better tools to look for SBOs in BSG and RSG stars, but with the \emph{GALEX} mission now completed we will need to wait for future space missions for this to happen \citep[e.g.][]{2002SPIE.4497..115F}. Interestingly, gravitational waves \citep{2016PhRvL.116f1102A} or neutrino detections \citep{1987PhRvL..58.1490H, 2013ApJ...778...81K} may be complementary methods for obtaining earth based high cadence observations of SBOs if candidates with a signature consistent with core collapse SNe in nearby galaxies within the detectors localization errors are targeted, but this will require either very wide field--of--view telescopes or arrays of robotic telescopes observing several galaxies simultaneously and able to reach the necessary absolute magnitudes shown in Figure~\ref{fig:SBO_LCs}. Finally, it should be noted that the current observation and data analysis strategy followed by HiTS allowed us to study the optical Universe with an unprecedented combination of total observed volume, high cadence for several continuous nights and real--time data reduction including visualization, emphasizing the importance of interdisciplinary collaborations for astronomy. These observations are currently being used as a verification dataset for the LSST software stack \footnote{http://dm.lsst.org/} and will likely become a legacy dataset for different scientific purposes. | 16 | 9 | 1609.03567 |
1609 | 1609.08169_arXiv.txt | We describe how to estimate the velocity dispersions of ultra diffuse galaxies, UDGs, using a previously defined galaxy scaling relationship. The method is accurate for the two UDGs with spectroscopically measured dispersions, as well as for ultra compact galaxies, ultra faint galaxies, and stellar systems with little or no dark matter. This universality means that the relationship can be applied without further knowledge or prejudice regarding the structure of a galaxy. We then estimate the velocity dispersions of UDGs drawn from two published samples and examine the distribution of total masses. We find, in agreement with the previous studies of two individual UDGs, that these systems are dark matter dominated systems, and that they span a range of at least $10^{10} < M_{200}/M_\odot < 10^{12} $. These galaxies are not, as an entire class, either all dwarfs or all failed $L_*$ galaxies. Estimates of the velocity dispersions can also help identify interesting subsets of UDGs, such as those that are likely to have the largest mass-to-light ratios, for subsequent spectroscopic study. | Deep imaging surveys are uncovering extensive samples of low surface brightness galaxies. Some low luminosity examples are satellites of our own Galaxy \citep{bechtol,drlica} and others, much larger and more luminous, lie well beyond the Local Group \citep{vdk,mihos,koda,munoz,roman}. These objects potentially hold key clues on open questions as diverse as the nature of dark matter \citep{ackerman} and the drivers of galaxy formation \citep{agretz}. To understand these galaxies and utilise them in addressing these broader questions, we must measure their masses. Are the distant low surface galaxies, broadly referred to as ultra diffuse galaxies or UDGs, ``failed" massive galaxies or spatially extended dwarf galaxies \citep{vdk,beaslyb,amorisco}? A recent study required 33.5 hrs of exposure time on a 10m telescope to obtain the integrated light spectrum from which the line of sight velocity dispersion of a single UDG was measured \citep{vdk16}. We are reaching the limit of what we can accomplish with current capabilities. Compiling large samples of such galaxies with measured internal kinematics is beyond what we can hope to do. How to choose the systems on which to spend our valuable resources most fruitfully? We propose exploiting galaxy scaling relations that relate photometric and kinematics parameters. Such relations also depend on the distance to the source, and so have most commonly been used in combination with measured photometry and kinematics to estimate distances \citep[for example,][]{tf}. However, recent large samples of UDGs \citep{vdk,munoz,koda} are confined to galaxies in galaxy clusters, which means that the distances are known. For galaxies with known or estimated distances we can use scaling relations to solve for the internal kinematics. This approach, utilising the Tully-Fisher relation, has been used for high surface brightness galaxies \citep{cole,gonzalez}. However, the Tully-Fisher relation does not apply to UDGs because their morphology suggests that they are not rotating disk galaxies. Below we describe the application of a scaling relation that does apply to low velocity dispersion, pressure supported systems. We will demonstrate that using this approach we recover the measured velocity dispersions of the two UDGs that have been spectroscopically measured so far. We then apply the method to two large, published samples of UDGs and conclude that the total masses of these systems range from those of dwarf galaxies to Milky-Way like systems. We conclude that UDGs, as a class, cannot be thought of as either all dwarfs or all failed $L_*$ galaxies. | We place ultra diffuse galaxies (UDGs) on an existing galaxy scaling relation to estimate their line of sight velocity dispersions. We find that the two UDGs with spectroscopically measured dispersions satisfy the scaling relation sufficiently well that our estimated velocity dispersions are within the spectroscopic $1\sigma$ uncertainties. Assuming that these two galaxies signify that UDGs as a class satisfy the scaling relation, we derive velocity dispersions, and the related enclosed dynamical masses, for the full set of UDGs in the Coma \citep{vdk} and Fornax \citep{munoz} clusters. We reach the following conclusions: $\bullet$ DF44 appears to lie in a massive ($\sim 10^{12} M_\odot$) as found by \cite{vdk16}, but is not representative of UDGs. It lies at the upper end in size and enclosed mass of the known UDGs. Therefore, while it appears to be what has been referred to as a failed L$_*$ galaxy, it is not typical of UDGs. $\bullet$ Different samples of UDGs can be quite different. The Coma and Fornax samples have little overlap in half light radii and that translates to little overlap in enclosed mass and total mass. The Fornax UDGs generally have $M_{200} < 10^{11} M_\odot$ and so can plausibly be referred to as dwarf galaxies (e.g. LMC-like and smaller). However, they are not fully representative of UDGs because the Fornax sample does not include an object as extreme as DF44. We conclude that UDGs are neither all dwarfs nor all failed L$_*$ galaxies. Subsequent spectroscopic measurements of velocity dispersions will help resolve whether the Fundamental Manifold can be used to reliably estimate velocity dispersions of UDGs that have known distances. If so, we will then be able to explore a number of questions regarding the nature of these objects without the overwhelming burden of obtaining spectroscopic velocity dispersions for large samples of these challenging galaxies. | 16 | 9 | 1609.08169 |
1609 | 1609.02388_arXiv.txt | Detection of periodicity in the broad-band non-thermal emission of blazars has so far been proven to be elusive. However, there are a number of scenarios which could lead to quasi-periodic variations in blazar light curves. For example, orbital or thermal/viscous period of accreting matter around central supermassive black holes could, in principle, be imprinted in the multi-wavelength emission of small-scale blazar jets, carrying as such crucial information about plasma conditions within the jet launching regions. In this paper, we present the results of our time series analysis of $\sim 9.2$ year-long, and exceptionally well-sampled optical light curve of the BL Lac OJ 287. The study primarily uses the data from our own observations performed at the Mt. Suhora and Krak\'ow Observatories in Poland, and at the Athens Observatory in Greece. Additionally, SMARTS observations were used to fill in some of the gaps in the data. The Lomb-Scargle Periodogram and the Weighted Wavelet Z-transform methods were employed to search for the possible QPOs in the resulting optical light curve of the source. Both the methods consistently yielded possible quasi-periodic signal around the periods of $\sim 400$ and $\sim 800$ days, the former one with a significance (over the underlying colored noise) of $\geq 99\%$. A number of likely explanations for such are discussed, with a preference given to a modulation of the jet production efficiency by highly magnetized accretion disks. This supports the previous findings and the interpretation reported recently in the literature for OJ 287 and other blazar sources. | } Quasi-periodic oscillations (QPOs) are routinely found in the X-ray light curves of Galactic binary systems \citep[e.g.,][for a review]{Remillard06}. In radio-quiet AGN, the first detection of short ($\sim 1$\,h) X-ray QPOs was reported by \citet{Gierlinski08} for the narrow-line Seyfert 1 galaxy REJ\,1034+396, followed by subsequent discoveries \citep[e.g.,][]{Alston15}. In the optical domain, 5.2\,yr periodicity has recently been found in the light curve of the radio-quiet quasar PG\,1302-102 \citep{Graham15a}, and some other analogous candidates have been identified (\citealt{Graham15b}; see also \citealt{Liu15} for the case of the high-redshift radio-loud quasar PSO\,J334.2028+01.4075, and \citealt{Zheng16} for the radio-quiet quasar SDSS\,J0159+0105). Short time scale quasi-periodic oscillations in optical or X-ray blazar light curves --- with characteristic timescales of the orders of hours or even minutes --- have been claimed in the past in several cases (e.g., \citealt{Lachowicz09} for PKS\,2155-304; e.g. \citealt{Rani10} for S5\,0716+714); these results have not been confirmed for the other periods in the same sources. On the other hand, analogous findings regarding quasi-periodicity in blazar light curves on longer timescales of months or years seem more robust, as summarized below. At radio frequencies, harmonics with periods ranging from about one year up to several years or a decade have been reported for AO\,0235 (\citealt{Liu06}; see also \citealt{Raiteri01}), PKS\,1510-089 \citep{Xie08}, NRAO\,530 \citep{An13}, and PKS\,1156+295 (\citealt{Wang14}, who confirmed the previous analyses by \citealt{Hovatta07,Hovatta08}). Also, \citet{King13} reported persistent $\sim 150$\,days periodicity in the 15\,GHz light curve of J1359+ 4011. Similar results have been reported in the optical domain, including \citet{Sandrinelli14,Sandrinelli16}, who confirmed the $\sim 315$\,days period in the light curve of PKS\,2155-304 claimed previously by \citet{Zhang14}, or \citet{Sandrinelli15,Sandrinelli16} who found marginally significant ($\sim 3\sigma$) periodicity in PKS\,0537-441, OJ\,287, 3C\,279, PKS\,1510-089 and PKS\,2005-489, on timescales ranging from tens of days up to a few years (often in harmonic relations). In the X-ray domain, \citet{Rani09} found the $\sim 17$\,days and $\sim 420$\,days quasi-periodicity in AO\,0235+164 and 1ES\,2321+419, respectively. In the high-energy $\gamma$-ray regime, \citet{Sandrinelli14} first detected the $\sim 635$\,day period in the light curve of PKS\,2155-304, most likely in a harmonic relation with the QPO found in the source at optical frequencies. More recently, \citet{Ackermann15} reported a periodic behavior of the blazar PG\,1553+113 with the characteristic timescale of 2.2\,yr (similar to the one found at optical frequencies, but not in the radio domain). BL Lac object OJ~287 (RA=08h\ 54m\ 48.87s, Dec= +20d\ 06m\ 30.64s and $z = 0.306$) is by now one of the best studied blazars in the history. Sparsely monitored for more than a century, and intensely observed during the last $40-50$ years, it has become the most famous case of a blazar periodicity, with its double optical outbursts repeating every $\sim 12$\,yr \citep{Sillanpaa88,Hudec13}. Besides this, a number of authors have also claimed the presence of QPOs in the source at various timescales: a $\sim 40-50$\,days quasi-periodicity for the source has been claimed at radio frequencies by \citet{Wu06}, and at optical frequencies by \citet{Pihajoki13}. Also \citet{Sandrinelli16} reported to have found a marginally-significant ($\sim 3\sigma$) signal in the optical light curve of OJ\,287 with the period of $\sim 435$\,days, in addition to a much less significant one of 203\,days. In this paper, we present our analysis and results of a search for quasi-periodical oscillations in the $9.2$\,year-long and exceptionally densely-sampled optical light curve of the blazar OJ 287, utilizing several widely-used statistical methods, including the Lomb-Scargle Periodogram, and Weighted Wavelet Z-transform. The results of the analysis are given below. | } Periodicity in the light curves of blazar sources (and, more generally, of other types of AGN), may be related to the period of a perturbing object in a binary black hole system \citep[eg][]{Lehto96,Graham15a}, or a jet precession caused by either closely orbiting binary black holes or warped accretion disks \citep[see the discussion in, e.g.,][]{Graham15b,Sandrinelli16}. Moreover, coherent helical/non-ballistic motions of relativistic blobs within blazar jets could be responsible for quasi-periodicity of blazar sources \citep[e.g.,][]{Camenzind92, Mohan15}. Finally, jet modulation by various instabilities developing within the innermost parts of accretion disks, in principle, could result in quasi-periodic modulation of the observed jet emission \citep[see the discussion in][]{Liu06,Wang14}. In the specific context of OJ\,287, previous claims of the detection of (quasi-)periodic oscillations on year-like timescales --- besides the 60-year modulation advocated by \citet{Valtonen06} and the 12-year cycle discussed by \citet{Sillanpaa88} --- have been made by a number authors. For example, \citet{Hughes98} identified a persistent oscillation in the radio light curve of the source, with a characteristic timescale of about $410$ days. This result was supported by the wavelet analysis attempted by \citet{Hovatta08}. More recently, using the data for the partly overlapping observing epoch considered in this paper, \citet{Sandrinelli16} claimed to have detected $\simeq 435\,$d and $\simeq 412\,$d period in the NIR-optical and $\gamma$-ray light curve of the blazar, respectively. The detection of the same signal in more than one band --- radio, optical, \emph{and} $\gamma$-rays --- enforce the physical relevance of the findings. It is therefore interesting to indicate at this point a possible analogy between a year-like periodicity we see in OJ~287 and the low frequency `C-type' QPOs commonly found in Galactic X-ray binaries. These QPOs come in harmonic peaks sticking out of a power spectrum at frequencies just above/around the transition from the ``white noise'' to the ``pink/red noise'' segments of the PSDs, and are typically linked to the Lense-Thirring precession of the innermost parts of accretion disks \citep[see, e.g.,][]{Stella98,Motta11}. Considering a possible explanation for the observed QPOs in OJ~287, we note that since binary black holes have been claimed to shape the 12\,yr periodicity of the optical light curve in the source (\citealt{Valtonen08}; but see also \citealt{Villforth10} for a critical review), the same scenario could not account for the $\sim 400$\,d period found in our analysis. Jet precession due to warped accretion disk, on the other hand, as well as a ``grand-design'' helical magnetic field in the OJ\,287 jet, seem both in conflict with the erratic jet wobbling observed on parsec scales by \citet{Agudo12} in the high-frequency radio domain at the time of our monitoring program. Hence we conclude that the likely explanation for the $\sim 400$\,d period in the blazar (with the possibly accompanying $\sim 800$\,d harmonic) could be a jet modulation by the innermost parts of the accretion disk. Interestingly, in the case of magnetically-arrested disks, the characteristic timescale of quasi-periodic oscillations in jet the production efficiency set by the rotating and unstable (``chocking'') magnetic field accumulated at the saturation level around the horizon of a spinning black holes, as seen in recent MHD simulations \citep{Tchekhovskoy11,McKinney12}, corresponds to tens/hundreds of the gravitational radius light-crossing times, $\sim 100 \, r_g/c \sim 10^6 \, (M_{BH}/10^{9} M_{\odot})$\,s. This is of the correct order of magnitude assuming $M_{BH} \simeq 2 \times 10^{10} \, M_{\odot}$ suggested by \citet{Valtonen08} for OJ~287. Even so, in such a case the year-like periodicity claimed for radio and optical light-curves of other blazars (most likely hosting smaller black holes) would remain a puzzle, unless lower black hole spins are considered, allowing to accommodate longer timescales for smaller black hole masses. | 16 | 9 | 1609.02388 |
1609 | 1609.09579_arXiv.txt | We present a consistent total flux catalogue for a $\sim$1 deg$^2$ subset of the COSMOS region (R.A. $\in [149.55\degr, 150.65\degr]$, DEC $\in [1.80\degr, 2.73\degr]$) with near-complete coverage in 38 bands from the far-ultraviolet to the far-infrared. We produce aperture matched photometry for 128,304 objects with $i < 24.5$ in a manner that is equivalent to the \citet{wright16a} catalogue from the low-redshift ($z < 0.4$) Galaxy and Mass Assembly (GAMA) survey. This catalogue is based on publicly available imaging from \textit{GALEX}, CFHT, Subaru, VISTA, \textit{Spitzer} and Herschel, contains a robust total flux measurement or upper limit for every object in every waveband and complements our re-reduction of publicly available spectra in the same region. We perform a number of consistency checks, demonstrating that our catalogue is comparable to existing data sets, including the recent COSMOS2015 catalogue \citep{laigle16}. We also release an updated \citet{davies15} spectroscopic catalogue that folds in new spectroscopic and photometric redshift data sets. The catalogues are available for download at \url{http://cutout.icrar.org/G10/dataRelease.php}. Our analysis is optimised for both panchromatic analysis over the full wavelength range and for direct comparison to GAMA, thus permitting measurements of galaxy evolution for $0 < z < 1$ while minimising the systematic error resulting from disparate data reduction methods. | \label{sec:intro} Wide-area multiwavelength surveys such as the Sloan Digital Sky Survey (SDSS; \citealt{york00}) and the UKIRT (UK Infrared Telescope) Deep Sky Survey \citep{lawrence07} have enabled the study of large, statistical samples of galaxies. However, such surveys are generally limited to low redshifts ($z ~< 0.3$), a single facility, and one region of the electromagnetic spectrum --- usually the ultraviolet, optical or near-infrared. To produce a comprehensive picture of galaxy evolution, one must observe galaxies over an extensive range of wavelengths to probe multiple physical properties. This requires the combination of multiple data sets across observatories and instruments, and thus the consolidation of disparate sensitivities, resolutions and data reduction techniques (see e.g. \citealt{driver15}). Obtaining a consistent, optically-motivated photometric catalogue for large multiwavelength datasets is highly non-trivial \citep{laidler07,wright16a}. Naively position matching existing catalogues gives rise to the possibility of table mismatches, especially when joining high-resolution (resolution $\sim$ 0.8\arcsec) optical data to low-resolution far-infrared data (resolution $\sim$18\arcsec). Disparate data reduction methods, even though they may represent the most appropriate photometry in each individual band, may use differently sized and shaped apertures for the same object and hence probe different physical scales. More subtly, the different means of calculating errors by different survey teams will affect the quality of spectral energy distribution (SED) fits for a particular galaxy. \citet{wright16a} show that the use of a multiwavelength catalogue derived using the same data reduction procedure across the full wavelength range results in reduced photometric inconsistency, and improves the accuracy of SED fits and star formation rate estimators compared to an equivalent catalogue constructed from table matching alone. One technique to construct a consistent multiwavelength catalogue is a variation of (forced) matched aperture photometry, proceeding initially with aperture definition on a high-resolution image. The apertures are then propagated to the lower resolution data after convolution with the point spread function and appropriate deblending. Software packages implementing this technique include TFIT \citep{laidler07} and the Lambda-Adaptive MultiBand Deblending Algorithm in R (\textsc{lambdar}; \citealt{wright16a}). One dataset that lends itself to the construction of such a catalogue is the Galaxy and Mass Assembly (GAMA; \citealt{driver11,liske15}) survey. GAMA is a highly complete low-redshift spectroscopic and multiwavelength imaging campaign that aims to characterise the distribution of energy, mass and structure from kiloparsec to megaparsec scales. The GAMA spectroscopic campaign targeted 230 degrees of sky using the AAOmega spectrograph on the 3.9~m Anglo-Australian Telescope, obtaining redshifts for $\sim$250,000 galaxies. This spectroscopy is complemented by ultra-violet imaging from the \textit{GALaxy Evolution eXplorer} (\textit{GALEX}; \citealt{martin05}), optical imaging from SDSS and the Kilo-Degree Survey (KiDS; \citealt{dejong15}), near-infrared imagery from the VISTA (Visible and Infrared Survey Telescope for Astronomy) Kilo-degree Infrared Galaxy (VIKING; \citealt{edge13}) survey, mid-infrared imagery from the \textit{Widefield Infrared Survey Explorer} \citep{wright10} and far-infrared imagery from \textit{Herschel}-Atlas \citep{eales10} --- see summary in \citet{driver15}. The project has examined a wide variety of topics, including the cosmic spectral energy distribution (e.g. \citealt{driver12}, Andrews et al. in prep), star formation rates \citep{davies16a}, large scale structure (e.g. \citealt{alpaslan14}), galaxy groups (e.g. \citealt{robotham11}), close pairs (e.g. \citealt{davies15b,davies16b}) and galaxy properties and structure (e.g. \citealt{taylor11,kelvin12,loveday15,moffett16}). The GAMA survey, while scientifically comprehensive, by design only probes the low redshift Universe ($z < 0.4$). It is therefore beneficial that an intermediate redshift ($0.3 < z < 1$) equivalent to GAMA is established in order to explore a broader time baseline. The Cosmological Evolution Survey (COSMOS; \citealt{scoville07}) region, covering 2~deg$^2$ of sky centred on R.A. = 10h00m28.6s, DEC = +02$\degr$12\arcmin21\arcsec.0 is suitable for this purpose. The program is anchored by F814W observations of the field using the \textit{Hubble Space Telescope} and has been expanded to include deep observations spanning from X-ray wavelengths to radio -- with observations conducted and released using \textit{Chandra}, \textit{GALEX}, the Canada-France-Hawaii Telescope (CFHT), Subaru, VISTA, \textit{Spitzer} and \textit{Herschel}. Spectroscopic surveys in the COSMOS region include zCOSMOS \citep{lilly07,lilly09}, the PRIsm MUlti-object Survey (PRIMUS; \citealt{coil11,cool13}), the VIMOS-VLT Deep Survey (VVDS; \citealt{garilli08}), the VIMOS Ultra Deep spectroscopic survey \citep{lefevre15}, the FMOS-COSMOS survey \citep{silverman15}, 3D-\textit{HST} \citep{brammer12} and SDSS DR10 \citep{ahn14}, complemented by large catalogues of photometric redshifts \citep{ilbert09,muzzin14,laigle16}. COSMOS has been used to study many aspects of galaxy formation and evolution, including the evolution of specific star formation rates (e.g. \citealt{ilbert15}), effects of environment on galaxy morphology (e.g. \citealt{capak07b}), high-redshift quasars (e.g. \citealt{masters12}) and dust obscured galaxies (e.g. \citealt{riguccini15}). However, the multiwavelength dataset was processed with different flux measurements and reduction methods resulting in a corresponding increase in systematic error. Here, we construct a catalogue of consistent total flux measurements spanning from the far-ultraviolet to the far-infrared for a subregion we shall refer to as G10 and based on existing COSMOS imagery. Our catalogue, when combined with the spectroscopic redshifts curated by \citet{davies15}, forms an intermediate redshift sample prepared in an identical way to and thus suitable for direct comparison to GAMA. The combined multiwavelength dataset is able to sample multiple processes occurring in the galaxy population across $0 < z < 1$, including (rest frame) ultraviolet light from star formation, optical and near-infrared emission from young and old stars, mid infrared emission from polycyclic aromatic hydrocarbons and warm dust (50~K), and far-infrared emission from cold dust (20~K). In Section \ref{sec:data}, we describe the multiwavelength dataset used. In Section \ref{sec:phot}, we construct a consistent 38 band photometric catalogue spanning the far-ultraviolet to the far-infrared in a subset of the COSMOS region using \textsc{lambdar}. In Section \ref{sec:con}, we demonstrate consistency with existing photometric catalogues in the region. Sections \ref{sec:release} and \ref{sec:conclusion} summarise the release content and present concluding remarks respectively. In four upcoming papers we use this data in conjunction with GAMA to examine stellar and dust masses (Driver et al. 2016 in prep, Wright et al. 2016 in prep.), the cosmic spectral energy distribution (Andrews et al. 2016 in prep) and star formation rates (Davies et al. 2016 in prep). We use AB magnitudes throughout this work. % | \label{sec:conclusion} We have produced a 38 band photometric catalogue in COSMOS spanning from far-ultraviolet to far-infrared wavelengths in a manner consistent with the equivalent \citet{wright16a} GAMA catalogue. We gathered multiwavelength imagery from the \textit{GALEX} Deep Imaging Survey, COSMOS, UltraVISTA, S-COSMOS, SPLASH, PEP and HerMES surveys. From this data, we obtained consistent total flux measurements for a sample of 185,907 sources using \textsc{lambdar}. This required the construction of an aperture catalogue by manually checking and adjusting raw \textsc{SExtractor} output, a process that is prohibitively labour-intensive for the next generation of galaxy surveys. We demonstrate that the resulting photometric catalogue has accurate astrometry, is consistent with existing photometric datasets and achieves adjacent colour distributions --- a proxy for photometric measurement error --- comparable to existing data sets. The released catalogue is complete for objects with $i < 24.5$~mag and partially complete to $i < 25$~mag due to a rigid flux cut made prior to the \textsc{lambdar} measurements. As our catalogue is designed for panchromatic analysis, including SED fitting, we tested it for a sample of 5619 galaxies using \textsc{magphys}. We found improved convergence and goodness of fit with our catalogue compared to table matching archival photometry. The catalogues and a cutout generator for the multiwavelength imagery used are available at \url{http://cutout.icrar.org/G10/dataRelease.php}. This sample will be used in future observations as an input catalogue for a spectroscopic survey to complete the G10 region. This catalogue will form the basis for a GAMA-equivalent multiwavelength database at intermediate redshifts. This database will enable the derivation of physical properties and structural parameters for the COSMOS region using the same techniques as GAMA and enable comparative studies of the cosmic spectral energy distribution (Andrews et al. in prep), galaxy structure and morphology, star formation rates (Davies et al. in prep), stellar masses (Wright et al. in prep) and panchromatic measurements of the extragalactic background light \citep{driver16b}. In combination with further spectroscopic observations of the G10 region, we will create catalogues of groups (akin to \citealt{robotham11}) and large scale structure. In addition to enabling comparisons to low-redshift galaxy evolution surveys, these catalogues will pave the way for future galaxy evolution surveys such as WAVES \citep{waves} and provide a basis for optically motivated stacking using 21~cm data from the COSMOS HI Large Extragalactic Survey (CHILES; \citealt{fernandez13}). | 16 | 9 | 1609.09579 |
1609 | 1609.00458_arXiv.txt | We have investigated the complex multiwavelength evolution of \SF\ during the rise of its 2005 outburst. We detected two hard X-ray flares, the first one during the transition from the soft state to the ultra-soft state, and the second one in the ultra-soft state. The first X-ray flare coincided with an optically thin radio flare. We also observed a hint of increased radio emission during the second X-ray flare. To explain the hard flares without invoking a secondary emission component, we fit the entire data set with the \emph{eqpair} model. This single, hybrid Comptonization model sufficiently fits the data even during the hard X-ray flares if we allow reflection fractions greater than unity. In this case, the hard X-ray flares correspond to a Comptonizing corona dominated by non-thermal electrons. The fits also require absorption features in the soft and ultra-soft state which are likely due to a wind. In this work we show that the wind and the optically thin radio flare co-exist. Finally, we have also investigated the radio to optical spectral energy distribution, tracking the radio spectral evolution through the quenching of the compact jet and rise of the optically thin flare, and interpreted all data using state transition models. | \label{sec:intro} Galactic black hole transients (GBHT) are systems that occasionally go into outburst, during which their X-ray luminosity may increase several orders of magnitude when compared to their quiescent levels. These objects are excellent laboratories to study the complex relationship between jets, winds and the accretion environment as the outbursts evolve on time scales of months. This rapid evolution allows for the detailed investigation of the properties of accretion states which are traced by X-ray spectral and timing properties, and the properties of outflows, where the jets are traced by the radio and optical/infrared (OIR) emission and the winds are traced by the properties of X-ray and optical absorption features. A detailed description of X-ray spectral and timing states of GBHTs can be found in \cite{McClintock06book} and \cite{Belloni10_jp}. At the start of a typical outburst, the GBHT is in the hard state (HS). In this state, the X-ray spectrum is dominated by a hard, power-law like component associated with Compton scattering of soft photons by a hot electron corona. Faint emission from a cool, optically-thick, geometrically-thin accretion disc may also be observed, which can be modelled by a multi-temperature blackbody \citep{Makishima86}. This state also exhibits strong X-ray variability (typically $>$20\% of the fractional rms amplitude). As the outburst continues and the X-ray flux increases, the GBHT usually transitions to a soft state (SS) in which the X-ray spectrum is now dominated by the optically-thick accretion disc, displaying low levels of X-ray variability ($<$ a few \%) and faint power-law emission. Between the HS and the SS, the source may transition through the hard and soft intermediate states (HIMS and SIMS, respectively) with properties in between the hard and soft states \citep[see][for further details]{Belloni10_jp}. Finally, some sources (e.g. \SF, Cyg X-3) may also show a so-called ultra-soft state (US), which has an extremely steep X-ray power-law (with a photon spectral index $\Gamma$ of $>$3) with a completely dominating disc contribution \citep{Szostek08, Zdziarski04}. \SF\ shows an even softer X-ray state, denoted the ``hyper-soft" state \citep{Uttley15}. GBHTs also show distinct multiwavelength characteristics throughout their outbursts. Radio and optical-infrared (OIR) observations indicate the presence of compact jets which exhibit flat to inverted radio spectrum (such that the radio spectral index $\alpha\gtrsim0$, $S_{\nu} \propto \nu^{\alpha}$ where $S_{\nu}$ is the radio flux density at frequency $\nu$) in the hard state \citep{Tananbaum72, Buxton04, Corbel13, Gallo10} which become quenched in the soft state \citep{Russell11, Fender99, Coriat11}. During the transition from the hard to soft state, the compact jets give way to relativistic and bright transient jets with an optically-thin radio spectrum \citep[$\alpha < 0$,][]{Fender09, Vadawale03}. Recent high-resolution grating observations of GBHTs and neutron stars revealed the presence of blue-shifted absorption features, especially \ion{Fe}{xxv} and \ion{Fe}{xxvi} lines, showing that these sources not only produce collimated jets, but can also drive winds \citep[][and references therein]{DiazTrigo14, Neilsen12}. Wind signatures are preferentially detected in soft states for high-inclination sources \citep{Done07, Ponti12}, where the inclination dependence indicates a thermal origin for the wind \citep{Begelman83}. A single observation of \SF\ on MJD~53461.5 revealed a rich series of absorption lines from a dense, highly-ionized wind, which was initially interpreted as magnetically-driven \citep{Miller06}. Most follow up studies have supported the magnetic origin of the wind \citep[][and references therein]{Neilsen12}. However, thermally-driven winds remain a possibility \citep{Netzer06}. \subsection{\SF} \label{sub:source} \SF\ was first discovered with the Burst and Transient Source Experiment (BATSE) on-board the Compton Gamma Ray Observatory \citep{Zhang94_iauc}. Subsequent radio observations revealed apparent-superluminal relativistic (0.92 $c$) jets \citep{Hjellming95, Tingay95}. Optical observations taken in quiescence indicate a FIII-FV giant or sub-giant with an orbital period of 2.62 days \citep{Orosz97}. In this study, we used primary and secondary masses of 6.3$\pm$0.5 $\rm M_{\odot}$ and 2.4$\pm$0.4 $\rm M_{\odot}$, respectively, which were obtained by modelling the ellipsoidal orbital modulations in quiescence \citep{Greene01}. The same model indicates a binary inclination of 70.2$^{\circ} \pm$1.9$^{\circ}$ which is consistent with deep absorption dips \citep{Kuulkers00} and strong wind emission \citep{Ponti12}. Alternative mass measurements exist \citep{Beer02, Shahbaz03}, but the differences are small and have no effect on our conclusions. The binary inclination angle of \SF\ is slightly different from the disc inclination angle obtained from radio imaging \citep{Orosz97, Maccarone02}. The distance to the source has been estimated via different methods, where the majority of works use a distance of 3.2$\pm$0.2 kpc, based on the analysis of \cite{Hjellming95}, which we also adopt. Between 1994 and 1997 BATSE detected several outbursts from \SF\ \citep{Zhang97_2}. These early outbursts showed a complex pattern between the hard X-ray flares and the optically-thin radio flares; the first three hard X-ray flares were very well correlated with superluminal radio flares \citep{Harmon95}. However, subsequent hard X-ray flares were not associated with any increased radio emission \citep{Tavani96, Zhang97_2}. This article investigates the multiwavelength evolution of \SF\ during its 2005 outburst rise, which was first detected on February 17 (MJD~53419) with the Proportional Counter Array (PCA) instrument on-board the Rossi X-ray Timing Explorer (\rxte) \citep{Markwardt05_atel}. The source was intensely monitored with \rxte\ throughout this outburst. There was also exceptional multiwavelength coverage during this outburst, which was followed daily in OIR with the Small and Moderate Aperture Research Telescope System (SMARTS; this work and \citealt{Kalemci13}), as well as frequent radio observations with the Very Large Array (VLA; this work and \citealt{Shaposhnikov07}). Grating observations with \chandra\ and \xmm\ taken during this outburst have also revealed wind features, the origin of which is still under debate \citep{Miller06, DiazTrigo14, Shidatsu16}. The structure of the paper is as follows. In \S\ref{sec:obs} we describe the multiwavelength observations and provide detailed information of the spectral extraction in all observing bands. In \S\ref{sub:states} we describe the source states and transitions during the rise, and then, for the first time, discuss the properties of the hard X-ray flares in \S\ref{sub:hard}. To explain the possible origin of hard X-ray flares, we conducted spectral fits with \emph{eqpair}, which are discussed in \S\ref{sub:eqpair}. The radio to OIR spectral energy distribution (SED) are shown in \S\ref{sub:SED}, with emphasis on the optically-thin radio flare. Finally we discuss our findings, focusing on the origin of hard X-ray flares and the relationship between the radio and wind emission. | \label{sec:discussion} \subsection{State transitions} While the general spectral and timing evolution, as well as the spectral state identification of this source have previously been discussed \citep{Shaposhnikov07, Motta12}, here, we discuss the evolution in terms of the \emph{eqpair} parameters and their relation to the multiwavelength evolution. First of all, the transition out of the hard state was marked by an increase in the optical depth ($\tau$) and an increase in the non-thermal electron distribution ($l_{nth}/l_{th}$). During the transition, the radio flux decreased and the radio spectrum became optically-thin as shown in Figs.~\ref{fig:seds} and \ref{fig:radind} (also see \citealt{Shaposhnikov07}). A similar evolution was observed in H1743$-$322 \citep{MillerJ12}, which showed radio spectral index softening and a slight drop in radio flux during the HIMS before the quenching of the compact jet and the launching of an optically-thin radio flare. While the H1743$-$322 radio flare peaked during the SS, VLBA observations indicated that the time of launch was close to the transition from the HIMS to the SIMS, which was a few days before the peak radio flux. It was not possible to do a similar analysis for \SF\, because the jet was not resolved. However, if we assume similar time scales it is possible that the optically-thin ejecta were launched earlier in this system. Such a delay would make the radio and hard X-ray flares out of synch, with the radio preceding the hard X-ray flare. % \cite{Corbel12} showed that the radio flux and radio spectral index gradually increasing (becoming flat and then inverted) as the compact jets were re-launched during the outburst decay of GX~339$-$4. MAXI~J1836$-$194 showed a similar evolution, where the radio spectrum softened as the source entered the HIMS from the HS, and then became highly inverted again in the hard state during the decay \citep{Russell13a, RussellT14}. A natural interpretation of this would be that the jets become more collimated and compact during the transition to the hard state during the decay, possibly as the magnetic flux accumulates close to the inner parts of the accretion flow. H1743$-$322 and this source show that perhaps the reverse evolution is taking place during the rise, that the magnetic flux diffuses out faster than it can be accumulated \citep{Begelman14}, reducing the power and collimation of the jet. However, the relatively strong flux at 24 $\mu$m during the HIMS (on MJD~53439.5) might suggest that this was not occurring in \SF. As discussed in \cite{Migliari07} and shown in Fig.~\ref{fig:seds}, it is difficult to explain the flux level as emission from the outer parts of the accretion disc or as dust from a circumbinary disc because the emission was variable and much stronger than what has been observed in other sources \citep{Muno06}. If it was coming from the jet, the radio spectrum cannot be fit with a single power-law, and may include multiple components as the compact jet was quenching. As the source made its transition to the SS, the non-thermal compactness ratio peaked at a level of 1 and remained steady while the optical depth decreases. As expected, the hard-to-soft compactness ratio decreased as well. At this time, the reflection fraction was $\sim$1, indicating a compact corona and an inner disc that was close to the black hole. In our \emph{power-law}+\emph{diskbb} fits, this transition showed an increase in the folding energy of the \emph{highecut} component, indicating higher and higher cut-off energies \citep{Joinet08} as the source moved towards the SS, which is in agreement with the increasing $l_{nth}/l_{th}$. Such behaviour has been observed during the state transitions of GX~339$-$4 and Swift~J1745$-$26 \citep{delSanto16} and can be explained by the presence of a dead-zone in the intermediate states in the elevated disc model of \cite{Begelman15}. The SS to US transition was coincident with an optically-thin radio flare (though the actual ejection may have preceded the transition) and an increase in the OIR flux. As the source evolved in the SS, the $l_{nth}/l_{th}$ decreased and the electrons thermalize. However, during the hard X-ray flare, the non-thermal compactness ratio increased up to unity, with a slight increase in hard-to-soft compactness ratio. Along with high reflection fraction, a single hybrid Comptonization component was adequate to represent the X-ray spectrum. We note that the reflection fraction was not well constrained because the lower energies of the reflection component is in the part of the spectrum with the iron absorption lines and edges (where the \rxte\ data has a higher effective area), and the resolution of \rxte\ makes it impossible to resolve each component. Nevertheless, it is not clear how the electrons became non-thermal and then thermal again in the US on those time scales. A possible explanation is the disc breaking scenario of \cite{Nixon14}. Since the inclination and spin angles are misaligned in \SF, the mechanism described in \cite{Nixon14} may be able to heat up the disc over the time scales observed here. \subsection{Hard X-ray flares and radio emission} We have identified two hard X-ray flares (see Fig.~\ref{fig:hxrevol}) during the rise of the 2005 outburst of \SF. The first flare occurred during the transition from the SS to the US. This transition also coincided with an optically-thin radio flare similar to the transient jets observed in many GBHTs \citep{Dhawan00, Fender06, Gallo10}. The association of the second hard X-ray flare with a radio flare is not as clear due to lack of radio observations between MJD~53455 and MJD~53460. Historically, \SF\ has shown many instances of luminous radio flares (relativistic ejections) coinciding with hard X-ray flares. But the association cannot be examined in detail in these older data due to limited spectral capability of the BATSE instrument on \emph{CGRO}. For example, between MJD~49550 and MJD~49700, three radio flares were observed to be coincident with X-ray flares (where the X-ray peaked earlier than the radio, \citealt{Harmon95}), but no radio flare was observed within the next year even though several hard X-ray flares took place \citep{Tavani96}. However, it is possible that radio ejection events did occur between MJD~49700 and MJD~50000 but were simply missed due to the timing of the radio observations and because the relation between spectral states and radio jets was not well known at the time and it was not easy to determine the X-ray spectral state. Aside from the historical note, the important observation here is the clear association of the radio jet with the hard X-rays. This association is obvious in the case of compact jets, which are always associated with the hard X-ray spectral state, and can only turn back on (following their quenching in the soft state) when the X-ray spectrum has hardened sufficiently during the outburst decay \citep{Kalemci13}. However, this association is less clear for the optically-thin flares during the outburst rise. \cite{Fender09} investigated the relationship between the spectral hardness and major radio ejections and found that while the association is complex, at least in XTE~J1859+226, a fast hardening is associated with a major flare event. This association can be related to the production, transport and dissipation of magnetic fields in the inner disc. The presence of hard X-rays emission indicates the presence of a geometrically-thick corona, which makes it easier to produce \citep{Begelman14, Kylafis15} and transport \citep{Beckwith09} magnetic flux. Similarly, with the compact jets, the presence of some form of hot, vertically-extended accretion flow may be a necessary condition for an optically-thin radio flares as well. \subsection{OIR evolution, radio flare and winds} \label{sub:disOIR} Explaining the behaviour of OIR emission from \SF\ is a difficult task due to sub-giant secondary contributing significant emission, especially in the HS and HIMS. For other well studied systems, the OIR emission is dominated by the compact jet during the hard state rise and decay (e.g. \citealt{Kalemci13}, but also see \citealt{Veledina13} for an alternative explanation based on a hot-flow model). In fact, for GX~339$-$4 in several outbursts \citep{Coriat09}, MAXI~J1836$-$194 \citep{Russell13a, RussellT14}, and XTE~J1550$-$564 (Kalemci et al. in preparation), the OIR emission drops down significantly as the source enters the HIMS. On the other hand, we observe no decrease in the OIR flux for \SF\ during the HIMS (although it is clear that the jet flux is decreasing, and perhaps becoming optically-thin at this time as shown in Fig.~\ref{fig:seds}), in fact it rises as the source enters the SS and then the US. This peculiar behaviour has also been discussed by \cite{Shidatsu16}. disc size cannot explain this difference as the binary separations of \SF, GX 339$-$4 and XTE~J1550$-$564 are similar. The only difference between them is the high inclination of \SF\, whereas the other sources are low inclination. A possible explanation is provided in \cite{Shidatsu16}, with the scattering in a strong wind increasing the irradiation and making the disc brighter. With the PCA observations, we infer the presence of winds from the beginning of the SS and beyond based on the detection of absorption lines. Our first detection is on MJD~53441.5, around the same time as the \chandra\ observation (Obsid 5460), which started on MJD~53441.9 \citep{Neilsen12}. The date of the \chandra\ observation is shown in Fig.~\ref{fig:eqp2} with a dashed line. In \cite{Neilsen12} this observation is described to be moving out of the hard state, while in \cite{Neilsen13}, it is simply described as an observation in the hard state. Our analysis, as well as earlier timing and spectral analysis, indicate that at this time, the source had already left the hard state and was completely in the soft state by MJD~53442.0 (Figs.~\ref{fig:sptr1} and \ref{fig:ttr1}). This observation is ``harder" than the other \chandra\ observation on MJD~53461.5, which was taken in the extremely soft hypersoft state. An interesting fact overlooked by earlier works is that the optically-thin radio flare, which peaked at around MJD~53447, coexisted with the disc wind detected by both \chandra\ \citep{Neilsen13} and \xmm\ \citep{DiazTrigo07}. The times of \chandra\ and \xmm\ observations are indicated by dashed and dotted lines, respectively, in Fig.~\ref{fig:eqp2}. While the compact jet / wind dichotomy is well documented \citep{Ponti12}, this is one of the rare cases that a wind and a jet of some form are observed together in a GBHT \citep[a recent case is the discovery of deep $H$ and $He$ $P-Cyg$ profiles existing along with radio emission from a compact jet in V404 Cyg,][]{MunozDarias16}. In the $\beta$ state of GRS~1915+105 (which generally show a soft X-ray spectrum) strong winds are observed, whereas no winds are observed in the hard states. Based on this, it was claimed that the intense mass loss due to winds were prohibiting the launching of the jets in this source by halting flow of matter into the compact jet \citep{Neilsen09}. Further analysis indicated that the winds were quenched during the dips (when the jets are presumably launched, \citealt{Mirabel98}), but were strong and fast in the flaring part of the $\beta$-state \citep{Neilsen12b}. Given that the winds are launched tens of thousands of gravitational radii from the black hole, it is more natural to assume that changes in inner accretion flow regulate the outflows, and it is not surprising to observe jets and winds together in transitional states. We note that \SF\ showed optically-thin radio flares in earlier outbursts for which the flux densities reach as high as 10 Jy, and were usually larger than 100 mJy, at 1.49 Ghz \citep{Hjellming95, Harmon95}. The 2005 outburst on the other hand only reached \wsim 6 mJy at its peak. Because the radio coverage was almost daily, it is unlikely that an order of magnitude larger radio peak was missed. Therefore, in the case of \SF, a weak wind was observed together with a weak optically-thin radio jet tapping the same accretion power reservoir. | 16 | 9 | 1609.00458 |
1609 | 1609.02207_arXiv.txt | Nuclear astrophysics aims to understand the heavens using nuclear physics as a tool. Progressively more capable satellites COBE, WMAP, Planck and powerful telescopes such as Subaru and Keck, with even more powerful telescopes coming on line such as the Thirty Meter Telescope made phenomena where the microphysics is nuclear and particle physics increasingly accessible. Aside from ``practical" consequences such as unraveling the origin of the elements around us, these efforts do lead to a deeper appreciation of the night sky in a centuries old tradition: In 1689 the haiku poet Basho noted the majesty of the Milky Way enveloping the Sado island off the Echigo coast (present day Niigata), where this meeting (2106 Nuclei in Cosmos conference) is auspiciously located. In many cases the bridge between observational astronomy and laboratory experiments is provided by neutrinos. These particles indeed play a very special role in the Cosmos: since they interact only weakly they can transport energy and entropy over large distances. In addition, since they can flip the third component of the strong isospin, they control the electron fraction in environments where nucleosynthesis takes place. More than half a century after their existence was first postulated, we finally seem to be getting closer to understanding the elusive physics of neutrinos. For a long time very little experimental information was available about neutrino properties, even though a minute neutrino mass has intriguing cosmological and astrophysical implications. This situation has changed in recent decades: intense experimental activity to measure many neutrino properties took place. For example it is now well established that weak-interaction eigenstates of the neutrinos do not coincide with the mass eigenstates, but are related by a unitary transformation, which can be parameterized as \begin{equation} \label{1a} \left( \begin{array}{ccc} 1 & 0 & 0 \\ 0 & C_{23} & S_{23} \\ 0 & -S_{23} & C_{23} \end{array} \right) \left( \begin{array}{ccc} C_{13} & 0 & S_{13} e^{-i\delta_{CP}} \\ 0 & 1 & 0 \\ - S_{13} e^{i \delta_{CP}} & 0 & C_{13} \end{array} \right) \left( \begin{array}{ccc} C_{12} & S_{12} & 0 \\ - S_{12} & C_{12} & 0 \\ 0 & 0 & 1 \end{array} \right) \end{equation} where $C_{ij} = \cos \theta_{ij}$, $S_{ij} = \sin \theta_{ij}$, and $\delta_{CP}$ is the CP-violating phase. Already a few years back the angles $\theta_{12}$ and $\theta_{23}$ were determined from the solar, atmospheric, and accelerator experiments, a fact that was recognized with the 2015 Nobel Prize in physics. Within the last decade, three reactor neutrino experiments, Daya Bay \cite{An:2012eh}, RENO \cite{Ahn:2012nd}, and Double Chooz \cite{Abe:2011fz}, were also able to measure the remaining mixing angle $\theta_{13}$. At short baselines the electron antineutrino survival probability is given as \begin{equation} P(\overline{\nu}_e \rightarrow \overline{\nu}_e) = 1 - \sin^2 2 \theta_{13} \sin^2 \left( \frac{\delta m_{ee}^2L}{4E} \right) - \cos^4 \theta_{13} \sin^2 2\theta_{12} \sin^2 \left( \frac{\delta m_{21}^2 L}{4E} \right) \end{equation} where the quantity $\delta m_{ee}^2$ is defined via the equality \begin{eqnarray} \sin^2 \left( \frac{\delta m_{ee}^2L}{4E} \right) &=& \cos^2 \theta_{12} \sin^2 \left( \frac{\delta m_{31}^2L}{4E} \right) + \sin^2 \theta_{12} \sin^2 \left( \frac{\delta m_{32}^2L}{4E} \right) \nonumber \\ &=& \sin^2 \left( \frac{\delta m_{31}^2L}{4E} \right) - \sin^2 \theta_{12} \sin^2 \left( \frac{\delta m_{21}^2L}{4E} \right) . \end{eqnarray} Note that at short distances the second term in the second line is extremely small. A recent measurement of the inverse beta decay with the full detector configuration at Daya Bay results in \cite{An:2015rpe} $|\delta m_{ee}^2| = (2.42 \pm 0.11) \times 10^{-3}$ eV$^2$ and $\sin^2 2 \theta_{13} = 0.084 \pm 0.005$. A simultaneous fit to the neutron capture on hydrogen and neutron capture on Gd at Daya Bay yields $\sin^2 2 \theta_{13} = 0.082 \pm 0.004$. \cite{An:2016bvr}. An independent measurement at RENO \cite{Kim:2016yvm} gives $|\delta m_{ee}^2| = [2.62 ^{+0.21}_{-0.23} (\rm stat.) ^{+0.12}_{-0.13} (\rm syst.) ] \times 10^{-3}$ eV$^2$ and $\sin^2 2 \theta_{13} = 0.082 \pm 0.009 (\rm stat.)\pm 0.006 (\rm syst.) $. These results are also in agreement with the measurements of the Double Chooz collaboration \cite{Abe:2015rcp}. Thus not only different experiments measuring $\theta_{13}$ converge on the same value, but they also single out this angle as the most precisely known neutrino mixing angle. | Their seemingly very small masses and feeble interactions with ordinary matter make neutrinos rather special. We made a lot of progress in understanding the elusive physics of neutrinos. Much experimental information is now available with more to come about neutrino properties enabling us to use them as valuable tools in exploring the Cosmos. These tools will be put into a good use with the upcoming powerful observatories and laboratory facilities. | 16 | 9 | 1609.02207 |
|
1609 | 1609.00728_arXiv.txt | text{ Like Hipparcos, Gaia is designed to give absolute parallaxes, independent of any astrophysical reference system. And indeed, Gaia's internal zero-point error for parallaxes is likely to be smaller than any individual parallax error. Nevertheless, due in part to mechanical issues of unknown origin, there are many astrophysical questions for which the parallax zero-point error $\sigma(\pi_0)$ will be the fundamentally limiting constraint. These include the distance to the Large Magellanic Cloud and the Galactic Center. We show that by using the photometric parallax estimates for RR Lyrae stars (RRL) within 8kpc, via the ultra-precise infrared period-luminosity relation, one can independently determine a hyper-precise value for $\pi_{0}$. Despite their paucity relative to bright quasars, we show that RRL are competitive due to their order-of-magnitude improved parallax precision for each individual object relative to bright quasars. We show that this method is mathematically robust and well-approximated by analytic formulae over a wide range of relevant distances. } \begin{document} \jkashead % | } Gaia will obtain astrometry for $>10^9$ stars, with parallax precisions down to $\sigma(\pi)\lesssim 6\muas$ for the brightest stars, $V\lesssim 12$. In contrast to traditional pre-Hipparcos astrometry, Gaia is designed to measure so-called ``absolute parallaxes''. However, since nothing in nature is truly ``absolute'', it behooves us to specify more precisely exactly what Gaia will measure. In traditional narrow angle astrometry, one measures the parallactic motion of some target star relative to a set of reference stars, and from this measures the ``relative parallax'' $\pi_\rel = \pi_{\rm target} - \pi_{\rm reference}$, where the last quantity is the mean parallax of the reference stars. One then estimates the distances of the reference stars, and hence $\pi_{\rm reference}$, by some non-astrometric method, usually photometric. If, for example, the reference stars are five times farther than the target star, and if their distances can be estimated to, say 30\% precision, and assuming $N=4$ reference stars, then the contribution of the error due to the reference frame is only $\sigma(\pi_{\rm reference})/\pi_{\rm target} = 30\%/5/\sqrt{4}=3\%$, which may well be lower than the contribution from the astrometric precision of the $\pi_\rel$ measurement. In a variant of this approach, one might use external quasars or galaxies as the reference frame, in which case $\pi_{\rm reference}=0$ to a precision adequate for most purposes. By contrast Hipparcos used wide-angle astrometry, which does not require any external reference frame for parallaxes (although it does for proper motions). To understand the basic principle of this approach, consider two telescopes that are rigidly separated by $90^\circ$. Let the first of these telescopes make two measurements of a star in the ecliptic, six months apart, both times at quadrature. That is, both of these measurements will suffer maximal parallactic deflection, but in opposite directions. Now let the second telescope measure the positions of a second star, also in the ecliptic at the same epochs. Since this second star is, by construction, aligned perfectly with the Sun, it will not suffer any parallactic deflection. Hence the relative change in position of these two stars directly gives the absolute parallax of the first. Now, of course, one of these two measurements of the second star could not be made in practice because it would lie directly behind the Sun. However, the point is that by simultaneously observing stars that are affected by parallax by substantially different (and easily calculable) amounts, one can extract the absolute parallax. For this method to work to a given specified precision, the ``basic angle'' between the two telescopes must remain fixed to the same precision. Or rather, any changes in the basic angle must be understood to this specified level of precision. If the basic-angle oscillations have power on timescales shorter than the rotation period of the two telescopes, then the amplitude of these oscillations can be derived (and so corrected) from the observations themselves. However, oscillations at the rotation period are indistinguishable from a zero-point offset of all parallaxes that are being measured. Uncertainty about this amplitude is therefore equivalent to introducing a ``$\pi_{\rm reference}$'' term, as in narrow-angle astrometry. For reasons that are not presently understood, the actual amplitude of these oscillations is about 1 mas, which is orders of magnitude higher than expected from the original design, and also orders of magnitude higher than the parallax precision of the best measurements. Happily, the great majority of this oscillation can be measured from engineering data, but an ultra-precise estimate of the Gaia system parallax zero point, $\pi_0$, will require external calibration. It may well be that for most applications a precision determination of $\pi_0$ is irrelevant. However, it is easy to imagine applications for which this is important. For example, the parallax of the Large Magellanic Cloud (LMC) is presently estimated to be $\pi_{\rm LMC}=20\,\muas$ (for a comparison of measurements see \citealt{degrijs14}). Consider measurements of 10,000 LMC stars at $V=16$, each with precision $\sigma(\pi)=40\,\muas$. Each measurement by itself would be ``useless'', having a 200\% error. Nevertheless, the combination of all of them would have an error $\sigma(\pi_{\rm LMC})=0.4\,\muas$, i.e., a 2\% error. However, if the zero-point error $\sigma(\pi_0)\sim 2\,\muas$, then this LMC distance measurement would be degraded by a factor 5. One method to measure $\pi_0$ is from quasars. There is about one such object per square degree to $V_0=18$ (e.g., \citealt{hewett01}). For the $\sim 3/4$ of these that are relatively unextincted, the Gaia precision\footnote{The Gaia site http://www.cosmos.esa.int/web/gaia/science-performance gives $\sigma(\pi) =(-1.631 + 680.766 · z + 32.732 · z2)1/2 · [0.986 + (1 - 0.986) · (V-I)$ where $z=\min(10^{0.4(G-15)},10^{-1.2})$} is anticipated to be $\sigma(\pi)\sim 140\,\muas$. Since these are each known a priori to have zero parallax (or rather $\pi\ll 1\,\muas$), the 30,000 that lie over $3\pi$ sterradians can be combined to yield $\sigma(\pi_0)\sim 0.8\,\muas$. There are $\sim 3$ times more quasars $(18<V_0<19)$ than $V_0<18$, but each contributes substantially less information. Including all quasars, we estimate $\sigma(\pi_0)\sim 0.6\,\muas$ from this technique. This estimate then sets the benchmark for other techniques. If these other methods can achieve a similar or better precision, then they can serve as an independent check on the quasars and improve the overall measurement of $\pi_0$. | \label{sec:conclude}} The Gaia mission data promises to transform our understanding of the Milky Way. In this work, we have shown that exploiting photometric parallax estimates for RRL within 8 kpc in conjunction with the precise IR P-L relation for these objects, one can measure the absolute parallax zero-point $\sigma(\pi_0)$ to precision of less than $0.5(\muas)$. Not only is this extremely precise, but it is also comparable to, and completely independent of, measurements of this quantity from quasars. We further show that once this is determined, one can refine the precision of the IR P-L zero point well beyond what is possible from photometric measurements alone. We anticipate this independent method will be of immediate use to the astronomical community. | 16 | 9 | 1609.00728 |
1609 | 1609.03037_arXiv.txt | Most future surveys designed to discover transiting exoplanets, including TESS and PLATO, will target bright ($V\lesssim 13$) and nearby solar-type stars having a spectral type later than F5. In order to enhance the probability of identifying transits, these surveys must cover a very large area on the sky, because of the intrinsically low areal density of bright targets. Unfortunately, no existing catalog of stellar parameters is both deep and wide enough to provide a homogeneous input list. As the first Gaia data release exploitable for this purpose is expected to be released not earlier than late 2017, we have devised an improved reduced-proper-motion method to discriminate late field dwarfs and giants by combining UCAC4 proper motions with APASS DR6 photometry, and relying on RAVE DR4 as an external calibrator. The output, named UCAC4-RPM, is a publicly-available, complete all-sky catalog of solar-type dwarfs down to $V\simeq 13.5$, plus an extension to $\logg > 3.0$ subgiants. The relatively low amount of contamination (defined as the fraction of false positives; $<30 \%$) also makes UCAC4-RPM a useful tool for the past and ongoing ground-based transit surveys, which need to discard candidate signals originating from early-type or giant stars. As an application, we show how UCAC4-RPM may support the preparation of the TESS (that will map almost the entire sky) input catalog and the input catalog of PLATO, planned to survey more than half of the whole sky with exquisite photometric precision. | Although a couple of thousand exoplanets have been discovered, the general picture on their structure, formation and evolution is still far from being complete. Putting things in context requires a more detailed analysis of individual systems than is currently achievable. This includes, for instance, the estimate of accurate stellar masses and ages through asteroseismology and a detailed characterization of the planetary atmospheres through spectroscopy. Both these tasks require high S/N, and are feasible only when targeting planets hosted by nearby and bright stars, which have been very rare indeed so far. In the next years, two cornerstone, space-based missions that will photometrically detect and partially characterize planetary systems around bright stars will be launched: TESS \citep{ricker2015}, a NASA Explorer mission selected for launch in 2017-2018; and PLATO (\citealt{rauer2014}), a medium-class mission selected for ESA's M3 launch opportunity (2022-2024). Aside from being bright ($V\lesssim13$), the most promising targets for exoplanetary science are solar-type stars, i.~e., main-sequence stars later than spectral type F5, which could include moderately-evolved subgiants following a broader definition. As a magnitude-limited sample of field stars is dominated by distant (and intrinsically bright) giants and early-type stars, the only way to access a large sample of bright, nearby solar-type stars (as required by transit exoplanet search surveys) is to dramatically increase the covered sky area up to a significant fraction of the whole celestial sphere. In this case, the selection of a homogeneous target list requires the deepest all-sky stellar classification ever attempted, able to assign at least a spectral type and luminosity class to every star brighter than $V = 13$. This should be regarded as a minimal requirement for carrying out the target selection task. Previous experience from the CoRoT and Kepler space missions \citep{deleuil2009,brown2011} has shown that detailed knowledge of the stellar parameters of the targets (such as effective temperature $\teff$, surface gravity $\logg$, metallicity $\meh$, stellar mass and radius $\rstar$, $\mstar$, age, etc.) along with the identification and characterization of the background stars \citep{deleuil2006} helps to prioritize the targets, and makes the follow-up and the rejection of false alarms much more efficient. The Gaia mission\footnote{http://sci.esa.int/gaia/} \citep{perryman2001}, launched in 2013 and currently collecting data, is expected to play a fundamental role in the target selection by performing an unprecedented ultra-high-precision astrometric survey of nearly every source brighter than $V\simeq 20$, along with low-resolution spectrophotometry and radial velocities. An intermediate catalog which includes stellar parameters from spectrophotometry is expected to be released at the end of 2017\footnote{http://www.cosmos.esa.int/web/gaia/release}. However, we should keep in mind that i) the Gaia stellar classification will be affected by crowding in the densest fields \citep{bailerjones2013,recioblanco2016}; ii) there will be some degeneracy among certain parameters, such as temperature and interstellar extinction \citep{straizys2006,bailerjones2010,bailerjones2013}; iii) other astronomical catalogs (for instance, X/UV/IR/narrow-band photometry and activity diagnostics) will be very complementary to the Gaia measurements; iv) space missions (e.g.\ TESS) and ground-based surveys which are presently in development may require a preliminary target list before 2017, for the performance analysis, to optimize the observing strategy, to fine-tune the spacecraft design, and to begin implementing the foreseen additional observations and coordinated follow-up programs. In this paper, we first introduce the basic problem of attempting a large-scale stellar classification by relying only on wide-band photometry and proper motions (Section~\ref{problem}), and review the existing techniques and catalogs designed for that purpose (Section~\ref{catalogs}). After showing that a new approach has to be devised, in Section~\ref{ucac4-rpm} we describe how we compiled a brand-new all-sky catalog of FGK dwarfs and subgiants, called UCAC4-RPM. We include a detailed description of the catalogs used as input, the grid-based algorithm exploited to define appropriate selection criteria, and the cross-matching procedure used to estimate the contamination and completeness of the resulting sample. Finally, in Sections~\ref{discussion} and \ref{conclusions} we discuss how UCAC4-RPM can be exploited by the ongoing or forthcoming space- and ground-based transit surveys, and how it could be extended and complemented in the future. | Throughout the previous sections, we described how we devised a new RPM-based algorithm to assign a luminosity class to field stars by knowing only their proper motions and two optical magnitudes. By applying this optimal algorithm on a new stellar catalog compiled by matching UCAC4, APASS~DR6 and Tycho-2, we ended up with UCAC4-RPM ---an all-sky sample of solar-type dwarf stars complete down to at least $V\simeq 13$. We demonstrated that the latter catalog, once complemented by subgiants within the same spectral type range, meets the requirements set by the PLATO team for the target selection of its main stellar samples. In particular, the relatively low level of contamination ($\lesssim$30\%) of UCAC4-RPM, together with a $\gtrsim$80\% completeness, is well suited to PLATO (but also TESS), whose telemetry allows us to select many more targets with respect to the nominal requirement of P1 stars, therefore compensating for the fraction lost due to contaminants. UCAC4-RPM proved to be helpful as a starting point to select the (provisional) coordinates of the long-duration pointing fields, which are needed at this stage to tune the observational strategy, to run engineering tests and to plan an optimal follow-up strategy for the object of interest to be delivered by PLATO. Also ongoing and forthcoming groud-based survey for exoplanet search may benefit by UCAC4-RPM catalog. It is possible to take future steps to improve UCAC4-RPM. The most obvious one is the inclusion of newer relases of APASS, to rely on a more accurate, complete and homogeneous source of $B$ and $V$ magnitudes. APASS~DR8 and DR9 are already available, but both suffer from photometric inhomogeneity due to the inclusion of more recent data and a filter change; it is expected that APASS~DR10, still to be released, will solve most of these problems. Additionally the training set used in the present work (RAVE) could be improved by including other wide-field spectroscopic surveys, such as SDSS/SEGUE \citep{yanny2009} and especially LAMOST/LEGUE \citep{deng2012}, which will enable us to calibrate our RPM selections on both hemispheres (RAVE is limited to $b\leq0$) and to fainter magnitudes, i.e., including more dwarfs of K and M spectral types. We performed a preliminary cross-check by mapping the surface gravity listed on the first public release of LAMOST (DR1) on the full UCAC4-RPM sample. As expected, the LAMOST parameters confirm the accuracy of our previous dwarf vs.~giant separation on the RPM diagram based on RAVE (Fig.~\ref{rpmcheck}). Once the first LAMOST complete release will be made available to the community, we expect to increase the size of our training set by a factor of ten. The Gaia final catalog, to be released no earlier than 2024, is expected to make UCAC4-RPM obsolete on most of the sky, thanks to its accurate spectrophotometric measurements (BP/RP instrument) and exquisitely precise geometrical parallaxes (ASTRO instrument). However, the algorithms on which UCAC4-RPM is based could help in exploiting the Gaia data to estimate luminosity classes at $V>11$ much earlier than 2024, starting from the second intermediate release (DR2; end of 2017), when just BP/RP integrated magnitudes and proper motions will be available, but not distances or surface gravities. A modified version of the algorithm presented in Section~\ref{ucac4-rpm} can be easily adapted and calibrated through LAMOST+RAVE to work in the ($G_\mathrm{BP}-G_\mathrm{RP}$, RPM($G$)) plane. | 16 | 9 | 1609.03037 |
1609 | 1609.06048_arXiv.txt | Primordial nucleosynthesis is one of the three historical strong evidences for the hot big bang model. Its last free parameter, the baryon--to--photon ratio of the Universe, is now deduced from observations of the anisotropies of the cosmic microwave background radiation (CMB), with a precision better than one percent \cite{Planck15}. There is a good agreement between the primordial abundances of \qua\ and \deu, deduced from observations, and from primordial nucleosynthesis calculations. However, the \sep\ calculated abundance is significantly higher than the one deduced from spectroscopic observations. Solutions to this problem that have been considered include stellar surface depletion of lithium, nuclear destruction during BBN or solutions beyond the standard model (see \cite{Fie11} for a review). Experiments have now excluded a conventional nuclear physics solution (see e.g. \cite{Ham13} and references therein), even though a few uncertain reaction rates could marginally affect Li/H predictions. This lithium problem has recently worsened. Most non--conventional solutions lead to an increase of deuterium production. However, recent deuterium observations have drastically reduced the uncertainty on primordial D/H abundance \cite{Coo14}, excluding such increase. With a precision of 1.6\% on the observed D/H value \cite{Coo14}, comparison with BBN predictions requires that the uncertainties on thermonuclear reaction rates governing deuterium destruction be reduced to a similar level. | As conclusions, we list below our comments regarding frequently asked questions concerning big bang nucleosynthesis. \begin{itemize} \item {\em There is no nuclear solution to the lithium problem.} Extensive sensitivity studies \cite{Coc12b} have not identified reactions, beyond those already known, that could have a strong impact on lithium nucleosynthesis. Unknown resonances that could sufficiently increase the cross sections of reactions that destroy $^7$Be were not found experimentally (see e.g. \cite{Ham13} and references therein) and in any case would have too low strengths \cite{Bro12} because of the Coulomb barrier. \item However, {\em without solving the lithium problem}, uncertainties affecting a few reaction rates, like \hag, may still affect the lithium production. The role of the $^7$Be(n,$\alpha)^4$He reaction is presently assumed to be negligible with respect to $^7$Be(n,p)$^7$Li. However, depending on the results of ongoing experiments (these proceedings), it could reduce the lithium production by a few percents. The up to now overlooked $^7$Be(n,p$\gamma$)$^7$Li channel could also have a similar effect \cite{Hay16}. \item {\em The effect of electron screening or modification of decay lifetime is negligible.} For reactions of interest to BBN, screening affects the laboratory cross sections at too low energies [e.g. $\lesssim$ 20~keV for \ddp\ or $^3$He(d,p)$^4$He] to affect measurement at BBN energies [$\approx$100~keV], on the one hand. On the other hand, the effect screening during BBN is completely negligible \cite{Wan11,Fam16}. It is well known that the lifetime of $^7$Be that decays by electron capture is modified in a plasma \cite{Sim13}. However, because of the Boltzmann suppression factor, at $T<$0.5 GK, the electron density is too low to provide the required reduction factor of $\approx$3000 on its 53 days half-life. \item Many exotic solutions to the lithium problem have been investigated (e.g. \cite{Yam14}), but most rely on extra neutron sources that {\em overproduce deuterium to levels now excluded by observations} \cite{Coo14,Coo16}. Few solutions beyond the Standard Model that do not suffer from this drawback are left, e.g. \cite{Gou16}. \item {\em Stellar physics solutions requires a uniform reduction of surface lithium over a wide range of effective temperature and metallicity.} With some fine--tuning, this could be achieved by the combined effects of atomic diffusion and turbulence in the outer layers of these stars \cite{Ric05}, or by lithium destruction, followed by a self--regulated re-enrichment of lithium by late time accretion \cite{Fu16}. \item {\em There is no \six\ problem anymore.} A few years ago, observations \cite{Asp06} of \six\ in a few metal poor stars had suggested the presence of a plateau, at typically \six/H $\approx10^{-11}$, orders of magnitude higher than the BBN predictions of \six/H $\approx1.3\times10^{-14}$ \cite{Ham10}. The uncertainties on the D$(\alpha,\gamma)^6$Li cross section have been experimentally constraint by a LUNA measurement \cite{And14} and by theory \cite{Muk16} confirming the BBN value. However, later, the observational \six\ plateau has been questioned due to line asymmetries which were neglected in previous abundance analyses. Hence, there is no remaining evidence for a plateau at very low metallicity \cite{Lin13} that can be used to derive a primordial \six\ abundance. \item With the high precision on D/H observations, the \dpg, \ddn\ and \ddp\ rates need to be known at the percent level! This demands accurate measurement at BBN energies where data are scarce (see Fig.~\ref{f:dpg}), to be compared with theories. The theoretical work of Arai \etal\ \cite{Ara11} was focused on low energies and does not correctly reproduce the \ddn\ and \ddp\ experimental data above $\approx$600~keV. It is highly desirable that these calculations be extended up to $\approx$2~MeV, to cover the range of experimental data and BBN energies. \end{itemize} | 16 | 9 | 1609.06048 |
|
1609 | 1609.06879_arXiv.txt | {In giant molecular clouds (GMCs), the fractional ionisation is low enough that the neutral and charged particles are weakly coupled. A consequence of this is that the magnetic flux redistributes within the cloud, allowing an initially magnetically supported region to collapse.} {We aim to elucidate the effects of ambipolar diffusion on the evolution of infinitely long filaments and the effect of decaying turbulence on that evolution.} {First, in ideal magnetohydrodynamics (MHD), a two-dimensional cylinder of an isothermal magnetised plasma with initially uniform density was allowed to evolve to an equilibrium state. Then, the response of the filament to ambipolar diffusion was followed using an adaptive mesh refinement multifluid MHD code. Various ambipolar resistivities were chosen to reflect different ratios of Jeans length to ambipolar diffusion length scale. To study the effect of turbulence on the ambipolar diffusion rate, we perturbed the equilibrium filament with a turbulent velocity field quantified by a rms sonic Mach number, $M_{\rm rms}$, of 10, 3 or 1.} {We numerically reproduce the density profiles for filaments that are in magnetohydrostatic and pressure equilibrium with their surroundings obtained in a published model and show that these equilibria are dynamically stable. If the effect of ambipolar diffusion is considered, these filaments lose magnetic support initiating cloud collapse. The filaments do not lose magnetic flux. Rather the magnetic flux is redistributed within the filament from the dense centre towards the diffuse envelope. The rate of the collapse is inversely proportional to the fractional ionisation and two gravitationally-driven ambipolar diffusion regimes for the collapse are observed as predicted in a published model. For high values of the ionisation coefficient, that is $X \geq 10^{-7}$, the gas is strongly coupled to the magnetic field and the Jeans length is larger than the ambipolar diffusion length scale. Then the collapse is governed by magnetically-regulated ambipolar diffusion. The gas collapses at velocities much lower than the sound speed. For $X \lesssim 10^{-8}$, the gas is weakly coupled to the magnetic field and the magnetic support is removed by gravitationally-dominated ambipolar diffusion. Here, neutrals and ions only collide sporadically, that is the ambipolar diffusion length scale is larger than the Jeans length, and the gas can attain high collapse velocities. When decaying turbulence is included, additional support is provided to the filament. This slows down the collapse of the filament even in the absence of a magnetic field. When a magnetic field is present, the collapse rate increases by a ratio smaller than for the non-magnetic case. This is because of a speed-up of the ambipolar diffusion due to larger magnetic field gradients generated by the turbulence and because the ambipolar diffusion aids the dissipation of turbulence below the ambipolar diffusion length scale. The highest increase in the rate is observed for the lowest ionisation coefficient and the highest turbulent intensity.} | Giant molecular clouds (GMCs) contain regions of enhanced density where star formation occurs. These regions often take the form of structures such as clumps and filaments \citep[e.g.][]{1994ApJ...425..641S,2003ApJS..149..343E}. These structures may be thermally supported, or in cases where the mass of the object exceeds the Jeans mass, a magnetic field can provide support against gravitational collapse, provided that it is sufficiently strong. Only if such structures are able to fragment or collapse into sufficiently dense cores, can protostellar objects form. Various mechanisms for initiating this collapse have been suggested, including collisions between clouds \citep[e.g.][]{2014ApJ...792...63T}, shock-cloud interactions \citep[e.g.][]{2006MNRAS.365...37B, 2013MNRAS.433.1258V}, and perturbation of cores by waves \citep[e.g.][]{1995ApJ...440..686M, 2007MNRAS.376..779V}. \citet{1956MNRAS.116..503M} suggested that ambipolar diffusion due to the relative motion between the neutrals and ions, can also fragment molecular clouds into dense cores. Detailed calculations of \citet{1976ApJ...207..141M, 1979ApJ...228..475M} subsequently showed that ambipolar diffusion does, indeed, leads to the self-initiated collapse of dense central regions while the cloud envelope remains magnetically supported as magnetic flux is redistributed. Additional numerical simulations of magnetically sub-critical self-gravitating sheets or layers further confirm the ambipolar-diffusion regulated fragmentation process \citep[e.g.][]{2011ApJ...728..123K,2014ApJ...794..127K}. In this paper we consider the effect of ambipolar diffusion and velocity perturbations on magnetically sub-critical filaments. In these models the magnetic field is perpendicular to the filament axis. Filamentary clouds threaded by magnetic fields (both parallel and perpendicular to their axis) are expected to form due to gravitational and thermal instabilities within thin dense layers \citep[e.g.][]{2007MNRAS.380..499K,2011MNRAS.414.2511V,2014ApJ...789...37V}. Many observed filaments in GMCs exhibit such a configuration \citep{2013MNRAS.436.3707L}, but as yet little theoretical study of such structures has been performed. In Sect.~\ref{Model} we describe the numerical code and our initial conditions based on the analytic work of \citet{2014ApJ...785...24T}. Then, in Sect.~\ref{ambipolar}, we investigate the effect of ambipolar diffusion on the evolution of the filaments. We also examine the interaction of velocity perturbations with the diffusion process in Sect.~\ref{sect:turbulent}. Finally, in Sect.~\ref{summary}, we discuss and summarise our results. | \label{summary} In this paper we have investigated the effect of ambipolar diffusion and decaying turbulence on infinitely long, isothermal, magnetically sub-critical filaments in two dimensions. Magnetohydrostatic equilibrium filaments in pressure-equilibrium with the external medium are generated numerically in ideal MHD as initial conditions. These equilibria reproduce the analytic profiles of \citet{2014ApJ...785...24T} and, by perturbing the equilibria with decaying velocity perturbations, we find that these equilibrium filaments are dynamically stable. By using a multifluid AMR MHD code, we then follow the response of the equilibrium filament to ambipolar diffusion. Due to the gradients of the magnetic field, ambipolar diffusion initiates the filament's contraction. For thermally sub-critical filaments this contraction is halted when a new equilibrium is reached. Magnetic support is lost, with flux loss rates increasing inversely proportional to the ionisation coefficient $X$, but thermal pressure gradients are enough to balance gravitational forces. The new equilibrium is the hydrostatic profile described by \citet{1964ApJ...140.1056O}. For thermally super-critical filaments the filament contains enough mass to overcome thermal pressure forces and to collapse gravitationally. The collapse rate depends on the flux loss rate and is, as for the sub-critical filaments, thus inversely proportional to the ionisation coefficient $X$. It is important to realise that ambipolar-diffusion regulated collapse solely depends on $X$ and no other variable such as for example density, magnetic field strength or external pressure. Two models with completely different properties, that is C3b and D2, show the same collapse times for the various values of $X$ when normalised to the collapse time for instantaneous magnetic flux removal. Two gravitationally-driven ambipolar diffusion regimes are observed: a magnetically-regulated one for $X \geq 10^{-7}$ and a gravitationally-dominated one for $X \lesssim 10^{-8}$ in agreement with \citet{1991ApJ...371..296M}. The former arises because the collision time between neutrals and ions is much shorter than the free-fall time (or $\lambda_{\rm AD} < \lambda_{\rm ff}$). Then the neutrals are strongly coupled to the magnetic field and the ambipolar diffusion is regulated by the magnetic field. The collapse is quasi-static with neutral velocities much smaller than the sound speed. In the latter regime, the ion-neutral collision time becomes comparable or longer than the free-fall time and $\lambda_{\rm AD} \geq \lambda_{\rm ff}$. The neutrals are then weakly coupled and high neutral velocites are attained. Ambipolar-diffusion regulated collapse is a slow process compared to free-fall collapse, especially in the magnetically-regulated regime. Numerical simulations of non-gravitating clouds show that turbulence enhances ambipolar diffusion by a factor of a few \citep{2004ApJ...603..165H, 2012ApJ...744...73L}. When the equilibrium filament is perturbed by adding a decaying turbulent velocity field, we find that the ambipolar-diffusion collapse times decrease when compared to the collapse time for instantaneous magnetic flux loss. The actual time scales increase as the turbulent motions provide additional support to the filament. The effect of the turbulence on the ambipolar diffusion is to speed up the diffusion rate as larger magnetic field gradients are generated, while ambipolar diffusion dissipates the turbulence below the ambipolar diffusion length scale. Because we only study the effect of decaying turbulence, we cannot disentangle the combined effect of turbulence and ambipolar diffusion, but the largest effect is observed for the lowest ionisation coefficient and the highest turbulent intensity. Other effects potentially enhance the collapse rate of the magnetised filament further. In addition to ambipolar diffusion, turbulent magnetic reconnection can be an efficient diffusion process for the magnetic field \citep[e.g.][]{1999ApJ...517..700L,2010ApJ...714..442S}. Also, our models are restricted to 2D. \citet{1991ApJ...373..169M} has shown that geometry, that is the dimensionality of the problem, plays an important role in the fragmentation process due to ambipolar diffusion. In a subsequent paper we will extend these filaments to three dimensions. | 16 | 9 | 1609.06879 |
1609 | 1609.04798_arXiv.txt | We explore whether close-in super-Earths were formed as rocky bodies that failed to grow fast enough to become the cores of gas giants before the natal protostellar disk dispersed. We model the failed cores' inward orbital migration in the low-mass or type I regime, to stopping points at distances where the tidal interaction with the protostellar disk applies zero net torque. The three kinds of migration traps considered are those due to the dead zone's outer edge, the ice line, and the transition from accretion to starlight as the disk's main heat source. As the disk disperses, the traps move toward final positions near or just outside 1~au. Planets at this location exceeding about 3~M$_\oplus$ open a gap, decouple from their host trap, and migrate inward in the high-mass or type II regime to reach the vicinity of the star. We synthesize the population of planets formed in this scenario, finding that some fraction of the observed super-Earths can be failed cores. Most super-Earths formed this way have more than 4~M$_\oplus$, so their orbits when the disk disperses are governed by type II migration. These planets have solid cores surrounded by gaseous envelopes. Their subsequent photoevaporative mass loss is most effective for masses originally below about 6 M$_\oplus$. The failed core scenario suggests a division of the observed super-Earth mass-radius diagram into five zones according to the inferred formation history. | \label{intro} The discovery of super-Earths has significantly enriched fundamental achievements of exoplanet observations. This was initially led by the radial velocity observations \citep[e.g.,][]{mq95,mb96,us07,mml11,fms14} and has subsequently been accelerated by the transit detections through {\it Kepler} mission \citep[e.g.,][]{bkb11,brb12,bbm14}. Super-Earths that have $1-20 M_{\oplus}$\footnote{Exoplanets that have $10-20 M_{\oplus}$ are sometimes called as (hot) Neptunes, rather than super-Earths.} by a working definition are very peculiar in the sense that there is no such analogue in the solar system. In fact, the mass range lies down between rocky planets like the Earth and the lower end of icy planets like the Neptune. The coupling of these two observations (the radial velocity and the transit) reveals that the mean density of super-Earths is quite diverse \citep[e.g.,][, see Figure \ref{fig1}]{wm14,mwp14,jfr16,gcd16}. These discoveries in turn stimulate active investigations of super-Earths, since their formation mechanisms are uncertain. \begin{figure}% \begin{center} \includegraphics[width=8cm]{fig1.eps} \caption{The mass-radius diagram for observed exoplanets. Among 65 exoplanets from the list of \citet{wm14}, 33 exoplanets smaller than 4 $R_{\oplus}$ that have 2$-\sigma$ mass determinations are extracted, following \citet{mwp14}. In the data, planetary masses are obtained by radial velocity observations as well as transit timing variation (TTV) while planet radii by transit observations such as the {\it Kepler} mission. Coupling of these two observations implies that there may be a transition in planet radius ($R_{tran} \simeq 1.5-1.6 R_{\oplus}$, also see the blue horizontal line) above which most planets are very likely to be made of gaseous envelopes atop rocky cores, below which most planets are very likely to be purely rocky \citep{r15}. The dividing line between "Rocky" (the yellow regime) and "Core+Envelope" (the green regime) is given by an empirical linear density-radius relation \citep[][]{wm14,mwp14}.} \label{fig1} \end{center} \end{figure} Currently, three kinds of formation scenarios are predominantly examined (see Table \ref{table1}). The first one is the so-called "in situ" formation scenario \citep[e.g.,][]{ml09,cl13}. In this picture, it is proposed that the observed super-Earths should have formed there through collisional growth of planetesimals and solids. While the orbital distribution of the observed super-Earths allows construction of a minimum-mass extrasolar nebula (MMEN) model \citep{cl13}, which can be compatible with the famous minimum-mass solar nebula (MMSN) model \citep[e.g.,][]{h81}, the origin of high solid densities in the inner part of disks remains to be explored \citep[e.g.,][]{s14}. \begin{table*} \begin{minipage}{17cm} \begin{center} \caption{Currently proposed scenarios to form Super-Earths} \label{table1} \begin{tabular}{lccc} \hline & in-situ & embryo assembly & failed core (this study) \\ \hline Main mode of solid transport & radial drift of dust particles & type I migration ($M_p \la 1 M_{\oplus}$) & type II migration ($M_p \ga 2-3 M_{\oplus}$) \\ Core formation & at $r < 1$ au in gas-rich disks & at $r < 1$ au in gas-poor disks & at $r > 1$ au in gas-rich disks \\ \hline \end{tabular} \end{center} \end{minipage} \end{table*} The other two kinds of scenarios require inward planetary migration that arises from tidal, resonant interactions between protoplanets and their gas disks \citep[e.g.,][]{kn12}. Planetary migration can be classified into two modes, depending on the mass of migrators. In general, type I migration is effective for low-mass planets such as terrestrial planets and cores of gas giants, while type II migration becomes in action for massive planets such as Jovian planets that can open up a gap in their gas disks. One of the remaining two scenarios is a scaled down version of gas giant formation \citep[e.g.,][]{abb06,mab09,rbl11,hp12,hp14}. In this picture, planets form via the core accretion scenario, wherein cores of planets form initially, followed by the subsequent gas accretion onto the cores \citep[e.g.,][]{p96,il04i}. The main difference with gas giant formation is that planetary cores that eventually become super-Earths cannot accrete the disk gas efficiently. This inefficient gas accretion can be explained by the slow growth rate of planetary cores which should occur at the later stage of disk evolution. At that time, the formation of planetary cores will take a longer time and the core mass itself would become less massive, both of which lengthen the timescale of gas accretion onto the core and hence prevent the core from growing up to gas giants even in gas disks.\footnote{Most of the currently existing studies implicitly assume that failed cores are dynamically isolated from others. This may be possible due to their masses that are $\ga 5 M_{\oplus}$ as shown in this study. Such planets may lead to dynamical clearing of the surrounding regions and hence suppress the subsequent giant impact stage (also see Section \ref{disc_2}).} This formation mechanism thus can be referred to as a "failed core" scenario. The second kind of scenarios with migration is an extension of rocky planet formation \citep[e.g.,][]{tp07,mn10,il10,hm12}. In this scenario, planetary embryos that are massive enough to undergo rapid, inward type I migration but not to initiate efficient gas accretion, migrate toward the inner edge of disks and accumulate there. Since gas disks are still present at that time, these embryos can establish a dynamically stable configuration. This arises from efficient damping of the embryos' eccentricities by the disk gas, which eventually suppresses orbital crossing with each other. Once the gas disks are gone, then these embryos collide together to form super-Earths. An advantage of this scenario is that massive protoplanets that potentially accrete gas in the disks, can form only after gas disks disperse significantly. Therefore, there is no need of considering possibilities that protoplanets will grow rapidly via runaway gas accretion in gas disks. In this paper, this scenario is called as "embryo assembly", since planets obtain most of their masses through collisions with other massive embryos and/or protoplanets. As discussed above, one may classify the currently proposed scenarios for super-Earths, depending on how solid materials are transported from the outer part to the inner part of disks, and on when and where planetary cores are formed (see Table \ref{table1}). It, however, should be noted that such a classification and each scenario are still under development. In this paper, we will explore the the failed core scenario. As already mentioned above, the scenario essentially allows the formation of massive cores in gas disks, which can potentially trigger runaway gas accretion there and hence may end up with gas giants, rather super-Earths. In addition, it is still a matter of debate how such massive cores are saved from rapid inward migration that leads to loss of the cores into the central star within the disk lifetime \citep[e.g.,][]{ward97,ttw02}. Currently, it is well known that the direction of type I migration is quite sensitive to the disk structure \citep[e.g.,][]{ward97,ttw02,pm06,bm08,pbck09,hp10c}. Nonetheless, the sensitivity causes further complexity of the migration, because its systematic effect both on planetary formation and on the resultant populations of planets becomes unclear. These considerations pose natural questions: what kind(s) of mechanisms would be needed to systematically understand the effect of planetary migration on the failed core scenario? In the end, is the failed core scenario plausible to explain (at least in part) the population of observed super-Earths? In this paper, we investigate the failed core scenario by taking account of planetary migration in detail and examine how important the scenario is to understand the observed properties of super-Earths. We will explore the effect of gas accretion in a subsequent paper. We, in particular, focus on planet traps - the stopping sites for rapid type I migration \citep{mmcf06,hp11}. Combing the traps with the core accretion growth of planets, we have so far examined a number of characteristic features of the exoplanet observations such as the mass-period relation \citep{hp12}, the orbital distribution of exoplanets \citep{hp13a}, and the planet-metallicity relation \citep{hp14}. Our previous studies suggest that planetary populations synthesized by the combination of the core accretion process with planet traps are reasonably consistent with the trends of the exoplanet observations. It is important that these results can be viewed as (partial) supportive evidence for the importance of planet traps. We here apply our model to the population of super-Earths and discuss how significant planet traps are to form them. We will demonstrate below that planet traps that can halt type I migration will play a crucial role in forming super-Earths, as pointed out by our earlier work. The movement of planet traps will terminate due to the dispersal of gas disks as the traps arrive roughly at the distance of 1 au from the central star. This suggests that when planet traps act as the birth sites of super-Earths, the cores must have dropped out of their host traps and migrated to their current location possibly by type II migration. The estimate of the gap-opening mass around $r=1$ au, therefore, provides a threshold value for trapped planets to be transported to the vicinity of the central star. Under the presence of dead zones in our model that have $\alpha \simeq 10^{-4}$ using the $\alpha$-prescription \citep{ss73}, the gap-opening mass becomes about $3M_{\oplus}$ at the disk temperature of about 200 K at 1 au. Furthermore, we perform a population synthesis analysis and show that most of super-Earths generated by our model obtain the mass of about $4-5 M_{\oplus}$ or higher, which is larger than the gap-opening mass at $r \simeq 1$ au. We thus demonstrate that switching of migration modes can provide profound insights as to origins of close-in super-Earths. The plan of this paper is as follows. In Section \ref{model}, a model used in this paper is described, wherein a recipe of planet traps and our population synthesis analysis are summarized. In Section \ref{resu}, we present our results and examine the role of planet traps on the formation of super-Earths. The results are derived from two distinct approaches: the one is to simply compute the closest orbital distance of planet traps from the central star, which is determined by disk evolution; the other is to perform population synthesis calculations. In Section \ref{disc}, we discuss other physical processes that may affect our findings, and provide some implications of our results for the mass-radius diagram of observed super-Earths. Section \ref{conc} is devoted to the conclusion of this work. | \label{conc} One of the key discoveries in exoplanet observations that are mainly done both by the radial velocity technique and by the transit method via the {\it Kepler} mission, is the rapidly growing population of close-in super-Earths. These samples are quite unique in the sense that their mass ($1 M_{\oplus} \la M_p \la 20 M_{\oplus}$) and semimajor axis ($r \la 1$ au) do not overlap with the parameter space covered by planets in the solar system (see Figure \ref{fig1}). This therefore triggers active investigations of super-Earths in which a number of formation mechanisms have currently been examined (see Table \ref{table1}). In this paper, we have considered a failed core scenario and examined planetary migration in detail. In the scenario, the formation of super-Earths is regarded as a scaled-down version of gas giant formation. As a result, the population of such planets should emerge at the late stage of disk evolution, wherein core formation takes a longer time due to a lower value of the solid density. Such slow growth of planetary cores can eventually prevent the cores from undergoing rapid gas accretion. In other words, the formation timing would become a central quantity to determine the feasibility of the failed core scenario. Also, the failed core scenario can lead to the formation of massive cores in gas disks, where rapid type I migration cannot be neglected. This provides another difficulty in examining the scenario in detail, because the direction of the migration is very sensitive to the local disk structures. Here, we have focused on planet traps - the specific location of protoplanetary disks at which rapid type I migration can be halted. We have investigated how plausible the failed core scenario is to account for the observed properties of close-in super-Earths. We have adopted a model that has been developed in our earlier work \citep[see Section \ref{model}]{hp11,hp12,hp13a}. We have first examined the movement of planet traps (see Figure \ref{fig2}). We have shown that, while planet traps can spread out over a wide range of disk radii at the early stage of disk evolution, they end up around $r \simeq 1$ au at the stage where gas disks are severely depleted. This occurs, because their movement is coupled directly with disk evolution (see equations (\ref{eq:r_dz}), (\ref{eq:r_il}), and (\ref{eq:r_ht})). In other words, their position can freeze out once disk evolution terminates due to photoevaporation of the disk gas. We have found that the final orbital distance of planet traps is beyond $r \simeq 1$ au even if a disk mass parameter varies considerably. Thus, our results indicate that when planet traps can play an important role in forming planets, planets should drop-out from their host trap to distribute within $r \simeq 1$ au. We have then computed a number of characteristic masses at planet traps. These include a threshold mass ($M_{mig,I}$) for planets to migrate fast enough to follow the moving of planet traps (see equation (\ref{eq:M_typeI})), a critical planet mass ($M_{sat}$) for saturation to become strong enough to null the effect of planet traps (see equation (\ref{eq:M_sat})), the gap-opening mass ($M_{gap}$) of planets (see equation (\ref{eq:M_gap})), and a planet mass ($M_{mig,II}$) that becomes comparable to the value of $2 \pi r^2 \Sigma_g$ (see equation (\ref{eq:M_typeII})). We have found that planet traps become effective for planets with certain masses ($M_{mig,I}< M_p < M_{gap}$) even if the effect of (partial) saturation is taken into account. This is because $M_{gap} \approx M_{sat}$. We have also confirmed that $M_{gap} < M_{mig,II}$, indicating that planets will undergo type II migration once their mass exceeds $M_{gap}$. Moreover, we have demonstrated that the value of $M_{gap}$ at the smallest orbital distance of planet traps is about $2-3 M_{\oplus}$ (see Figure \ref{fig2}). This implies that when planets can grow up to this threshold value or higher, they can drop-out from their planet traps and be transported to the vicinity of their central star via inward type II migration. We have confirmed that the value of $M_{gap}$ is not so sensitive to the disk mass parameter ($\eta_{acc}$). We have also performed a population synthesis analysis to investigate how useful the failed core scenario is to understand the properties of observed close-in super-Earths. We have initially examined the resultant PFFs of Jovian planets, and demonstrated that the coupling of planet traps with the core accretion scenario can lead to planetary populations, which are in a good agreement with the trend of exoplanet observations (see Table \ref{table4}); the PFFs of exo-Jupiter are much higher than those of hot Jupiters. On the other hand, we have found that the PFFs of the Low-mass planets are not fully consistent with the observations which reveal that close-in super-Earths are currently the most dominant. This difference probably suggests that a number of formation mechanisms would be needed to reproduce the entire population of observed super-Earths. We have then turned our attention to the population of the Low-mass planet regime (see Table \ref{table3}), and examined how these planets are formed in our model. To proceed, we have plotted a cumulative fraction of the normalized PFFs as a function of planetary mass (see Figure \ref{fig3}). We have found that the fraction rises rapidly when the planetary mass exceeds $4-5 M_{\oplus}$ (that is defined as $M_{min}^{CA}$) (see Table \ref{table5}). As discussed above, the value is characterized well by the statistical average value of $M_{gap}$ ($\braket{M_{gap}}$). Our results therefore indicate that switching of migration modes (from trapped type I to type II migration) is needed for planets that grow predominantly at planet traps to distribute in the vicinity of the central star. It is important that the value of $M_{min}^{CA}$ is independent of planet formation parameters such as the critical core mass ($M_{c,crit}$). Thus, we can conclude that planets with the mass of $M_{min}^{CA}$ or higher can be made of solid cores surrounded by gaseous atmospheres (see Figure \ref{fig4}). In addition, we have examined the effect of photoevaporative mass loss from planets (see equation (\ref{eq:f_lost})). We have found that the mass loss decreases the cumulative fraction of planets, which corresponds to the range of planetary mass ($1 M_{\oplus} \la M_p \la 20 M_{\oplus}$) (see Figure \ref{fig3}). This is a direct reflection that planets in the range undergo the partial or the complete removal of their envelope. As a result, a threshold mass of planets at which the fraction suddenly increases, goes up to $6-7 M_{\oplus}$ (that is defined by $M_{min}^{CA+PE}$) (see Table \ref{table5}). Thus, our results imply that planets with the mass of $M_{min}^{CA+PE}$ or lower may not currently have gaseous atmosphere due to photoevaporative mass loss, when the planets are formed via the failed core scenario coupled with planet traps. We have discussed physical processes that are simplified and/or not included in our present model (Section \ref{disc}). These include planetary migration, planet-planet interactions, gas accretion and the subsequent mass loss, and dust physics at ice lines. While some of them may have a considerable impact on our results, our findings may still be valuable in the sense that they may provide a conservative estimate on the resultant planetary populations. Finally, we have provided some implications that can be obtained from our results (see Figure \ref{fig4}). Since our analysis may be useful to derive the threshold value of planetary mass (see Table \ref{table5}), the coupling of our results with the mass-radius diagram may act as an exoplanet "phase" diagram in which some formation and evolution processes can be speculated based on the "location" of planets in the diagram. Overall, the failed core scenario coupled with planet traps can reproduce some fraction of observed close-in super-Earths, and may be important in the sense that it can serve as one of the crucial mechanisms to fully understand the diversity of observed properties of such planets. In a subsequent paper, planet-planet interactions should be incorporated into our model, in order to examine how our current results can be affected. | 16 | 9 | 1609.04798 |
1609 | 1609.08145_arXiv.txt | We present the light curves of the hydrogen-poor superluminous supernovae (SLSNe-I) \dam\ and \dcc, discovered by the (intermediate) Palomar Transient Factory. Both show excess emission at early times and a slowly declining light curve at late times. The early bump in \dam\ is very similar in duration ($\sim$10~days) and brightness relative to the main peak (2--3~mag fainter) compared to those observed in other SLSNe-I. In contrast, the long-duration ($>$30~days) early excess emission in \dcc, whose brightness competes with that of the main peak, appears to be of a different nature. We construct bolometric light curves for both targets, and fit a variety of light-curve models to both the early bump and main peak in an attempt to understand the nature of these explosions. Even though the slope of the late-time light-curve decline in both SLSNe is suggestively close to that expected from the radioactive decay of $^{56}$Ni and $^{56}$Co, the amount of nickel required to power the full light curves is too large considering the estimated ejecta mass. The magnetar model including an increasing escape fraction provides a reasonable description of the \dam\ observations. However, neither the basic nor the double-peaked magnetar model is capable of reproducing the \dcc\ light curve. A model combining a shock breakout in an extended envelope with late-time magnetar energy injection provides a reasonable fit to the \dcc\ observations. Finally, we find that the light curves of both \dam\ and \dcc\ can be adequately fit with the circumstellar medium (CSM) interaction model. | \label{sec:introduction} Supernovae that reach an absolute magnitude brighter than the (arbitrary) limit of $M=-21$ are labelled superluminous \citep{2012Sci...337..927G}. Even though they are very rare, several tens of them have been discovered over the past decade thanks to the ever-increasing survey speed of optical telescopes. They are observationally separated into two classes based on the detection of hydrogen in their spectra, similar to classical supernovae \citep[see][]{1997ARA&A..35..309F}: hydrogen-rich Type II SLSNe show clear Balmer features \citep[e.g.,][]{2007ApJ...659L..13O,2007ApJ...666.1116S}, while the hydrogen-poor Type I SLSNe do not \citep{2012Sci...337..927G}. The latter commonly exhibit a distinct W-shaped feature identified as \ion{O}{2} at rest-frame wavelengths 4000--4500~\AA\ \citep{2011Natur.474..487Q,2016MNRAS.458.3455M}. At late times, the spectra of these SLSNe-I evolve to appear like those of normal Type Ic SNe \citep{2010ApJ...724L..16P}, leading many authors to refer to this class as SLSN-Ic. Sometimes this Type I/II distinction is not so obvious; for example, spectra of CSS121015:004244+132827 (hereafter referred to as \css) have similarities to both Type II and I objects \citep{2013arXiv1310.1311B}, while iPTF~13ehe, classified as Type I, shows the emergence of broad \ha\ emission at late times (\citealp{2015ApJ...814..108Y}; see also \citealp{2015A&A...584L...5M,2016ApJ...828...87W}). There is some evidence that the energy source powering the hydrogen-rich SLSNe is interaction of the SN ejecta with optically-thick material at a large distance ($\sim$10$^{15}$~cm), as they typically reveal Balmer emission lines indicative of interaction with a hydrogen-rich circumstellar medium \citep[CSM; e.g.,][]{1994ApJ...420..268C,1994MNRAS.268..173C,2014ApJ...788..154O}. Because of the similarity with normal SNe of Type IIn, this class is also referred to as SLSNe-IIn. However, some SLSNe-II do not exhibit narrow emission lines, while they are of Type II as they reveal broad hydrogen features during the photospheric phase (\citealp{2016arXiv160401226I}; see also \citealp{2012ApJ...747..118M}). The energy source of the hydrogen-poor SLSNe is still under debate, with the most promising candidates being (1) additional energy input from a central engine, such as a spinning-down magnetar \citep[e.g.,][]{2010ApJ...717..245K,2010ApJ...719L.204W,2013ApJ...770..128I} or an accreting black hole \citep{2013ApJ...772...30D}; (2) interaction of the ejecta with a hydrogen-poor shell expelled by the progenitor star some time before the explosion \citep{2010arXiv1009.4353B,2011ApJ...729L...6C}; and (3) the radioactive decay of a large amount of nickel produced in the explosion, potentially due to pair-instability conditions \citep[see][]{2009Natur.462..624G} though this is still being debated \citep{2010ApJ...717L..83M,2010A&A...512A..70Y,2013Natur.502..346N}. The light curves of some hydrogen-poor SLSNe, such as the two PTF sources presented in this paper, decay very slowly at late times, with a slope similar to that expected from the decay of radioactive nickel and cobalt. These SLSNe are part of the hydrogen-poor class, but are sometimes referred to as Type R \citep[``radioactive''; see][]{2012Sci...337..927G}. The host galaxies of SLSNe are found to be irregular, compact, low-mass galaxies with high specific star formation rates \citep{2011ApJ...727...15N,2013arXiv1311.0026L,2015ApJ...804...90L,2015MNRAS.449..917L}. A comprehensive study of 32 host galaxies of all PTF-discovered SLSNe (until the end of 2012) found hydrogen-poor SLSNe to have a preference for environments in hosts with a metallicity upper bound of about half solar, while the hydrogen-rich SLSNe do not show such a preference \citep{2016arXiv160408207P}. A very similar conclusion was reached independently by \citet{2016arXiv160504925C}. In emission, the galaxies hosting SLSNe have broadly similar characteristics as the hosts of gamma-ray bursts (GRBs). \citet{2015MNRAS.449..917L} find the host-galaxy emission-line strengths of redshift $z<1$ SLSNe-I to be significantly stronger than in GRB hosts, but \citet{2016arXiv160701045J} do not confirm this result over the range $0.3<z<0.7$. In {\it absorption}, the environments of SLSNe-I appear to be significantly poorer in their neutral gas content, as traced by \ion{Mg}{1} and \ion{Mg}{2}, than those of GRBs \citep{2014ApJ...797...24V}. Focusing on the hydrogen-poor class, several of these have shown evidence for an early-time light-curve ``bump'' or excess emission before the onset of the main peak. Examples are \oz\ \citep{2012A&A...541A.129L}, \bdq\ \citep{2015ApJ...807L..18N}, and \des\ \citep{2015arXiv151206043S}. In fact, \citet{2016MNRAS.457L..79N} suggest that early bumps such as the ones above may be ubiquitous in hydrogen-poor SLSNe. Early bumps have also been observed in normal stripped-envelope SNe, and recently in a normal SN Ic from a massive progenitor \citep{2016A&A...592A..89T}. To date, such early excess emission has not been reported for any hydrogen-rich SLSN-II. This early excess emission is of particular interest, as it may provide a clue regarding what is powering these explosions. In the case of \oz, \citet{2012A&A...541A.129L} propose that the precursor bolometric plateau might be related to a recombination wave in a H-poor CSM. The study by \citet{2012ApJ...756L..22M} has shown that a dip in the light curve is naturally expected when shock breakout occurs within a dense CSM. \citet{2015ApJ...807L..18N} propose that the initial peak in \bdq\ may arise from the post-shock cooling of extended stellar material, while reheating by a central engine is driving the main peak. The high kinetic energy inferred from fitting the \citet{2011ApJ...728...63R} model to the initial peak ($E_{\rm k}\sim 2\times10^{52}$~erg) of \bdq\ may favor a black hole accretion engine \citep{2013ApJ...772...30D} rather than a magnetar. The early-time excess emission in the case of \des\ shows rapid cooling from 22,000~K to 8,000~K over the course of 15 rest-frame days. The authors find that a shock-cooling model of CSM at a distance of $\sim$400~\Rsun, followed by a magnetar causing the main peak of the light curve can adequately explain the entire light curve. \begin{deluxetable*}{lccccccccc} \tablecaption{Log of Spectroscopic Observations of \dcc\label{tab:13dcc_logspec}} \tablehead{ \colhead{UTC Date} & \colhead{Telescope} & \colhead{Instrument} & \colhead{Exp.~Time} & \colhead{Grating/Grism/Filter} & \colhead{Slit~Width} & \colhead{$\lambda$ Coverage} & \colhead{Res.\tablenotemark{a}} & \colhead{I.Q.\tablenotemark{b}} & \colhead{Airm.} \\ & & & (min.) & & \arcsec & (\AA) & (\AA) & \arcsec & } \startdata 2013 Nov. 26 & P200 & DBSP & 20 & 600/4000, 316/7500 & 1.5 & 3400--10,400 & 9.3 & 2.9 & 1.3 \\ 2013 Dec. 3 & Keck~I & LRIS & 21 & 400/3400, 400/8500 & 1.0 & 3200--10,240 & 6.0 & 1.6 & 1.4 \\ 2013 Dec. 4 & Keck~I & LRIS & 10 & 600/4000, 400/8500 & 1.0 & 3140--10,240 & 5.8 & 1.2 & 1.4 \\ 2013 Dec. 31 & Magellan Baade & IMACS & 25 & Gra-300-4.3 & 0.9 & 3700--9,700 & 6.1 & 0.9 & 1.3 \\ 2014 Jan. 6 & P200 & DBSP & 60 & 600/4000, 316/7500 & 1.5 & 3300--10,400 & 8.5 & 1.8 & 1.6 \enddata \tablenotetext{a}{The resolution of the spectra was determined from the width of the [O~I] $\lambda$5577 night-sky line.} \tablenotetext{b}{The image quality, or effective seeing, was measured directly from the width of the object's spatial profile around 6000~\AA.} \end{deluxetable*} In this paper, we present two SLSNe-I discovered by the (intermediate) Palomar Transient Factory \citep{2009PASP..121.1334R,2009PASP..121.1395L} that also show evidence for early excess emission: \dam\ and \dcc. The modest bump in \dam\ is very similar in duration and brightness relative to the main peak to the cases discussed above. However, the long-duration early excess emission in \dcc, whose brightness competes with that of the main peak, appears to be of a different nature. This paper is organized as follows. In Sec.~\ref{sec:observations} we present the photometric observations that we obtained for \dam\ and \dcc, as well as the spectroscopic sequence of \dcc. We construct the bolometric luminosity evolutions in Sec.~\ref{sec:photometry}, which we confront with models in Sec.~\ref{sec:modelling}. The models \citep[see][]{2012ApJ...746..121C,2013ApJ...773...76C,2015ApJ...808L..51P} assume different energy sources (radioactive decay, magnetar heating, and CSM interaction) and predict the ensuing light curve based on a number of parameters; we infer estimates of the best-fit values by fitting the semi-analytical models to the bolometric light curves. We discuss our results and briefly conclude in Sec.~\ref{sec:discussion}. Unless noted otherwise, the uncertainties listed in this paper are at the 1$\sigma$ confidence level. We adopt the cosmological parameters as derived by the Planck collaboration in 2015 \citep[H$_0=68$~km~s$^{-1}$~Mpc$^{-1}$, $\Omega_{\rm m}=0.31$, $\Omega_{\Lambda}=0.69$;][]{2015arXiv150201589P}. | \label{sec:discussion} Figure~\ref{fig:Mg} shows a comparison of the absolute $g$-band light curves of \dcc, \bdq, \oz, \des, and \dam. The \dcc\ data were derived by converting the observed $r$-band measurements to rest-frame $g$, where we computed the K-corrections adopting the temperature evolution of \dam\ (see Fig.~\ref{fig:bol13dcc}). Since the effective wavelengths of the observed $r$-band and rest-frame $g$-band filters match well at the redshift of \dcc\ ($z=0.431$), these K-corrections are very close to $-$2.5~log~(1+$z$) and do not depend very much on the assumed temperature evolution. For \dam\ ($z=0.107$), we instead computed a continuous rest-frame $g$-band light curve by adopting the temperature and bolometric luminosity evolution as derived in Fig.~\ref{fig:bol12dam}. For each SLSN shown in Fig.~\ref{fig:Mg}, except for \dcc, we shifted the \dam\ light curve, in both time and magnitude, to match the data points. This is simply to show how the observed early excess emission of each SLSN compares in duration and magnitude with that of the others. For both \oz\ and \des\ the match with \dam\ at early epochs is quite good, whereas at late times the \des\ light curve is dropping much more rapidly than that of \dam. For \bdq\ the early excess emission is longer and brighter than that of \dam, but at late times the two match surprisingly well. The obvious outlier is the light curve of \dcc: its early emission is of a much longer duration and a few magnitudes brighter than those of the others. \citet{2016ApJ...821...36K} suggest that the early excess emission observed for \bdq\ could be caused by a magnetar central engine whose energy input drives a shock through the pre-exploded SN ejecta, resulting in a burst of shock-breakout emission several days after the explosion. The radiation is expected to be released in the optical/UV wavelengths and to have a duration of several days, with the emission being dimmer than the main light-curve peak resulting from continued magnetar heating. As can be seen in Fig.~7 of \citet{2016ApJ...821...36K}, for standard magnetar and ejecta parameters, the shock-breakout emission produces only a kink in the overall light curve. In principle, this model is capable of explaining the early excess emission in \dam. However, the early emission that we observe for \dcc\ is too bright and too extended to be accommodated in this model. Even when pushing the ejecta mass and kinetic energy to large values \citep[see Fig.~8 of][]{2016ApJ...821...36K}, this model is not able to reproduce the \dcc\ observations. Based on the model fits presented in Sec.~\ref{sec:modelling}, the CSM interaction model seems to provide better fits to the data. Moreover, the magnetar model is unable by itself to fit the precursor bumps seen in SLSNe but requires the addition of the Piro model (or similar) and the presence of extended material at large radii to explain the first peak. The combination of Piro and magnetar provides a reasonable description of the \dcc\ light curve, even though it does not fit the data well around peak and appears to overshoot the data beyond $+40$~days. Moreover, the explosion energy for this model reaches the upper end of the range that we allowed: $E_{\rm SN}=10^{53}$~erg. We find that the CSM interaction model, which in fact contains one free parameter less than the number used in the combined Piro and magnetar model, can provide a good fit to the \dcc\ light curve, even without the addition of the Piro model (see Table~\ref{tab:csm} and Fig.~\ref{fig:model13dcc}). The first minimum in the light curve cannot be accurately reproduced (see the left panel of Fig.~\ref{fig:model13dcc}), but a time-dependent opacity (we have assumed a constant opacity throughout this paper) could be invoked to explain the early light-curve shape. Also, this \dcc\ light-curve dip is reminiscent of the picture proposed by \citet{2012ApJ...756L..22M}, where a drop in the light curve arises naturally. These authors suggest that such a dip is a solid prediction from the strong interaction scenario regardless of the power source for the early emission. The ability of the CSM interaction model to reproduce the light curves of many SLSNe can be explained by the fact that it includes many free parameters. In addition, the CSM model by \citet{2012ApJ...746..121C,2013ApJ...773...76C} includes a number of simplifying assumptions of which the most important are those of a central power source and a stationary photosphere, allowing for the use of Arnett-style diffusion by using analytical equations. However, it is not clear if these assumptions hold in real CSM interaction. Chatzopoulos et al. verify their model against a more sophisticated hydrodynamical model developed for the H-rich SN~2006gy \citep{2013MNRAS.428.1020M} and obtain a reassuringly similar light curve for similar parameters. Extending the use of the model to H-poor SLSNe, such as \dam\ and \dcc, is not trivial, however, owing to the different treatment of the opacity when hydrogen is absent. Nevertheless, models that successfully reproduce the light curves of H-poor SLSNe by using CSM interaction have now been reproduced in radiation hydrodynamics simulations by \citet{2015arXiv151000834S}, although these authors do report discrepancies in some parameters of their models in comparison to those obtained with the Chatzopoulos models for the same SLSNe. The CSM interaction semi-analytical models remain a valuable tool, and it is clear that H-poor CSM interaction can reproduce the light curves of SLSNe, more naturally explaining features such as premaximum bumps in a self-consistent way. However, these caveats show that caution should be used in interpreting the best-fit model parameters based on $\chi^2$ minimization of the semi-analytical CSM interaction models. On a final note, an argument often invoked against the CSM interaction model for hydrogen-poor SLSNe is the lack of narrow emission lines, as observed for Type IIn SNe. However, to date, the spectroscopic signature(s) of interaction with a H-deficient dense CSM has not been investigated through spectral synthesis modelling, due to its complexity. \begin{deluxetable}{crccc} \tablecaption{Log of observations of \dam.\label{tab:12dam_logphotometry}} \tablehead{ \colhead{MJD} & \colhead{Phase\a} & \colhead{Telescope} & \colhead{Filter} & \colhead{Magnitude\b} \\ (days) & (days) & & & AB } \startdata 56009.34 &$-$78.92 & P48 & R & $>$ 22.26 \\ 56013.35 &$-$75.30 & P48 & R & $>$ 22.10 \\ 56017.29 &$-$71.73 & P48 & R & $>$ 22.13 \\ 56020.41 &$-$68.91 & P48 & R & $>$ 21.94 \\ 56023.28 &$-$66.33 & P48 & R & $>$ 20.69 \\ 56025.33 &$-$64.47 & P48 & R & 20.11 $\pm$ 0.17 \\ 56027.26 &$-$62.73 & P48 & R & 20.02 $\pm$ 0.12 \\ 56027.29 &$-$62.70 & P48 & R & 20.11 $\pm$ 0.11 \\ 56027.32 &$-$62.68 & P48 & R & 20.04 $\pm$ 0.10 \\ 56033.49 &$-$57.10 & P48 & R & 19.89 $\pm$ 0.08 \\ 56034.16 &$-$56.49 & P48 & R & 19.96 $\pm$ 0.08 \\ 56034.19 &$-$56.46 & P48 & R & 19.96 $\pm$ 0.08 \\ 56034.26 &$-$56.41 & P48 & R & 19.83 $\pm$ 0.07 \\ 56036.17 &$-$54.68 & P48 & R & 19.77 $\pm$ 0.08 \\ 56036.21 &$-$54.64 & P48 & R & 19.68 $\pm$ 0.07 \\ 56036.37 &$-$54.50 & P48 & R & 19.70 $\pm$ 0.07 \\ 56038.35 &$-$52.71 & P48 & R & 19.37 $\pm$ 0.04 \\ 56038.38 &$-$52.69 & P48 & R & 19.39 $\pm$ 0.04 \\ 56040.43 &$-$50.83 & P48 & R & 19.28 $\pm$ 0.04 \\ 56040.46 &$-$50.81 & P48 & R & 19.22 $\pm$ 0.04 \\ 56046.40 &$-$45.44 & P48 & R & 18.55 $\pm$ 0.03 \\ 56046.43 &$-$45.41 & P48 & R & 18.57 $\pm$ 0.03 \\ 56048.43 &$-$43.60 & P48 & R & 18.34 $\pm$ 0.03 \\ 56048.47 &$-$43.57 & P48 & R & 18.43 $\pm$ 0.02 \\ 56048.50 &$-$43.54 & P48 & R & 18.36 $\pm$ 0.03 \\ 56058.47 &$-$34.53 & P48 & R & 17.76 $\pm$ 0.02 \\ 56059.28 &$-$33.81 & P48 & R & 17.69 $\pm$ 0.01 \\ 56061.33 &$-$31.95 & P48 & R & 17.64 $\pm$ 0.01 \\ 56063.29 &$-$30.18 & P48 & R & 17.53 $\pm$ 0.01 \\ 56063.32 &$-$30.15 & P48 & R & 17.51 $\pm$ 0.01 \\ 56063.35 &$-$30.12 & P48 & R & 17.57 $\pm$ 0.01 \\ 56065.44 &$-$28.24 & P48 & R & 17.46 $\pm$ 0.02 \\ 56066.38 &$-$27.39 & P48 & R & 17.44 $\pm$ 0.01 \\ 56066.41 &$-$27.36 & P48 & R & 17.46 $\pm$ 0.01 \\ 56066.45 &$-$27.33 & P48 & R & 17.44 $\pm$ 0.01 \\ 56068.48 &$-$25.50 & P48 & R & 17.38 $\pm$ 0.01 \\ 56070.40 &$-$23.76 & P48 & R & 17.28 $\pm$ 0.01 \\ 56070.44 &$-$23.73 & P48 & R & 17.28 $\pm$ 0.01 \\ 56070.47 &$-$23.70 & P48 & R & 17.29 $\pm$ 0.01 \\ 56076.37 &$-$18.36 & P48 & R & 17.12 $\pm$ 0.01 \\ 56076.40 &$-$18.33 & P48 & R & 17.16 $\pm$ 0.01 \\ 56362.48 & 240.09 & P48 & R & 20.70 $\pm$ 0.13 \\ 56362.51 & 240.12 & P48 & R & 20.42 $\pm$ 0.10 \\ 56362.54 & 240.14 & P48 & R & 20.59 $\pm$ 0.13 \\ 56366.49 & 243.71 & P48 & R & 20.40 $\pm$ 0.17 \\ 56366.50 & 243.72 & P48 & R & 20.56 $\pm$ 0.10 \\ 56366.51 & 243.73 & P48 & R & 20.78 $\pm$ 0.13 \\ 56366.52 & 243.74 & P48 & R & 20.60 $\pm$ 0.11 \\ 56369.54 & 246.47 & P48 & R & 20.51 $\pm$ 0.13 \\ 56374.51 & 250.96 & P48 & R & 20.48 $\pm$ 0.12 \\ 56376.51 & 252.76 & P48 & R & 21.01 $\pm$ 0.20 \\ 56385.51 & 260.90 & P48 & R & 20.48 $\pm$ 0.27 \\ 56068.22 &$-$25.73 & P60 & i & 17.52 $\pm$ 0.01 \\ 56068.22 &$-$25.73 & P60 & r & 17.42 $\pm$ 0.01 \\ 56068.22 &$-$25.73 & P60 & B & 17.14 $\pm$ 0.01 \\ 56068.22 &$-$25.72 & P60 & g & 17.13 $\pm$ 0.01 \\ 56068.23 &$-$25.72 & P60 & i & 17.54 $\pm$ 0.01 \\ 56068.23 &$-$25.72 & P60 & r & 17.34 $\pm$ 0.01 \\ 56068.23 &$-$25.72 & P60 & B & 17.21 $\pm$ 0.01 \\ 56068.23 &$-$25.72 & P60 & g & 17.13 $\pm$ 0.01 \\ 56075.23 &$-$19.39 & P60 & i & 17.31 $\pm$ 0.01 \\ 56075.24 &$-$19.39 & P60 & r & 17.22 $\pm$ 0.01 \\ 56075.24 &$-$19.39 & P60 & B & 16.80 $\pm$ 0.01 \\ 56075.24 &$-$19.39 & P60 & g & 17.17 $\pm$ 0.01 \\ 56080.20 &$-$14.90 & P60 & i & 17.19 $\pm$ 0.01 \\ 56080.20 &$-$14.90 & P60 & r & 17.02 $\pm$ 0.01 \\ 56080.20 &$-$14.90 & P60 & B & 16.87 $\pm$ 0.01 \\ 56080.21 &$-$14.90 & P60 & g & 16.82 $\pm$ 0.01 \\ 56080.21 &$-$14.90 & P60 & i & 17.19 $\pm$ 0.01 \\ 56080.21 &$-$14.90 & P60 & r & 17.05 $\pm$ 0.01 \\ 56080.21 &$-$14.90 & P60 & B & 16.90 $\pm$ 0.01 \\ 56080.21 &$-$14.89 & P60 & g & 16.76 $\pm$ 0.01 \enddata \tablenotetext{a}{Calculated using MJD$_{r,\rm peak} = 56096.7$ and $z=0.107$.} \tablenotetext{b}{The magnitudes have {\it not} been corrected for Galactic extinction.} \end{deluxetable} \addtocounter{table}{-1} \begin{deluxetable}{crccc} \tablecaption{(continued) Log of observations of \dam.} \tablehead{ \colhead{MJD} & \colhead{Phase\a} & \colhead{Telescope} & \colhead{Filter} & \colhead{Magnitude\b} \\ (days) & (days) & & & AB } \startdata 56085.21 &$-$10.38 & P60 & i & 17.11 $\pm$ 0.01 \\ 56085.21 &$-$10.38 & P60 & r & 16.98 $\pm$ 0.01 \\ 56085.21 &$-$10.38 & P60 & B & 16.83 $\pm$ 0.01 \\ 56085.21 &$-$10.38 & P60 & i & 17.21 $\pm$ 0.01 \\ 56085.21 &$-$10.38 & P60 & r & 17.05 $\pm$ 0.01 \\ 56085.21 &$-$10.37 & P60 & B & 16.83 $\pm$ 0.01 \\ 56085.22 &$-$10.37 & P60 & g & 16.78 $\pm$ 0.01 \\ 56090.20 & $-$5.87 & P60 & i & 17.05 $\pm$ 0.01 \\ 56090.20 & $-$5.87 & P60 & r & 16.92 $\pm$ 0.01 \\ 56090.20 & $-$5.87 & P60 & g & 16.74 $\pm$ 0.01 \\ 56090.20 & $-$5.87 & P60 & i & 17.19 $\pm$ 0.01 \\ 56090.21 & $-$5.87 & P60 & r & 16.91 $\pm$ 0.01 \\ 56090.21 & $-$5.87 & P60 & B & 16.89 $\pm$ 0.01 \\ 56090.22 & $-$5.85 & P60 & g & 16.78 $\pm$ 0.01 \\ 56098.29 & 1.44 & P60 & i & 17.10 $\pm$ 0.01 \\ 56098.29 & 1.44 & P60 & r & 16.98 $\pm$ 0.01 \\ 56098.29 & 1.44 & P60 & B & 16.90 $\pm$ 0.01 \\ 56098.30 & 1.44 & P60 & i & 17.12 $\pm$ 0.01 \\ 56098.30 & 1.44 & P60 & r & 16.95 $\pm$ 0.01 \\ 56098.30 & 1.44 & P60 & B & 16.89 $\pm$ 0.01 \\ 56098.30 & 1.44 & P60 & g & 16.78 $\pm$ 0.01 \\ 56103.21 & 5.88 & P60 & i & 17.08 $\pm$ 0.01 \\ 56103.21 & 5.88 & P60 & r & 16.96 $\pm$ 0.01 \\ 56103.22 & 5.89 & P60 & g & 16.97 $\pm$ 0.01 \\ 56103.22 & 5.89 & P60 & i & 17.13 $\pm$ 0.01 \\ 56103.23 & 5.90 & P60 & r & 17.12 $\pm$ 0.01 \\ 56103.23 & 5.90 & P60 & B & 16.86 $\pm$ 0.01 \\ 56103.23 & 5.90 & P60 & g & 16.84 $\pm$ 0.01 \\ 56108.21 & 10.40 & P60 & i & 17.11 $\pm$ 0.01 \\ 56108.21 & 10.40 & P60 & r & 17.02 $\pm$ 0.01 \\ 56108.21 & 10.40 & P60 & B & 17.00 $\pm$ 0.01 \\ 56108.21 & 10.40 & P60 & g & 16.80 $\pm$ 0.01 \\ 56108.21 & 10.40 & P60 & i & 17.12 $\pm$ 0.01 \\ 56108.21 & 10.40 & P60 & r & 17.02 $\pm$ 0.01 \\ 56108.21 & 10.40 & P60 & B & 16.99 $\pm$ 0.01 \\ 56108.21 & 10.40 & P60 & g & 16.87 $\pm$ 0.01 \\ 56113.20 & 14.91 & P60 & i & 17.13 $\pm$ 0.01 \\ 56113.20 & 14.91 & P60 & r & 17.09 $\pm$ 0.01 \\ 56113.20 & 14.91 & P60 & B & 17.08 $\pm$ 0.02 \\ 56113.21 & 14.91 & P60 & g & 16.86 $\pm$ 0.01 \\ 56113.21 & 14.91 & P60 & i & 17.15 $\pm$ 0.01 \\ 56113.21 & 14.91 & P60 & r & 17.02 $\pm$ 0.01 \\ 56113.21 & 14.91 & P60 & B & 17.10 $\pm$ 0.02 \\ 56113.21 & 14.91 & P60 & g & 16.89 $\pm$ 0.01 \\ 56118.25 & 19.47 & P60 & i & 17.17 $\pm$ 0.01 \\ 56118.25 & 19.47 & P60 & r & 17.11 $\pm$ 0.01 \\ 56118.25 & 19.47 & P60 & B & 17.18 $\pm$ 0.01 \\ 56118.25 & 19.47 & P60 & g & 16.99 $\pm$ 0.01 \\ 56118.26 & 19.47 & P60 & i & 17.21 $\pm$ 0.01 \\ 56118.26 & 19.47 & P60 & r & 17.08 $\pm$ 0.01 \\ 56118.26 & 19.47 & P60 & B & 17.16 $\pm$ 0.01 \\ 56118.26 & 19.48 & P60 & g & 16.99 $\pm$ 0.01 \\ 56127.28 & 27.62 & P60 & g & 17.47 $\pm$ 0.02 \\ 56129.20 & 29.36 & P60 & i & 17.37 $\pm$ 0.01 \\ 56129.20 & 29.36 & P60 & r & 17.24 $\pm$ 0.01 \\ 56129.20 & 29.36 & P60 & B & 17.41 $\pm$ 0.01 \\ 56129.20 & 29.36 & P60 & g & 17.22 $\pm$ 0.01 \\ 56159.20 & 56.46 & P60 & i & 17.62 $\pm$ 0.02 \\ 56161.18 & 58.24 & P60 & r & 17.56 $\pm$ 0.01 \\ 56161.18 & 58.24 & P60 & B & 18.11 $\pm$ 0.02 \\ 56161.18 & 58.25 & P60 & g & 17.87 $\pm$ 0.01 \\ 56165.17 & 61.85 & P60 & i & 17.54 $\pm$ 0.02 \\ 56177.15 & 72.67 & P60 & i & 17.67 $\pm$ 0.02 \\ 56177.15 & 72.67 & P60 & r & 17.78 $\pm$ 0.01 \\ 56177.15 & 72.67 & P60 & B & 18.46 $\pm$ 0.03 \\ 56177.15 & 72.67 & P60 & g & 17.98 $\pm$ 0.01 \\ 56185.16 & 79.91 & P60 & i & 17.90 $\pm$ 0.02 \\ 56185.16 & 79.91 & P60 & r & 17.92 $\pm$ 0.02 \\ 56185.16 & 79.91 & P60 & B & 18.54 $\pm$ 0.03 \\ 56185.16 & 79.91 & P60 & g & 18.18 $\pm$ 0.02 \\ 56185.17 & 79.91 & P60 & i & 17.80 $\pm$ 0.02 \\ 56185.17 & 79.92 & P60 & r & 17.92 $\pm$ 0.02 \enddata \tablenotetext{a}{Calculated using MJD$_{r,\rm peak} = 56096.7$ and $z=0.107$.} \tablenotetext{b}{The magnitudes have {\it not} been corrected for Galactic extinction.} \end{deluxetable} \addtocounter{table}{-1} \begin{deluxetable}{crccc} \tablecaption{(continued) Log of observations of \dam.} \tablehead{ \colhead{MJD} & \colhead{Phase\a} & \colhead{Telescope} & \colhead{Filter} & \colhead{Magnitude\b} \\ (days) & (days) & & & AB } \startdata 56185.17 & 79.92 & P60 & B & 18.59 $\pm$ 0.04 \\ 56185.17 & 79.92 & P60 & g & 18.15 $\pm$ 0.02 \\ 56201.11 & 94.32 & P60 & g & 18.46 $\pm$ 0.03 \\ 56202.11 & 95.22 & P60 & i & 17.92 $\pm$ 0.02 \\ 56202.11 & 95.22 & P60 & r & 18.15 $\pm$ 0.02 \\ 56214.10 & 106.05 & P60 & i & 18.30 $\pm$ 0.03 \\ 56214.10 & 106.06 & P60 & r & 18.31 $\pm$ 0.02 \\ 56214.10 & 106.06 & P60 & B & 19.09 $\pm$ 0.04 \\ 56214.11 & 106.06 & P60 & g & 18.58 $\pm$ 0.02 \\ 56237.55 & 127.23 & P60 & i & 18.73 $\pm$ 0.06 \\ 56245.54 & 134.45 & P60 & i & 18.73 $\pm$ 0.13 \\ 56245.54 & 134.45 & P60 & B & 19.84 $\pm$ 0.18 \\ 56245.54 & 134.45 & P60 & g & 19.29 $\pm$ 0.09 \\ 56250.51 & 138.95 & P60 & r & 19.21 $\pm$ 0.04 \\ 56251.51 & 139.85 & P60 & r & 19.17 $\pm$ 0.03 \\ 56252.51 & 140.75 & P60 & i & 19.10 $\pm$ 0.05 \\ 56252.51 & 140.75 & P60 & B & 19.97 $\pm$ 0.11 \\ 56252.51 & 140.75 & P60 & g & 19.59 $\pm$ 0.05 \\ 56262.54 & 149.81 & P60 & r & 19.42 $\pm$ 0.09 \\ 56265.48 & 152.46 & P60 & g & 19.56 $\pm$ 0.21 \\ 56266.47 & 153.36 & P60 & r & 19.16 $\pm$ 0.04 \\ 56266.47 & 153.36 & P60 & g & 19.62 $\pm$ 0.07 \\ 56268.47 & 155.16 & P60 & B & 20.11 $\pm$ 0.10 \\ 56273.45 & 159.67 & P60 & r & 19.53 $\pm$ 0.05 \\ 56273.45 & 159.67 & P60 & g & 19.86 $\pm$ 0.06 \\ 56281.43 & 166.87 & P60 & i & 19.64 $\pm$ 0.15 \\ 56281.43 & 166.88 & P60 & r & 19.50 $\pm$ 0.09 \\ 56283.43 & 168.68 & P60 & r & 19.55 $\pm$ 0.03 \\ 56283.43 & 168.68 & P60 & B & 20.21 $\pm$ 0.09 \\ 56283.43 & 168.68 & P60 & g & 19.84 $\pm$ 0.05 \\ 56297.39 & 181.29 & P60 & i & 19.66 $\pm$ 0.07 \\ 56297.39 & 181.29 & P60 & r & 19.67 $\pm$ 0.04 \\ 56297.39 & 181.29 & P60 & B & 20.13 $\pm$ 0.10 \\ 56297.39 & 181.29 & P60 & g & 20.00 $\pm$ 0.06 \\ 56323.31 & 204.71 & P60 & i & 20.03 $\pm$ 0.13 \\ 56323.32 & 204.71 & P60 & r & 20.45 $\pm$ 0.15 \\ 56323.32 & 204.71 & P60 & g & 20.39 $\pm$ 0.16 \\ 56324.31 & 205.61 & P60 & r & 20.30 $\pm$ 0.07 \\ 56345.33 & 224.60 & P60 & i & 20.18 $\pm$ 0.13 \\ 56345.34 & 224.60 & P60 & r & 20.41 $\pm$ 0.11 \\ 56345.34 & 224.61 & P60 & B & 20.75 $\pm$ 0.27 \\ 56345.34 & 224.61 & P60 & g & 20.83 $\pm$ 0.16 \\ 56345.34 & 224.61 & P60 & i & 20.22 $\pm$ 0.12 \\ 56345.34 & 224.61 & P60 & r & 20.36 $\pm$ 0.10 \\ 56345.35 & 224.61 & P60 & B & 20.63 $\pm$ 0.24 \\ 56345.35 & 224.61 & P60 & g & 20.61 $\pm$ 0.13 \\ 56346.25 & 225.43 & P60 & i & 20.55 $\pm$ 0.20 \\ 56346.25 & 225.43 & P60 & r & 20.18 $\pm$ 0.12 \\ 56346.26 & 225.44 & P60 & g & 20.43 $\pm$ 0.17 \\ 56347.43 & 226.50 & P60 & i & 20.17 $\pm$ 0.21 \\ 56347.43 & 226.50 & P60 & r & 20.59 $\pm$ 0.22 \\ 56347.44 & 226.50 & P60 & g & 20.62 $\pm$ 0.26 \\ 56349.24 & 228.13 & P60 & i & 19.78 $\pm$ 0.21 \\ 56350.24 & 229.04 & P60 & r & 20.49 $\pm$ 0.25 \\ 56350.25 & 229.04 & P60 & g & 20.55 $\pm$ 0.31 \\ 56352.24 & 230.84 & P60 & r & 20.47 $\pm$ 0.16 \\ 56352.24 & 230.84 & P60 & g & 20.90 $\pm$ 0.28 \\ 56353.23 & 231.74 & P60 & r & 20.46 $\pm$ 0.09 \\ 56353.24 & 231.74 & P60 & B & 21.23 $\pm$ 0.26 \\ 56353.24 & 231.74 & P60 & g & 20.62 $\pm$ 0.10 \\ 56356.23 & 234.44 & P60 & i & 20.44 $\pm$ 0.12 \\ 56362.27 & 239.90 & P60 & r & 20.55 $\pm$ 0.09 \\ 56362.28 & 239.91 & P60 & B & 21.01 $\pm$ 0.21 \\ 56362.28 & 239.91 & P60 & g & 20.82 $\pm$ 0.10 \\ 56362.28 & 239.91 & P60 & r & 20.56 $\pm$ 0.09 \\ 56362.28 & 239.91 & P60 & B & 20.97 $\pm$ 0.19 \\ 56362.28 & 239.91 & P60 & g & 20.82 $\pm$ 0.10 \\ 56364.20 & 241.65 & P60 & i & 20.25 $\pm$ 0.09 \\ 56370.31 & 247.16 & P60 & r & 20.70 $\pm$ 0.13 \\ 56370.31 & 247.16 & P60 & B & 20.84 $\pm$ 0.23 \\ 56371.25 & 248.01 & P60 & g & 21.28 $\pm$ 0.29 \\ 56372.20 & 248.87 & P60 & r & 20.44 $\pm$ 0.22 \enddata \tablenotetext{a}{Calculated using MJD$_{r,\rm peak} = 56096.7$ and $z=0.107$.} \tablenotetext{b}{The magnitudes have {\it not} been corrected for Galactic extinction.} \end{deluxetable} \addtocounter{table}{-1} \begin{deluxetable}{crccc} \tablecaption{(continued) Log of observations of \dam.} \tablehead{ \colhead{MJD} & \colhead{Phase\a} & \colhead{Telescope} & \colhead{Filter} & \colhead{Magnitude\b} \\ (days) & (days) & & & AB } \startdata 56373.18 & 249.76 & P60 & i & 20.71 $\pm$ 0.25 \\ 56373.18 & 249.76 & P60 & r & 20.51 $\pm$ 0.21 \\ 56374.18 & 250.66 & P60 & r & 20.20 $\pm$ 0.25 \\ 56375.29 & 251.66 & P60 & i & 20.64 $\pm$ 0.14 \\ 56375.29 & 251.66 & P60 & r & 21.33 $\pm$ 0.18 \\ 56375.29 & 251.67 & P60 & g & 21.01 $\pm$ 0.09 \\ 56376.17 & 252.46 & P60 & r & 20.89 $\pm$ 0.25 \\ 56382.17 & 257.88 & P60 & B & 21.04 $\pm$ 0.21 \\ 56386.17 & 261.49 & P60 & g & 21.25 $\pm$ 0.13 \\ 56388.50 & 263.59 & P60 & r & 21.19 $\pm$ 0.24 \\ 56388.50 & 263.59 & P60 & g & 21.37 $\pm$ 0.15 \\ 56390.28 & 265.20 & P60 & r & 20.72 $\pm$ 0.23 \\ 56392.27 & 267.00 & P60 & r & 21.20 $\pm$ 0.26 \\ 56409.27 & 282.36 & P60 & i & 20.78 $\pm$ 0.19 \\ 56412.46 & 285.24 & P60 & i & 21.50 $\pm$ 0.24 \\ 56413.45 & 286.13 & P60 & r & 21.58 $\pm$ 0.26 \\ 56416.39 & 288.79 & P60 & r & 21.91 $\pm$ 0.30 \\ 56416.40 & 288.79 & P60 & g & 21.92 $\pm$ 0.17 \\ 56422.43 & 294.25 & P60 & i & 21.36 $\pm$ 0.20 \\ 56422.44 & 294.25 & P60 & g & 22.35 $\pm$ 0.28 \\ 56425.34 & 296.88 & P60 & r & 21.94 $\pm$ 0.27 \\ 56431.39 & 302.34 & P60 & i & 21.40 $\pm$ 0.25 \\ 56432.32 & 303.18 & P60 & i & 21.57 $\pm$ 0.23 \\ 56452.32 & 321.25 & P60 & r & 21.97 $\pm$ 0.29 \\ 56459.24 & 327.50 & P60 & i & 21.37 $\pm$ 0.18 \\ 56460.29 & 328.45 & P60 & i & 21.45 $\pm$ 0.28 \\ 56470.27 & 337.46 & P60 & i & 21.83 $\pm$ 0.31 \\ 56471.21 & 338.31 & P60 & r & 22.19 $\pm$ 0.25 \\ 56507.22 & 370.84 & P60 & r & 22.10 $\pm$ 0.23 \\ 56070.26 &$-$23.88 & LCOGT & r & 17.23 $\pm$ 0.02 \\ 56070.26 &$-$23.88 & LCOGT & r & 17.22 $\pm$ 0.02 \\ 56083.36 &$-$12.05 & LCOGT & r & 17.00 $\pm$ 0.02 \\ 56086.44 & $-$9.26 & LCOGT & r & 16.91 $\pm$ 0.02 \\ 56086.45 & $-$9.26 & LCOGT & r & 16.94 $\pm$ 0.02 \\ 56089.33 & $-$6.65 & LCOGT & r & 16.89 $\pm$ 0.02 \\ 56089.34 & $-$6.65 & LCOGT & r & 16.90 $\pm$ 0.02 \\ 56128.38 & 28.62 & LCOGT & r & 17.20 $\pm$ 0.02 \\ 56128.38 & 28.62 & LCOGT & r & 17.22 $\pm$ 0.03 \\ 56149.24 & 47.47 & LCOGT & r & 17.41 $\pm$ 0.02 \\ 56149.25 & 47.47 & LCOGT & r & 17.42 $\pm$ 0.01 \\ 56152.24 & 50.17 & LCOGT & r & 17.48 $\pm$ 0.02 \\ 56152.24 & 50.17 & LCOGT & r & 17.46 $\pm$ 0.03 \\ 56156.24 & 53.79 & LCOGT & r & 17.55 $\pm$ 0.01 \\ 56156.24 & 53.79 & LCOGT & r & 17.62 $\pm$ 0.02 \\ 56181.23 & 76.36 & LCOGT & r & 17.87 $\pm$ 0.04 \\ 56181.23 & 76.36 & LCOGT & r & 17.91 $\pm$ 0.03 \\ 56187.24 & 81.79 & LCOGT & r & 17.97 $\pm$ 0.02 \\ 56196.21 & 89.90 & LCOGT & r & 18.18 $\pm$ 0.02 \\ 56196.22 & 89.90 & LCOGT & r & 18.03 $\pm$ 0.03 \\ 56197.22 & 90.81 & LCOGT & r & 18.11 $\pm$ 0.02 \\ 56197.22 & 90.81 & LCOGT & r & 18.23 $\pm$ 0.02 \\ 56200.21 & 93.51 & LCOGT & r & 18.19 $\pm$ 0.02 \\ 56200.21 & 93.51 & LCOGT & r & 18.23 $\pm$ 0.02 \\ 56360.00 & 237.85 & Keck~I & r & 20.44 $\pm$ 0.30 \\ 56630.00 & 481.75 & Keck~I & r & 22.46 $\pm$ 0.19 \enddata \tablenotetext{a}{Calculated using MJD$_{r,\rm peak} = 56096.7$ and $z=0.107$.} \tablenotetext{b}{The magnitudes have {\it not} been corrected for Galactic extinction.} \end{deluxetable} \begin{deluxetable}{crccc} \tablecaption{Log of observations of \dcc.\label{tab:13dcc_logphotometry}} \tablehead{ \colhead{MJD} & \colhead{Phase\a} & \colhead{Telescope} & \colhead{Filter} & \colhead{Magnitude\b} \\ (days) & (days) & & & AB } \startdata 56533.36 &$-$59.36 & P48 & R & 19.68 $\pm$ 0.12 \\ 56533.42 &$-$59.32 & P48 & R & 19.73 $\pm$ 0.06 \\ 56533.45 &$-$59.29 & P48 & R & 19.76 $\pm$ 0.05 \\ 56537.35 &$-$56.57 & P48 & R & 19.66 $\pm$ 0.06 \\ 56537.47 &$-$56.49 & P48 & R & 19.86 $\pm$ 0.04 \\ 56537.50 &$-$56.46 & P48 & R & 19.70 $\pm$ 0.04 \\ 56538.45 &$-$55.80 & P48 & R & 19.80 $\pm$ 0.04 \\ 56539.34 &$-$55.18 & P48 & R & 19.66 $\pm$ 0.07 \\ 56539.38 &$-$55.15 & P48 & R & 19.79 $\pm$ 0.07 \\ 56539.42 &$-$55.12 & P48 & R & 19.80 $\pm$ 0.06 \\ 56540.44 &$-$54.41 & P48 & R & 19.98 $\pm$ 0.19 \\ 56541.34 &$-$53.78 & P48 & R & 19.73 $\pm$ 0.08 \\ 56541.37 &$-$53.76 & P48 & R & 19.55 $\pm$ 0.17 \\ 56541.41 &$-$53.73 & P48 & R & 19.38 $\pm$ 0.17 \\ 56542.34 &$-$53.08 & P48 & R & 19.77 $\pm$ 0.08 \\ 56542.45 &$-$53.00 & P48 & R & 19.69 $\pm$ 0.06 \\ 56542.49 &$-$52.98 & P48 & R & 19.71 $\pm$ 0.06 \\ 56543.39 &$-$52.35 & P48 & R & 19.84 $\pm$ 0.06 \\ 56543.46 &$-$52.30 & P48 & R & 19.91 $\pm$ 0.06 \\ 56543.50 &$-$52.27 & P48 & R & 19.74 $\pm$ 0.05 \\ 56545.33 &$-$51.00 & P48 & R & 19.55 $\pm$ 0.07 \\ 56546.32 &$-$50.30 & P48 & R & 19.76 $\pm$ 0.08 \\ 56546.43 &$-$50.22 & P48 & R & 19.86 $\pm$ 0.05 \\ 56546.46 &$-$50.20 & P48 & R & 19.86 $\pm$ 0.05 \\ 56547.37 &$-$49.57 & P48 & R & 19.79 $\pm$ 0.07 \\ 56547.49 &$-$49.48 & P48 & R & 19.80 $\pm$ 0.06 \\ 56548.32 &$-$48.90 & P48 & R & 19.82 $\pm$ 0.07 \\ 56548.47 &$-$48.80 & P48 & R & 19.76 $\pm$ 0.04 \\ 56548.51 &$-$48.77 & P48 & R & 19.71 $\pm$ 0.06 \\ 56549.41 &$-$48.14 & P48 & R & 19.88 $\pm$ 0.04 \\ 56549.45 &$-$48.11 & P48 & R & 19.82 $\pm$ 0.04 \\ 56549.49 &$-$48.09 & P48 & R & 19.76 $\pm$ 0.03 \\ 56550.39 &$-$47.45 & P48 & R & 19.77 $\pm$ 0.07 \\ 56550.48 &$-$47.39 & P48 & R & 19.83 $\pm$ 0.04 \\ 56550.51 &$-$47.37 & P48 & R & 19.81 $\pm$ 0.03 \\ 56551.42 &$-$46.74 & P48 & R & 19.78 $\pm$ 0.08 \\ 56551.46 &$-$46.71 & P48 & R & 19.86 $\pm$ 0.07 \\ 56551.49 &$-$46.69 & P48 & R & 19.83 $\pm$ 0.04 \\ 56552.40 &$-$46.06 & P48 & R & 19.76 $\pm$ 0.09 \\ 56552.44 &$-$46.02 & P48 & R & 19.88 $\pm$ 0.08 \\ 56552.47 &$-$46.00 & P48 & R & 19.80 $\pm$ 0.09 \\ 56561.28 &$-$39.84 & P48 & R & 19.81 $\pm$ 0.16 \\ 56561.36 &$-$39.79 & P48 & R & 19.87 $\pm$ 0.10 \\ 56561.40 &$-$39.77 & P48 & R & 19.85 $\pm$ 0.09 \\ 56562.45 &$-$39.03 & P48 & R & 19.90 $\pm$ 0.08 \\ 56562.48 &$-$39.01 & P48 & R & 19.81 $\pm$ 0.07 \\ 56562.51 &$-$38.98 & P48 & R & 19.72 $\pm$ 0.08 \\ 56563.42 &$-$38.35 & P48 & R & 19.74 $\pm$ 0.09 \\ 56563.49 &$-$38.30 & P48 & R & 19.75 $\pm$ 0.08 \\ 56563.52 &$-$38.28 & P48 & R & 19.91 $\pm$ 0.12 \\ 56564.43 &$-$37.65 & P48 & R & 19.80 $\pm$ 0.07 \\ 56564.48 &$-$37.61 & P48 & R & 19.95 $\pm$ 0.07 \\ 56564.51 &$-$37.59 & P48 & R & 19.87 $\pm$ 0.08 \\ 56565.41 &$-$36.96 & P48 & R & 19.78 $\pm$ 0.07 \\ 56565.47 &$-$36.92 & P48 & R & 19.88 $\pm$ 0.06 \\ 56565.50 &$-$36.90 & P48 & R & 19.86 $\pm$ 0.08 \\ 56566.40 &$-$36.27 & P48 & R & 19.82 $\pm$ 0.06 \\ 56566.45 &$-$36.23 & P48 & R & 19.79 $\pm$ 0.05 \\ 56566.48 &$-$36.21 & P48 & R & 19.85 $\pm$ 0.07 \\ 56567.39 &$-$35.58 & P48 & R & 19.77 $\pm$ 0.05 \\ 56567.44 &$-$35.54 & P48 & R & 19.76 $\pm$ 0.06 \\ 56567.47 &$-$35.52 & P48 & R & 19.86 $\pm$ 0.06 \\ 56568.38 &$-$34.89 & P48 & R & 19.93 $\pm$ 0.06 \\ 56568.45 &$-$34.84 & P48 & R & 19.83 $\pm$ 0.07 \\ 56568.48 &$-$34.81 & P48 & R & 19.83 $\pm$ 0.07 \\ 56569.39 &$-$34.18 & P48 & R & 19.77 $\pm$ 0.08 \\ 56569.46 &$-$34.13 & P48 & R & 19.84 $\pm$ 0.06 \\ 56569.49 &$-$34.11 & P48 & R & 19.76 $\pm$ 0.07 \\ 56570.39 &$-$33.48 & P48 & R & 19.83 $\pm$ 0.22 \\ 56570.47 &$-$33.43 & P48 & R & 19.39 $\pm$ 0.16 \\ 56570.50 &$-$33.40 & P48 & R & 19.94 $\pm$ 0.22 \\ 56571.40 &$-$32.77 & P48 & R & 20.01 $\pm$ 0.13 \enddata \tablenotetext{a}{Calculated using MJD$_{r,\rm peak} = 56618.3$ and $z=0.431$.} \tablenotetext{b}{The magnitudes have {\it not} been corrected for Galactic extinction.} \end{deluxetable} \addtocounter{table}{-1} \begin{deluxetable}{crccc} \tablecaption{(continued) Log of observations of \dcc.} \tablehead{ \colhead{MJD} & \colhead{Phase\a} & \colhead{Telescope} & \colhead{Filter} & \colhead{Magnitude\b} \\ (days) & (days) & & & AB } \startdata 56571.46 &$-$32.73 & P48 & R & 19.72 $\pm$ 0.08 \\ 56571.49 &$-$32.71 & P48 & R & 20.08 $\pm$ 0.18 \\ 56572.39 &$-$32.08 & P48 & R & 19.83 $\pm$ 0.08 \\ 56572.46 &$-$32.04 & P48 & R & 19.85 $\pm$ 0.07 \\ 56573.40 &$-$31.38 & P48 & R & 19.83 $\pm$ 0.07 \\ 56573.44 &$-$31.35 & P48 & R & 19.98 $\pm$ 0.07 \\ 56573.47 &$-$31.33 & P48 & R & 19.99 $\pm$ 0.11 \\ 56576.39 &$-$29.29 & P48 & R & 19.83 $\pm$ 0.06 \\ 56576.49 &$-$29.22 & P48 & R & 19.90 $\pm$ 0.08 \\ 56576.53 &$-$29.19 & P48 & R & 19.78 $\pm$ 0.10 \\ 56588.21 &$-$21.03 & P48 & R & 19.76 $\pm$ 0.13 \\ 56588.35 &$-$20.93 & P48 & R & 20.05 $\pm$ 0.16 \\ 56588.38 &$-$20.91 & P48 & R & 19.76 $\pm$ 0.12 \\ 56589.28 &$-$20.28 & P48 & R & 19.62 $\pm$ 0.07 \\ 56589.31 &$-$20.26 & P48 & R & 19.93 $\pm$ 0.10 \\ 56589.33 &$-$20.24 & P48 & R & 19.76 $\pm$ 0.06 \\ 56590.36 &$-$19.53 & P48 & R & 19.74 $\pm$ 0.07 \\ 56590.38 &$-$19.51 & P48 & R & 19.72 $\pm$ 0.08 \\ 56591.28 &$-$18.88 & P48 & R & 19.75 $\pm$ 0.08 \\ 56591.32 &$-$18.85 & P48 & R & 19.64 $\pm$ 0.04 \\ 56591.35 &$-$18.83 & P48 & R & 19.68 $\pm$ 0.04 \\ 56592.30 &$-$18.17 & P48 & R & 19.61 $\pm$ 0.04 \\ 56592.39 &$-$18.11 & P48 & R & 19.67 $\pm$ 0.06 \\ 56592.42 &$-$18.09 & P48 & R & 19.71 $\pm$ 0.04 \\ 56596.19 &$-$15.45 & P48 & R & 19.75 $\pm$ 0.11 \\ 56596.30 &$-$15.37 & P48 & R & 19.66 $\pm$ 0.08 \\ 56596.33 &$-$15.36 & P48 & R & 19.67 $\pm$ 0.07 \\ 56597.23 &$-$14.73 & P48 & R & 19.54 $\pm$ 0.09 \\ 56597.31 &$-$14.67 & P48 & R & 19.70 $\pm$ 0.08 \\ 56597.34 &$-$14.65 & P48 & R & 19.55 $\pm$ 0.07 \\ 56598.24 &$-$14.02 & P48 & R & 19.59 $\pm$ 0.10 \\ 56598.31 &$-$13.97 & P48 & R & 19.62 $\pm$ 0.06 \\ 56598.33 &$-$13.95 & P48 & R & 19.67 $\pm$ 0.07 \\ 56599.24 &$-$13.32 & P48 & R & 19.58 $\pm$ 0.07 \\ 56599.28 &$-$13.29 & P48 & R & 19.71 $\pm$ 0.06 \\ 56599.30 &$-$13.28 & P48 & R & 19.61 $\pm$ 0.06 \\ 56600.25 &$-$12.62 & P48 & R & 19.59 $\pm$ 0.05 \\ 56600.36 &$-$12.54 & P48 & R & 19.54 $\pm$ 0.07 \\ 56600.39 &$-$12.52 & P48 & R & 19.47 $\pm$ 0.16 \\ 56602.17 &$-$11.27 & P48 & R & 19.55 $\pm$ 0.12 \\ 56602.23 &$-$11.23 & P48 & R & 19.57 $\pm$ 0.15 \\ 56602.25 &$-$11.21 & P48 & R & 19.54 $\pm$ 0.11 \\ 56603.17 &$-$10.57 & P48 & R & 19.37 $\pm$ 0.09 \\ 56603.24 &$-$10.53 & P48 & R & 19.44 $\pm$ 0.07 \\ 56603.29 &$-$10.49 & P48 & R & 19.51 $\pm$ 0.06 \\ 56604.29 & $-$9.79 & P48 & R & 19.60 $\pm$ 0.04 \\ 56604.41 & $-$9.71 & P48 & R & 19.57 $\pm$ 0.04 \\ 56604.44 & $-$9.69 & P48 & R & 19.58 $\pm$ 0.06 \\ 56605.34 & $-$9.06 & P48 & R & 19.53 $\pm$ 0.06 \\ 56605.38 & $-$9.03 & P48 & R & 19.62 $\pm$ 0.06 \\ 56605.41 & $-$9.01 & P48 & R & 19.54 $\pm$ 0.05 \\ 56607.16 & $-$7.79 & P48 & R & 19.82 $\pm$ 0.11 \\ 56607.29 & $-$7.69 & P48 & R & 19.54 $\pm$ 0.07 \\ 56607.32 & $-$7.67 & P48 & R & 19.47 $\pm$ 0.07 \\ 56608.22 & $-$7.04 & P48 & R & 19.34 $\pm$ 0.06 \\ 56608.26 & $-$7.02 & P48 & R & 19.60 $\pm$ 0.12 \\ 56608.29 & $-$6.99 & P48 & R & 19.48 $\pm$ 0.10 \\ 56609.19 & $-$6.36 & P48 & R & 19.64 $\pm$ 0.10 \\ 56609.34 & $-$6.26 & P48 & R & 19.45 $\pm$ 0.07 \\ 56537.33 &$-$56.58 & P60 & r & 19.81 $\pm$ 0.04 \\ 56637.12 & 13.15 & P60 & i & 19.95 $\pm$ 0.13 \\ 56637.12 & 13.15 & P60 & r & 19.88 $\pm$ 0.13 \\ 56637.12 & 13.15 & P60 & g & 20.14 $\pm$ 0.11 \\ 56645.12 & 18.74 & P60 & r & 19.93 $\pm$ 0.18 \\ 56645.12 & 18.74 & P60 & g & 20.13 $\pm$ 0.14 \\ 56647.13 & 20.15 & P60 & i & 19.75 $\pm$ 0.16 \\ 56647.13 & 20.15 & P60 & r & 20.83 $\pm$ 0.24 \\ 56647.14 & 20.15 & P60 & g & 20.46 $\pm$ 0.11 \\ 56648.11 & 20.83 & P60 & i & 20.15 $\pm$ 0.10 \\ 56648.12 & 20.84 & P60 & r & 20.02 $\pm$ 0.04 \\ 56648.12 & 20.84 & P60 & g & 20.43 $\pm$ 0.05 \\ 56655.20 & 25.79 & P60 & i & 19.99 $\pm$ 0.12 \enddata \tablenotetext{a}{Calculated using MJD$_{r,\rm peak} = 56618.3$ and $z=0.431$.} \tablenotetext{b}{The magnitudes have {\it not} been corrected for Galactic extinction.} \end{deluxetable} \addtocounter{table}{-1} \begin{deluxetable}{crccc} \tablecaption{(continued) Log of observations of \dcc.} \tablehead{ \colhead{MJD} & \colhead{Phase\a} & \colhead{Telescope} & \colhead{Filter} & \colhead{Magnitude\b} \\ (days) & (days) & & & AB } \startdata 56655.20 & 25.79 & P60 & r & 20.25 $\pm$ 0.11 \\ 56655.21 & 25.79 & P60 & g & 20.85 $\pm$ 0.15 \\ 56656.17 & 26.46 & P60 & i & 20.17 $\pm$ 0.10 \\ 56656.17 & 26.46 & P60 & r & 20.25 $\pm$ 0.08 \\ 56656.18 & 26.47 & P60 & g & 20.66 $\pm$ 0.09 \\ 56657.13 & 27.14 & P60 & i & 20.14 $\pm$ 0.05 \\ 56657.13 & 27.14 & P60 & r & 20.21 $\pm$ 0.03 \\ 56657.14 & 27.14 & P60 & g & 20.62 $\pm$ 0.04 \\ 56658.13 & 27.83 & P60 & i & 20.15 $\pm$ 0.05 \\ 56658.13 & 27.83 & P60 & r & 20.25 $\pm$ 0.04 \\ 56658.13 & 27.84 & P60 & g & 20.66 $\pm$ 0.05 \\ 56659.20 & 28.58 & P60 & i & 20.21 $\pm$ 0.06 \\ 56660.15 & 29.25 & P60 & i & 20.12 $\pm$ 0.06 \\ 56660.16 & 29.25 & P60 & r & 20.29 $\pm$ 0.05 \\ 56660.16 & 29.25 & P60 & g & 20.65 $\pm$ 0.05 \\ 56662.21 & 30.68 & P60 & i & 20.38 $\pm$ 0.13 \\ 56662.21 & 30.68 & P60 & r & 20.30 $\pm$ 0.08 \\ 56662.22 & 30.69 & P60 & g & 20.82 $\pm$ 0.08 \\ 56671.19 & 36.96 & P60 & i & 20.41 $\pm$ 0.15 \\ 56671.19 & 36.96 & P60 & r & 20.60 $\pm$ 0.14 \\ 56672.19 & 37.66 & P60 & i & 20.88 $\pm$ 0.26 \\ 56672.20 & 37.66 & P60 & r & 20.59 $\pm$ 0.21 \\ 56673.18 & 38.35 & P60 & r & 20.60 $\pm$ 0.19 \\ 56676.13 & 40.41 & P60 & i & 20.70 $\pm$ 0.08 \\ 56676.13 & 40.41 & P60 & r & 21.33 $\pm$ 0.10 \\ 56676.14 & 40.42 & P60 & g & 21.59 $\pm$ 0.09 \\ 56683.25 & 45.39 & P60 & i & 21.13 $\pm$ 0.16 \\ 56683.25 & 45.39 & P60 & r & 21.37 $\pm$ 0.19 \\ 56685.13 & 46.70 & P60 & i & 21.41 $\pm$ 0.14 \\ 56685.13 & 46.70 & P60 & r & 21.65 $\pm$ 0.20 \\ 56697.14 & 55.10 & P60 & i & 21.43 $\pm$ 0.21 \\ 56697.15 & 55.10 & P60 & r & 21.51 $\pm$ 0.19 \\ 56710.18 & 64.20 & P60 & i & 21.25 $\pm$ 0.21 \\ 56701.20 & 57.93 & DCT & r & 22.20 $\pm$ 0.24 \\ 56701.20 & 57.93 & DCT & i & 21.74 $\pm$ 0.21 \\ 56908.50 & 202.79 & DCT & r & 23.20 $\pm$ 0.28 \\ 57070.00 & 315.65 & HST & F625W& 25.00 $\pm$ 0.13 \enddata \tablenotetext{a}{Calculated using MJD$_{r,\rm peak} = 56618.3$ and $z=0.431$.} \tablenotetext{b}{The magnitudes have {\it not} been corrected for Galactic extinction.} \end{deluxetable} | 16 | 9 | 1609.08145 |
1609 | 1609.03899_arXiv.txt | Diffusive shock acceleration (DSA) at relativistic shocks is widely thought to be an important acceleration mechanism in various astrophysical jet sources, including radio-loud active galactic nuclei such as blazars. Such acceleration can produce the non-thermal particles that emit the broadband continuum radiation that is detected from extragalactic jets. An important recent development for blazar science is the ability of {\it Fermi}-LAT spectroscopy to pin down the shape of the distribution of the underlying non-thermal particle population. This paper highlights how multi-wavelength spectra spanning optical to X-ray to gamma-ray bands can be used to probe diffusive acceleration in relativistic, oblique, magnetohydrodynamic (MHD) shocks in blazar jets. Diagnostics on the MHD turbulence near such shocks are obtained using thermal and non-thermal particle distributions resulting from detailed Monte Carlo simulations of DSA. These probes are afforded by the characteristic property that the synchrotron $\nu F_{\nu}$ peak energy does not appear in the gamma-ray band above 100 MeV. We investigate self-consistently the radiative synchrotron and inverse Compton signatures of the simulated particle distributions. Important constraints on the diffusive mean free paths of electrons, and the level of electromagnetic field turbulence are identified for three different case study blazars, Mrk 501, BL Lacertae and AO 0235+164. The X-ray excess of AO 0235+164 in a flare state can be modelled as the signature of bulk Compton scattering of external radiation fields, thereby tightly constraining the energy-dependence of the diffusion coefficient for electrons. The concomitant interpretations that turbulence levels decline with remoteness from jet shocks, and the probable significant role for non-gyroresonant diffusion, are posited. | Extragalactic jets of collimated relativistic outflows are some of the most powerful emitters of radiation in the Universe. Two contrasting types of jets are found in transient gamma-ray bursts (GRBs), and persistent but highly-variable active galactic nuclei (AGN). Both are considered prime candidates for the production of ultra-high energy cosmic rays above \teq{10^{18}}eV, and perhaps also the high energy neutrinos recently detected by IceCube \citep{Aartsen14}. Yet they are very different in their origin. GRBs probably result from the explosion of massive progenitor stars in {\it hypernovae}, or the merger of compact binary neutron stars, whereas AGN are continually powered over eons by material accreted onto supermassive black holes (SMBHs) in the center of distant galaxies. The masses of such SMBHs are typically in the range \teq{10^{6.5}-10^9 M_{\odot}} \citep{Bentz-Katz15},\footnote{For an accessible Web database of SMBH mass listings, see {\tt http://www.astro.gsu.edu/AGNmass}.} deduced in part from their large luminosities, typically \teq{10^{42}-10^{47}}erg/sec, and also from reverberation mapping techniques \citep[e.g.,][]{Peterson04}. In this paper, the focus is on the interpretation of the jet environments of AGN, in particular on the specific subset known as {\it blazars}, which exhibit flares with short timescale variations in radio, optical, X-ray and $\gamma$-ray wavebands. The class of blazars was identified following the discovery \citep{Hartmann92,Lin92} of transient gamma-ray emission in 3C 279 and Mrk 421 by the EGRET instrument on the Compton Gamma-Ray Observatory in the 100 MeV -- 1 GeV range. This was around the same time that the Whipple atmospheric \v{C}erenkov telescope (ACT) discerned that Mrk 421 also emitted at TeV energies \citep{Punchetal92}. Blazars generally evince the key identifying feature of their parent BL Lac objects like {\it BL Lacertae} of a general absence of emission lines in their non-thermal optical spectra. Their radio continua often possess quite flat spectra. These characteristics suggest synchrotron radiation from non-thermal electrons gyrating in magnetic fields as the origin of their emission in these two wavebands. Such a hypothesis is supported by detections of polarization in both radio and optical bands. Most remarkable among the ensemble of blazar polarimetry studies is the observation of rapid optical polarization angle swings in 3C 279 \citep[seen also for other blazars, including PKS 1510-089;][]{Marscher10}, contemporaneous with a gamma-ray flare seen by {\it Fermi}-LAT \citep{Abdo10b}. This indicates a large-scale ordering of the magnetic field in the 3C 279 jet, and spatial coincidence of the optical and gamma-ray emission zones. In many blazars, this putative synchrotron component extends to X-rays, where polarization measurements in the not-to-distant future define a hope for further constraining the emission mechanism of blazars \citep[e.g.,][]{Krawczynski11}. The prevailing paradigm is that blazars' gamma-ray signals are generated by inverse Compton scattering of synchrotron photons by the same relativistic electrons that emit this radio-to-X-ray signal, so-called synchrotron-self-Compton (SSC) models \citep[e.g.,][]{MCG92,MK97,CB02}. An alternative possibillity for the target photons seeding this upscattering is an IR/optical/UV photon source of possibly disk or ambient origin \citep[e.g.,][]{DSM92,SBR94} not too distant from the black hole. The main competing scenario for $\gamma$-ray production in blazars is the hadronic model \citep[e.g.,][]{MB89,RM98,MP01} where non-thermal protons collide with disk and jet-associated synchrotron photons, generating pair cascades via photo-pion/muon production processes. Proton synchrotron radiation can also appear in the $\gamma$-ray band. High energy neutrinos are a signature product of such cascades, forging a connection of hadronic blazar models to the observation of energetic neutrinos up to around \teq{10^{15}}eV by the IceCube experiment \citep{Aartsen14}. Blazar spectra observed above 300 GeV by ACTs are typically quite steep, defining a turnover. These photons are subject to strong \teq{\gamma\gamma} pair absorption in propagating to Earth due to the presence of intergalactic infra-red and optical starlight background fields. Correcting for such attenuation can yield inferences of very flat intrinsic source spectra \citep[e.g.,][]{SBS07}, yet uncertainties in the background fields permeate such protocols. An important recent development for blazar science has been the improvement of sensivity in the 100 MeV-100 GeV window, afforded by the {\it Fermi}-Large Area Telescope (LAT). Over the last eight years, LAT data has enabled measurements of the power-law index of spectra from numerous blazars \citep{Abdo09,Abdo10c}. This provides a more robust measure of the underlying non-thermal particle population, since the LAT band is generally subject to only small \teq{\gamma\gamma} pair attenuation, and negligibly so below around 3 GeV. This is a new probe enabled for multiwavelength blazar studies, a tool that we exploit to some extent here. Yet more particularly, one of our case studies, AO 0235+164, is not yet detected with current ACT facilities, so that the {\it Fermi}-LAT data prove critical to our interpretation of its jet environment. The rapid flux variability seen in radio, X-ray and GeV--TeV flares \citep[e.g., see][for Mrk 421]{MFTetal99,TKMetal00} drives the prevailing picture for the blazar environment: their jets are relativistic, and compactly structured on small spatial scales that are unresolvable by present gamma-ray telescopes. For a survey of radio and X-ray jet properties including spatial morphology, see \citet{HK06}. The supersonic outflows in these jets naturally generate relativistic shocks, and these can form the principal sites for dissipation of the ballistic kinetic energy via acceleration of electrons and perhaps also ions to the ultrarelativistic energies demanded by the X-ray and $\gamma$-ray data. Diffusive Fermi acceleration at such jet shocks is believed to be the main candidate mechanism for energizing such charges \citep{BE87,Drury83,JE91}. This is motivated by the fact that Fermi acceleration is both extremely efficient and very fast, precipitating acceleration rates \teq{{\dot \gamma}} of the order of the gyrofrequency \teq{\omega_g = eB/mc}. This follows because the diffusive collisions of electrons and ions with MHD turbulence in jets is putatively dominated by gyroresonant interactions. The efficiency of particle-turbulence interactions is an issue that we will visit in this paper. We note that shocks also may span a large surface cross section of the jet that the outward-propagating plasma may encounter with high probability. Another possibility is that shear layers encapsulating sharp velocity gradients transverse to the net flow may precipitate Fermi-type acceleration \citep{Ostrowski90} due to transport of charges straddling the shear ``discontinuity.'' Observational evidence supporting such transverse velocity structure comes from parsec-scale limb-brightening of blazar and radio galaxy jets revealed in VLBI observations \citep[e.g., see ][for Mrk 501]{GGFetal04}. MHD/hydrodynamic simulations of spine-sheath jets and their launching \citep[e.g.,][]{MK07,MN07,MHN07,TMN08} indicate that the sheath, in combination with a poloidal magnetic field, aids in stabilizing the jet. Rayleigh-Taylor-type instabilities may develop at the spine-sheath interface, and such turbulence offers another promising avenue for accelerating relativistic particles. Indeed, particle-in-cell \teq{e^{\pm}} plasma simulations of relativistic shear flows \citep{LBS13} indicate substantial energization of pairs in the boundary layers. Other paradigms such as acceleration by reconnection of magnetic fields embedded in Poynting-flux dominated outflows can be envisaged. Whether shocks or shear boundaries or reconnection zones provide the dominant energization site for non-thermal particles in blazars defines a major objective for future theoretical efforts within the blazar community. Our considerations here will focus on acceleration at blazar jet shocks. To explore the blazar shock paradigm, the standard practice has been to develop multiwavelength (MW) spectral models, spanning radio to TeV gamma-ray wavelengths \citep{DSM92,SBR94,MCG92,MK97,CB02,BDF08}. This generally gives a broad-brush assessment of a range of jet environment parameters, but the spectral fits in any one band are of limited quality. The high statistics spectroscopy afforded by the {\it Fermi-LAT} data demands a closer look at the constraints observations can elucidate for the shocked environs of blazar jets and on the shock acceleration process itself. In particular, theoretical studies of shock acceleration have determined \citep[e.g.,][]{EJR90,KH89,ED04,SB12} that a wide variety of power-law indices for accelerated charges are possible for a given velocity compression ratio across the discontinuity. These indices are sensitive to the orientation of the mean magnetic field, and the character of the {\it in situ} MHD turbulence. In addition, it has been understood for nearly two decades \citep[see][]{IT96} that diffusive shock acceleration is so efficient that low levels of field turbulence are required in blazar jets to accommodate synchrotron spectral peaks appearing in the X-ray band. This serves as a central issue to the discourse of this paper. In the multi-wavelength blazar models developed herein, detailed results from comprehensive Monte Carlo simulations of diffusive acceleration at relativistic shocks are employed, building upon the exposition of \citet{SB12}. These simulations capture the relationship between turbulence parameters and the power-law index, and also the connection between the thermal bulk of the population (a hot Maxwellian-like component) and the power-law tail of the accelerated species. Such an approach goes beyond an elemental description of the distribution of accelerated charges, usually a power-law truncated at some minimum and maximum particle Lorentz factors, that is commonly invoked in blazar spectral modeling. Our shock simulation approach is outlined in Section~\ref{sec:shock_accel}. We fold these simulation results through the one-zone SSC/external Compton models of \citet{Boettcher13} for radiation emission and transport in blazars. The radiation modeling protocols are summarized in Section~\ref{sec:radmodels}, and subsequently results from this analysis are presented therein. Diagnostics are obtained for the particle mean free path and the level of field turbulence for our three chosen blazars, namely Mrk 501, BL Lacertae, and the BL Lac object AO 0235+164. The non-thermal emission components {\it are primarily generated by leptons that undergo repeated drift acceleration and interspersed upstream reflections in the shock layer}, the result being extremely flat distributions emerging from the acceleration zones. By exploring a range of dependences of the diffusive mean free path \teq{\lambda} on the momentum \teq{p} of accelerated charges, the possible interpretation that turbulence levels decline with remoteness from a shock is identified, perhaps signalling a significant role for non-gyroresonant diffusion in the vicinity of blazar jet shocks, an interpretation discussed in Section~\ref{sec:discussion}. This determination is impelled by the fact that gyroresonant diffusive acceleration and energization in magnetic reconnection zones are just too fast to accommodate the multi-wavelength spectra of blazars, in particular the positioning of the synchrotron \teq{\nu F_{\nu}} peak in the optical or X-ray bands. \newpage \section[]{Simulations of Diffusive and Drift Acceleration at Shocks} \label{sec:shock_accel} To construct a more precise assessment of multiwavelength emission models for blazars, and derive a deeper understanding of their jet physics, it is necessary to go beyond simplistic truncated power-law distributions for the radiating leptons and hadrons. This is our approach in this paper, where detailed modeling of particle distribution functions, anisotropies, and their relationship to plasma turbulence are afforded through simulations of shock acceleration processes. Simulations are the most encompassing technique to apply to the blazar jet problem because the shocks are inherently relativistic: faster and slower regions of the jet flow travelling near \teq{c} have relative velocities that are mildly-relativistic. This domain will form the focus of our exposition here. \subsection[]{Background on Shock Acceleration Theory} \label{sec:shock_theory} Various approaches have been adopted to model diffusive shock acceleration (DSA) at relativistic discontinuities. These include analytic methods \citep[e.g.,][]{Peacock81,KS87,KH89,Kirk00}, and Monte Carlo simulations of convection and diffusion \citep[e.g.,][]{EJR90,BO98,ED04,NO04,SB12}. Particle-in-cell (PIC) plasma simulations have also been employed to study this problem via modeling the establishment of electromagnetic turbulence self-consistently \citep[e.g.,][]{Gallant92,Silva03,Nishikawa05,Spitk08,SS09} in conjunction with the driver charges and their currents. Important insights are gleaned from each of these techniques. All exhibit the core property that in collisionless shocks, non-thermal, charged particles gain energy by scattering between MHD turbulence (for example magnetic ``islands'' observed in PIC simulations) that is ``anchored'' in the converging upstream and downstream plasmas, the so-called {\it Fermi mechanism}. This defines a fundamental association between turbulence contained in shock environs, diffusion and ultimate acceleration. Each of these approaches to the shock acceleration problem has both merits and limitations. Analytic techniques provide core insights into global character, though are often restricted to treating particles well above thermal energies where the acceleration process possesses no momentum scale. Monte Carlo simulations with prescribed injected turbulence \citep[e.g.,][]{NO04} explore wave-particle interactions and diffusive acceleration well above thermal energies, but do not treat the electrodynamics of wave generation and dissipation self-consistently. Particle-in-cell simulations provide the greatest depth in modeling shock layer microphysics by solving Maxwell's equations and the Newton-Lorentz force equation, though their macro-particle approximation does eliminate electrodynamic information on the smallest scales. With increased box sizes, PIC codes have now realized the establishment of truly non-thermal components, but still cannot explore acceleration beyond modest non-thermal energies \citep[e.g.,][]{SS09}; their focus is still on the thermal dissipation and injection domains in shock layers. The Monte Carlo approach employed here subsumes the microphysics in a parametric description of diffusion, and so does not investigate the feedback between charges, currents and hydromagnetic waves. Yet it is ideal for describing the diffusive and shock drift acceleration of particles from thermal injection scales out to the highest energies relevant to blazar jet models, and so is the preferred tool for interfacing with astronomical emission datasets. A key feature of both relativistic and non-relativistic shock acceleration theory is that the acceleration process possesses no momentum scale, and the resulting particle distribution takes the form \teq{dn/dp \propto p^{-\sigma}}. For non-relativistic shocks, since their speeds \teq{v\approx c} far exceed \teq{u_{1x}} (\teq{u_{2x}}), the upstream (downstream) flow speed component in the co-ordinate direction \teq{x} normal to the shock, the energetic particles are nearly isotropic in all fluid frames. The acceleration process then establishes \citep[e.g.,][]{Drury83,JE91} a power-law distribution with index \teq{\sigma = (r+2)/(r-1)}, where \teq{r=u_{1x}/u_{2x}} is the shock's velocity compression ratio. The index \teq{\sigma} in this \teq{u_1\ll c} limit is independent of the shock speed, \teq{u_1}, the upstream field obliquity angle \teq{\ThetaBfone} (to the shock normal, which is in the \teq{x}-direction throughout this paper), and any details of the scattering process. This canonical result has propelled the popularity of shock acceleration as a key element of paradigms promoted over the last four decades for the generation of cosmic rays. In contrast, it is widely understood that because plasma anisotropy is prevalent in relativistic shocks, when \teq{u_1\sim c}, the index \teq{\sigma} of the power-law distribution is a function of the flow speed \teq{u_1}, the field obliquity angle \teq{\ThetaBfone}, and the nature of the scattering. Test-particle acceleration in parallel (\teq{\ThetaBfone =0}) relativistic shocks evinces the essential property that a so-called ``universal'' spectral index \teq{\sigma\sim 2.23} exists in the two limits of \teq{\Gamma_1 \gg 1} and small angle scattering, i.e., \teq{\delta \theta \ll 1/\Gamma_1}, for a shock compression ratio of \teq{r=u_{1x}/u_{2x}=3}. This was showcased in the seminal work \citet{Kirk00}, which employed semi-analytic methods to solve the diffusion-convection equation, and was also generated in the Monte Carlo analyses of \citet{BO98} and \citet{Baring99}. Here \teq{\delta\theta} is the average angle a particle's momentum vector deviates in a scattering ``event,'' and $\Gamma_1 = (1 - [u_1/c]^2)^{-1/2}$ is the Lorentz factor of the upstream flow in the shock rest frame. For all other parameter regimes in relativistic shocks, a wide range of departures of \teq{\sigma} from this special index is observed \citep[see][and references therein]{ED04,SB12}. While this is a complication, it enables powerful spectral diagnostics on the large scale electromagnetic structure of shocks and also the MHD turbulence in their environs. \subsection[]{The Monte Carlo Simulational Method} \label{sec:MC} The Monte Carlo technique employed here to model acceleration at blazar shocks solves the Boltzmann equation with a phenomenological scattering operator \citep[][]{JE91,EJR90,SB12}. The simulation space is divided into grid sections with boundaries parallel to the planar shock interface. Sections can possess different flow velocities, mean magnetic field vectors, etc., defining the MHD structure of the shock. For electron-ion shocks, the vastly different inertial scales can be easily accommodated, though the present application is for pair plasma shocks. Charges are injected far upstream and diffuse and convect in the shock neighborhood. Their flux-weighted contributions to the momentum distribution function are logged at any position \teq{x}. Each particle is followed until it leaves the system by either convecting sufficiently far downstream, or exceeding some prescribed maximum momentum \teq{\pmax}. To enhance the speed of the code, a {\it probability of return} boundary is introduced \citep[following][]{EJR90} at several diffusion lengths downstream of the shock, beyond which the decision to retain or discard a charge subject to convection and diffusion is made statistically: see \citet{SB12}. The details of particle transport are given in \citet{EJR90,ED04,SB12}: the simulation can model both scatterings with small angular deflections \teq{\delta\theta \ll 1/\Gamma_1} (seeding pitch angle diffusion) and large ones (\teq{\delta\theta \gtrsim 1/\Gamma_1}) using several parameters. These deflections can be viewed as a mathematical discretization of charge trajectory segments in electromagnetic turbulence: the shorter the segment, the smaller \teq{\delta\theta} is. The scatterings are assumed to be quasi-elastic in the local fluid frame, an idealization that is usually valid because in blazar jets and other astrophysical systems the flow speed far exceeds the Alfv\'{e}n speed (true for the models in Table~1), and contributions from stochastic second-order Fermi acceleration are small. In this paper, the focus will be on small angle \teq{\delta\theta \ll 1/\Gamma_1} domains, where {\it pitch angle diffusion} is realized; results for large angle scattering are illustrated in \citet{EJR90} and \citet{SBS07}. The simulation assumes that the complicated plasma physics of wave-particle interactions can be described by a simple scattering relation for particles, viz. \begin{equation} \lambda_{\parallel} \; =\; \lambda_1 \, \dover{\rho_1}{\rho} \left ( \dover{r_g}{r_{g1}} \right )^\alpha \;\equiv\; \eta_1\, r_{g1} \left( \dover{p}{p_1} \right)^{\alpha} \;\; , \quad \kappa_{\parallel} = \lambda_{\parallel} v/3\ , \label{eq:mfp_alpha} \end{equation} where \teq{v=p/m} is the particle speed in the local upstream or downstream fluid frame, \teq{r_g =pc/(QeB)} is the gyroradius of a particle of charge $Qe$, and $\rho$ is the plasma density, with a far upstream value of \teq{\rho_1}. Here, \teq{\lambda_\parallel} (\teq{\kappa_{\parallel}}) is the mean free path (spatial diffusion coefficient) in the local fluid frame, parallel to the field {\bf B}, with \teq{\lambda_\parallel\gtrsim r_g} being a fundamental bound (Bohm limit) for physically meaningful diffusion. Scalings of the momentum \teq{p_1=mu_{1x}=m\beta_{1x}c} and the mean free path \teq{\lambda_1= \eta_1 r_{g1}} for \teq{r_{g1} =p_1c/(QeB_1)} are introduced to simplify the algebra in Eq.~(\ref{eq:mfp_alpha}). In the local fluid frame, the time, \teq{\delta t_f}, between scatterings is coupled \citep{EJR90} to the mean free path, \teq{\lambda_{\parallel}}, and the maximum scattering (i.e. momentum deflection) angle, \teq{\delta\theta} via \teq{ \delta t_f\approx \lambda_{\parallel} \,\delta\theta^{2}/(6v)} for particles of speed \teq{v\approx c}. Scattering according to Eq.~(\ref{eq:mfp_alpha}) is equivalent to a kinetic theory description \citep[e.g.,][]{FJO74,Jokipii87} where the diffusion coefficients perpendicular to (\teq{\kappa_{\perp}}) and parallel to (\teq{\kappa_{\parallel}}) the local field vector {\bf B} are related via \begin{equation} \kappa_{\perp} \; =\; \dover{\kappa_{\parallel} }{1 + (\lambda_{\parallel}/r_g)^2 } \quad . \label{eq:kinetic_theory} \end{equation} The parameter \teq{\eta \equiv \lambda_{\parallel}/r_g \propto p^{\alpha-1}} then characterizes the ``strength" of the scattering and the importance of cross-field diffusion: when \teq{\eta \sim 1} at the Bohm diffusion limit, \teq{\kappa_{\perp} \sim \kappa_{\parallel}} and particles diffuse across magnetic field lines nearly as quickly as along them. This Bohm case corresponds to extremely turbulent fields, whose fluctuations satisfy \teq{\delta B/B\sim 1}. Each scattering event is an elastic deflection of the fluid frame momentum vector {\bf p} through angle \teq{\sim \delta\theta}, so that the number of deflections constituting \teq{\lambda_\parallel}, a large angle deflection scale (i.e., a turnaround through angle \teq{\sim \pi /2}), is proportional to \teq{(\delta\theta)^{-2}}. Values in the range \teq{1/2 \lesssim \alpha \lesssim 3/2} for this diffusion index are inferred from observations of turbulence in disparate sites in the solar wind, and also from hybrid plasma simulations, a context that will be discussed later in Section~\ref{sec:discussion}. We remark here that the story that will unfold in this paper is that blazar jets may present a picture with some similarities to the solar wind environment, with higher values of \teq{\alpha} in some cases that are associated with the large dynamic ranges of momenta and mean free paths in their relativistic jets. \begin{figure} \centerline{\includegraphics[width=9.5cm]{mn16bbs_fig1.eps}} \vspace*{-5pt} \caption{Particle distributions \teq{n_s(\gamma\beta ) \equiv m_s c\, dn_s/dp} (normalized; \teq{s=e,p}) in momentum space for acceleration simulation runs in the small angle scattering limit, corresponding to strong mildly-relativistic shocks of upstream flow speed \teq{\beta_{1x}\equiv u_{1x}/c =0.71}. Here the de Hoffmann-Teller (HT) frame upstream flow speed was set at \teq{\betaoneHT =0.84=\beta_{1x}/\cos\ThetaBfone}, with \teq{\ThetaBfone \approx 32.3^{\circ}} being the upstream field obliquity to the shock normal in the HT frame. Distributions are displayed for six different forms for the momentum dependence of the diffusive mean free path \teq{\lambda\equiv\lambda_{\parallel}}, namely \teq{\lambda/r_g \propto p^{\alpha -1}} with \teq{\alpha -1 = 0, 1/2} and \teq{2}, as labelled --- see Eq.~(\ref{eq:mfp_alpha}) and associated discussion. The shock velocity compression ratio was fixed at \teq{r=u_{1x}/u_{2x}=3.71}, and the upstream temperature corresponded to a sonic Mach number of \teq{{M}_{\rm S}\sim 4}. See Fig.~10 of \citet{SB12} for more \teq{\alpha =1} cases. At high momenta \teq{p\gg p_1}, many of the distributions are close to the flat \teq{n_s(\gamma\beta )\propto 1/(\gamma\beta )} asymptote that is highlighted in dark green. \label{fig:accel_dist}} \end{figure} In relativistic shocks, the distribution functions of accelerated charges are sensitive to the choices of both the \teq{\eta_1=\lambda_1/r_{g1}} and \teq{\alpha} diffusion parameters. This property is illustrated in Fig.~\ref{fig:accel_dist} for the mildly-relativistic shock domain that is germane to our blazar study. This sensitivity is addressed at length in \citet{SB12}, where the \teq{\alpha = 1} restriction was adopted for simplicity, so that \teq{\lambda_{\parallel} \propto r_g} and a single value of \teq{\eta} applies for all particle momenta. In subluminal shocks, for which \teq{u_{1x}/\cos\ThetaBfone <c}, i.e., those where a de Hoffman-Teller (HT) frame \citep{HT50} can be found, distributions \teq{dn/dp \propto p^{-\sigma}} generally possess indices in the range \teq{1 < \sigma < 2.5}. The HT frame is that where the shock is at rest and the fluid flows along {\bf B} at all positions upstream and downstream. Since \teq{\ThetaBfone} is the upstream field obliquity angle to the shock normal in the HT frame, and because mildly-relativistic shocks are {\it de rigueur} for structures in blazar jets, this angle can be quite modest: subluminal shocks are a real possibility --- see Table 1 for the \teq{\ThetaBfone} values used in our multiwavelength spectral fitting studies. In contrast, superluminal shocks that possess higher obliquities have steep distributions with \teq{\sigma > 2.5} when diffusion is not near the Bohm limit; in these circumstances, rapid convection downstream of the shock spawns high loss rates from the acceleration process \citep[e.g.,][]{BK90,ED04,SB12}. \subsubsection{Representative $e^-$ Distributions for Blazar Studies} Examples of both constant and momentum-dependent \teq{\lambda /r_g} are depicted in Fig.~\ref{fig:accel_dist}, where the power-law tails smoothly blend into the upper end of the thermal distribution. These distribution results were normalized to unit integrated density in all cases but one, with the \teq{\lambda /r_g = 100\, (p/p_1)^{1/2}} example being normalized to \teq{n_s=1/2} so as to aid clarity of the Figure. For much of the fairly restricted range of obliquities corresponding to \teq{u_1/\cos\ThetaBfone <c}, when \teq{\eta} is a constant for all momenta, the larger the value of \teq{\eta = \eta_1}, the smaller is \teq{\sigma}. This flattening trend is clearly represented by the comparison of the Bohm case (black) and the \teq{\eta =100} case (purple) in the Figure. In fact, when \teq{\eta\gg 1}, then indices \teq{\sigma\sim 1} are generally observed \citep{SB12} and the \teq{p\gg p_1} distribution is extremely flat. The origin of this behavior was found to be {\bf shock drift acceleration} (SDA), where charges with select gyrational phases incident upon the shock from upstream, are trapped and repeatedly reflected back upstream by the shock discontinuity. \citet{SB12} observed that the shock drift process is quickly disrupted when the turbulence levels increase and \teq{\eta} drops below around 100 or so. Shock drift acceleration is a phenomenon that has been well-studied in the context of non-relativistic heliospheric shocks \citep{Jokipii82,DV86}. It has also been observed in recent PIC simulations of non-relativistic electron-proton shocks \citep{Park12,Park15}, albeit primarily as a source of pre-injection of protons into the Fermi process that is manifested therein on large scales that sample the turbulent magnetic structures. The asymmetry of the drift electric fields straddling the discontinuity leads to a net energization during a gyrational reflection event. Successive episodes of upstream excursions and reflection at the shock precipitate efficient acceleration and postpone convective loss downstream for many shock interaction cycles. This extensive confinement to the shock layer automatically generates flat distributions with \teq{\sigma\sim 1-1.5} \citep{SB12}, in a sense analogous to the blazar jet shear layer studies of \citet{Ostrowski90}, and more or less commensurate with distribution indices realized in some magnetic reconnection models \citep[e.g.][]{Cerutti12}. Only when the field obliquity \teq{\ThetaBfone} is high enough that the shock is superluminal, and \teq{u_{1x}/\cos\ThetaBfone >c}, does the powerful convective action of the flow overwhelm reflection, and the acceleration process effectively shuts off \citep{SB12,BK90,ED04}, with the index \teq{\sigma} increasing rapidly above \teq{3-4}. Such distributions are nominally too steep to accommodate spectra from {\it Fermi}-LAT blazar data \citep[e.g., see][]{Abdo10a}, and so it is concluded that blazar jets must possess subluminal or ``marginally luminal'' relativistic shocks. This is the MHD shock phase space that is explored exclusively in this paper. These acceleration distribution results are extended here to cases where \teq{\eta (p)} is an increasing function of momentum \teq{p}, as in Eq.~(\ref{eq:mfp_alpha}), with representative \teq{\alpha >1} examples being displayed in Fig.~\ref{fig:accel_dist}. The subluminal shock parameters used therein are mostly those in Fig.~10, left panel, of \citet{SB12}. The distributions all display the generic trait of a dominant thermal population with a power-law tail that extends as high as the geometric scale of the diffusive acceleration zone permits: this environmental parameter is addressed in Section~\ref{sec:radmodels}. Each distribution possesses an injection efficiency from thermal into the Fermi process, i.e. for slightly supra-thermal momenta \teq{p\sim 0.3p_1-3p_1}, that reflects a combination of the values of \teq{\eta} and \teq{\alpha} in Eq.~(\ref{eq:mfp_alpha}). The ultimate power-law index at high momenta \teq{p\gg p_1} is determined by large values for \teq{\eta (p)} in all these simulation runs, i.e., realizing \teq{\sigma\sim 1} domains due to efficient shock drift acceleration in putatively weak turbulence. The key new spectral signature presented here is the appearance of ``flattening'' breaks in the high energy tail that are manifested when \teq{\eta (p)} begins to exceed values around \teq{10-30}; see the \teq{\eta_1=3}, \teq{\alpha =3/2} and \teq{\eta_1=3}, \teq{\alpha =2} cases in Fig.~\ref{fig:accel_dist}. Hence, \teq{\alpha >1} models can possess relatively inefficient injection at thermal energies, with dominant thermal components, but exhibit efficient acceleration at momenta \teq{p\gtrsim 100p_1}. The properties of plasma turbulence that might generate \teq{\alpha >1} circumstances are discussed below in Section~\ref{sec:discussion}, where the blazar spectral modeling that is presented in Section~\ref{sec:radmodels} is interpreted. \section[]{Multiwavelength Radiation Emission Modeling} \label{sec:radmodels} In this section, the focus is on modeling the broadband continuum emission of the blazars Mrk 501, BL Lacertae and AO 0235+164 using specific thermal plus non-thermal electron/pair distributions generated by the Monte Carlo technique. There are quite a few shock parameters that can be varied, so to maximize our insights, we keep the MHD Rankine-Hugoniot structure of the shock identical for all runs, and vary mostly the magnitude and direction of the magnetic field, and the diffusion parameters \teq{\eta_1} and \teq{\alpha} highlighted in Eq.~(\ref{eq:mfp_alpha}). Accordingly, flow speeds, compression ratios and upstream plasma temperatures are fixed so as to reduce the myriad of model possibilities. The objective is to assess how the global character of the radio to X-ray to gamma-ray data can provide insights into the nature of the accelerator and its plasma environment. To facilitate this goal, we concentrate solely on leptonic models \citep{MCG92,MK97,CB02,Boettcher13}, those with synchrotron emission predominantly at frequencies below \teq{10^{19}}Hz, and with inverse Compton emission dominating the gamma-ray signal. Hadronic models are also of great interest, but then introduce more parameters, and so don't tend to provide constraints that are as restrictive as those explored here. We remark that the critical synchrotron constraints on the values of \teq{\eta_1} and \teq{\alpha} will apply to either leptonic or hadronic scenarios for the gamma-ray signals. \subsection{Radiation Code and Geometry Essentials} To evaluate the light emission from the complete thermal+non-thermal electron distributions generated by the Monte Carlo simulation described in the previous subsection, we adopt radiation modules from the one zone blazar radiation transfer code of \citet{Boettcher13} and \citet{BC02}. This code treats synchrotron emission, synchrotron self-absorption at low radio frequencies, and also inverse Compton scattering of both synchrotron and external (nominally of disk/torus origin) seed photons. Bremsstrahlung is modeled in the code, but is usually insignificant in blazars: it is so for all our case study blazars, Mrk 501, BL Lac and AO 0235+164. The emission components are computed for isotropic distributions in the jet frame and then blue-shifted and Doppler-boosted to the observer frame using the bulk Lorentz factor \teq{\Gamma_{\rm jet}} of the jet, and the observing angle \teq{\theta_{\rm obs}} relative to its axis. Note that the distributions generated in the shock acceleration simulations are not isotropic in all reference frames. The dominant, compressed field is located downstream of a shock, and this is the zone where one anticipates the major contribution to both synchrotron and SSC emission originates. Since the diffusion process approximately isotropizes the electron distribution in the downstream fluid frame, this is technically the reference frame for which the radiation modules apply. Therefore the applicable boost would be by a Lorentz factor \teq{\Gamma \sim \Gamma_{\rm jet}} that is mildly-relativistic relative to the mean jet motion represented by \teq{\Gamma_{\rm jet}}. For simplicity, here we ignore this subtle distinction to reduce the number of model parameters, and set \teq{\Gamma\to \Gamma_{\rm jet}}. To complete the radiation modeling, the gamma-ray portion of the spectrum was then modified to treat line-of-sight attenuation of photons in propagation from source to Earth. This absorption derives from \teq{\gamma\gamma\to e^+e^-} pair opacity due to the intervening extragalactic background light (EBL) in the optical and infra-red; this is of stellar and dust origin in galaxies. Such \teq{\gamma\gamma} opacity due to the EBL has been extensively studied over the past two decades, and many models for its space density and its impact on VHE gamma-ray spectra of blazars and gamma-ray bursts exist. Here, we adopt the attenuation correction model of \citet{FRD10}. Note that the radiation models also treat \teq{\gamma\gamma\to e^+e^-} pair opacity internal to the blazar jets, which can lead to attenuation of $\gamma$-rays by IR and optical synchrotron fields. In practice, the jet Lorentz factors \teq{\Gamma_{\rm jet}} are chosen sufficiently high for all three blazars studied that the jets are internally transparent to this pair opacity. Opting for a {\bf one-zone radiative model} here is motivated by simplicity. In reality, one expects multi-zone structure to the emission region, encompassing gradients in the field magnitude, electron density and other system parameters. Such complexity is the natural inference from high-resolution imaging of jets in radio, optical and X-ray wavebands. Multi-zone constructions effectively introduce more modeling variables, not necessarily precipitating a gain in insight. The most noticable change introduced by upgrading to a multi-zone scheme is the broadening of the \teq{\nu F_{\nu}} peaks in both the synchrotron and inverse Compton components: this muting of the spectral curvature of these turnovers is a consequence of distributing the cooling rates for each of the processes over a broader range of values. The spatial structure conceived in a multi-zone model is a sequence of segments subdividing the jet where acceleration and emission are produced in each segment. For our one-zone construction, we consider just one such segment. The acceleration zone is the most confined portion of this region, being a planar slab of thickness \teq{2R_{\rm acc}} with a planar shock dividing the zone roughly in halves. In terms of the acceleration, since the flow is relativistic, most of the diffusive transport occurs in the downstream region, while the shock drift contribution straddles the shock and encroaches primarily into the upstream region. On larger scales, there is a radiation volume whose linear dimension \teq{R_{\rm rad}} perpendicular to the shock layer is defined by the approximate equality between the radiative cooling time (synchrotron and/or inverse Compton, whichever is the shortest), and the jet fluid convection time. Since the field is compressed by the shock, the downstream region is where the synchrotron emission is most prevalent, and if the gamma-rays are generated by the SSC mechanism, then here also is where the inverse Compton signal is the strongest. MHD turbulence will permeate the entire radiating volume, but is most likely more concentrated near the shock, i.e. in the acceleration zone. The reason for this is that turbulence is naturally generated in the shock layer by dissipation of the ballistic kinetic energy of the upstream flow as it decelerates, and this turbulence should abate with distance from the shock due to damping. Such is the nature of interplanetary shocks and their environs embedded in the solar wind. A schematic of this model geometry is given in Figure~\ref{fig:model_geom}, wherein the depicted turbulent fields are actually a two-dimensional projection of Alfv\'en-like fluctuations simulated using a Kolmogorov power spectrum of modest inertial range. An exponential damping of the fluctuations is imposed on scales \teq{\vert x\vert \gtrsim 5} for the purposes of illustration. \begin{figure} \vspace*{-10pt} \centerline{\includegraphics[width=9.0cm]{blazar_turb_fig.eps}} \vspace*{-20pt} \caption{Schematic of the blazar model geometry under consideration, consisting of a region proximate to the shock that is the acceleration zone within a slab of approximate thickness \teq{2R_{\rm acc}}, which is embedded in a much larger radiation zone of size \teq{R_{\rm rad}} (not to scale). The MHD background magnetic field structure of the shock is indicated by the blue straight line segments. Superposed on this is a turbulent field, signified by the red field line projections that are computed from a Kolmogorov power spectrum of finite inertial range spanning around a decade in wavenumber \teq{k}. This perturbation is exponentially damped on a scale of \teq{R_{\rm acc}} in this depiction so as to highlight confinement of turbulence to the shock layer. \label{fig:model_geom}} \end{figure} In the absence of information on the physical scale of the acceleration zone, the Monte Carlo simulations do not describe the high-energy cut-off of the electron distribution. We therefore extend the initial electron spectra like those in Fig.~\ref{fig:accel_dist} out to a maximum energy \teq{\gamma_{\rm max}mc^2 \approx \pmax c}. Such an extrapolation is acceptable because the asymptotic power-law has been realized in the Monte Carlo simulation results with the displayed dynamic range in momenta. This maximum energy is constrained in two ways: \vspace{-0pt} \begin{enumerate} \item particles will not be accelerated beyond an energy for which the (synchrotron plus inverse Compton) radiative cooling time scale, \teq{t_{\rm rad} = 3 m_e c^2 / (4\gamma\, c \, \sigt \, {\cal U} )}, where \teq{{\cal U} = \UB + {\cal U}_{\rm rad}} is the sum of the jet frame energy densities in the magnetic field and the photon field, is shorter than the acceleration time scale, \teq{t_{\rm acc} = \gamma m_e c \, \eta( p) / (e B)}, which will be discussed in greater detail below; \item particles will not continue to be accelerated once their energy establishes a diffusive mean free path \teq{\lambda_{\parallel} = \eta(p) r_g(\gamma)\equiv \eta_1\, (p/p_1)^{\alpha}} in Eq.~(\ref{eq:mfp_alpha}) that exceeds the size \teq{R_{\rm acc}} of the acceleration zone. Such energetic charges are assumed to escape upstream or downstream from the shock environs. \end{enumerate} Here \teq{\gamma = \{ 1+ (p/mc)^2 \}^{1/2} \approx p/(mc)} for energetic leptons. With \teq{\gamma_{\rm max}} determined as the smaller of the limiting values from the two constraints above, the high-energy portion of the electron distribution {\it injected} at the acceleration site can be written as \teq{n_{\rm acc} (\gamma) \propto \gamma^{-\sigma (p )} \, \exp\left(-\gamma / \gamma_{\rm max} \right)} for \teq{\gamma\gg 1}, where \teq{\sigma} is the high-energy index of the {\it simulated} Monte Carlo distribution \teq{dn/dp \propto p^{-\sigma}} of accelerated particles. The radiation module then uses as input an injected distribution, of thermal particles plus charges accelerated up to \teq{\gamma_{\rm max}} that are redistributed over the entire radiation zone, from which they escape on a time scale \teq{t_{\rm esc} = \eta_{\rm esc}\, R_{\rm rad}/c}. This escape parameter sets the normalization of the effective equilibrium density of electrons in concert with the jet luminosity. If \teq{t_{\rm rad} < t_{\rm esc}} when \teq{\gamma = \gammax}, a break in the electron spectra is expected at a Lorentz factor \begin{equation} \gamma_{\rm br} \; \sim\; \dover{3 m_e c^2 }{4 \eta_{\rm esc}\, R_{\rm rad} \, \sigt \, {\cal U} } \label{eq:gamma_br_cool} \end{equation} where the radiative cooling time equals the escape timescale. At this \teq{\gamma}, the distribution steepens by an index increment \teq{\Delta\sigma = 1}. As a reminder, \teq{{\cal U} = \UB + {\cal U}_{\rm rad}} is the total field plus radiation energy density in the co-moving frame of the jet, and this result applies for inverse Compton scattering in the Thomson limit, a fairly representative situation. This cooling break at \teq{\gamma_{\rm br} <\gamma_{\rm max}} is in addition to acceleration flattening breaks illustrated in Fig.~\ref{fig:accel_dist}, and generates a corresponding break in the photon spectrum by an index of \teq{1/2}, the well-known signature of the transition from slow to fast cooling at high Lorentz factors. To couple the emission signal to the acceleration and cooling information, the partially-cooled electron distributions are normalized to a total kinetic luminosity of electrons in the blazar's two jets, \begin{equation} L_{\rm kin} \; =\; 2\pi R_{\rm rad}^2 c \, \Gamma^2_{\rm jet} \, m_e c^2 \int_{1}^{\infty} n_e (\gamma) \, \gamma \, d\gamma \label{eq:Lkin} \end{equation} which is a free input parameter for fitting purposes. The factor of 2 accounts for both the bright approaching jet that we observe, and the vastly fainter receding one. As most of acceleration simulations illustrated in this paper produce distributions that are extremely flat, approaching \teq{n_e(\gamma )\propto \gamma^{-1}} (i.e., \teq{\sigma\sim 1}), a substantial and sometimes dominant contribution to \teq{L_{\rm kin}} comes from the non-thermal electrons up to the cooling break \teq{\gamma_{\rm br}}. This is the case for the distributions exhibited in Fig.~\ref{fig:blazar_edist} below. Observe that since changes in the electron spectrum and density will lead to changes in the co-moving synchrotron photon field, our code determines the equilibrium electron distribution through an iterative process. It starts by considering only the magnetic \teq{\UB} and external radiation \teq{{\cal U}_{\rm rad}} energy densities to determine a first-order equilibrium \teq{e^-} distribution. Then \teq{\UB} is used to evaluate the co-moving synchrotron photon field, which is added to the energy densities to re-determine the electron cooling rates and to re-evaluate the break and maximum electron energies. The process is repeated until convergence is achieved, in just a few iterations. The magnetic field is specified by means of a magnetic partition fraction \teq{\epsilonB \equiv \LB /L_{\rm kin}}, where \begin{equation} \LB\; =\; \pi R_{\rm rad}^2 \, c \, \Gamma_{\rm jet}^2 \, \UB \label{eq:mag_power_jet} \end{equation} is the power carried by the magnetic field along the jet, partly in the form of Poynting Flux. Two powers of \teq{\Gamma_{\rm jet}} appear here, defining the energy boost and the time dilation factors. From the perspective of the radiation modules, the field is effectively assumed to be tangled in the co-moving jet frame so that isotropic emissivities for both synchrotron and inverse Compton processes are employed. In reality, the acceleration considerations capture both turbulent and quasi-laminar fields at different relative levels on different spatial scales. Furthermore, the shocks are moving at mildly-relativistic speeds in the jet frame, thereby engendering ``beaming'' anisotropies in the electron distribution; these will be neglected for the radiation modeling since a number of the other parameters have a more important impact on the results. The radiative output from the equilibrium electron distribution is evaluated using the radiation transfer modules of \citet{Boettcher13}. Two parameters derived from the spectroscopic models define the characteristic rates for acceleration processes. These are the non-relativistic electron gyrofrequency, \teq{\omegaB}, and the electron plasma frequency \teq{\omega_p}: \begin{equation} \omegaB \; =\; \dover{eB}{m_ec} \quad ,\quad \omega_p \; =\; \sqrt{ \dover{4\pi n_e e^2}{m_e} }\quad . \label{eq:wB_wp} \end{equation} The gyrofrequency specifies the rate of acceleration in gyroresonant plasma mechanisms such as shock drift acceleration, while the plasma frequency governs the rate of energization in electrodynamic mechanisms such as magnetic reconnection. The ratio of these two frequencies is captured in the non-relativistic plasma magnetization parameter \begin{equation} \sigma \; =\; \dover{\omegaB^2}{\omega_p^2} \; =\; \dover{B^2}{4\pi n_e m_ec^2} \quad , \label{eq:sigma_param} \end{equation} familiar in kinetic plasma and MHD simulation studies. The electron density \teq{n_e} is obtained from the kinetic luminosity \teq{L_{\rm kin}} in Eq.~(\ref{eq:Lkin}). We highlight these parameters here as they will facilitate the interpretative elements later on. \subsection{Generic Features of Synchrotron plus Inverse Compton Blazar One-Zone Models} Before progressing to the blazar modeling, it is instructive to provide an outline of the central elements of the multiwavelength fitting diagnostics of the plasma environment in blazar jets. The key constraints derived are representative values of \teq{\eta (p) = \lambda_{\parallel}/r_g} at the low injection momenta, \teq{p\sim p_1\sim mu_1}, and at the maximum momentum \teq{\pmax} in the acceleration zone. These two mean free path parameters couple via the index \teq{\alpha} of the diffusion law in Eq.~(\ref{eq:mfp_alpha}). Foremost in our considerations is the value of \teq{\eta (\pmax )}, which far exceeds unity for our chosen blazars. The origin of this property is that large \teq{\eta} are required to fit the frequency of the synchrotron \teq{\nu F_{\nu}} peak just below the spectral turnover of this emission component. For strong cooling in the acceleration zone, the turnover is created when the cooling rate \begin{equation} -\, \dover{d\gamma}{dt}\biggl\vert_{\rm cool} \; =\; \dover{4 \sigt}{3m_ec} \, {\cal U} \, \gamma^2 \label{eq:cool_rate} \end{equation} in the jet frame is approximately equal to the electron acceleration rate for the diffusive Fermi process \begin{equation} \dover{d\gamma}{dt}\biggl\vert_{\rm acc} \;\sim\; \left( \dover{u_{1x}}{c} \right)^2 \, \dover{eB}{\eta\, m_ec} \quad\hbox{for}\quad \eta (p) \; =\; \eta_1 \left( \dover{p}{p_1} \right)^{\alpha-1}\; . \label{eq:acc_rate} \end{equation} Note that this acceleration rate is an approximate scale, and more precise forms that isolate upstream and downstream contributions, and field obliquity influences can be found in reviews like \citet{Drury83} and other papers such as \citet{Jokipii87}. Now introduce a cooling parameter \teq{\epsilon_{\rm syn} = \UB/{\cal U}} that represents the fractional contribution of the synchrotron process to the electron cooling in the jet frame. The contribution \teq{{\cal U}_{\rm rad}} to \teq{{\cal U}} is determined from the comoving radiation field (radio-to-gamma-ray) from the radiation codes as outlined above. This synchrotron parameter can be approximately expressed in terms the ratio of optical/X-ray ``peak luminosities'' to the pseudo-bolometric luminosity: \teq{\epsilon_{\rm syn}\approx L_{O-X}/(L_{O-X} + L_{\gamma})}. For Mrk 501, this has a value of about \teq{\epsilon_{\rm syn} \sim 0.5-0.6}, as can be inferred from the broadband spectrum in Fig.~\ref{fig:MW_spec_Mrk501}. In contrast, for the synchrotron-dominated blazar BL Lacertae, \teq{\epsilon_{\rm syn} \sim 0.8-0.9}, which is somewhat higher than the values inferred for Mrk 421 and Mrk 501. For the flare state of AO 0235+164, the strong gamma-ray signal sets \teq{\epsilon_{\rm syn} \sim 0.25}. The acceleration/cooling equilibrium then establishes a maximum electron Lorentz factor at \begin{equation} \gammax \; =\; \left( \dover{9\epsilon_{\rm syn}}{4\eta_1} \, \Bigl( \dover{u_{1x}}{c} \Bigr)^2 \, \dover{e}{B\, r_0^2} \right)^{1/(1 + \alpha )} \quad , \label{eq:gamma_max} \end{equation} assuming that \teq{p_1\sim m_ec} in a mildly relativistic shock with \teq{u_{1x}\sim c}. Here \teq{r_0 = e^2/m_ec^2} is the classical electron radius. If \teq{\epsilon_{\rm syn}\sim 1}, then synchrotron emission dominates the cooling of particles, and when \teq{\alpha =1}, corresponding to \teq{\eta} being independent of electron momentum, the well-known result that \teq{\gammax \propto B^{-1/2}} follows. Since \teq{e/Br_0^2 = B_{\rm cr}/(\fsc B)} for a fine structure constant \teq{\fsc =e^2/(\hbar c)} and \teq{B_{\rm cr} = m_e^2 c^3/e\hbar \approx 4.41\times 10^{13}}Gauss as the Schwinger field, the scale of \teq{\gammax} in the range \teq{\sim 10^3-10^6} is established for blazars with fields \teq{B\sim 1-10}Gauss in the jet frame, depending on the value of \teq{\alpha}. This is readily seen by recasting Eq.~(\ref{eq:gamma_max}) in the form \begin{equation} \gammax \; \sim\; \sqrt{ \dover{2 {\cal E}_s(\alpha )}{3} } \left( \dover{6\times 10^{15}}{\eta_1\, B} \right)^{1/(1 + \alpha )} \quad , \label{eq:gammax_syn} \end{equation} for \teq{B} in units of Gauss, and \begin{equation} {\cal E}_s(\alpha ) \; =\; \dover{3}{2} \left( \dover{9 \epsilon_{\rm syn}}{4} \, \Bigl[ \dover{u_{1x}}{c} \Bigr]^2 \right)^{2/(1 + \alpha )} \label{eq:calE_s_def} \end{equation} being of the order of unity for blazar shocks. Increasing \teq{\alpha} above unity lowers \teq{\gammax} because the acceleration becomes less efficient (slower) in competing with cooling. The synchrotron turnover/cutoff energy \teq{E_{\rm syn}\propto \gammax^2 B} corresponding to \teq{\gammax} is \begin{equation} E_{\rm syn} \; \sim\; \dover{\delta_{\rm jet}\, {\cal E}_s(\alpha )}{\eta_1\, (1+z)} \, \left( \dover{\fsc B}{B_{\rm cr}} \eta_1 \right)^{(\alpha -1)/(\alpha + 1)} \dover{m_ec^2}{\fsc} \quad , \label{eq:Esyn_max} \end{equation} The blueshift factor \teq{\delta_{\rm jet} = 1/[\Gamma_{\rm jet}( 1-\beta_{\rm jet}\cos\theta_{\rm obs} )]} due to Doppler beaming has been included, where \teq{\theta_{\rm obs}\ll 1} is the observer's viewing angle with respect to the bulk velocity of the jet in the emission region. Also, \teq{z} is the redshift of the distant blazar, so that \teq{E_{\rm syn}} represents an observer's measurement of the energy. For the Bohm limiting case of \teq{\eta (\gammax )=1}, imposing both \teq{\alpha=1} and \teq{\eta_1=1}, and setting \teq{\delta_{\rm jet} = 1}, the \teq{\Gamma_1 =1} and \teq{u_{1x}=c} evaluation yields \teq{E_{\rm syn}\sim 100} MeV, independent of the magnetic field strength. This widely-known result was highlighted in \citet{DeJager96} for considerations of $\gamma$-ray emission at relativistic pulsar wind nebular shocks, but was derived much earlier by \citet{GFR83} in the context of AGN. This special result holds provided that the acceleration process is gyroresonant, which is the prevailing paradigm for both non-relativistic and relativistic shocks. For Mrk 501 and our other case study blazars, \teq{\gammax} becomes weakly dependent on \teq{B}, and the turnover energy \teq{E_{\rm syn}} quickly drops below the \teq{m_ec^2/\fsc} scale. For representative jet fields \teq{B\sim 1}Gauss, one quickly estimates using Eq.~(\ref{eq:Esyn_max}) that for the synchrotron \teq{\nu F_{\nu}} peak energy \teq{E_{\rm syn}} to appear in the optical, one requires values \teq{\alpha \sim 2.5-3}, and often also values \teq{\eta_1\gtrsim 100}. This is a central feature of our case studies below: in order to move the synchrotron peak into either the X-ray or optical band, large values of \teq{\eta (\gammax )} are required. In such circumstances, the acceleration process may not be gyroresonant, a realization that we will discuss in Sec.~\ref{sec:discussion}. The compelling need for \teq{\eta (\gammax )\gg 1} in the shock acceleration interpretation of blazar activation was first emphasized by \citet{IT96}, who chose a {\it momentum-independent} \teq{\eta \sim 10^5} when exploring multi-wavelength modeling of Mrk 421 spectra. In our more versatile construct here, \teq{\alpha} is not fixed to unity. To realize such large \teq{\eta} at \teq{\gammax}, without forcing \teq{\eta_1} to similar large values that would suppress efficient injection of charges from the thermal population into the Fermi mechanism \citep{SB12}, values \teq{\alpha > 1} are necessitated. Synchrotron self-Compton models possess an additional constraint on system parameters via the observer's frame energy of the SSC peak \teq{\ESSC} in the hard gamma-rays. This is just positioned according to the energy gained by synchrotron photons of around the turnover energy \teq{E_{\rm syn}} subjected to inverse Compton scattering. For the Thomson limit, this energy enhancement ratio is \teq{4\gammax^2/3}, so that \begin{equation} \gammax \;\sim\; \sqrt{\dover{3 \ESSC}{4E_{\rm syn}} } \quad . \label{eq:gammax_SSC} \end{equation} In conjunction with Eq.~(\ref{eq:Esyn_max}), this restricts the values of \teq{B} and \teq{\alpha}. This constraint on \teq{\gammax} is not applicable for external inverse Compton models where the target photon field is not the synchrotron population, but perhaps originating from a proximate disk; this will actually be the case for our blazars BL Lac and AO 0235+164. It is also modified somewhat when the upscattering samples the Klein-Nishina regime, a domain that is barely encroached upon in most of our models, since they satisfy \teq{4\gammax E_{\rm syn}/(\delta_{\rm jet}\, m_ec^2) \lesssim 1}. This ensemble of diagnostics for our model parameters is summarized in Fig.~\ref{fig:spec_schematic}. In it, an array of observational data from the multiwavelength campaign on Mrk 421 of January -- May 2009 is depicted: this serves as a selection from Fig.~8 in \citet{Abdo11b}. Mrk 421 was chosen for the Figure as it will not be studied here, though we note that its broadband spectrum is qualitatively similar to that of our case study blazar Mrk 501. Schematic model synchrotron and SSC spectra are exhibited in Fig.~\ref{fig:spec_schematic}, though specifically as would be computed in the jet rest frame. This then permits identification of the Doppler factor \teq{\delta_{\rm jet}} via measurement of the frequency ratio between peaks in the observed and modeled spectra for each of the emission components. Additionally, the ratio of observed to theoretical peak heights measures the flux enhancement \teq{\delta_{\rm jet}^4} between model and data, noting that this element of the schematic is not to scale. The separation of the component peaks (observed or modeled) defines the inverse Compton scattering enhancement factor \teq{4\gammax^2/3} (Thomson limit), and so constrains the value of \teq{\gammax} approximately according to Eq.~(\ref{eq:gammax_SSC}). For the illustrated case, \teq{\gammax\sim 10^4} and the synchrotron peak energy is around \teq{10^{-4}m_ec^2}, so that the Compton scattering near the \teq{\sim}TeV peak is marginally in the Klein-Nishina regime. These are standard paths to computing spectral fitting parameters for SSC models of blazars. What is more unique to the present study is the determination of the diffusion parameter \teq{\eta (\gammax )} by comparing the synchrotron model peak energy with the fundamental bound \teq{m_ec^2/\fsc}, as depicted in Fig.~\ref{fig:spec_schematic}. This is essentially inverting Eq.~(\ref{eq:Esyn_max}) to yield a combination of \teq{\eta_1} and \teq{\alpha} that is captured in Eq.~(\ref{eq:eta1_alpha_reln}) below. These two diffusion parameters will form a central focus for our case studies below. \begin{figure} \vspace*{-10pt} \centerline{\hskip 10pt\includegraphics[width=9.5cm]{Mrk421_MW_schematic.eps}} \vspace*{-5pt} \caption{ Generic construction of multiwavelength synchrotron self-Compton (SSC) spectral modeling for blazars in the \teq{\nu F_{\nu}} representation. The illustration includes a selection of data for Mrk 421 published in Fig.~8 of \citet{Abdo11b} for a campaign during the first half of 2009. This includes the ochre band signifying various radio detections, purple squares for optical data from different facilities, yellow-green (Swift-XRT) and green circles (Swift-BAT) for X-rays, red dots for {\it Fermi}-LAT measurements, and black triangles for MAGIC TeV-band data. The solid curves are the schematic model spectra applicable to the comoving jet frame, with radio-to-optical/UV (orange) curve constituting the synchrotron component, and the X-ray-to-\teq{\gamma}-ray (blue) curves denoting the SSC spectrum. The frequency offsets between the model peak and the data peak for these two components (marked by solid and dashed green vertical lines, respectively) approximately define the value of the jet Doppler factor \teq{\delta_{\rm jet}}. The separation of the peaks of the synchrotron and SSC components is by a factor of the order of \teq{\gammax^2}. In the jet frame, the synchrotron model peak lies at an energy that is a factor \teq{\eta (\gammax )} below the fundamental bound of \teq{m_ec^2/\fsc} (see text). \label{fig:spec_schematic}} \end{figure} Observe the appearance of a thermal SSC component in the X-rays due to bulk Comptonization of shock-heated thermal electrons; this actually emerges smoothly above the non-thermal inverse Compton contribution at flux levels below those of the plot scale. A similar thermal synchrotron component is not exhibited (below 1 GHz) in the Figure, since it would be suppressed by synchrotron self-absorption in the radio. Each component also exhibits a cooling break, at \teq{E_{\rm br,syn} \sim 3\times 10^{12}}Hz and \teq{E_{\rm br,SSC} \sim 3\times 10^{21}}Hz, respectively, where the photon spectrum steepens by an index of \teq{1/2} at higher frequencies. The approximate value of \teq{\gamma_{\rm br}} can be deduced since \teq{\gamma_{\rm br}^2} roughly represents the ratio of the energy of the break to the thermal peak energy in the SSC component. This follows from the presumption that the shock heated electrons are at most mildly-relativistic (true for mildly-relativistic blazar shocks), and one would conclude that \teq{\gamma_{\rm br}\sim 10^2} from the depiction in Fig.~\ref{fig:spec_schematic}. The precise value of \teq{\gamma_{\rm br}} in Eq.~(\ref{eq:gamma_br_cool}) yields a measure of the relative sizes of the acceleration and cooling zones. It is evident that constraints imposed by the energies of the synchrotron and SSC \teq{\nu F_{\nu}} peaks on jet environmental parameters lead to strong couplings between them. The most obvious of these is the calibration of the magnetic field strength using the ratio of the SSC and synchrotron peak frequencies. Combining Eqs.~(\ref{eq:gammax_syn}) and~(\ref{eq:gammax_SSC}), one arrives at the coupling between \teq{B} and \teq{\alpha}: \begin{equation} B \;\sim\; \dover{6\times 10^{15}}{\eta_1} \left( \dover{8E_{\rm syn}}{9 \ESSC} \right)^{(1 + \alpha )/2} \label{eq:B3_SSC_syn_ratio} \end{equation} for \teq{{\cal E}_s(\alpha )\sim 1}. Evidently, for synchrotron peaks in the X-rays (HBLs), fields \teq{B\sim 1}Gauss require \teq{\alpha \sim 1-1.5}, whereas, for blazars with synchrotron peaking in the optical, higher values of \teq{\alpha} are needed to establish similar jet fields. The ratio of the SSC to synchrotron peak luminosities also provides a modest constraint on \teq{B} via the \teq{\epsilon_{\rm syn}} parameter. The second diffusion/acceleration parameter, \teq{\eta_1}, can be introduced by inverting Eq.~(\ref{eq:Esyn_max}), again for the case \teq{{\cal E}_s(\alpha )\sim 1}. Set \teq{\nu_{\hbox{\fiverm S,14}}} to be the synchrotron peak frequency in the comoving frame in units of \teq{10^{14}}Hz, and \teq{\nu_{\hbox{\fiverm SSC,24}}} as the SSC peak frequency in the jet frame in units of \teq{10^{24}}Hz. Eliminating \teq{B} using Eq.~(\ref{eq:B3_SSC_syn_ratio}), one can then ascertain the fiducial relationship between \teq{\eta_1} and \teq{\alpha} for representative observed blazar synchrotron and SSC peak energies: \begin{equation} \eta_1^{2/(\alpha +1)} \;\sim\; \dover{1.5 \times 10^{-2} \delta_{\rm jet}}{(1+z)\, \nu_{\hbox{\fiverm SSC,24}}} \, \left( 10^{-10}\, \dover{\nu_{\hbox{\fiverm S,14}} }{\nu_{\hbox{\fiverm SSC,24}} } \right)^{(\alpha -3 )/2}\quad . \label{eq:eta1_alpha_reln} \end{equation} With \teq{\delta_{\rm jet}\sim 10-30}, this clearly suggests that \teq{\alpha >2} cases are needed for synchrotron peaks to appear in the optical, and concomitantly \teq{\eta_1 \gtrsim 30-100}. The physical implications of such diffusion parameters will be addressed in Section~\ref{sec:discussion}. Lower values of \teq{\alpha} and \teq{\eta_1} can be entertained if the synchrotron peak appears in the X-rays, as it does for Mrk 501. \subsection{Case Study: Mrk 501} An obvious first choice for study is the well-monitored BL Lac blazar Markarian 501, located at a redshift of \teq{z=0.034}. Multiwavelength campaigns are readily available for this source \citep[e.g.,][]{Catanese97,Albert07}, thereby eliminating problems associated with non-contemporaneous data in different wavebands. The more recent campaign conducted during the March--July 2009 epoch serves amply to illustrate the advances in understanding offered by combining shock acceleration simulation results with radiation emission codes. The broadband spectral character of Mrk 501 differs only modestly from that depicted in Fig~\ref{fig:spec_schematic} for its sister blazar Mrk 421, and is also fairly similar to the SED of more distant blazars such as PKS 2155-304 \citep[at \teq{z=0.116}; see][]{Aharonian09}, yet it is still suitable to an SSC model construction with an added optical component. An extensive summary of the mid-March--early August, 2009 campaign involving the VERITAS and MAGIC telescopes above 100 GeV, {\it Fermi}-LAT in the 20MeV to 300 GeV energy window, RXTE-PCA (2--60 keV) and Swift XRT+BAT (0.3-150 keV) in the X-rays, combined with UV (Swift-UVOT), optical and radio measurements, is provided in \citet{Abdo11a}. For most of this epoch, the source was in a low/moderate state, for which broadband spectral data are depicted in Fig.~\ref{fig:MW_spec_Mrk501}. During this campaign, in the short window May 1--5, 2009, Mrk 501 underwent a sizable flare in gamma-rays, evincing a high state at around 5 times the Crab nebula flux in the VERITAS/MAGIC/{\it Fermi}-LAT bands. This was an enhancement of a factor of 3-5 over the quieter portions of the campaign, and received considerable attention in \citet{Abdo11a} and also \citet{Aliuetal16}. The flare exhibited a spectrum that was much harder above \teq{10^{23}}Hz than for the extended state. To aid clarity, the VERITAS data were not depicted in Fig.~\ref{fig:MW_spec_Mrk501}, being similar to the MAGIC points above \teq{10^{25}}Hz. Moreover, the multi-wavelength data represent the long-term average over the 4.5 month interval, and for radio, optical, UV and X-rays, were taken from the spectrum in Fig.~8 of \citet{Abdo11a}. The gamma-ray points, for {\it Fermi}-LAT and MAGIC, omit an interval of 30 days that brackets the strong flaring episode, and are taken from Fig.~9 of \citet{Abdo11a}. Note that the units for the $y$-axis in this and subsequent multi-wavelength spectral figures employ the Jansky-Hertz choice familiar to radio astronomers, and can be converted to the form in Fig.~\ref{fig:spec_schematic} that is often used by high energy astrophysicists via 1 Jy Hz \teq{=10^{-23}} erg cm$^{-2}$ sec$^{-1}$. \begin{figure} \vspace*{-10pt} \centerline{\hskip 10pt\includegraphics[width=9.5cm]{Mrk501_spec_draft.eps}} \vspace*{-5pt} \caption{Multiwavelength \teq{\nu F_{\nu}} spectra (points), together with model fits as described in the text, for the extended March-July 2009 monitoring of the blazar Markarian 501. The campaign data are taken from Figures 8 and 9 of Abdo et al. (2011), and constitute a ``low'' state with gamma-ray data bracketing the flare of early May, 2009 being omitted (see text). The gamma-ray detections and upper limits are from {\it Fermi}-LAT (purple triangles) and MAGIC (black circles), and the X-ray points are {\it Swift}-BAT (muddy orange squares), {\it Swift}-XRT (light blue triangles) and RXTE-PCA data (open red circles). Optical, UV and radio measurements are detailed in Abdo et al. (2011); since a variety of radio fluxes were recorded for various regions much larger than the compact X-ray/$\gamma$-ray zones, they are marked by a representative band, and were not fit. The broadband models consist of a synchrotron component (dashed green curve) up to the X-ray band, an SSC component in the X-rays and gamma-rays (dashed red curve), and a separate thermal host galaxy emission component (dotted black). The full jet model spectra for the extended low state includes a correction \citep{FRD10} for \teq{\gamma\gamma} absorption by the extragalactic background light, and combined with the data, constrain the diffusion parameters to \teq{\eta_1=100} and \teq{\alpha =1.5} (see text). \label{fig:MW_spec_Mrk501}} \end{figure} In the SSC interpretation, the extended state is very slightly synchrotron-dominated, contrasting the May 1--5, 2009 high state, which is marginally inverse Compton-dominated; throughout, the synchrotron peak is positioned in the X-ray band. The low-variability UV/optical data is not attributed exclusively to the jet, but mostly to the host galaxy: at least 2/3 of the flux is believed \citep[e.g.,][]{Nilssonetal07} to originate in the host, and is herein modeled via a separate thermal component. The radio data were from various (mostly single-dish) facilities \citep[see the list in Table 1 of][]{Abdo11a} that generally do not resolve the jet, and accordingly provide upper bounds to the signal expected from the highly-variable $\gamma$-ray emission zone. For this reason, they are depicted in Fig.~\ref{fig:MW_spec_Mrk501} as a swath, positioned at the fluxes measured at the various radio frequencies. Using the radiation modeling protocols and coding described in Section~\ref{sec:radmodels}, multiwavelength spectra from radio to VHE gamma-rays were generated for a suite of particle distributions generated by the Monte Carlo shock acceleration simulation. The turbulence parameters \teq{\alpha} and \teq{\eta_1} were adjusted for the various simulation runs to hone in on a candidate ``best case,'' the spectral results for which are exhibited in Fig.~\ref{fig:MW_spec_Mrk501}. The spectra so generated are dominated by two components, synchrotron emission in the radio-to-X-ray band, and a synchrotron self-Compton signal in the gamma-ray range; the observer frame computations of these components at the source are isolated in this Figure as dashed curves. The overall spectrum at Earth is subjected to \teq{\gamma\gamma} attenuation from the line-of-sight EBL, as described above, and is depicted as the solid M/W curve in Fig.~\ref{fig:MW_spec_Mrk501} that is the nominal fit to the data; this attenuation is omitted in the SSC component depiction so as to illustrate its magnitude. Simulation and radiation model input parameters, and also those derived from the input ones, are listed in Table~1; some of them can be compared with model fitting parameters listed in Table~2 of \citet{Abdo11a}. Since in this broadband spectrum the thermal and non-thermal components are obviously distinct, the detailed shape of the host galaxy optical spectrum is irrelevant for our modeling goals. Therefore, no attempt was made to perform a detailed fit to the optical (e.g., with a spectral template for an elliptical galaxy); this component was simply represented with a separate blackbody spectrum for illustrative purposes. \begin{table*} \centering \begin{minipage}{140mm} \caption{Input and Derived Parameters for Blazar Models} \begin{tabular}{@{}lrcccc@{}} \hline \multicolumn{2}{c}{Parameter} & Mrk 501 & BL Lacertae & AO 0235+164 \\ Name & Symbol/units & Extended State & Extended State & High State \\ \hline \multicolumn{5}{c}{Jet+Source Parameters} \\ \hline Redshift & $z$ & 0.034 & 0.069 & 0.94 \\ Jet Lorentz factor & $\Gamma_{\rm jet}$ & 25 & 15 & 35 \\ Observing angle & $\theta_{\rm obs}$ & $2.29^{\circ}$ & $3.82^{\circ}$ & $1.7^{\circ}$ \\ Emission region size & $R_{\rm rad}$ [cm] & $1.2\times 10^{16}$ & $2.5\times 10^{15}$ & $1.0\times 10^{16}$ \\ $e^-$ injection luminosity\footnote{Expressed in Eq.~(\ref{eq:Lkin}), wherein the electron number density $n_e$ is normalized.} & $L_{\rm kin}$ [erg/sec] & $1.5\times 10^{44}$ & $6.7\times 10^{43}$ & $4.8\times 10^{46}$ \\ Escape time scale\footnote{The escape time is given by \teq{t_{\rm esc} = \eta_{\rm esc}\, R_{\rm rad}/c}.} & $\eta_{\rm esc}$ & 100 & 5 & 300 \\ Magnetic partition & \teq{\epsilonB \equiv \LB /L_{\rm kin}} & $3\times 10^{-4}$ & $0.5$ & 0.06 \\ Dusty torus luminosity & [erg/sec] & --- & $6\times 10^{44}$ & $3.4\times 10^{44}$ \\ Dusty torus temperature & [K] & --- & $2\times 10^3$ & $10^3$ \\ \hline \multicolumn{5}{c}{Shock Parameters} \\ \hline Upstream speed & $u_{1x}/c$ & 0.71 & 0.71 & 0.71 \\ Field obliquity & $\ThetaBfone $ & $32.3^{\circ}$ & $32.3^{\circ}$ & $52.4^{\circ}$ \\ Upstream temperature & $T_1$ (K) & $5.45\times 10^7$ & $5.45\times 10^7$ & $5.45\times 10^7$ \\ Compression ratio & $r$ & 3.71 & 3.71 & 3.71 \\ Injection mean free path & $\eta_1=\eta(p_1)$ & 100 & 20 & 225 \\ Diffusion index & $\alpha$ & 1.5 & 3 & 3 \\ \hline \multicolumn{5}{c}{Derived Parameters} \\ \hline Magnetic field & $B$ [Gauss] & $1.15\times 10^{-2}$ & $2.52$ & $2.50$ \\ Field luminosity & $\LB$ [erg/sec] & $4.5\times 10^{40}$ & $3.35\times 10^{43}$ & $2.9\times 10^{45}$ \\ Synchrotron partition\footnote{This is given by \teq{\epsilon_{\rm syn}\approx L_{O-X}/(L_{O-X} + L_{\gamma})}.} & $\epsilon_{\rm syn}$ & 0.6 & 0.83 & 0.25 \\ Cyclotron frequency & $\omegaB$ & $2.0 \times 10^5$ & $4.44\times 10^7$ & $4.4\times 10^7$ \\ $e^-$ plasma frequency & $\omega_p$ & $1.95\times 10^5$ & $1.4\times 10^7$ & $4.0\times 10^7$ \\ Shock magnetization\footnote{Eq.~(\ref{eq:sigma_param}) applies to non-relativistic cases; it must be divided by the Lorentz factor $\Gamma_1$ for relativistic shocks.} & $\sigma$ & $1.08$ & $10.0$ & $1.2$ \\ Alfv\'en speed & $ B/\sqrt{4\pi \rho} $ & $2.34\times10^{-3}\,$c & $0.074\,$c & $0.026\,$c \\ Cooling break & $\gamma_{\rm br}$ & $3.2\times 10^3$ & $1.75\times 10^2$ & 2.0 \\ Maximum $e^-$ Lorentz factor & $\gamma_{\rm max}$ & $8.5\times 10^5$ & $3.64\times 10^3$ & $1.61\times 10^3$ \\ Maximum $\lambda_{\parallel}/r_g$ & $\eta (\gammax )$ & $9.2 \times 10^4$ & $2.6 \times 10^8$ & $5.8 \times 10^8$ \\ Maximum mean free path\footnote{This is $\lambda_{\parallel}(\gammax) = \gammax\eta (\gammax )\, c/\omegaB$, and is realized in the acceleration region and beyond.} & $\lambda_{\rm max}$ [cm] & $1.17 \times 10^{16}$ & $6.5 \times 10^{14}$ & $6.4 \times 10^{14}$ \\ Synchrotron cutoff & $\nu_{\rm syn}$ [Hz] & $4.2 \times 10^{17}$ & $3.1 \times 10^{14}$ & $1.7 \times 10^{14}$ \\ SSC cutoff frequency & $\nu_{\hbox{\fiverm SSC}}$ [Hz] & $8.8 \times 10^{25}$ & $7.2 \times 10^{20}$ & $5.9 \times 10^{20}$ \\ \hline \end{tabular} \end{minipage} \end{table*} The principal constraints derived from the Mrk 501 multiwavelength modeling are that the mean free path parameter \teq{\eta\sim \eta_1} is modest for mildly-relativistic electrons injected into the shock acceleration process, and that at the highest energies, \teq{\eta (\gammax)\gtrsim 10^5} in order to place the synchrotron turnover (i.e., \teq{\nu F_{\nu}} peak) in the hard X-ray band. The values we obtained for the extended state were \teq{\eta_1=100} and \teq{\alpha =3/2} (i.e. \teq{\eta (p) \propto p^{1/2}}). These parameters generate moderately efficient injection from thermal energies into the Fermi mechanism (see Fig.~\ref{fig:accel_dist}), and mean free paths that possess a stronger momentum dependence than Bohm diffusion (\teq{\alpha =1}), suggesting interactions with weaker turbulence for more energetic particles diffusing on larger spatial scales. The interpretive elements of this parametric fit will be embellished upon in Section~\ref{sec:discussion}. It should be borne in mind that this fitting ``solution'' is representative of the appropriate parameter space, but is by no means a unique choice. It can be generally stated that there is a tolerance of about \teq{\pm 0.2} in \teq{\alpha}, and a tolerance of a factor of \teq{\sim 1-5} in \teq{\eta_1} permitting M/W spectral fits of similar quality. The uncertainties in the gamma-ray fluxes, particularly in the {\it Fermi}-LAT band, preclude greater precision for modeling. Moreover, in this analysis, the shock MHD parameters were kept fixed at the values chosen for Figure~10 of \citet{SB12}, namely a shock speed of \teq{u_{1x}=0.71c}, a velocity compression ratio of \teq{r=2.71}, a field obliquity of \teq{\ThetaBfone = 32.3^{\circ}}, and a low sonic Mach number of around 4. Adjusting the MHD structure of the shock would introduce changes to the optimal choices for the shock layer turbulence parameters \teq{\alpha} and \teq{\eta_1} in a data fitting protocol, so the ones for spectra exhibited in the Figure serve as general guides, and should not be presumed sacrosanct. Likewise, the choice of the observational epoch should always be borne in mind: Mrk 501 is highly variable above the optical band, and so sets of fitting parameters will change when different epochs are modeled. This is evident for strong flaring episodes such as the May 1--5, 2009 event not treated here. Yet it also applies for other observational campaigns, such as the recent NuSTAR-led multi-wavelength monitoring spanning the period April 1 -- August 10, 2013 \citep{Furniss15}. The X-ray spectrum therein is often harder and more luminous than that depicted here, though the gamma-ray spectrum (MAGIC+VERITAS+{\it Fermi}-LAT) is more or less commensurate with that in Fig~\ref{fig:MW_spec_Mrk501} --- see Fig.~9 of \citet{Furniss15} for details. To appreciate how extreme Mrk 501 can become during flares, one need look no further than the depiction of the April 1997 flare in Fig.~6 of \citet{Acciari11}, wherein the synchrotron peak moves to above \teq{3\times 10^{19}}Hz (i.e. nearly two decades higher than the listing in Table~1), and its flux level is slightly over a decade higher than that illustrated in Fig.~\ref{fig:MW_spec_Mrk501}. Modeling this extreme HBL behavior would lower the inferred value of \teq{\eta (\gammax )} by a factor of 30--50. Yet for this flare, and also for the NuSTAR M/W campaign, the principal conclusion of a strong momentum dependence for the diffusion coefficient and large departures from the Bohm limit at \teq{\gamma\sim\gammax} would be upheld, and is anticipated to apply to a broad selection of observing campaigns. The electron distribution corresponding to the spectrum in Fig.~\ref{fig:MW_spec_Mrk501} is illustrated in Figure~\ref{fig:blazar_edist}, in a manner analogous to the \teq{\nu F_{\nu}} representation: \teq{\gamma^2 n_e(\gamma )} distributions give an approximate scaling of the emission power of synchrotron and SSC signals from electrons of a given Lorentz factor \teq{\gamma}. The thermal Maxwellians are fairly prominent, yet they couple efficiently and smoothly into the non-thermal accelerated population. The exponential turnovers near the maximum energies are precipitated by efficient radiative cooling in the compact acceleration zone, depicted in Fig.~\ref{fig:model_geom}. This shuts off the acceleration, and the electron distribution evolves through continued cooling in the larger radiation zone. Cooling ceases when the radiation cooling length becomes comparable to the size of the radiation region, which is specified by the parameter \teq{R_{\rm rad}} in Table 1, and this introduces the break at \teq{\gamma\sim \gamma_{\rm br}}, the approximate value of which satisfies Eq.~(\ref{eq:gamma_br_cool}). For the extended state of Mrk 501, since the synchrotron component slightly dominates the SSC one in terms of the energy flux, the inverse Compton mechanism contributes the minority of the cooling; for any considerations of the high state, the situation would be reversed. For each process, efficient cooling steepens the electron distribution by an index of unity above the cooling break, and this induces the well-known steepening by an index of \teq{1/2} in the photon spectrum. Such a break is seen at around \teq{10^{12.5}}Hz for the synchrotron component in Fig.~\ref{fig:MW_spec_Mrk501}, and is barely discernible at around \teq{10^{19}}Hz for the SSC contribution therein. We note in passing that the injected shock acceleration distributions that led to those depicted for the larger cooling zone in this Figure have most of their energy allocated to the particles near \teq{\gammax}. In such cases, the acceleration simulations need to be upgraded to account for non-linear modifications to the shock MHD structure imposed by the pressure of the energetic non-thermal particles. This feedback phenomenon is well understood in non-relativistic shock theory \citep[e.g.][]{EE84,EBJ96,BERGG99}, and motivates future refinements to blazar studies of the genre herein, using extensions of this feedback theory to relativistic shocks along the lines of \citet{EWB13}. \begin{figure} \vspace*{-10pt} \centerline{\hskip 10pt\includegraphics[width=9.3cm]{blazar_edist.eps}} \vspace*{-5pt} \caption{Complete thermal plus non-thermal electron distributions employed in the multiwavelength modeling results for Mrk 501 depicted in Fig.~\ref{fig:MW_spec_Mrk501}, and BL Lacertae as shown in Fig.~\ref{fig:MW_spec_BLLac}, as functions of the four-velocity or dimensionless electron momentum \teq{\gamma\beta}. These distributions, applicable to the jet frame, are derived from the Monte Carlo simulation results that are injected into the radiating volume of the blazar jet, and are modified by cooling losses above the cooling break energy \teq{\gamma_{\rm br}\, m_ec^2}, and eventually by the suppression of acceleration by rapid radiative cooling at \teq{\gamma\sim \gamma_{\rm max}}, thereby precipitating the exponential tails. Given that \teq{n_e(\gamma )} is the distribution of electrons differential in Lorentz factor, this \teq{\gamma^2 n_e(\gamma )} representation highlights the fact that most of the synchrotron and SSC radiative power is generated by electrons with the Lorentz factors near and above \teq{\gamma_{\rm br}}. The Mrk 501 case (blue histogram) had \teq{\eta_1=100} and \teq{\alpha = 1.5}, while the BL Lac one (red) had \teq{\eta_1=20} and \teq{\alpha = 3}. The comparison \teq{n_e(\gamma )\propto \gamma^{-1}} differential distribution exhibited in green is the expected power law for high \teq{\eta_1} cases (see Section~\ref{sec:shock_accel}). \label{fig:blazar_edist}} \end{figure} Beholding the multiwavelength spectrum, at first sight, the character of the Mrk 501 SSC model fits here resemble those in numerous expositions on this source, including \citet{Acciari11} and \citet{Furniss15}. However there is a key development here beyond phenomenological electron distributions, namely broken power laws that are truncated at a minimum electron Lorentz factor, which are usually invoked in other studies of Mrk 501 \citep[e.g., see][]{Aleksic15} and various other blazars. These forms are unphysical from the standpoint of shock acceleration theory. Here, {\it we have complete thermal-plus-non-thermal distributions at our disposal}, and this serves to advantageously define the thermal plasma density normalization relative to that of the accelerated electron population (see Fig.~\ref{fig:blazar_edist}). The plasma density connects to determination of the shock structure, and the number density of accelerated electrons benchmarks the radiative output and also the jet kinetic energy \teq{L_{\rm kin}} via Eq.~\ref{eq:Lkin}. In addition, knowing the total number density \teq{n_e} fairly precisely permits the determination of the electron plasma frequency \teq{\omega_p}, and the plasma magnetization parameter \teq{\sigma = \omegaB^2/\omega_p^2} (not possible for phenomenological power-law distributions). Inspection of Table~1 reveals that the magnetization is \teq{\sigma\sim 1} for our Mrk 501 fitting, a result that will be interpreted in Section~\ref{sec:discussion} as implying similar speeds of acceleration for reconnection and for the diffusive Fermi mechanism in the Bohm limit. Also listed in Table 1 are determinations of the hydrogenic Alfv\'en speed, in all cases virtually non-relativistic and substantially inferior to the shock speed. These rise to relativistic speeds with the addition of a dominant pair component to the jet. Accordingly, here we forge an interconnection between photon emission and underlying shock plasma properties in blazar jets, in a substantial advance beyond previous works. From the spectroscopic point of view, it is apparent from Fig.~\ref{fig:MW_spec_Mrk501} that the thermal portion of the electron distribution functions played no role in constraining parameter space. The mildly-relativistic electrons in their thermal component are inefficient generators of both synchrotron and SSC radiation. Thermal synchrotron would appear in low frequency radio windows and is heavily self-absorbed by the inverse synchrotron process, which is treated in our radiation codes: the signatures of such attenuation generally appear at frequencies below \teq{10^{10}}Hz when \teq{B\lesssim 10}Gauss. Thermal SSC appears in the optical-X-ray window, and is swamped by the strong X-ray synchrotron emission for Mrk 501. Both contributions are off-scale in the Figure. Hence, in the case of Mrk 501, spectroscopy associated with the thermal jet electrons is irrelevant. This contrasts the interesting case of AO 0235+164, our final object for study. \subsection{Case Study: BL Lacertae} To offer a picture distinct from that of Mrk 501, the focus turns to BL Lacertae, the prototype BL Lac object, which has a much softer synchrotron component that peaks in the optical (i.e., an LBL) and dominates the hard X-ray and gamma-ray signals. It is also a nearby source, at \teq{z=0.069}, and various multiwavelength campaigns have been staged to help elucidate its character. A survey undertaken in the early days of the {\it Fermi} mission is the focal point here. This was for the extended period of August -- October 2008, and while the source was variable and possessed episodes of relatively enhanced radiative output, its broadband emission was not that of a characteristically high state. The data compilation for this three month epoch is detailed in the {\it Fermi}-LAT collaboration paper on a multitude of so-called LBAS blazars \citep{Abdo10c}. The purpose of this compendium was a first quick-look survey of the general character of bright blazar spectra in the {\it Fermi} era. Therefore, for BL Lacertae, the mutliwavelength data presented in that paper are depicted in Fig.~\ref{fig:MW_spec_BLLac}, serving as our modeling benchmark. The data in color are from {\it Fermi}-LAT in the 100 MeV to 30 GeV energy window (blue), Swift XRT (0.3-10 keV, green) and Swift BAT (15-200 keV, purple) in the X-rays, Swift-UVOT in the near ultra-violet, radio measurements from Effelsberg, Owens Valley and RATAN, and data from other optical, infra-red and radio facilities. Since not all were precisely contemporaneous with {\it Fermi}-LAT observations and each other, modest mismatches between model fits and data should not be over-interpreted. Synchrotron emission from BL Lac is bright enough to outshine its host galaxy, and so no galaxy contribution is discernible in the optical, unlike Mrk 501. Fig.~\ref{fig:MW_spec_BLLac} also presents archival MAGIC data \citep{Abdo10c} in grey that were not used in the fitting protocol, and serve as a guide for the variable character of gamma-ray signals from BL Lac. Fig.~23 of \citet{Abdo10c} can be consulted for a more extensive presentation of archival spectroscopy for BL Lac. The radiation modeling protocols were implemented as described above, using a similar suite of particle distributions generated by the Monte Carlo shock acceleration simulation. Radiation model and simulation input parameters, and also those derived from the input ones, are listed in Table~1. The shock MHD parameters were identical to the Mrk 501 study. The turbulence parameters \teq{\alpha} and \teq{\eta_1} for the simulations were again adjusted to optimize the fit, for which spectral results are presented in Fig.~\ref{fig:MW_spec_BLLac}. The radio-to-optical data dictate that the synchrotron component is dominant, and turns over below the X-ray band. This imposes a higher value of the magnetic field and its associated portion of the total energy budget than for the case of Mrk 501. Concomitantly, synchrotron cooling of electrons in the acceleration zone is more rampant, thereby reducing the maximum Lorentz factor to \teq{\gammax\sim 3.5\times 10^3} (see Fig.~\ref{fig:blazar_edist}). Accordingly, the synchrotron self-Compton \teq{\nu F_{\nu}} peak energy appears in the soft gamma-rays, below 10 MeV, as is evident in Fig.~\ref{fig:MW_spec_BLLac}. This circumstance is essentially unavoidable in LBLs that are synchrotron-dominated. Hence, the {\it Fermi}-LAT data require another component to be present. To describe such, here we employ an external Compton signal seeded by emission in the near infra-red (NIR) from a dusty torus associated with the accretion flow onto the central black hole. These constitute an isotropic thermal photon field with the luminosity and temperature as listed in Table~1. It is a simple matter to discern why this EC component emerges in the {\it Fermi}-LAT band. While the torus appears in the NIR for an observer, as an external isotropic field proximate to the jet, it is boosted in the jet frame into the far-UV. This is at considerably higher frequency than the synchrotron emission, which clearly peaks in the IR band in the jet frame. Accordingly, the EC peak is substantially bluer than that of the SSC in both the jet frame, and our observer frame, by over a factor of 100 higher in energy. Moreover, the Doppler-boosted flux enhancement of the torus seed photons in the jet frame can generate a gamma-ray \teq{\nu F_{\nu}} flux comparable to the SSC without the seed NIR signal being discernible above the very strong, beamed, synchrotron component at those frequencies. The breadth of the EC peak is restricted by the intrinsic width of the Planck spectrum from the torus. \begin{figure} \vspace*{-10pt} \centerline{\hskip 10pt\includegraphics[width=9.5cm]{BLLac_spec.eps}} \vspace*{-5pt} \caption{Multiwavelength \teq{\nu F_{\nu}} spectra (points) spanning the radio, optical, X-ray and gamma-ray bands, together with model fits as described in the text, for the August--October 2008 extended {\it Fermi}-LAT observation of the blazar BL Lacertae. The data are taken from the LBAS study of Abdo et al. (2010), and are briefly detailed in the text: the colored data are approximately contemporaneous with the {\it Fermi}-LAT collection. The grey MAGIC points above \teq{10^{25}}Hz are archival, being shown to illustrate the spectral variability between different observational epochs; they were not accommodated in the model fit. The broadband models consist of a synchrotron component (dashed green curve) up to the optical band, an SSC component in the X-rays and gamma-rays (dashed purple curve), and an external inverse Compton emission component (dash-dot grey) seeded by IR photons from a dusty accretion torus. The full jet model spectrum (orange) includes a correction for \teq{\gamma-\gamma} absorption by the extragalactic background light, and combined with the data, constrain the model diffusion parameters to \teq{\eta_1=20} and \teq{\alpha =3} (see text). \label{fig:MW_spec_BLLac}} \end{figure} While this introduces an extra subset of model parameters, there appears no likely alternative in one-zone leptonic models. Similar broad spectral structure from X-rays to hard gamma-rays could be generated by a multi-zone blazar emission model: such was the protocol adopted by \citet{ZW16} in the modeling of AP Librae, an LBL object whose multiwavelength spectroscopy is very similar to that of BL Lacertae. A possible discriminator between these two competing pictures is the appearance of a flat bridge in the spectrum from our protocol between the two inverse Compton components. Such a feature might not appear in a multi-zone construction, which may evince broader spectral curvature. This provides motivation for future sensitive gamma-ray detectors below the {\it Fermi}-LAT window, perhaps employing Compton telescope technology. At the very highest energy gamma-rays, as for Mrk 501, our modeling incorporated a \teq{\gamma\gamma\to e^+e^-} pair opacity correction \citep{FRD10} due to the intervening EBL. However, since the EC peak in BL Lac lies at 10 GeV and well inside the LAT window, this correction is quite small, and did not influence our fitting protocol. It is anticipated that for this source, when it is in a hard state as captured in the non-contemporaneous MAGIC data, such attenuation compensations will come into play. It is notable that BL Lacertae is singular in our ensemble of blazars in that it has measurements of the jet bulk velocity. \citet{Jorstad05} outline VLBI imaging, polarimetry and variability studies of a number of radio jet sources, including BL Lac. Combining imaging and flux variability information at 7mm, they were able to isolate both the apparent superluminal speed of the jet and the Doppler factor. This yielded a determination of the bulk Lorentz factor of \teq{\Gamma_{\rm jet}\sim 7\pm 2} for BL Lac: see Table~11 of \citet{Jorstad05} and associated discussion. In a later investigation of TeV flaring in BL Lac using VERITAS, \citet{Arlen13} employed internal \teq{\gamma\gamma\to e^+e^-} pair transparency arguments based on \citet{DG95} to constrain the jet bulk motion to Lorentz factors \teq{\Gamma_{\rm jet}\gtrsim 13-17}. \citet{Arlen13} argued that this determination was possibly not in conflict with the lower value obtained in the VLBI study, with the gamma-ray emitting region being distinct from the radio ones and perhaps lying closer to the central black hole, thereby invoking a decelerating jet picture. Since these gamma-ray constraints are most germane to our multiwavelength investigation here, they guided our adoption of \teq{\Gamma_{\rm jet}=15} for the models of BL Lac. This guaranteed pair transparency for the jet for gamma-rays below 1 TeV. Higher values of \teq{\Gamma_{\rm jet}} that permit jet pair transparency can also lead to viable fits with the same particle distributions from acceleration theory, yet require reasonable adjustments of environmental parameters such as the field strength and electron number density. As with Mrk 501, the upshot of the M/W modeling is that a large value of \teq{\alpha} is required to generate relatively efficient injection of electrons from thermal energies into the acceleration process, and inefficient acceleration at the highest Lorentz factors. Notably, since BL Lac is an LBL blazar, a value of \teq{\alpha \sim 3} was required to increase \teq{\eta (\gammax )} above \teq{10^7} and move the synchrotron peak into the optical band. The values we obtained, \teq{\eta_1=20} and \teq{\alpha =3} (i.e. \teq{\eta (p) \propto p^2}), indicate diffusion well away from the Bohm limit at all momenta, and that the level of turbulence must decline with distance from the shock even more profoundly than for Mrk 501. As with our first case study, there is a tolerance of about \teq{\pm 0.2} in the range of \teq{\alpha}, and a tolerance of a factor of \teq{\sim 1-5} in \teq{\eta_1} permitting M/W spectral fits of similar quality. The value of \teq{\eta_1=20} obtained for BL Lac is noticeably lower than for our two other case study blazars. Yet, it is also somewhat higher than the value of \teq{\eta_1\sim 5} determined from fitting thermal+non-thermal distributions of protons and alpha particles that are measured {\it in situ} at interplanetary shocks in the non-relativistic solar wind (Baring et al. 1997). This suggests that mildly relativistic shocks in blazar jets are less efficient at injecting charges into the acceleration process than are traveling shocks in the solar wind. Note that the contributions to each radiation component from the thermal electron population are small enough to be off scale in the Figure. The electron distribution corresponding to the fit is depicted in Fig.~\ref{fig:blazar_edist}. An issue that has not yet been addressed concerns the variability timescale, which must be large enough to accommodate the time for acceleration up to the largest particle energies required to fit the spectral data. The acceleration rate is given by \teq{d\gamma /dt\sim \omegaB/\eta} in Eq.~(\ref{eq:acc_rate}), assuming a relativistic shock. Thus, the most energetic charges are accelerated on timescales of \begin{equation} \tau_{\rm acc} \;\sim\; \dover{\eta (\gamma_{\rm max}) \, \gamma_{\rm max}}{\omegaB} \label{eq:} \end{equation} in the rest frame of the jet. Each of these quantities on the right hand side is listed in Table 1, yielding derived values of the characteristic cumulative timescale for acceleration \teq{\tau_{\rm acc}}, also listed in the Table. Thus we infer \teq{\tau_{\rm acc} \sim 4\times 10^5}sec for Mrk 501, i.e. around 4.5 days. Similarly, \teq{\tau_{\rm acc} \sim 2\times 10^4}sec for Bl Lac. To connect to the observer's frame, these are shortened by time dilation factors of \teq{1/\Gamma_{\rm jet}} for transverse variability, i.e to around \teq{1.6\times 10^4}sec or 4 hours for Mrk 501, and around 20 minutes for Bl Lac. Such dilation applies to fluctuations in the direction normal to the line of sight to the observer, i.e. approximately the jet axis. Longitudinal variability along the jet, such as are expected from colliding shells, reduces the acceleration timescales by another factor of \teq{1/\Gamma_{\rm jet}}, rendering them very short in the observer frame. These are then less than the variability timescales for both sources in the hard gamma-rays (e.g., see Aliu et al., 2016, for Mrk 501 data in the high state that was excluded from the epoch considered in this paper), for both extended states and flaring episodes within the observing epochs under consideration, and so no internal inconsistency arises. \subsection{Case Study: AO 0235+164} To provide another contrasting case, we turn to the LBL blazar AO 0235+164, which has a radio-optical component that turns over below the UV band. It is normally weak in X-rays, however, in flare state, it possesses a significant X-ray excess. Such enhancements have been discussed \citep{SMMP97,BSMM00} as a possible signature of a bulk Comptonization effect \citep[see also][]{SBR94}, where hot thermal jet electrons upscatter an ambient, quasi-thermal radiation field that manifests itself in the infra-red, optical or UV bands. This seed field could be optical/UV radiation from the broad-line or narrow-line regions, or IR emission from warm dust in the greater AGN environment; it is distinct from the non-thermal jet synchrotron emission. One of the objectives here is to explore whether the X-ray ``excess'' can be attributed to the substantial thermal electron pool evinced in shock acceleration distributions such as those depicted in Fig.~\ref{fig:accel_dist}. In the fast-moving jet, this population may produce significant radiative signatures due to Comptonization of an external radiation field, taking advantage of the same Doppler boosting enhancement that was captured in our modeling of BL Lac. The focus here is on modeling data collected during the flaring episode of September -- November 2008. In particular, we employ the multi-wavelength campaign observations discussed in \citet{Ackermann12}; see also \citet{Agudo11} for additional flare light curves and polarization data. The broadband spectrum for the high state portion of the outburst, during the interval MJD 54761-54763 (October 22-24, 2008), is depicted in Fig.~\ref{fig:MW_spec_AO0235}. This appears in Fig.~7 of \citet{Ackermann12}, wherein it can be compared with a low state spectrum obtained about six weeks later that does not contain a {\it Swift} XRT detection. The list of observatories and missions participating in this campaign is extensive --- it spans radio and optical, with {\it Swift} XRT and RXTE providing X-ray monitoring, and {\it Fermi}-LAT providing the principal gamma-ray observations. The LAT spectrum displays a break or possible turnover at energies around \teq{3-4}GeV, with a power-law index of about \teq{2.1} below this feature; the source is not detected in the TeV band \citep{Errando11}. \begin{figure} \vspace*{-10pt} \centerline{\hskip 10pt\includegraphics[width=9.5cm]{AO0235_spec.eps}} \vspace*{-5pt} \caption{Multiwavelength \teq{\nu F_{\nu}} spectra (points) spanning the radio, optical, X-ray and gamma-ray bands, together with model fits as described in the text, for the October 2008 high state {\it Fermi}-LAT observation of the blazar AO 0235+164. The campaign data are taken from \citet{Ackermann12}, corresponding to the high state epoch during October 22-24. The gamma-ray detections and upper limits are from {\it Fermi}-LAT, while the X-ray ``butterfly'' block (blue) represents {\it Swift} XRT data. The broadband models consist of a synchrotron component (dashed green curve) up to the optical band, a two-order SSC contribution in the optical, X-rays and gamma-rays (dashed purple curve), and bulk Comptonization emission (EC: dotted black curve) between \teq{\sim 0.1}keV and \teq{\sim 10}GeV. The total model spectrum (orange) includes a very small correction for \teq{\gamma\gamma} absorption by the extragalactic background light, and combined with the data, constrains the model diffusion parameters to \teq{\eta_1=225} and \teq{\alpha =3} (see text). \label{fig:MW_spec_AO0235}} \end{figure} The major model components and total multi-wavelength spectrum in our fit of the flare data are also exhibited in Fig.~\ref{fig:MW_spec_AO0235}. The jet and disk parameters used to derive the model fit are listed in Table 1, and we remark that they are somewhat different from those in our earlier introductory exposition \citep{BBS14} on this source. Notably, the magnetic field at \teq{2.5}Gauss was sub-equipartition by about a factor of 16. To place the synchrotron turnover frequency in the optical band (\teq{\sim 10^{14}}Hz), the values of \teq{\eta (\pmax )} and \teq{B} have been chosen so that the maximum electron Lorentz factor in the comoving frame of the jet is \teq{\gammax \approx 1.6\times 10^3}. The shock obliquity was modestly higher than for the BL Lac and Mrk 501 studies, a choice that renders the thermal electrons comparatively more numerous. As with BL Lac, the model fit required \teq{\alpha =3} and \teq{\eta_1=225}, so that the diffusive scales of energetic particles are very large. The tolerances in these parameters that permit M/W fits of similar calibre are about \teq{\pm 0.2} in the range of \teq{\alpha}, and a factor of \teq{\sim 1-3} in \teq{\eta_1}. The moderate value of \teq{\gammax} then positions the peak SSC frequency at around \teq{\gammax^2 10^{14}\hbox{Hz}\sim 6\times 10^{20}}Hz, i.e. around 5 MeV, as is evident in Fig.~\ref{fig:MW_spec_AO0235}. This is not the maximum extension of SSC component: a second-order IC image of the synchrotron continuum appears up to about 10 GeV. A cooling break is also apparent in the SSC spectrum at around a few keV. Observe that the SSC component is of insufficient luminosity to model the {\it Swift} XRT and {\it Fermi}-LAT signals --- another component is needed to explain them. These two high state detections are modeled here using bulk-Comptonization/inverse Compton scattering of an external radiation field. This field is presumed to be quasi-isotropic in the observer's frame, so that the jet material moves at \teq{\Gamma_{\rm jet}=35} relative to these seed photons, and so can inverse Compton scatter them, i.e. increase their frequency by a factor of \teq{\sim 4 \Gamma_{\rm jet}^2}. This choice of \teq{\Gamma_{\rm jet}} lies between that adopted in some of the supporting modeling in the {\it Fermi}-LAT paper of \citet{Ackermann12}, namely \teq{\Gamma_{\rm jet}=20}, and the much higher value of \teq{\Gamma_{\rm jet}\sim 50} indicated by structure variability in the VLBA imaging over extended epochs that was presented in \citet{Jorstad01}. We note that these values exceed the typical range of \teq{\Gamma\sim 10-20} quoted in \citet{Ghisellini14} for BL Lac objects where Compton dominance (or not) in SSC models is used to constrain \teq{\Gamma} --- clearly, employing external Compton scenarios will modify these estimates. The Lorentz factor of the thermal electrons in the distributions of Fig.~\ref{fig:accel_dist} is less than about 2, thereby modifying this IC energy enhancement by at most a factor of a few. Hence, to explain the {\it Swift} XRT flaring flux, the seed background radiation needed a temperature of \teq{T\approx 10^3}K. This corresponds to infra-red radiation, probably from a dusty torus. The thermal portion of the electron population boosts the ambient photons up to the X-ray range, and roughly produces the steep {\it Swift} XRT spectrum if it corresponds to the higher momentum side of the Maxwell-Boltzmann distribution, as is clearly evident in Fig.~\ref{fig:MW_spec_AO0235}. This then connects to a very flat external Compton tail, generated by the broken power-law tail portions of distributions like those depicted in Fig.~\ref{fig:accel_dist}, modulo the cooling breaks discussed above. The IC tail possesses a flat spectral index of around 3/2, rising in the \teq{\nu F_{\nu}} representation to meet the {\it Fermi}-LAT flux above 100 MeV. The large IC flux indicates very strong Compton cooling of electrons, and this generates a very low cooling break energy at \teq{\gamma_{\rm br}\sim 2}. Note that despite the high redshift (\teq{z=0.94}) of this source relative to those of the other two blazars, the EBL absorption correction is minimal since the spectrum extends only up to around 20 GeV, i.e., similar to the situation for BL Lac. A peculiarity of this blazar is that sensitive optical/UV spectroscopy reveals the presence of an intervening damped Lyman-$\alpha$ absorbing system at a redshift of \teq{z=0.524} \citep[e.g.,][and references therein]{Junkk04}. This absorber manifests itself via the appearance of a broad, redshifted 2175\AA\ feature that is attributed to the presence of graphitic dust grains nearer than AO 0235+164. Such a concentrated repository of dust (perhaps in a galaxy) along the line of sight can also provide \teq{\gamma\gamma} pair absorption for the AO 0235+164 spectrum. Yet this source of opacity would again be confined to the energy band greater than 20 GeV, and so has minimal impact on the spectral modeling here. This system can, however, serve to attenuate soft X-rays via photoelectric absorption. In modeling the broadband data, the ephemeral {\it Swift} X-ray component is considered to be a bulk Comptonization signal, which typically span a narrow waveband. \citet{SMMP97} demonstrate that such a component needs to be generated at perhaps 100 gravitational radii or more in order for the kinetic luminosity of the jet to not overproduce X-rays. The maximum mean free paths \teq{\lambda_{\parallel}(\gammax)} listed in Table~1 are consistent with this constraint in systems where the central black hole exceeds around \teq{10^8M_{\odot}} in mass. \citet{BSMM00} note that such Comptonization X-ray signals are rare in the OVV quasar population. AO 0235+164 is peculiar in exhibiting such a steeply declining X-ray spectrum, contrasting BL Lac as depicted in Fig.~\ref{fig:MW_spec_BLLac}. It is natural to ask whether the X-ray flare in AO 0235+164 could alternatively be a distinct synchrotron component from a different region in a multi-zone construction. This requires more model parameters, with no additional constraints. To render this explanation viable, either the magnetic field of the second zone should be much larger, or the maximum Lorentz factor \teq{\gammax} should be much higher, or both. Either of these requirements would also demand a smaller cooling volume so as to reduce the X-ray luminosity below the optical one. If \teq{\gammax} is much larger than the value of \teq{1.6\times 10^3} listed in Table 1, one anticipates the appearance of a significant SSC component extending up into the TeV band. At MeV energies, this will necessarily lie below the one displayed in Fig.~\ref{fig:MW_spec_AO0235} that corresponds to a synchrotron fit in the radio to optical. Accordingly its extension out to the LAT band would likely lie below the observational data, and one envisages difficulties for a picture with two distinct synchrotron (and also SSC) components without an external Compton one. While it is possible, in principle, to construct a viable two-zone synchrotron plus SSC plus EC model for the AO 0235+164 multiwavelength spectrum, we view the interpretation of the steep X-ray continuum in the high state as a bulk Comptonization signature as the simpler explanation. \vspace{-10pt} \section[]{Discussion: Turbulence \& Diffusion} \label{sec:discussion} The recurring theme in all of these multi-wavelength spectral fits is that the diffusion of charges in blazar acceleration zones should be well removed from the Bohm limit \teq{\lambda \sim r_g} at nearly all energies, and furthermore that the diffusion scale \teq{\lambda} should lengthen fairly rapidly with increasing particle momentum. This indicates that the physical environment of the blazar shocks should possess a diminished level of turbulence at larger distances from the discontinuity, precipitating the longer diffusive mean free paths for larger momenta. This circumstance is not entirely without parallel. In the heliosphere, the solar wind contains traveling interplanetary (IP) shocks, at which accelerated ions are clearly detected, and MHD turbulence is measured using magnetometers mounted on {\it in situ} spacecraft. The turbulence near such IP shocks is observed to satisfy \teq{\langle\delta B/B\rangle \sim 0.1}, and the diffusive transport of suprathermal H and He ions at \teq{\sim 10} keV energies is inferred to be not very different from the Bohm limit, with mean free paths of \teq{\lambda_{\parallel} \lesssim 10^4}km \teq{\sim 10^{-4}}AU \citep[e.g.,][]{BOEF97}. In contrast, away from these shocks, the turbulence is much more benign, and the diffusive mean free paths \teq{\lambda_{\parallel}} of 10 MeV--1 GeV ions are much larger \citep{FJO74,Palmer82,Bieber94,ZMBM98}, on the scale of \teq{0.1-3}AU or so. \subsection{Turbulence, Diffusion and the Solar Wind} From such contrasting information, \teq{\alpha >1} circumstances must clearly exist. Determinations of the diffusion index in the solar wind are non-trivial, and not all that precise. The reason for this is primarily that magnetometer measurements consist of three-dimensional vector fields {\bf B} at one particular space location. This restricts knowledge of the global field character, and prohibits assessment of true diffusive trajectories of charges on extended spatial scales. For examples of ion diffusion characteristics inferred from older solar wind data, the reader can consult \citet{EMP90} for observations at the Earth's bow shock, which indicate that \teq{1/2 < \alpha < 3/2}. In contrast, lower values \teq{\alpha\sim 1/2} are obtained from turbulence in the interplanetary magnetic field on larger scales at heliodistances around 20 AU \citep{Moussas92}, and in addition \teq{1/2 < \alpha < 4/5} can be deduced from ions accelerated in solar particle events \citep{MGH83} close to the sun. It is immediately obvious that these values of \teq{\alpha} do not match the larger values inferred from the blazar spectral modeling. Accordingly, different conditions must apply to the energetic particles in the blazar jet environment, as will become apparent shortly. More recent magnetometer measurements of the quiet solar wind \citep{PRG07} by the {\it Wind} spacecraft over a period of several years captures a broader dynamic range for the power spectrum. The {\it Wind} data indicate that the inertial range for turbulence \teq{\langle (\delta B)^2/8\pi \rangle} spans 3-4 decades in frequency above \teq{\nu_{\rm stir}\sim 3\times 10^{-5}}Hz, which marks the stirring scale of the system, i.e., \teq{v_{\rm sw}/\nu_{\rm stir} \sim 10^7}km, where \teq{v_{\rm sw}\sim 400}km/sec is the approximate speed of the solar wind. The spectrum for this inertial range is fairly close to the standard Kolmogorov \teq{\nu^{-5/3}} form. Below this frequency, the turbulence spectrum flattens, resembling the shape depicted in the schematic of Figure~\ref{fig:powerspec}, and one anticipates that diffusion in solar wind plasma turbulence must decline. The normalization of the power spectrum establishes the diffusive scale of ions in such turbulence, and estimates using quasi-linear diffusion theory lead to the aforementioned diffusive scale of the order of \teq{0.1-3}AU in the undisturbed interplanetary medium. Note that departures from pure Kolmogorov forms for solar wind field turbulence are observed in high heliographic latitude (i.e. well out of the ecliptic) {\it Ulysses} magnetometer determinations \citep{Goldstein95}. This excursion should alter the momentum dependence of the mean free path somewhat. An abundance of data in the heliosphere indicates that the solar wind is very structured, with active regions containing shocks, enhanced turbulence and the production of energetic ions and electrons. Perhaps the popular picture that shocks generate turbulence through their viscous dissipation, and that this is muted or damped in regions remote from them, applies to both the solar wind outflow and also to blazar jets. This is the interpretative message encapsulated in this paper. We now explore some elements of diffusion theory in the solar wind, bearing in mind that while commonalities between the structured, non-relativistic heliospheric wind and turbulent relativistic blazar jets are expected to exist, there should also be some inherent differences between turbulence in these two contrasting cosmic plasma environments. Appraising the assembly of solar wind data in terms of the stochastic impact of magnetic turbulence on charge motion requires a theoretical construct for modeling diffusion. The historical tool dating from the 1960s has been the quasi-linear theory (QLT) of diffusion in magnetodynamic turbulence, i.e. that embedded in non-relativistic flows, for which electric field fluctutations are small. There are many expositions of such in the literature, and here we capture their essence through a compact set of equations. The choice here of outlining elements of QLT in non-relativistic environs is motivated by the ability to efficaciously define certain simple and important characteristics, principally the momentum dependence of the diffusive mean free path, and how this connects to the power spectrum of turbulence. It is anticipated that the global character of \teq{\lambda_{\parallel}(p)} being a strongly increasing function of momentum will carry over into the mildly relativistic plasma regime that is germane to our blazar jet frame context. \begin{figure} \vspace*{-10pt} \centerline{\hskip 10pt\includegraphics[width=9.3cm]{blazar_powerpec.eps}} \vspace*{-5pt} \caption{Schematic plot of a representative power spectrum \teq{\langle \delta \mathbf{B}\rangle^2/B^2} for cascading MHD turbulence in astrophysical systems like the solar wind, interpreted as approximately what should exist in blazar jets. The turbulence is seeded at low wavenumbers {\bf k} (long wavelengths) around a stirring scale of \teq{\vert \mathbf{k}\vert\sim k_{\rm stir}} and cascades in a roughly scale-independent manner to high frequencies or short wavelengths. Eventually it reaches a dissipation scale, here chosen to be \teq{\vert \mathbf{k}\vert \sim 10^{2.5}k_{\rm stir}}, where absorption of MHD waves by the plasma heats the charges. The power spectrum in the intermediate wavenumber domain, the inertial range that is usually what is captured in solar wind magnetometer data, is shown here for a Kolmogorov scaling \teq{\langle \delta \mathbf{B}\rangle^2/B^2\propto k^{-5/3}} in one-dimensional MHD turbulence (see text); other approximate scalings are realized in Nature. At extremely long wavelengths \teq{\lambda \gg 2\pi /k_{\rm stir}} beyond the stirring scale, the precise form of the power spectrum is usually unknown, and is represented by the two cases shown, namely flat and declining. \label{fig:powerspec}} \end{figure} In the diffusive limit of particle transport, the spatial development of particle distribution functions along the mean magnetic field scales with the spatial diffusion coefficient \teq{\kappa_{\parallel} = \langle (\Delta x)^2\rangle/(2\Delta t)}, which appears in the second-order derivative terms in the Fokker-Plank form of the transport equation. In the limit of transport subject to small, chaotic field fluctuations, charges diffuse in their pitch angle \teq{\mu = \mathbf{v} \cdot \mathbf{B} / (\vert \mathbf{v}\vert . \vert \mathbf{B}\vert )}, so that \teq{\kappa_{\parallel}} can be related \citep[e.g.,][]{Jokipii66,KP69,HW70,Schlick89} to the pitch-angle diffusion coefficient via \begin{equation} \dover{\lambda_{\parallel}v}{3}\; =\; \kappa_{\parallel} \;\equiv\; \dover{\langle (\Delta x)^2\rangle}{2\Delta t} \; =\; \dover{v^2}{8} \int_{-1}^1 \dover{(1-\mu^2 )^2}{D_{\mu\mu}}\, d\mu\quad , \label{eq:diffusion_coeff} \end{equation} for charges of speed \teq{v}. The relationship between \teq{\kappa_{\parallel}}, \teq{D_{\mu\mu}} and the mean free path \teq{\lambda_{\parallel}} parallel to the mean field presumes isotropy of the diffusion process. Using QLT, the pitch angle coefficient can be approximately related to the power spectrum of field turbulence: \begin{equation} D_{\mu\mu} \;\equiv\; \dover{\langle (\Delta \mu)^2\rangle}{2\Delta t} \; \approx\; \dover{\pi v}{r_g^2B^2} \, \dover{1-\mu^2}{\vert \mu\vert} \, {\cal P} \left( k_{\parallel} = \dover{\omegaB}{\gamma v\,\vert \mu\vert} \right)\quad . \label{eq:Dmumu_QLT} \end{equation} In this expression, \teq{r_g =pc/(eB) = \gamma v/\omegaB} is the gyroradius of the charge, as employed throughout this paper, and \teq{B} is the strength of the mean unperturbed field. Note that slightly different constructions for \teq{D_{\mu\mu}} appear in the literature. The function \teq{{\cal P}(k_{\parallel})} is the power spectrum of one-dimensional ``{\it slab}'' turbulence with propagation wavenumbers \teq{k_{\parallel}} along the mean field direction; it is essentially a Fourier transform of the spatial fluctuations in \teq{(\delta B)^2}. The coupling between the power spectrum and the diffusion coefficient implicitly presumes that the turbulence is propagating at non-relativistic speeds, i.e., is magnetostatic, so that the applicable diffusion tensor is three-dimensional. Details of this derivation are presented in \citet{Jokipii66,FJO74,Lee82,ZMBM98}, and numerous other papers in the literature. A key element of this result is that it presumes resonant interactions between charges and waves at the Doppler-shifted cyclotron frequency, i.e. \teq{k_{\parallel}v\, \vert \mu\vert = \omegaB/\gamma}. This is a circumstance that is not always realized in plasma turbulence, as will become apparent below. Departures from this idealized gyroresonant case have been presented in \citet{Schlick89}, for example, wherein embellishments such as incorporating the influence of finite Alfv\'en wave speeds can ameliorate the problem of the inherent suppression of diffusion through pitch angles \teq{\arccos \mu = \pi /2}. Notwithstanding, the combination of Eqs.~(\ref{eq:diffusion_coeff}) and~(\ref{eq:Dmumu_QLT}) serves to define the generic character of diffusion in turbulent fields. To this purpose, one can follow Appendix A of \citet{ZMBM98} and approximate the power spectrum of one-dimensional turbulence with a modified Kolmogorov form: \begin{equation} P(k) \; \propto\; \dover{1}{(1+k^2/k^2_{\rm stir})^{5/6}}\quad . \label{eq:power_spec_Kolmogorov} \end{equation} This possesses true Kolmogorov dependence \teq{P(k)\propto k^{-5/3}} in the inertial range where \teq{k\gg k_{\rm stir}}, and is constant at wavenumbers below the stirring or field coherence scale \teq{k_{\rm stir}}. Such a form resembles the solid curve in the schematic plot in Fig.~\ref{fig:powerspec}, specifically in the inertial range below the dissipation scale at small wavelengths. Inserting it into the core QLT equations above yields \begin{equation} \lambda_{\parallel} \;\propto\; \dover{3B^2}{4\pi} \, r_g^2 \int_0^1 \mu\, (1-\mu^2) \left\{1 + \dover{1}{(r_g \mu k_{\rm stir})^2} \right\}^{5/6} \, d\mu\; . \label{eq:mfp_QLT} \end{equation} This is essentially the form given in Eq.~(A4) of \citet{ZMBM98}. For low momentum charges, those with \teq{r_g k_{\rm stir}\lesssim 1}, the inertial range is accessible, and the dominant contribution to the integral comes from small pitch angle cosines \teq{\mu}, so that \teq{\lambda_{\parallel}\propto r_g^{1/3} \propto p^{1/3}}, approximately. This behavior, which is commensurate with the values discussed above for heliospheric shocks, is predicated on the presumption of 1D Kolmogorov turbulence, and adjusts according to changes in the power spectrum and dimensionality of the turbulence. The conclusion of generally weak momentum dependence of the diffusion mean free path at low momenta is not restricted to QLT considerations. Complementary tools for modeling diffusion in MHD turbulence are plasma simulations, of which there are numerous analyses in the heliospheric literature. Hybrid plasma simulations \citep{GBS92,GJ99} in which electrons are treated as a background fluid, but ions are modeled kinetically, suggest mean free paths of non-relativistic ions obeying Eq.~(\ref{eq:mfp_alpha}) with \teq{\alpha \sim 2/3} in one-dimensional Alfv\'en-like turbulence. Similiar diffusion scales are realized in the solar wind turbulent diffusion analysis of \citet{ZMBM98} that combines slab and 2D turbulence. Therein, the actually momentum dependence is not well-constrained, though it is roughly consistent with QLT predictions. More recently \citet{CS14} computed diffusion coefficients using hybrid plasma simulations of non-relativistic shocks with strong turbulence. These isolate ion diffusion that is fairly close to the Bohm limit \citep[contrasting][]{ZMBM98}, with a momentum dependence \teq{\kappa_{\parallel}\propto p^{\alpha +1}} over roughly two orders of magnitude that is slightly weaker (namely \teq{\alpha \sim 0.8-1}) than for the Bohm case, i.e., \teq{\kappa_{\parallel}\propto \lambda_{\parallel}v \propto p^2}. While the forgoing exposition pertains to non-relativistic systems, {\it a priori} one anticipates that its general character should also apply to relativistic plasmas, where electric field fluctuations are no longer small. Returning to inferences from formal QLT, in stark contrast to this low energy domain, for charges with high momenta and therefore large gyroradii, \teq{r_g k_{\rm stir}\gtrsim 1}, the Doppler resonance condition \teq{k_{\parallel}v\, \vert \mu\vert = eB/(\gamma mc)} is not satisfied by turbulence in the inertial range, only by low frequency waves, which are represented by the flat portion of the solid curve in Fig.~\ref{fig:powerspec}. The result from Eq.~(\ref{eq:mfp_QLT}) is then \teq{\lambda_{\parallel}\propto p^2}, i.e. \teq{\alpha\sim 2} and \teq{\eta \propto p}, corresponding to a strong momentum dependence. This nuance is well-known in the heliospheric physics community. The diffusion of {\it energetic charges} then receives a substantial contribution from non-gyroresonant collisions with inertial turbulence. Such is expected to be the situation for electrons with large gyroradii being accelerated in blazar shocks. \subsection{Turbulence and Diffusion: Blazar Contexts} The broadband spectral modeling presented for the three blazars clearly indicates that plasma turbulence does not persist at the same levels out to very large distances from blazar shocks, but instead declines in intensity. This applies in the inertial range, but can also be more pronounced at wavelengths longer than the stirring scale. Such is the circumstance represented by the dashed curve in Fig.~\ref{fig:powerspec}, which is similar to the power spectra exhibited in the theoretical modeling of Alfv\'en-like and compressible, non-relativistic MHD turbulence that is presented in \citet{CL03}. For this situation, one can anticipate that \teq{\alpha} can effectively rise above two. Mathematically, this domain could be described by a power spectral form \teq{P(k) \propto (k_{\rm stir}/k +k^2/k^2_{\rm stir})^{-5/6}}, which would yield \teq{\alpha > 2} in Eq.~(\ref{eq:mfp_QLT}). In other words, \teq{\alpha >2} scenarios are entirely realistic for blazar multi-wavelength models. Low levels of long wavelength turbulence can be enough to eventually turn around the highest energy charges that travel upstream in shock drift reflection cycles. Note that it is not necessary that such turbulence be generated within the shock zone: some portion of it may in fact be convected into the shock from the remote regions of the jet, somewhat akin to the turbulence fed into interplanetary shocks in the solar wind. The discussion so far has focused on Alfv\'enic turbulence in non-relativistic magnetized plasmas. It is natural to ask whether the properties of relativistic turbulence are similar. Studies of relativistic MHD turbulence are both more limited in number and more recent, with the non-relativistic domain being deeply investigated because of its relevance to the heliosphere, cosmic ray studies, the interstellar medium and supernova remnants. The upshot of recent work on relativistic Alfv\'en-like turbulence \citep[e.g., see][]{Cho05,BL09,RR13,ZMacF13} is that a consensus appears to be emerging that the power spectrum resembles the schematic depicted in Fig.~\ref{fig:powerspec}, and that the inertial range sustains a form fairly close to the classic Kolmogorov spectrum. There may be some differences in the anisotropy of turbulence from the non-relativistic regime, and this is naturally expected due to the wave speeds approaching \teq{c}. In particular, the 3D wavenumber profiles of Weibel instability-driven turbulence near steep shocks inherently differ from those for Alfv\'en-like turbulence propagating remotely from shocks, and so the two should yield different anisotropies to the diffusion relative to the mean field direction. The details of particle diffusive transport in relativistic turbulence has not yet been fully explored. It is anticipated that scale-free power laws \teq{\lambda_{\parallel}\propto p^{\alpha}} may well be realized for some isolated ranges of momenta. The spatial diffusion tensor will indeed be anisotropic, yet the basic conclusion in our study of much greater turbulence near shocks than remote from them will not be sensitive to the microdetails of the tensor characteristics, nor more complicated forms for the \teq{\lambda_{\parallel}(p)} function. Our indication that blazar multiwavelength modeling leads to inferences of weak turbulence and long diffusive mean free paths for the most energetic leptons is not a consequence of the specific shock parameters chosen in our case studies. Varying the shock speed, magnetic field obliquity, temperature and compression ratio from the values in Table 1 will lead mostly to adjustments in \teq{\eta_1}, and thereby the efficiency of injection from the thermal population into the acceleration process. The derived values of \teq{\alpha >1} will not be impacted by varying these shock properties to any significant degree, since generating synchrotron turnovers in the optical or X-rays imposes such \teq{\lambda_{\parallel}/r_g\gg 1} values at \teq{\gammax}. Note that envisaging magnetic reconnection as an alternative scenario does not alter this circumstance. Reconnection is generally invoked to provide an acceleration mechanism (with \teq{E>B}) that is faster than the cyclotron frequency. While perhaps desirable when modeling the GeV gamma-ray flares detected by the {\it Fermi}-LAT in the Crab nebula, this is the opposite of what is needed for blazars. One can gauge the relative efficacy of reconnection versus gyroresonant acceleration using derived model parameters listed in Table~1. Since the trapping of electrons in relativistic reconnection zones in the PIC simulation of \citet{SS14} is effected by the merging of magnetic islands traveling at nearly \teq{c}, and the effective trapping ``radius'' is the inertial scale \teq{\sim c/\omega_p}, the acceleration rate for relativistic reconnection is of the order of \teq{\omega_p}. This is to be compared with the Bohm-limited gyroresonant acceleration rate \teq{\omegaB}. For non-relativistic cases, their ratio is just \teq{\sqrt{\sigma}}, where \teq{\sigma} is the plasma magnetization parameter in Eq.~(\ref{eq:sigma_param}), which is around unity for our Mrk 501 and AO 0235+164 fits, and somewhat larger than unity for BL Lac. Such an estimate is only slightly modified by the Lorentz factor \teq{\Gamma_1\lesssim 2} for our mildly-relativistic shocks. This tells us that magnetic reconnection is at most only slightly less efficient at accelerating charges in our blazar shocks than gyroresonant acceleration, an inference that would be impossible without our virtually unique ability to benchmark the electron number density relative to the non-thermal population. It is salient to comment upon potential diagnostics on field turbulence in blazars that can be afforded by radio and optical polarimetry. Polarization data require sufficient count statistics that generally sample only larger spatial regions, unless intense flaring provides sufficient flux to probe smaller scales. In the case of BL Lac, \citet{Gaur14} report monitoring in the V band spanning around 19 months from May 2008 until January 2010 that exhibits variable polarization degrees in the range \teq{\sim 5-30}\%. Similar, but slightly smaller values were observed in the R Band by \citet{Marscher08} during late 2005. Comparable polarization degrees have been discerned for other blazars, indicating the presence of moderately coherent fields. This ties more directly to an average MHD structure than it does to the micro-turbulence being considered here. The scales for the diffusion of all charges simulated in this study, even the most energetic ones, and also those of the jet turbulence, are too small to be captured via polarimetry, being far inferior to the sizes of causally-connected regions inferred from polarization variability estimates. No conflict emerges between the physical conclusions here of small levels of turbulence far from shocks, and the moderate field coherence deduced from the optical polarimetry of blazars. As a final remark, this picture of large diffusion mean free paths for the most energetic electrons present in blazar emission regions is not confined to our subset of case study sources, which were selected to span a significant range of blazar character. The conclusion of \teq{\eta (\gammax ) \gg 1} should be widely applicable to blazars of both the HBL and LBL synchrotron spectral varieties in their compact central regions. It is notable that a similar conclusion of such large \teq{\eta (\gammax )}, well above the Bohm regime, is made much further out for the radio quasar 4C 74.26. \citet{Araudo15} report that MERLIN radio interferometer imaging in the proximity of the southern radio lobe of this source has been able to map the bow shock zone and discern a very restricted spatial scale. The narrowness of the emission zone is far inferior to the radiation-reaction limited synchrotron cooling length one infers for the shock zone by equating estimates like Eqs.~(\ref{eq:cool_rate}) and~(\ref{eq:acc_rate}). This leads \citet{Araudo15} to conclude both that wave damping must inhibit the spatial scale of amplified, turbulent magnetic field, and that energetic charges at \teq{\gamma \sim \gammax} will possess diffusive mean free paths \teq{\lambda_{\parallel}\gtrsim 10^6r_g}, an inference concordant in character with that for our blazar studies here. | This paper has offered a deep exploration of multi-wavelength spectral fits to flaring emission from the blazars Mrk 501, BL Lacertae and the Bl Lac object AO 0235+164. The modeling uses complete thermal plus non-thermal electron distribution functions obtained directly from Monte Carlo simulations of diffusive acceleration at relativistic shocks that span over six decades in energy, a material advantage over plasma codes. These simulations presume a planar shock interface and phenomenologically describe the diffusive mean free path of charges' interactions with MHD turbulence in a two-dimensional (tensor) construction, omitting contributions to particle transport that might emerge from complicated magnetic structures in the shock layer, for example rippling of the interface. The importance of treating the intimate connection between non-thermal and thermal leptons in the jet is highlighted by the successful modeling of the ephemeral X-ray component in AO0235+164 as a bulk Comptonization signature of the thermal population. Our integrated application of simulated relativistic electron populations to blazar spectral modeling provides fundamentally new insights into the plasma environment of extragalactic jets. In particular, the multiwavelength model fits specify the values of the magnetic field, plasma density, plasma frequency and Alfv\'en speed concomitantly. Restriction to one-zone constructions is made in order to reduce the number of model parameters. We believe that any extension of this protocol to encapsulate multiple radiation zones, perhaps desirable for treating the BL Lac gamma-ray signal, will not alter the principal inferences we forge concerning the blazar jet environment. The main conclusions are as follows. In order to model both the low synchrotron turnover energy in the optical and X-ray bands, and the prominent hard gamma-ray flux seen by {\it Fermi} and Atmospheric \v{C}erenkov Telescopes at even higher energies, leptonic models need the diffusive mean free path to increase rapidly with electron momentum, \teq{\lambda_{\parallel}\propto p^{\alpha}} with \teq{\alpha \gtrsim 3/2}, starting with fairly modest values \teq{\lambda_{\parallel}/r_g \sim 20-100} at low, quasi-thermal momenta. This establishes mean free paths orders of magnitude larger than electronic gyroradii at the Lorentz factors that generate the synchrotron and SSC \teq{\nu F_{\nu}} peak emission. This determination is robust in that dividing the jet region into several zones will not ameliorate this constraint on the \teq{\lambda_{\parallel}(p)} function, since it is principally dictated by the synchrotron component. We argue here that this strong momentum dependence of \teq{\lambda_{\parallel}/r_g} is not unexpected, due to somewhat constrained inertial ranges of plasma turbulence, and associated non-gyroresonant diffusion. It is perhaps also due to a decline in the level of such turbulence away from shocks that inject energetic particles into the much larger blazar jet emission zones. Such a spatial non-uniformity in jet plasma turbulence spanning small injection sites to larger dissipative zones in some ways parallels the much more intimately-observed environment of the active solar wind. Future investigative programs naturally indicated by this study include detailed assessment of the particle diffusion characteristics, and their variation with particle momenta, in relativistic MHD turbulence. \newpage | 16 | 9 | 1609.03899 |
1609 | 1609.04688_arXiv.txt | The current status of the phenomenology of short-baseline neutrino oscillations induced by light sterile neutrinos in the framework of 3+1 mixing is reviewed. | 16 | 9 | 1609.04688 |
||
1609 | 1609.01252_arXiv.txt | In this work, we present results from a long-term precipitable water vapor (PWV) study in the Chajnantor area, in northern Chile. Data from several instruments located at relevant sites for sub-millimeter and mid-infrared astronomy were processed to obtain relations between the atmospheric conditions among the sites. The data used for this study can be considered the richest dataset to date, because of the geographical sampling of the region, including sites at different altitudes, a time span from 2005 to 2014, and the different techniques and instruments used for the measurements. We validate a method to convert atmospheric opacity from 350 $\mu$m tipper radiometers to PWV. An average of 0.68 PWV ratio between Cerro Chajnantor and Llano of Chajnantor was found. | \hspace{4mm} The Chajnantor area is considered one of the best sites in the world for millimeter, submillimeter and mid-infrared astronomy due to the combination of dryness, i.e. low precipitable water vapor (PWV) and the notable high altitude above $5000$ meters. The area has been selected as the location for world-class observatories, such as ALMA (\cite{wotten09}), APEX (\cite{gusten06}), CBI (\cite{padin02}), TAO (\cite{motohara11}), and others. A study of the tropospheric distribution of PWV over the area is of relevance to determine the optimum altitude for new observatories that can impact project costs significantly. Another reason to support such a study is to observe temporal drifts and cycles in the atmospheric conditions for the site, as well as contributing with the understanding of calibrations and cross-comparisons between instruments located in the area. There are a number of publications in the literature referring to PWV and weather conditions in the Chajnantor area and their implications to astronomy. For example, in \cite{delgado99}, several empirical relations for atmospheric variables are derived to understand the climatology in the Chajnantor area and how to obtain the amount of PWV from radiometric measurements. Long-term studies using climatological data for the Chajnantor area was presented by \cite{bustos00} and \cite{otarola05}. The phase correction by using the PWV at the line of sight in interferometric mode at ALMA was studied by \cite{nikolic13}. PWV estimates and forecasts for the Chajnantor area using GOES satellites were produced by \cite{marin15}. An atmospheric measurement campaign was performed by \cite{turner10} on Cerro Toco in 2009. The atmospheric transparency has been intensively studied by \cite{radford00}, \cite{giovanelli01}, \cite{radford11}, and \cite{radford16}. In addition, the differences in PWV between the Chajnantor Plateau and Cerro Chajnantor was shown in \cite{bustos14} for a time span of 5 days. The amount of PWV obtained by each instrument located in the Chajnantor area is measured independently and is only valid for that location and time. There are no studies involving the spatial distribution of PWV over the Chajnantor area and its temporal variations, nor long term relations that connect the amount of PWV at different locations and measured with different instruments. In this work, we have derived a new method to convert atmospheric opacity (optical depth) from $350 \,\mu m$ tipper radiometers to PWV, using the AM spectroscopy/atmospheric model (\cite{paine16}). We have aimed this study to understand the ratio of PWV between the Chajnantor Plateau (5080\,m of altitude), Cerro Toco (5320\,m), and the summit of Cerro Chajnantor (5612\,m). \subsection{Instruments} \hspace{4mm} Different instruments have been located on the Chajnantor area over the years. These instruments differ in their technical characteristics and specifications, operations approach, location, altitude, and also in the observables they measure. The instruments used in this study, their locations, and time span are detailed in Table \ref{tab-data-inst}. Water vapor radiometers (WVR) provide a measure of the atmospheric brightness temperature of a selected water molecule vibro-rotational absorption/emission band. Then, an atmospheric model is used to compute PWV. This technique is used by APEX, RHUBC-II and UdeC, as shown in Table \ref{tab-data-inst}. The APEX telescope includes a 183 GHz water vapor radiometer installed in the Cassegrain cabin. The instrument measures PWV whenever the telescope shutter is open. We have decided to use the data from this instrument as the main reference for comparisons among the selected sites. On the other hand, the PWV at mid-infrared was estimated from atmospheric extinction using photometry, as for the TAO data. The submillimeter tippers measure sky brightness temperature at several zenith angles and fit for optical depth, at a certain bandwidth defined by the optical filter at the input of the instrument. \subsection{Software} The computational tools used to produce all graphs, as well as derive the physical relationships and results include: AM 9.0 (Atmospheric Model, \cite{paine16}), Python version 2.7.2, and TOPCAT (Tool for OPerations on Catalogues And Tables) version 4.1. \begin{table*}[HB!] \caption{Instruments used in this study and their location, altitude, observable and time span. }% \centering \begin{tabular}{|c |c | c | c | c | c | c |} \hline Instrument & Location & Altitude (m) & Observable & Time span & ID \\ & & & & & \\ \hline \hline APEX & Chajnantor & 5107 & PWV & 2006 to 2014 & APEX \\ radiometer & plateau & & from WVR at 183 GHz& & \\ \hline TAO & Cerro Chajnantor & 5640 & PWV from mid-infrared & 2009 to 2011 & TAO \\ & summit & & astronomical observations & & \\ \hline RHUBC-II & Cerro & 5320 & PWV from WVR & August to October & TOCO \\ & Toco & & at 183 GHz & 2009 & \\ \hline & Chajnantor & 5080 & Opacity & 1997 to 2005 & TA-1 \\ & Plateau (NRAO)& & at 350 $\mu m$ & & \\ & & & & & \\ Tipper & Chajnantor & 5080 & Opacity & 2005 to 2010 &TA-2 \\ radiometer A & plateau (CBI)& & at 350 $\mu m$ & & \\ & & & & & \\ & Chajnantor & 5107 & Opacity & 2011 to now & TA-3\\ & plateau (APEX) & & at 350 $\mu m$ & & \\ \hline & Chajnantor & 5080 & Opacity & 2000 to 2005 & TB-1 \\ & Plateau (NRAO) & & at 350 $\mu m$ & & \\ & & & & & \\ Tipper & Chajnantor & 5080 & Opacity & 2005 to 2009 & TB-2 \\ radiometer B & plateau (CBI) & & at 350 $\mu m$ & & \\ & & & & & \\ & Cerro & 5612 & Opacity & 2009 to now &TB-3 \\ & Chajnantor & & at 350 $\mu m$ & & \\ & summit & & & & \\ \hline & Cerro & 5612 & PWV & 2011 & UdeC \\ UDEC & Chajnantor & & at 183 GHz & & \\ & summit & & & & \\ \hline \end{tabular} \label{tab-data-inst} \end{table*} | This work collected $> 3\cdot 10^{6}$ data coming from different instruments and locations over the Chajnantor area and for long periods of time. The AM v9.0 has been recently updated and the use of the modified version prove crucial to get the results shown here. The method used in this work to convert atmospheric opacity from the sub-mm tippers to PWV, has been validated by the instruments at these sites and turned out to be a useful tool for atmospheric studies in the Chajnantor area. The tipper converted PWV data was calibrated using APEX PWV measurements. When co-located at the Plateau, and once the data was properly calibrated, both tipper radiometers measured the same amount of PWV. This fact allowed us to use the tipper measurements when the instruments were taken to different locations in the area. The CBI site shows 7\% excess of PWV compared to the APEX site. This difference is due to the altitude difference between the sites. For Cerro Chajnantor summit, the PWV ratio with the Plateau is 0.68, which indicates 32\% less PWV compared to the Plateau and for our long term study. Although this value is close to previous results, i.e. cite{bustos14} and \cite{radford16}, our study is longer term and therefore averages out temporal variabilities in the atmospheric quantities. Cerro Toco is a interesting site to study the amount of PWV since it is located between the Chajnantor Plateau and Cerro Chajnantor summit. Our result for PWV ratio between Toco and the Plateau is subject to selection bias, as explained in the text. More data for this site would be needed to estimate and clarify the observed PWV ratios between Toco and APEX. For the interested reader, the AM configuration files used in this paper can be requested by e-mail to F. Cort\'es ([email protected]). | 16 | 9 | 1609.01252 |
1609 | 1609.01314_arXiv.txt | \noindent The Qatar-2 transiting exoplanet system was recently observed by the {\it Kepler} as part of {\it K2} Campaign 6 in short-cadence mode. We identify dozens of starspot-crossing events, when the planet eclipsed a relatively dark region of the stellar photosphere. The observed patterns of these events demonstrate that the planet always transits over the same range of stellar latitudes, and therefore that the stellar obliquity is less than about 10$^\circ$. We support this conclusion with two different modeling approaches: one based on explicit identification and timing of the events, and the other based on fitting the light curves with a spotted-star model. We refine the transit parameters and measure the stellar rotation period ($18.5 \pm 1.9$~days), corresponding to a 'gyrochronological' age of $1.4 \pm 0.3$ Gyr. Coherent flux variations with the same period as the transits are well modeled as the combined effects of ellipsoidal light variations ($15.4 \pm 4.8$~ppm) and Doppler boosting ($14.6 \pm 5.1$~ppm). The magnitudes of these effects correspond to a planetary mass of $2.6 \pm 0.9~M_{\text{Jup}}$ and $3.9 \pm 1.5~M_{\text{Jup}}$, respectively. Both of these independent mass estimates agree with the mass determined by the spectroscopic Doppler technique ($2.487 \pm 0.086~M_{\text{Jup}}$). No occultations are detected, giving a 2$\sigma$ upper limit of 0.06 on the planet's visual geometric albedo. We find no evidence for orbital decay, although we are only able to place a weak lower bound on the relevant tidal quality factor: $Q'_\star > 1.5\times 10^4$~(95\% confidence). | The obliquity of a planet-hosting star (the angle between the star's rotation axis and orbit normal) may bear information about a planet's formation, migration and tidal evolution history \citep{Queloz2000, Ohta2005, Gaudi2007, Winn2010}. For example, dynamically hot scenarios for hot Jupiter formation, such as planet-planet scattering \citep{Chatterjee2008} and Kozai-Lidov mechanism \citep{Fabrycky2007}, should often produce large obliquities. Dynamically cold scenarios such as disk migration \citep{Lin1996} and {\it in situ} formation \citep{Batygin2015} should preserve low obliquities, unless there are mechanisms for exciting obliquities independently of hot-Jupiter formation \citep[e.g.,][]{Bate2010, Batygin2012}. One way to determine the stellar obliquity --- or, to be more precise, to recognize when the obliquity is low --- is to observe a sequence of flux anomalies that occur when a transiting planet repeatedly passes in front of a starspot. The analysis of these ``starspot-crossing anomalies'' takes advantage of the precise time-series photometry that is available for the systems that have been observed by the {\it Kepler} and {\it CoRoT} spacecraft. This method does not require intensive time-series spectroscopy, unlike the more traditional method based on the Rossiter-McLaughlin effect, which is often difficult to apply to relatively faint or slowly-rotating stars. \citet{Silva2003} anticipated the observable signal of a transiting planet crossing over a starspot: the loss of light is temporarily reduced, because the starspot has a lower intensity than the surrounding photosphere. This produces a brief flux enhancement or ``bump'' in the transit light curve. It soon became clear that spot-crossing anomalies can be used to study the properties of starspots \citep{Silva-Valio2010}, demonstrate the presence of active latitudes \citep{Sanchis-Ojeda2011Hat} and constrain the stellar obliquity \citep{Sanchis-Ojeda2011Wasp,Nutzman2011}. Qatar-2b is a hot Jupiter with a mass of 2.5~$M_{\text{Jup}}$, a radius of 1.1~$R_{\text{Jup}}$, and an orbital period of 1.34~days. It was discovered by the Qatar Exoplanet Survey \citep[QES,][]{Bryan2012}. The host star Qatar-2A is a relatively bright K dwarf ($V = 13.3$, $M_\star = 0.740 \pm 0.037 ~M_{\odot}$). Radial velocity follow-up revealed the presence of a long-term trend which was attributed to a more distant companion. \citet{Mancini2014} constrained the obliquity of Qatar-2b using spot-crossing anomalies seen in the ground-based multi-color transit observations. However, the stellar rotation period was unknown at the time of their analysis. Without the ability to calculate the rotational phase of each transit, \citet{Mancini2014} had to make the assumption that two particular spot-crossing anomalies they observed were caused by eclipses of the same spot. With this assumption, they found the stellar rotation period to be $14.8 \pm0.3$ days \citep[after the correction described by] []{Mancini2016}, and the sky-projected obliquity (the angle between the sky projections of the stellar rotation axis and the orbit normal) to be $\lambda = 4.3 \pm 4.5^{\circ}$. Qatar-2 was within the field of view of {\it K2} Campaign 6. Being a confirmed planet, Qatar-2 was selected for 1~min (``short-cadence'') time sampling, instead of the usual 30~min sampling. The precise, continuous and well-sampled {\it K2} photometric data provides an opportunity to study Qatar-2b in greater detail. As we will show, the {\it K2} data reveal the stellar rotation period to be $18.5 \pm 1.9$ days, at odds with the period determined by \citet{Mancini2014}. Moreover, the {\it K2} data show evidence for numerous spot-crossing anomalies caused by different spot groups. This leaves little room for doubt in the interpretation of these events, and the conclusion that the stellar obliquity is low. The short-cadence data also allow for better resolution of the ingress/egress phases of the transit, leading to improved estimates of the basic transit parameters. The data can also be searched for occultations, which would reveal the planet's dayside brightness; and for ellipsoidal variations (ELV) and the effects of Doppler boosting (DB), the amplitudes of which can be used to make independent estimates of the planetary mass. Finally, the continuous sequence of transit times permits a search for any variations in the intervals between transits, which could be caused by additional orbiting bodies or tidal effects. The paper is organized in the following way. Section \ref{sec: photo} describes our reduction of {\it K2} data. Section \ref{sec: refine} lays out the analysis of the light curve and the refinement of transit parameters. Section \ref{sec: ttv} presents a search for changes in the transit period. Section \ref{sec: rot} discusses the measurement of the stellar rotation period, and the associated ``gyrochronological'' age. Section \ref{sec: elv} presents the search for occultations, ELV, and DB effects. Section \ref{sec: anomalies} presents the analysis of spot-crossing anomalies and the implications for the stellar obliquity. Finally, Section \ref{sec: dis} summarizes and discusses all our findings. While this work was in the final stages of preparation, we became aware of the work of \citet{Mocnik}, who performed a similar analysis of the same data. Our study and their study have reached similar conclusions regarding the stellar obliquity, stellar rotation period, transit-timing results, and flux modulation outside of transits. Some small differences exist in the quantitative results, which we describe in the appropriate sections. | \label{sec: dis} In this work, we presented the analysis of the {\it K2} short-cadence observation of Qatar-2. The continuous monitoring, high precision and high cadence of the {\it K2} data helped to refine the transit parameters. In addition, the data quality was high enough to facilitate the identification and exclusion of data points affected by spot-crossing anomalies, leading to a less biased set of transit parameters. We measured the stellar rotation period of Qatar-2A $18.5 \pm 1.9$ days based on the out-of-transit flux variation of the {\it K2} light curve. Using the technique of gyrochronology, the rotation period led to an independent estimate of the stellar age, $1.4 \pm 0.3$ Gyr. The rotation period also played a crucial role in our obliquity determination; the lack of an independently measured rotation period had been a missing piece of the puzzle in a previous effort to determine the stellar obliquity. The nondetection of the secondary eclipse allowed us to place a constraint on the planet's geometric albedo in the {\it Kepler} bandpass: $A_g < 0.06$, with 95\% confidence. This is consistent with previous investigations that showed "hot Jupiters'' often have low albedos \citep{Esteves2015, Gandolfi2013, Kipping2011}. We detected the ellipsoidal light variation and Doppler boosting effects in the {\it K2} light curve, after filtering out long-term stellar variability and systematic effects. The magnitudes of these two effects imply planetary masses of $2.6 \pm 0.9~M_{\text{Jup}}$ and $3.9 \pm 1.5~M_{\text{Jup}}$, both of which are consistent with the mass determined from the spectroscopic Doppler technique \citep{Bryan2012}. We have updated the ephemerides of Qatar-2b with the new mid-transit times observed by {\it K2}. There is no evidence for orbital decay, leading to a lower bound on the stellar tidal quality factor $Q'_\star >1.5\times 10^{4}$~(95\% confidence). We identified dozens of spot-crossing anomalies in the {\it K2} light curve. These anomalies revealed the presence of active regions on the host star along the planet's transit chord. This suggests that Qatar-2 is magnetically active, as one would expect for a star with the relatively young age that was determined from gyrochronology. We used the observed spot-crossing anomalies to demonstrate that the the obliquity of Qatar-2 is very likely smaller than 10$^\circ$. We did this in two different ways. First we identified individual spot-crossing anomalies and measured their properties, including their times of occurrence. We then used a simple geometric model for which the parameters were determined by requiring spatial coincidences of the spot and the planet at the times of observed anomalies. In a separate approach, we fitted a photometric model to a portion of the light curve, based on the premise of a planet transiting a limb-darkened star with a circular starspot. Neither model can be relied upon for precise quantitative results, because of the strong assumptions that were made, such as the circular shape of the spots and the lack of spot evolution. Nevertheless the qualitative results leave little room for doubt that the obliquity is lower than 10$^\circ$. A low obliquity for Qatar-2 is consistent with a pattern that has been prevoiusly noted: the hot Jupiter hosts with photospheres cooler than about 6100-6300~K tend to have low obliquities \citep{Winn2010}. | 16 | 9 | 1609.01314 |
1609 | 1609.00915_arXiv.txt | {\small {\bf Summary:}\ Research field of celestial mechanics. Historical overview: apparent motion of planets, and solar and lunar eclipse as impetus for celestial mechanics. Ancient celestial mechanics. Appolonius and the idea of epicyclic motion. Ptolemy and the geocentric system. Copernicus and the heliocentric system. Kepler and the three Kepler laws. Galileo: satellites of Jupiter as a model for the Solar system, the begin of mechanics. Newton: mathematical formulation of mechanics, gravitational force. Einstein: the problem of perihelion advance of Mercury and the general theory of relativity. Three aspects of celestial mechanics: physics of motion, mathematics of motion and (numerical) calculation of motion. The astronomical objects and specific goals and problems of the modelling of their motion: artificial satellites, the Moon, major planets, asteroids, comets, Kuiper belt objects, satellites of the major planets, rings, interplanetary dust, stars in binary and multiple systems, stars in star clusters and galaxies. }\bigskip \chapter{Two-body Problem} \label{chapter-two-body-problem} | {\small {\bf Summary:}\ Calculation of position and velocity from the Kepler elements. Calculation of the Kepler elements from the position and velocity. Orbit determination (an overview). }\bigskip We have seen above that there are two equivalent ways to represent a particular two body motion: (1) to specify the initial conditions for the equation of motion, i.e. the position and velocity vectors $\ve{r}=(x,y,z)$ and ${\dot \ve{r}}=(\dot x,\dot y,\dot z)$ together with the corresponding moment of time $t_0$ and the parameter $\kappa$, and (2) to fix the whole set of the six Kepler elements $a$, $e$, $i$, $\omega$, $\Omega$, $M_0=M(t_0)$ again together with the moment of time $t_0$ for which the mean anomaly $M_0$ is supposed to be known and the parameter $\kappa$. Very often in the practical calculation one wants to switch between these two representations, that is to transform the position and velocity into the corresponding Kepler elements or vice verse. Here we give the set of formulas enabling one to perform these transformations for the case of elliptic motion. The transformation from the Kepler elements to the position and velocity vectors can be done in the following way: \begin{enumerate} \item calculate mean motion as $n=\kappa\,a^{-3/2}$ and mean anomaly as $M=n\,(t-t_0)+M_0$ (here the position and velocity vectors can be calculated for any arbitrary moments of time $t$, not necessarily for the moment $t_0$ for which the mean anomaly $M_0$ is specified), \item calculate eccentric anomaly $E$ from $E-e\sin E=M$, \item calculate position and velocity in the orbital plane: \begin{eqnarray} X&=&a\,(\cos E-e), \nonumber\\ Y&=&a\,\sqrt{1-e^2}\,\sin E, \nonumber\\ \dot X&=&-{a\,n\,\sin E\over 1-e\,\cos E}, \nonumber\\ \dot Y&=&{a\,n\,\sqrt{1-e^2}\,\cos E\over 1-e\,\cos E}, \nonumber \end{eqnarray} \item calculate the position and velocity vectors in space as \begin{eqnarray} \pmatrix{x\cr y\cr z}&=&P\,\pmatrix{X\cr Y\cr 0}, \nonumber\\ \nonumber\\ \nonumber\\ \pmatrix{\dot x\cr \dot y\cr \dot z}&=&P\,\pmatrix{\dot X\cr \dot Y\cr 0}, \nonumber \end{eqnarray} \noindent where $P=\mat{A}^T_z(\Omega)\,\mat{A}^T_x(i)\,\mat{A}^T_z(\omega)$ defined by (\ref{P-matrix}). \end{enumerate} The transformation from the position and velocity vectors to the Kepler elements is a bit more complicated and can be done as follows: \begin{enumerate} \item from the integrals of the areas $\ve{r}\times{\dot \ve{r}}=\ve{c}$ one gets $\ve{c}=(c_x,c_y,c_z)$ \begin{eqnarray} c_x&=&y\,\dot z-\dot y\, z, \nonumber\\ c_y&=&z\,\dot x-\dot z\, x, \nonumber\\ c_z&=&x\,\dot y-\dot x\, y, \nonumber\\ c&=&|\ve{c}|=\sqrt{c_x^2+c_y^2+c_z^2}. \nonumber \end{eqnarray} \noindent Then the semi-latus rectum can be calculated as \begin{equation} p={c^2\over \kappa^2}, \nonumber \end{equation} \noindent and from \begin{eqnarray} \pmatrix{c_x\cr c_y\cr c_z}=\mat{A}^T_z(\Omega)\,\mat{A}^T_x(i)\,\pmatrix{0\cr 0\cr c} = \left(\begin{array}{rrr} c\sin i\sin\Omega\\ -c\sin i\cos\Omega\\ c\,\cos i \end{array} \right) \nonumber \end{eqnarray} \noindent which gives us three equations. The third equation can be written as \begin{equation} \cos i = {c_z\over c} \end{equation} \noindent and since $0\le i\le\pi$ this one equation is sufficient to calculate the inclination $i$. The two other equations read \begin{eqnarray} \sin\Omega&=&{c_x\over\sqrt{c_x^2+c_y^2}}, \nonumber\\ \cos\Omega&=&-{c_y\over\sqrt{c_x^2+c_y^2}} \end{eqnarray} \noindent and allow one to calculate $\Omega$. Note that if $c_x^2+c_y^2=0$, the inclination $i=0$ and $\Omega$ is not defined. \item From Eq. (\ref{conic-section}) and $v=u-\omega$ one gets two equations \begin{eqnarray} e\,\cos v&=&{p\over r}-1, \nonumber\\ e\,\sin v&=&{\sqrt{p}\over \kappa}\,{\ve{r}\cdot{\dot \ve{r}}\over r} \end{eqnarray} \noindent which can be used to calculate both the eccentricity $e$ and the true anomaly $v$. Then using (\ref{tan-v-E}) one can calculate the eccentric anomaly $E$, and from (\ref{Kepler-in-E-int}) the mean anomaly $M$. All these values of anomalies $v$, $E$ and $M$ correspond to time $t_0$ for which the position and velocity of the body is specified. Finally, from $p$ and $e$ it is easy to calculate the semi-major axis as $a=p\,{(1-e^2)}^{-1}$; \item From \begin{eqnarray} \pmatrix{x\cr y\cr z}=\mat{A}^T_z(\Omega)\,\mat{A}^T_x(i)\, \pmatrix{r\cos (v+\omega)\cr r\sin (v+\omega)\cr 0} \nonumber \end{eqnarray} \noindent one gets \begin{eqnarray} \label{cos-sin-v+omega} \cos (v+\omega)&=&{x\over r}\,\cos\Omega+{y\over r}\,\sin\Omega, \nonumber\\ \sin (v+\omega)&=&\left(-{x\over r}\,\sin\Omega+{y\over r}\,\cos\Omega\right)\, \cos i+{z\over r}\,\sin i. \end{eqnarray} \noindent From these two equations one calculates the angle $v+\omega$ and since $v$ is known, the argument of perihelion $\omega$. \end{enumerate} | 16 | 9 | 1609.00915 |
|
1609 | 1609.05192_arXiv.txt | We present an X-ray photometric analysis of six gravitationally lensed quasars, with observation campaigns spanning from $5$ to $14$ years, measuring the total ($0.83-21.8$ keV restframe), soft ($0.83-3.6$ keV), and hard ($3.6-21.8$ keV) band image flux ratios for each epoch. Using the ratios of the model-predicted macro-magnifications as baselines, we build differential microlensing light curves and obtain joint likelihood functions for the average X-ray emission region sizes. Our analysis yields a Probability Distribution Function for the average half-light radius of the X-Ray emission region in the sample that peaks slightly above $1$ gravitational radius and with nearly indistinguishable $68\%$ confidence (one-sided) upper limits of $17.8$ and $18.9$ gravitational radii for the soft and hard X-ray emitting regions, assuming a mean stellar mass of $0.3$ $M_{\odot}$. We see hints of energy dependent microlensing between the soft and hard bands in two of the objects. In a separate analysis on the root-mean-square (RMS) of the microlensing variability, we find significant differences between the soft and hard bands but the sign of the difference is not consistent across the sample. This suggests the existence of some kind of spatial structure to the X-ray emission in an otherwise extremely compact source. We also discover a correlation between the RMS microlensing variability and the average microlensing amplitude. | X-ray emission is one of the defining characteristics of active galactic nuclei (AGN). However, most properties of the X-ray corona are obtained only through spectral analyses, since neither current nor near-future instrumentation can resolve the X-ray emitting regions of AGN. Reverberation mapping and quasar microlensing provide the only probes of the spatial structure of the different AGN components, with the latter better suited to the more compact regions like the X-ray corona or the accretion disk. Reverberation mapping studies have succeeded in mapping the more spatially extended regions such as the broad line regions \citep[e.g., ][]{Bentz2009,Zu2011,Bentz2013,Kollatschny2014}, the dust torus \citep[e.g., ][]{Koshida2014}, and with limited results for accretion disks \citep[e.g., ][]{Shappee2014, Fausnaugh2016}. Microlensing refers to the micro-arcsecond effects produced by light ray deflections of emision from a background source by foreground stars. It has the advantage over reverberation mapping in that the signal only becomes stronger as the source becomes more compact. The Einstein radius gives a typical scale of \begin{equation} \label{eq:AE} R_{E}=D_{ol}\sqrt{\frac{4GM}{c^{2}}\frac{D_{ls}}{D_{ol}D_{os}}}, \end{equation} where $M$ denotes the deflector mass and $D_{ol}$, $D_{os}$, and $D_{ls}$ are the angular diameter distances between the observer, lens, and source respectively. If the apparent size of the source is comparable or smaller in size than the Einstein radius, typically a few light-days, the observed flux varies because the magnification changes as the source, lens, and observer move relative to each other \citep[see, e.g., the review by][]{Wambsganss2006}. This makes extragalactic microlensing a unique tool for probing the spatial structure of the central region of quasars, because most AGN components are comparable in size to the Einstein radius or smaller. As a result, microlensing has been successfuly used to obtain size estimates for the broad line region spanning several tens of light days \citep[e.g., ][]{Sluse2012,Guerras2013a}, the accretion disk spanning $\sim$10 light-days \citep[e.g., ][]{Morgan2008,Mediavilla2011a,Jorge2014}, and the X-ray corona of $\sim$1 light-day \citep[e.g., ][]{Dai2010,Pooley2012,Morgan2012,Mosquera2013,Blackburne2014,Blackburne2015,Macleod2015}. Here, we present updated X-ray light curves for a sample of $6$ lensed quasars with redshifts between $z_{s}=1.3$ and $z_{s}=2.3$. We derive total, soft, and hard energy band light curves in Section~\ref{sec:photometries} and examine them for evidence of microlensing. We perform a simple analysis of several aspects of the microlensing variability in Section~\ref{sec:microlensing_analysis}. In Section~\ref{sec:source_size_estimates} we derive a Probability Distribution Function (PDF) for the average size of the X-ray emitting region in the sample. Section~\ref{sec:discussion} presents a summary of the results. We assume a flat $\Lambda$CDM cosmology with $H_0 = 70~{\rm km~s^{-1} Mpc}^{-1}$, $\Omega_m = 0.3$, and $\Omega_\Lambda=0.7$.% | \label{sec:discussion} We have measured full, soft and hard band X-ray light curves for $6$ lensed quasars to look for microlensing by comparing the observed flux ratios with the ratios predicted by macro lens modeling. We have tested for energy-dependent variability in several ways: a $\chi^2$ fit to the light curves of quadruple lenses, a comparison between the microlensing amplitude RMS of the soft and hard bands, and estimates of the average source size in the full, hard and soft X-ray energy bands. Our $\chi^2$ test for energy-dependent microlensing shows a lack of correlation between soft and hard band in 2 of the 4-image lenses. This can be explained by a size difference between emitting regions, but also by a lack of correlation in the time domain. The RMS of the microlensing variability between the hard and soft bands is significantly different for a number of image pairs, but the sign of the difference varies and shows no consistent pattern. If a higher RMS is interpreted as arising from a more compact hard X-ray source, then this picture is consistent with the recent review of a sample of $8$ lensed quasars by \cite{Chartas2016}, where for some objects the hard X-ray emission regions seems to be more compact than the soft and in others the soft appears to be smaller. However, a physical interpretation of a higher RMS level of microlensing variability in one band might not be straightforward, as suggested by the inconsistencies shown in Table~\ref{tab:mean_rms_table} among image pairs of the same quadruple objects. An analysis on this question is in preparation (Guerras et al. 2017). Our estimates of the average source size indicate that any size difference between the hard and soft emitting regions must be modest. This is in good agreement with the general picture that emerges from fully time-dependent studies of individual objects. \cite{Blackburne2015} found the same upper limit for the size of the hard and soft X-ray emitting regions in HE 1104$-$1805, as did \cite{Morgan2012} for QJ 0158$-$4325. \cite{Mosquera2013} could find only ``weak evidence'' that the hard X-ray emitting region in Q 2237+030 was more compact than the soft X-ray emitting region, and in HE 0435$-$1223 \cite{Blackburne2014} found no evidence for a size difference. The physical structure of what is believed to be a hot corona responsible for the X-ray continuum in quasars is poorly understood, and there are other astrophysical examples where hotter does not necessary equal smaller (e.g., the Solar corona). There is even some evidence \citep{Chartas2012} suggesting that the soft emitting region could be more compact at least in one case. The X-ray light curves span time intervals comparable to typical microlensing time scales, Treating them as single-epoch observations would result in overestimated source sizes. One way to address this problem is with fully time dependent calculations \citep{Kochanek2004}, but these are very computationally expensive. Here we introduce a simple approximation which at least avoids the bias of the single epoch method. This approximation essentially consists of introducing probability distribution modelled from time averaged tracks across the magnification patterns instead of isolated data points. The size estimates are consistently smaller using this approach. If we do not include the effects of time averaging, the source size estimate increases by a factor of $1.7$, which is a significant bias in the size estimate. We also introduced a radial model for the microlensing optical depth based upon the best fit to real data given by \cite{Oguri2014}. This model is simpler than the de Vaucouleurs stellar distribution plus NFW dark-matter halo used in previous time-dependent studies on individual objects \citep[e.g.,][]{Morgan2008, Dai2010}, yet it is an improvement with respect to using a uniform value in previous single-epoch studies over a heterogeneous sample of quasars \citep[e.g.,][]{Guerras2013b, Jorge2015}, where it is desirable to do the least possible assumptions on the lens galaxies to give a uniform treatment to all objects in the sample. We also find a functional relationship between the RMS and the average value of the microlensing amplitude. This suggest that both observables carry physical information (e.g., about the quasar source size or the optical depth in galactic halos), and a more detailed analysis of this correlation is in preparation (Guerras et al. 2017). These two observables could be used complementarily to constraint physical properties from microlensing variability for a better understanding of lensed quasars. | 16 | 9 | 1609.05192 |
1609 | 1609.09469_arXiv.txt | We present a description of cosmic neutrinos as a dispersive fluid. In this approach, the neutrino phase space is reduced to density and velocity fields alongside a scale-dependent sound speed. This sound speed depends on redshift, the initial neutrino phase space density and the cold dark matter gravitational potential. The latter is a new coupling between neutrinos and large scale structure not described by previous fluid approaches. We compute the sound speed in linear theory and find that it asymptotes to constants at small and large scales regardless of the gravitational potential. By comparing with neutrino N-body simulations, we measure the small scale sound speed and find it to be lower than linear theory predictions. This allows for an explanation of the discrepency between N-body and linear response predictions for the neutrino power spectrum: neutrinos are still driven predominantly by the cold dark matter, but the sound speed on small scales is not stable to perturbations and decreases. Finally, we present a calibrated model for the neutrino power spectrum that requires no additional integrations outside of standard Boltzmann codes. | Recently, \citet{bib:Banerjee2016} performed numerical simulations treating neutrinos as a fluid. They evolved the density and velocity fields using the continuity and Euler equations and estimated the full position-dependent stress tensor from N-body neutrino particles. While this is the most accurate way to close the neutrino hierarchy equations, other possibilities exist including those discussed here. In Fig. \ref{fig:data_cs} we demonstrated that, on small scales, the non-linearity in the sound speed is due to the neutrino density, not its stress. In addition, on large scales the behaviour becomes more linear and the sound speed is unimportant as $j_0(k c_s (s-s')) \simeq 1$. We speculate that it may be sufficient to close the hierarchy equations using an approach analagous to \cite{bib:AliHaimoud2012} but using linear response to compute $\Pi$ instead of $\delta$. The benefits of such a scheme are significant compared to N-body. For instance, in our particle implementation there are $N^3$ neutrinos and $(N/2)^3$ CDM particles, each requiring six 4-byte floats. Hence, neutrinos are allocated $8/9$ of the available memory. On the other hand, in a grid based implementation there could be $N^3$ CDM particles, and two grids with $(N/2)^3$ cells each requiring one 4-byte integer. In this case neutrinos only require $4/100$ of the memory available. In addition, less computational time could be spent on neutrinos (due to the simplified hydrodynamic structure) and more time on the CDM. Finally, as in \cite{bib:Banerjee2016}, the neutrinos could be simulated starting at a high redshift (compared to our N-body implementation which starts them at $z\le10$). If the redshift is too high (e.g. above the neutrino relativistic to non-relativistic transition), this approach does not accurately describe neutrinos (which would be significantly relativistic) but the calculation would at least be self-consistent and such high redshift discrepencies are unlikely to propagate to late time effects. Despite these benefits, a dispersive fluid approach would require extensive comparisons to N-body results to calibrate the sound speed and validate the results. \end{section} \begin{section}{Conclusion} \label{sec:conclusion} We have considered neutrinos as a dispersive fluid and found that this provides additional physical insights into their clustering behaviour. We have computed the sound speed and shown that it depends on the intial neutrino velocity distribution and also the non-linear cold dark matter. We find that the excess in power observed in the N-body neutrino power spectrum compared to linear response can be explained via a higher-order modification to the sound speed. Based on this, we have provided a simple model for the neutrino power spectrum that requires no additional integration beyond standard Boltzmann code outputs. Finally, we speculate that treating neutrinos as a dispersive fluid could allow for them to be simulated efficiently in both memory and processing time. \end{section} | 16 | 9 | 1609.09469 |
|
1609 | 1609.07474_arXiv.txt | The Moon migrated to $r_{\leftmoon}\simeq3.8\times10^{10}$ cm over a characteristic time $r/v=10^{10}$ Gyr by tidal interaction with the Earth's oceans at a present velocity of $v=3.8$ cm yr$^{-1}$. We derive scaling of global dissipation that covers the entire history over the past 4.52 Gyr. Off-resonance tidal interactions at relatively short tidal periods in the past reveal the need for scaling {with amplitude}. The global properties of the complex spatio-temporal dynamics and dissipation in broad spectrum ocean waves is modeled by damping $\epsilon = h F/(2Q_0)$, where $h$ is the tidal wave amplitude, $F$ is the tidal frequency, and $Q_0$ is the $Q$-factor at the present time. It satisfies $Q_0\simeq 14$ for consistency of migration time and age of the Moon consistent with observations for a near-resonance state today. It shows a startingly fast eviction of the Moon from an unstable near-synchronous orbit close to the Roche limit, probably in a protolunar disk. Rapid spin down of the Earth from an intial $\sim30\%$ of break-up by the Moon favored early formation of a clement global climate. Our theory suggests moons may be similarly advantageous to potentially habitable exoplanets. | Over the past eons, tidal interaction produced a major evolution of the Earth-Moon system, producing a relatively slowly rotating planet down from about one-third of centrifugal break-up at birth, very similar then to Jupiter's state today. For a detailed discussion on the break-up angular velocity $\Omega_b\simeq \sqrt{{GM}/{R^3}}$ of a planet of mass $M$ and radius $R$, see, e.g., \cite{dav99}. Tidal interactions are inherently dissipative, determined by a phase-lag between tidal deformation and position of the perturber, i.e., a misalignment of tidal bulge relative to the Earth-Moon direction. For the Earth, dissipation is primarily in the ocean tidal flows, more so than viscoelastic deformation of the Earth's mantle \citep[e.g.][]{mun68,mun71,lam77,dic94,ray94,ray96,efr09,efr12,efr15}, whose seismic frequencies are relative high compared to the tidal frequency \citep{lov11,mun60,dah74} with the exception of those driven by ocean waves \citep{web07}. However, a detailed quantitative account for the overall Moon's migration time due to various nonlinear dissipation channels \citep[][]{sto48,sto57,mun68,egb01}, remains to be identified. Tidal dissipation \citep[][]{mun71,web82,dic94} has various mechanisms in shallow water wave theory, some of which have recently been highlighted in detailed numerical simulations on ocean dynamics and dissipation covering relatively short initial and present epochs \citep[e.g][]{tou94,tou98,egb04,sta14}. Dissipation generally occurs when nonlinear steepening exceeds the mitigating effect of dispersion. Determining the net global result from detailed modeling of tidal dissipation is particularly challenging by the diversity of oceans and coastal regions, where most of the dissipation is expected to occur \citep{mil66,mun71}. Here, we focus on scaling in both amplitude and frequency of global tidal dissipation to account for the Earth-Moon history over the past 4.52 Gyr. {This approach aims at providing an effective description of an otherwise complex spatio-temperal distribution of dissipation in broadband ocean waves.} For a confrontation with data, the Moon's migration time is computed by numerical integration of angular momentum transfer backwards in time, to the instant of its formation from the Earth or a surrounding proto-lunar disk. This approach enables taking into account variations in tidal implitude over a few orders of magnitude, the effect of which seems not to have been computed before. The evolution by coupling to the Earth's spin \citep[e.g.][]{efr09} is conveniently described by the orbital angular momentum $m\sqrt{GM(1-e^2)a}$ with semi-major axis $a$, where $m=7.35\times 10^{22}$ g and $M=5.97\times 10^{27}$ g denote the mass of the Moon and, respectively, Earth. The orbital ellipticity $e$ is presently about $5.5\%$. A large dynamic range in tidal interaction strength arises from the tidal amplitude $h\propto r^{-3}$ by which the Moon's specific angular momentum {at radius $r\simeq a$} evolves according to \begin{eqnarray} \frac{dj}{dt} \propto \frac{h}{r^3} \sin(2\Delta \varphi), \label{EQN_j} \end{eqnarray} where $\Delta\varphi$ is the phase-lag of the Moon's orbit relative {to the tide raised on the Earth}, and $r^{-3}$ is the mutual interaction strength for a given $h$. As a result, $dj/dt\propto r^{-6}$ \citep[cf.][]{efr09}. To begin, we first recall some general conditions for dissipation in ocean tidal flows (\S2). In \S3, we formulate our scaling of dissipation in tidal amplitude and frequency. It serves to parameterize damping in our model based on (\ref{EQN_j}) and the pendulum equation (\S4). This model is explored numerically in \S5. In \S6, we summarize the results. | The migration time of the Moon is determined by the rate of tidal dissipation, predominantly in tidal waves. In the past, forcing was mostly off-resonance in the inertial range above the resonance frequency $\omega_0$ in (\ref{EQN_PE}). {Linear theory of dissipation described by the Stokes' limit $\epsilon=\epsilon_0$} falls short of explaining a migration time equal to the Moon's age. {Nonlinear dissipation mechanisms pointed to by (\ref{EQN_Ursell}) are important}, now and even more so during off-resonance in the past. Dissipation in bores is similar to that in shocks of compressible gas dynamics. The same is expected to feature spectral broadening and hardening by dispersion \citep[e.g.][]{sea85,bar04}. Damping representing total tidal dissipation hereby is expected to scale both with tidal amplitude and frequency proposed in (\ref{EQN_epsD}). Total dissipation in the ocean tides contains also a component of internal dissipation by flows over non-smooth surfaces \citep{egb00,egb01}. It scales effectively with the cube of horizontal tidal flows \citep{gem15}, akin to high Reynolds number flows past solid objects. According to the theory of shallow water wave equations, column height and height-averaged horizontal velocity satisfy the same wave equation in the linearized limit, the latter with amplitude \begin{eqnarray} u=\sqrt{\frac{g}{d}}h. \label{EQN_uh} \end{eqnarray} For harmonic perturbations, internal dissipation is equivalently described by damping proportional to $u$, and hence by $h$ as a consequence of (\ref{EQN_uh}). In Table 2, this is included in the scaling by tidal amplitude. In scaling damping of ocean waves by $hF$, the same would be included approximately to within 30\%. {In considering (\ref{EQN_r4e}) as an effective description of the dominant dissipation in semidiurnal tidal interactions, higher order tidal modes \citep{doo21} are neglected with generally complex dependence on ocean basin geometry \citep{mun66}. In a linearized approximation, one might contemplate including damping by higher harmonics. However, the latter are difficult to constrain observationally based on data of the Moon's migration velocity. Even if known, evolution of ocean basins on the geological time scale of $10^2$ Myr \citep[e.g.][]{wil66,wil75} makes estimation of time-averages over a recent epoch highly uncertain. For this reason and the expectation that dissipation in the semi-diurnal tides are dominant, these higher order perturbations fall outside the scope of the present approach.} Numerical integration of (\ref{EQN_r4e}) with (\ref{EQN_epsD}) reproduces a Moon migration time equal to its age for a present $Q$-factor (\ref{EQN_QER}) in excellent agreement with the observational constraint (\ref{EQN_QE}). By (\ref{EQN_epsD}), it introduces a dynamic $Q$-factor \begin{eqnarray} Q= \frac{1}{2\min\left\{1, \epsilon_0 hF \right\},} \end{eqnarray} that was substantially below $Q_0=1/(2\epsilon_0)$ in the past. Its trajectory features a startingly fast eviction of the Moon at or close to the initial, unstable synchronous orbit at $j_0$. A protolunar disk with the same composition as the Earth \citep[e.g.][]{bar14} might require even higher Earth spin rates prior, to facilitate its ejection by a giant impact \citep{can01,cuk12}. If so, the Moon is even more pertinent as a deposit of the Earth's initial spin angular momentum. {Some uncertainty in the formulation of our model arises from the fact that the ocean eigenfrequency $\omega_0$ (closest to the semidiurnal tidal frequency) is that of the Atlantic ocean. By continental drift, it evolves on aforementioned geological time scale that, in particular, might include intermittent closure \citep{wil66,wil75}. In our numerical formulation, the present state then represents a time-average of a recent epoch, small relative to the Moon's age of 4.5 Gyr, probably satisfying \begin{eqnarray} \left<\frac{\omega^\prime}{\omega_0}\right> > \frac{\omega^\prime_0}{\omega_0}, \label{EQN_ave} \end{eqnarray} where the right hand side refers to today's instantaneous value. Here, we assume that during closure, $\omega_0$ is determined by the remaining oceans, also since $Q$ of the Atlantic near closure is probably suppressed by inversion \citep[uplift of the ocean basin;][]{wil75}). In Fig. \ref{figKE}, (\ref{EQN_ave}) corresponds to moving along the ordinate to the right. By the relative flatness of the curves shown, our main conclusions hold, provided (\ref{EQN_ave}) remains in the window \begin{eqnarray} 0.95<\left< \frac{\omega^\prime}{\omega_0}\right> <1.35 \label{EQN_ave2} \end{eqnarray} for solutions to exist. Implications of larger perturbations fall outside the scope of the present formulation.} Once oceans form, eviction of the Moon was swift with steep spindown of the Earth from what was very similar to that of Jupiter today. In contrast to Jupiter's extreme weather patterns \citep{som88}, our dramatically reduced spin with small Coriolis forces today facilitates a clement global climate. It is tempting to consider application of (\ref{EQN_epsD}) further to exoplanet-moon systems \citep{kip09,for11,sch11,ea16}, whose selection is currently based on temperature, mass and size \citep{men06}. Relatively distant moons may be indirect evidence for effective tidal interaction with oceans, further favoring a potentially clement climate conducive to advanced life. {\bf Acknowledgments.} The author gratefully thanks the reviewer for constructive comments and pointing out several pertinent references. This report was supported in part by the National Research Foundation of Korea under grant No. 2015R1D1A1A01059793 and 2016R1A5A1013277. | 16 | 9 | 1609.07474 |
1609 | 1609.08553_arXiv.txt | We present {\it Chandra} ACIS-S and ATCA radio continuum observations of the strongly lensed dusty, star-forming galaxy SPT-S J034640-5204.9 (hereafter SPT0346-52) at $z$ = 5.656. This galaxy has also been observed with ALMA, {\it HST}, {\it Spitzer}, {\it Herschel}, APEX, and the VLT. Previous observations indicate that if the infrared (IR) emission is driven by star formation, then the inferred lensing-corrected star formation rate ($\sim$ 4500 $M_{\sun}$ yr$^{-1}$) and star formation rate surface density $\Sigma_{\rm SFR}$ ($\sim$ 2000 $M_{\sun} {\rm yr^{-1}} {\rm kpc^{-2}}$) are both exceptionally high. It remained unclear from the previous data, however, whether a central active galactic nucleus (AGN) contributes appreciably to the IR luminosity. The {\it Chandra} upper limit shows that SPT0346-52 is consistent with being star-formation dominated in the X-ray, and any AGN contribution to the IR emission is negligible. The ATCA radio continuum upper limits are also consistent with the FIR-to-radio correlation for star-forming galaxies with no indication of an additional AGN contribution. The observed prodigious intrinsic IR luminosity of (3.6 $\pm$ 0.3) $\times$ 10$^{13}$ $L_{\sun}$ originates almost solely from vigorous star formation activity. With an intrinsic source size of 0.61 $\pm$ 0.03 kpc, SPT0346-52 is confirmed to have one of the highest $\Sigma_{\rm SFR}$ of any known galaxy. This high $\Sigma_{\rm SFR}$, which approaches the Eddington limit for a radiation pressure supported starburst, may be explained by a combination of very high star formation efficiency and gas fraction. | \label{sec:intro} A population of gravitationally lensed dusty star-forming galaxies (DSFGs) has been discovered by the South Pole Telescope (SPT) survey \citep{Vieira2010} and facilitated our understanding of the stellar, gas, and dust content of the high-redshift Universe. One of the sources stands out as the most extraordinary discovered so far in the 2500 deg$^2$ survey: SPT-S J034640-5204.9 (hereafter SPT0346-52) at $z$ = 5.656, among the highest-redshift DSFGs known. It has been the focus of a multi-wavelength observational campaign with {\it HST}, {\it Spitzer}, {\it Herschel}, the Atacama Large Millimeter/submillimeter Array (ALMA), the Atacama Pathfinder EXperiment (APEX), and the Very Large Telescope (VLT). Our lens model obtained from ALMA 870 $\micron$ imaging shows that SPT0346-52 is magnified by the foreground lensing galaxy (at $z$ $\sim$ 1.1) a factor of 5.6 $\pm$ 0.1 with an intrinsic 870 $\micron$ flux of 19.6 $\pm$ 0.5 mJy, and has an intrinsic size of $R_{\rm eff}$ = 0.61 $\pm$ 0.03 kpc ($R_{\rm eff}$ being half-light radius; \citealt{Hezaveh2013, Spilker2016}). Multi-band spectral energy distribution (SED) fitting gives an intrinsic infrared (IR; 8-1000${\micron}$) luminosity $L_{\rm IR}$ of (3.6 $\pm$ 0.3) $\times$ 10$^{13}$ $L_{\sun}$ and a star formation rate (SFR) of 4500 $\pm$ 1000 $M_{\sun}$ yr$^{-1}$ \citep{Ma2015}. Given its size, SPT0346-52 turns out to have one of the highest IR luminosity surface density and SFR surface density $\Sigma_{\rm SFR}$ of any known galaxy (\citealt{Hezaveh2013,Spilker2015,Spilker2016}). The central question is whether this high luminosity surface density arises solely from intense star formation, or if there is an obscured active galactic nucleus (AGN). The dust temperature of SPT0346-52, 52.4 $\pm$ 2.2 K \citep{Gullberg2015}, is higher than that of typical DSFGs and reaches into the territory of AGN-dominated sources (Figure \ref{fig:LciiLfir_Td}). SPT0346-52 has an $L_{\rm [CII]}$/$L_{\rm FIR}$ ratio consistent with FIR-luminous quasars at $z$ $\sim$ 6 and also shows an $L_{\rm [CII]}$ deficit relative to $L_{\rm FIR}$ and $L_{\rm CO(1-0)}$, which has been observed in AGN-dominated sources (\citealt{Stacey2010,Sargsyan2014,Gullberg2015}). It shows strong H$_2$O emission lines \citep{Weiss2013} similar to that of the strongly lensed quasars H1413+1143 and APM 08279+5255 \citep{Bradford2011}. SPT0346-52 is also optically obscured and does not show any indications of type-1 or type-2 AGN in deep VLT optical spectroscopy \citep{Hezaveh2013}. DSFGs are in a unique phase of galaxy formation and evolution where the assembly of the stellar and supermassive black hole (SMBH) masses are believed to be closely coupled \citep{Alexander2012}. To test if \mbox{SPT0346-52} hosts an AGN and determine whether it is star formation-dominated or AGN-dominated, we resort to {\it Chandra}. Hard X-ray emission (rest-frame energies $>$ 2 keV) is the best indicator of AGN activity. These high-energy photons can penetrate through heavy obscuration, revealing the signature of the accreting black hole. A significant fraction of X-ray detected DSFGs have been found to be AGN-dominated in the X-ray, while some are powered by pure star formation (e.g., \citealt{Laird2010,Georgantopoulos2011,Johnson2013}). X-ray observations of the well-studied DSFG samples from the ALMA LABOCA E-CDF-S Submillimeter Survey (ALESS; \citealt{Wang2013}) reveal that 17\% of DSFGs appear to host an AGN. We here compare SPT0346-52 with these DSFGs and starbursts and quasars in the literature to understand the nature of the most extraordinary source found so far in the SPT survey. In addition to X-ray, radio also can be used to distinguish star-forming galaxies from AGN. Radio continuum emission from galaxies arises due to both thermal and non-thermal processes in massive star formation. These same massive stars also provide the primary sources of dust heating in the interstellar medium and the FIR emission is primarily due to the re-emitted starlight by dust. Star-forming galaxies that are not radio-loud AGN are observed to follow a tight FIR-to-radio correlation that holds over five orders of magnitude in galaxy luminosity (e.g., \citealt{Yun2001}). In contrast, radio-loud AGN will exhibit elevated radio emission above this relation (e.g., \citealt{Yun2001, Condon2002}). We utilize the Australia Telescope Compact Array (ATCA) to probe the radio continuum emission of SPT0346-52 to examine whether or not it is consistent with the FIR-to-radio correlation. In this paper, we present the results from X-ray observations with the {\it Chandra} Observatory and radio continuum observations with ATCA to constrain the AGN activity, in combination with our existing multi-wavelength data. The Galactic column density towards \mbox{SPT0346-52} is $N_H$ =1.8 $\times$ 10$^{20}$ cm$^{-2}$. We assume a $\Lambda$CDM cosmology with $H_0$ =69.3 km s$^{-1}$ Mpc$^{-1}$, $\Omega_{m}$ = 0.286, and $\Omega_{\Lambda}$ = 0.713 (WMAP9; \citealt{Hinshaw2013}). We adopt the definition of $L_{\rm IR}$ to be integrated over rest-frame 8-1000 $\micron$ and $L_{\rm FIR}$ integrated over rest-frame 42.5-122.5 $\micron$. We assume a \cite{Chabrier2003} initial mass function (IMF) throughout the paper. \begin{figure}[ht] \centering {\includegraphics[width=9cm, height=7.5cm]{LCII_Lfir_Td_forpaper_16April2016.eps}} \caption{$L_{\rm [CII]}/L_{\rm FIR}$ vs. $T_{\rm dust}$. The red circles are SPT DSFGs. The low-z and high-z samples are compiled by \cite{Gullberg2015} (see references therein). SPT0346-52 (red star) lies in the region surrounded by AGN-dominated galaxies. This diagnostic provided the motivation to search for X-ray signatures of an AGN in this extreme source.} \label{fig:LciiLfir_Td} \end{figure} | \end{table} \subsection{{\it Chandra} X-ray data} SPT0346-52 was observed with the Advanced CCD Imaging Spectrometer (ACIS; \citealt{Garmire2003}) on board {\it Chandra} on 2015 July 29. The source was placed at the aim point of the back-illuminated ACIS-S3 chip. The data were taken in Very Faint mode and were initially processed by the {\it Chandra} X-ray Center (CXC) using software version 10.4.1 and CalDB version 4.6.8. We reprocessed the data with the {\it Chandra} Interactive Analysis of Observations (CIAO; version 4.7) tool $chandra\_repro$. All the bad pixels were removed and the standard grade (0,2,3,4,6), status, and good-time filters were applied. The net exposure time for the observation is 49.52 ks. We performed energy filtering on events into three {\it Chandra} bands: the soft (SB; 0.5-2.0 keV), hard (HB; 2-8 keV), and full (FB; 0.5-8.0 keV) bands. The soft and hard bands probe rest-frame energies 3.3-13.3 keV and 13.3-53.2 keV for $z$ = 5.656, respectively. No detectable X-ray emission is expected from the foreground lens. Existing optical and radio data show no evidence for an AGN in the lens. Moreover, the foreground lens is an elliptical galaxy (based upon the light profile fitting by \citealt{Ma2015}), and thus should also have negligible X-ray emission from star formation (more than three orders of magnitude below the detection threshold). \begin{figure*} \centering {\includegraphics[width=6.08cm,height=6.08cm]{0346Thumbnail_Jan18.eps} \includegraphics[width=6.06cm,height=6.06cm]{0346chandra_15July2016.eps}} \caption{$Left:$ 20$\arcsec$ $\times$ 20$\arcsec$ cutouts of SPT0346-52 showing the HST/WFC3 (gray), Spitzer/IRAC (blue contours), and ALMA band 7 (red contours) data. $Right:$ {\it Chandra} 0.5-8 keV full-band data. The green circle shows the source extraction aperture enclosing the ALMA contours (red). The energies of the three photons are 1.284 (upper right), 2.585 (lower left), and 2.009 (lower right) keV.} \label{fig:data} \end{figure*} We matched the source to the position of ALMA (Figure \ref{fig:data}) and used a source extraction radius of 1.75$\arcsec$ enclosing all the lensed images, which is $\sim$1.3 times the 90\% encircled-energy aperture radius (at 0.3$\arcmin$ off-axis angle). The aperture was chosen such that it is large enough to enclose the ALMA contours without including too much background. The background counts were estimated by placing 78 circular apertures with the same size at random positions in the field. We only detected 3.19 net source counts in the full band (0.73 in the SB and 2.47 in the HB). Due to the low count level, we utilize the tool $aprates$ in the $srcflux$ script in CIAO to place a proper upper limit on the X-ray flux. This tool employs Bayesian statistics to compute the background-marginalized posterior probability distribution for source counts/flux. The posterior distribution can be used to determine flux value and confidence intervals or upper limits. The resultant FB flux is consistent with a non-detection with a 3 $\sigma$ upper limit of 6.0 $\times$ 10$^{-15}$ \flux. We derive the rest-frame 0.5-8 keV apparent luminosity $L_{0.5-8 {\rm keV}}$ (without absorption correction) using the following equation, \begin{equation} L_{0.5-8 {\rm keV}} = 4\pi d^2_{\rm L} f_{0.5-8 {\rm keV}} (1+z)^{\Gamma_{\rm eff} - 2} \label{eq:f2L} \end{equation} \noindent where $d_{\rm L}$ is the luminosity distance at $z$=5.656 and $\Gamma$ is the effective power-law photon index. In principle, the photon index can be derived from the hardness ratio, which is the ratio of the photon count rates in the HB and the SB. However we cannot derive a reliable hardness ratio based on the upper limits in both bands. Instead, $\Gamma_{\rm eff}$ is fixed to 1.4 following \cite{Xue2011} and \cite{Wang2013}. We then derive the rest-frame 0.5-8 keV absorption-corrected luminosity $L_{0.5-8 {\rm keV, unabs}}$ by replacing $\Gamma_{\rm eff}$ with intrinsic photon index $\Gamma_{\rm int}$ and $f_{0.5-8 {\rm keV}}$ with unabsorbed flux $f_{0.5-8 {\rm keV, unabs}}$ in Equation \ref{eq:f2L}. We assume $\Gamma_{\rm int}$ = 1.8, a typical value for AGNs. The unabsorbed flux is estimated using the tool $modelflux$ within the $srcflux$ script. We run simulations adopting Sherpa \citep{Freeman2001} models $xspowerlaw$$\times$$xszphabs$$\times$$xsphabs$ with fixed $\Gamma$ = 1.8 for the power-law model and hydrogen column density $N_{\rm H}$ for the (intrinsic and Galactic) absorption models. We scale the measured 3$\sigma$ upper limit on the flux by a fractional correction for absorption based on the typical intrinsic $N_{\rm H}$ (2.3 $\times$ 10$^{23}$ cm$^{-2}$) from the ALESS SMG sample \citep{Wang2013}. The absorption-corrected 3 $\sigma$ upper limit is $f_{0.5-8 {\rm keV, unabs}}$ $<$ 7.6 $\times$ 10$^{-15}$ \flux. Since SPT0346-52 is gravitationally lensed, we further correct the X-ray flux and luminosity for lensing magnification assuming there is no differential magnification between the FIR (i.e., ALMA) and the X-ray emission \citep{Hezaveh2012}. The magnification-corrected FB flux and luminosity upper limits are listed in \mbox{Table \ref{tab:chandra}}. \begin{table*} \centering \caption{{\it Chandra} X-ray properties of SPT0346-52} \begin{tabular}{lcccccccc} \hline \hline Source Name & Redshift & Exptime & Full-band & Background & \multicolumn{2}{c}{Full-band Flux (3$\sigma$) } & \multicolumn{2}{c}{Full-band Luminosity (3$\sigma$) } \\ & & (ks) & (count) & (count) & \multicolumn{2}{c}{ ($\times$ 10$^{-15}$ \flux) } & \multicolumn{2}{c}{($\times$ 10$^{44}$ \lum)} \\ & & & & &$f_{0.5-8 {\rm keV}}$ & $f_{0.5-8 {\rm keV, unabs}}$ & $L_{0.5-8 {\rm keV}}$ & $L_{0.5-8 {\rm keV, unabs}}$ \\ \hline SPT0346-52 & 5.656 & 49.52& 3.19 & 0.81 & $<$ 1.07 & $<$ 1.36 & $<$ 1.20 & $<$ 3.23 \\ \hline \end{tabular} \label{tab:chandra} \end{table*} \begin{figure*} \centering {\includegraphics[width=18.5cm, height=16.2cm]{lfir_lx_forpaper_5April2016.eps}} \caption{$L_{\rm FIR}$ vs. $L_{\rm X}$ for SMGs and literature galaxies. SPT0346-52 is the red arrow representing the 3 $\sigma$ upper limit on the absorption-corrected $L_{\rm X}$. We have corrected $L_{\rm FIR}$ and $L_{\rm X}$ for lensing magnification. We show the ALESS SMGs from \cite{Wang2013}. The X-ray detected SMGs are marked as starburst-classified SMGs and AGN-classified SMGs. The dividing line between star formation-dominated and AGN-dominated appears to be $L_{\rm X}$ = 0.004 $\times$ $L_{\rm FIR}$, which is the median ratio of SCUBA SMGs found by \cite{Alexander2005a}. The squares are literature galaxies that are classified as star formation-dominated (labeled with blue `S') or AGN-dominated (labeled with orange `A'). These data points are compiled by \cite{Alexander2005a} (see references therein) and we have re-fit the FIR photometry in a consistent manner as described in \cite{Gullberg2015}. We assume 30\% error on $L_{\rm X}$ for the literature galaxies. The dashed white line and the gray region represent the fiducial $L_{\rm FIR}$ vs. $L_{\rm X}$ relation and its standard deviation for the quasars in \cite{Elvis1994}. The $L_{\rm X}$-to-$L_{\rm FIR}$ ratio for pure starbursting systems (black dashed line) is about 2 orders of magnitude lower.} \label{fig:LFIR_Lx} \end{figure*} \subsection{ATCA radio data} SPT0346-52 was observed with ATCA for 3960s at 5.5 and 9.0 GHz and 4068s at 2.1 GHz on 2012 January 25 in the 6A array configuration using the CABB in the 1M-0.5k mode. The data was reduced in the same manner as in \cite{Aravena2013}. The resultant synthesized beam sizes are 7.7$\arcsec$ $\times$ 5.2$\arcsec$ at 2.1 GHz, 3.3$\arcsec$ $\times$ 2.2$\arcsec$ at 5.5 GHz and 2.1$\arcsec$ $\times$ 1.3$\arcsec$ at 9.0 GHz. The continuum was not detected in any band and we place 3 $\sigma$ upper limits (i.e., 3 $\times$ rms noise values calculated within a 1$\arcmin$ region around the source position) of 0.213 mJy at 2.1 GHz, 0.114 mJy at 5.5 GHz and 0.138 mJy at 9.0 GHz on the radio emission from the source. | 16 | 9 | 1609.08553 |
1609 | 1609.03477_arXiv.txt | {} {We have recently reported the discovery of five low redshift Lyman continuum (LyC) emitters (LCEs, hereafter) with absolute escape fractions \fesclyc\ ranging from 6 to 13\%, higher than previously found, and which more than doubles the number of low redshift LCEs. We use these observations to test theoretical predictions about a link between the characteristics of the Lyman-alpha (\lya) line from galaxies and the escape of ionising photons. } {We analyse the \lya\ spectra of eight LCEs of the local Universe observed with the Cosmic Origins Spectrograph onboard the Hubble Space Telescope (our five leakers and three galaxies from the litterature), and compare their strengths and shapes to the theoretical criteria and comparison samples of local galaxies: the Lyman Alpha Reference Survey, Lyman Break Analogs, Green Peas, and the high-redshift strong LyC leaker {\it Ion2}.} {Our LCEs are found to be strong \lya\ emitters, with high equivalent widths, EW(\lya)$> 70$ \AA, and large \lya\ escape fractions, \fesclya $>20\%$. The \lya\ profiles are all double-peaked with a small peak separation, in agreement with our theoretical expectations. They also have no underlying absorption at the \lya\ position. All these characteristics are very different from the \lya\ properties of typical star-forming galaxies of the local Universe. A subset of the comparison samples (2-3 Green Pea galaxies) share these extreme values, indicating that they could also be leaking. We also find a strong correlation between the star formation rate surface density and the escape fraction of ionising photons, indicating that the compactness of star-forming regions plays a role in shaping low column density paths in the interstellar medium of LCEs. } { The \lya\ properties of LCEs are peculiar: \lya\ can be used as a reliable tracer of LyC escape from galaxies, in complement to other indirect diagnostics proposed in the literature.} | \label{s_intro} Cosmic reionisation is a major event in the history of the Universe, which took place during the first billion years: after the dark ages, the first sources of light released ionising radiation in the intergalactic medium, reheating the gas, and changing its opacity. This \emph{cosmic feedback} of the first astrophysical objects is thought to have influenced the formation and evolution of galaxies in their first stages \citep{Ocvirk16} and determines their detectability \citep{Stark11, Schenker14, Dijkstra14, Dijkstra15}. The open question in actual studies of cosmic reionisation is to understand the relative contribution of the two types of sources likely involved: quasars and star-forming galaxies. Quasars inject a copious amount of ionising photons into the intergalactic medium (IGM), but were thought to be too rare at high redshift to have contributed significantly \citep{Fontanot12,Fontanot14}. A new population of faint Active Galactive Nuclei (AGNs) has been recently reported \citep{Giallongo15}, which could sustain reionisation, assuming that their escape fraction of ionising photons is high \citep[e.g. 100\%\ in ][]{Madau15}. The escape fraction of ionising photons from AGNs has to be investigated in more details, recent measurements report much lower values than assumed in the previous studies \citep{Micheva16}. Our present study focuses on the second type of sources of the cosmic reionisation: massive stars in galaxies. As for AGNs, the main uncertainty to quantify the role of star formation in the cosmic reionisation is the escape fraction of ionising photons from galaxies. Young and massive stars produce ionising radiation in situ, the difficult question is to observe, measure, and quantify how many of these photons, if any, escape the interstellar medium (ISM) of galaxies. So far, one clear detection has been reported at high redshift, with a high Lyman continuum absolute escape fraction \citep[\fesclyc $>50$\%][]{Vanzella16, deBarros16}, another ``clean'' Lyman Continuum Emitter (LCE) was recently found at $z=3.15$ by \cite{Shapley16}, and few low-redshift detections with low escape fractions ($<5\%$) have been found during the last decade \citep[][Puschnig et al., submitted]{Bergvall06,Leitet13,Borthakur14,Leitherer16}. Significant progress at low redshifts has recently been achieved with the COS spectrograph onboard {\sl HST} by \cite{Izotov16a, Izotov16b}, who found five LCEs at $z \sim 0.3$ with \fesclyc $\sim 6-13$\%. These sources are the basis of the present study. Given the extremely low success rate of observational searches for LCEs, several pre-selection methods for good LCE candidates have been proposed, involving the alteration of nebular emission line strengths \citep{Zackrisson13}, high \oiii/\oii\ ratios potentially tracing density-bounded \hii\ regions \citep{Jaskot13, Nakajima14}, or the non-saturation of the metallic low-ionisation absorption lines \citep{Heckman11, Alexandroff15, Vasei16} tracing a low covering fraction of the absorbing gas along the line of sight. We recently proposed a method based on the Lyman-alpha (\lya) spectral shape of LCEs: strong and narrow \lya\ profiles with a shift of the main peak smaller than $\sim150$ \kms, or double peaks closer than $\sim300$ \kms would indicate LyC leakage \citep[see ][ for more details]{Verhamme15}. Time is ripe to test these simple predictions. In the present work we assess the specificity of the \lya\ properties of LCEs by comparison to other samples of low-redshift star-forming galaxies whose \lya\ properties have been measured and analysed: the Lyman Alpha Reference Survey (LARS), Lyman-Break Analogs (LBAs), and ``Green Pea'' galaxies (GPs). The LARS sample consists of 14 star-forming galaxies at redshifts $z = 0.02 - 0.2$. They were selected from the cross-matched SDSS and {\sl GALEX} catalogues to sample the range of far-UV luminosities observed in $z\sim3$ Lyman Break Galaxies (LBGs). Only galaxies with active star formation were included, requiring EW(Ha)$ > 100$\,\AA. The aim of LARS is to study the mechanisms governing \lya\ escape from galaxies \citep[see][for an overview]{Ostlin14}. Most of LARS galaxies are dwarf irregulars. LARS14 is also a GP (see below). The LBAs were selected from the {\sl GALEX} catalogue for their far-UV luminosity, high surface brightness, and compactness \citep{Heckman05}. They resemble the high-redshift LBGs in physical size, stellar mass, star-formation rate, metallicity, dust extinction, and gas velocity dispersion. UV and optical morphologies of 30 LBAs were studied by \citet{Overzier09}. The UV shows massive star-forming clumps, while evidence of interaction is seen in the optical images. In 20\% of the sample, all of the UV light comes from a single, compact star-forming clump, usually characterized by strong outflows. Far-UV spectra of 22 LBAs were obtained with {\sl HST}/COS \citep{Heckman11, Alexandroff15}. Most of their \lya\ lines are observed in emission, with a variety of profiles including P-Cygni, as well as broad and narrow double-peaks. The escape of ionising photons from one LBA, J0921+4509, was reported in \citet{Borthakur14}, with an absolute escape fraction of $1\%$. The probability of LyC leakage from the other LBAs was discussed in \citet{Alexandroff15} by comparing the indirect diagnostics for LyC escape discussed above. Green Pea galaxies were noticed in the Galaxy Zoo SDSS images for their bright green colour and compactness. The colour is due to strong \oiiil\,\AA\ emission, reaching equivalent widths as high as 1000\,\AA. \citet{Cardamone09} identified a sample of $\sim250$ GPs in SDSS DR7, at $z=0.11-0.36$, while \citet{Izotov11} extended the number to $\sim800$ over a larger redshift range $z=0.02-0.63$ (the colour changes with redshift, but their properties remain similar). GPs share many properties with high-redshift LBGs and LAEs. Twelve archival {\sl HST}/COS spectra of GPs (outside our sample) are available \citep{Henry15,Yang16}, drawn from \citet{Cardamone09}. They are all strong \lya\ emitters, most of them with double-peaked line profiles. The escape of ionising photons from GPs was discussed on the basis of their \lya\ spectral shapes \citep{Verhamme15}. Some are best fitted by synthetic \lya\ spectra emergent from low column density geometries, indicating that they could be leaking \citep[][Orlitova et al. in prep.]{Yang16}. As reported in \citet{Izotov16a, Izotov16b}, we detected LyC emission from five compact, strongly star-forming galaxies of the local Universe, with high \oiii/\oii\ ratios ($>4$). These galaxies belong to the GP category. To our knowledge, this was the first attempt to test if high \oiii/\oii\ ratios are linked to LyC leakage. The major result of this study has been that five out of the five observed objects are leaking ionising radiation, and the finding of a correlation between \fesclyc\ and \oiii/\oii\ \cite[see Fig.\ 14 in][]{Izotov16b}. Therefore, a high \oiii/\oii\ ratio appears as a potential signature of LyC escape. Since nothing was known about \lya\ for these galaxies, it is interesting to study the \lya\ properties of these LCEs, and to use these sources to test our theoretical predictions regarding the relation between LyC leakage, \lya\ escape, and the detailed \lya\ line profiles \citep{Verhamme15}. This is the main objective of this paper. In Sect.~\ref{s_data} we describe the observational data used in our study. In Sect.~\ref{s_LyaLyC} we discuss the \lya\ properties of LCEs and comparison samples, showing that LCEs have a strong \lya\ emission. We then present and discuss the detailed \lya\ line profiles of these galaxies (Sect.~\ref{s_profiles}). The results are briefly put into perspective and discussed in Sect.~\ref{s_discuss}. Our main conclusions are summarised in Sect.~\ref{s_conclude}. | \label{s_conclude} We have analyzed the properties of the \lya\ line of five low-redshift Lyman continuum leaking galaxies observed with the COS spectrograph onboard {\sl HST} that have recently been reported by \cite{Izotov16a,Izotov16b}. The $z \sim 0.3$ sources were selected for compactness and for showing a high emission line ratio \oiii/\oii\, which has previously been suggested as a possible diagnostics for Lyman continuum escape \citep{Jaskot13,Nakajima14}. For comparison we have also included the other confirmed low-$z$ Lyman continuum emitters (LCEs), and {\sl HST}/COS observations from samples of Green Pea (GPs) galaxies and Lyman break analogs (LBAs), which have been suggested to be LCEs (see Sect.\ \ref{s_data}). We have presented the behaviour of the \lya\ equivalent width, EW(\lya), the \lya\ escape fraction (\fesclya), and several measures of the \lya\ line profile, and discussed how these quantities depend on the Lyman continuum escape fraction, \fesclyc. We found that strong LCEs, defined by an escape fraction \fesclyc$>5$ \%, are characterised by strong \lya\ emission, as measured by EW(\lya) and \fesclya. Furthermore, their \lya\ line profiles show a ``double-peak" structure with two emission peaks blue- and redward of the central wavelength with a small separation between the peaks ($\vsep \sim 300-400$ \kms), confirming qualitatively the predictions of \cite{Verhamme15} from radiation transfer models. More precisely, our main results can be summarised as follows: \begin{itemize} \item All strong LCEs show EW(\lya)$>70$ \AA\ and \fesclya $>0.2$, suggesting that these quantities could be used as efficient selection criteria to find potential sources (or analogs) of cosmic reionisation (cf.\ Fig.~\ref{LyaStrength}). \item The restframe EW(\lya) increases with the Lyman continuum escape fraction and reaches EW(\lya)$\sim 70-130$ \AA\ (Fig.~\ref{LyaStrength_vs_LyC} left). \item Empirically the \lya\ escape fraction correlates with the Lyman continuum escape fraction (Fig.~\ref{LyaStrength_vs_LyC} right). On average \fesclya/\fesclyc$\sim 2-3$, indicating that \lya\ photons escape more efficiently than the ionising photons, as expected due to radiation transfer effects (Section \ref{s_LyaLyC}). \item The separation between the blue and red peaks decreases with increasing \fesclyc, as expected for decreasing \hi\ column densities \citep{Verhamme15}. Whereas the red peak of most LCEs is $V_{\mathrm{peak}}^{\mathrm{red}} \sim 80-150$ \kms, the blue peak can be shifted by larger amounts ($V_{\mathrm{peak}}^{\mathrm{blue}} \sim$ -300 to -150 \kms). The shift of the blue peak is found to correlate with \fesclyc (Fig.~\ref{peaks_vs_fescLyC}). \item No clear correlation is found for LCEs between the ratio of the blue and red equivalent widths and \fesclyc. \item Comparison of the \lya\ line properties of GPs and LBAs with those of the LCEs shows that a subset of these galaxies could indeed be emitters in the Lyman continuum (see Fig.~\ref{peaks_vs_EWLya}). \item We find a correlation between \fesclyc\ and the SFR surface density (cf.\ Fig.~\ref{sigma_vs_fescLyC}), indicating that the compactness of star-forming regions could play a significant role in shaping low density channels through the ISM of LCEs. \end{itemize} Although larger samples of LCEs are needed to establish more robust correlations between Lyman continuum escape and other properties, the observations reported here should provide useful guidance to find and select more efficiently the sources of cosmic reionisation or their analogs. | 16 | 9 | 1609.03477 |
1609 | 1609.08848_arXiv.txt | \clustername\ is a galaxy cluster at $z=\zcluster$, strongly lensing a quasar at $z=\zQSO$ into six widely separated images. In recent \hst\ imaging of the field, we identify additional multiply lensed galaxies, and confirm the sixth quasar image that was identified by Dahle et al. (2013). We used the Gemini North telescope to measure a spectroscopic redshift of $z=\zarcB$ of one of the secondary lensed galaxies. These data are used to refine the lens model of \clustername, compute the time delay and magnifications of the lensed quasar images, and reconstruct the source image of the quasar host and a second lensed galaxy at $z=$\zarcA. This second galaxy also appears in absorption in our Gemini spectra of the lensed quasar, at a projected distance of \qsoAdist\ kpc. Our model is in agreement with the recent time delay measurements of Dahle et al. (2015), who found \tab=\timeABobs\ days and \tac=\timeACobs\ days. We use the observed time delays to further constrain the model, and find that the model-predicted time delays of the three faint images of the quasar are \tad=\timeADpred\ days, \tae=\timeAEpred\ days, and \taf=\timeAFpred\ days. We have initiated a follow-up campaign to measure these time delays with Gemini North. Finally, we present initial results from an X-ray monitoring program with \Swift, indicating the presence of hard X-ray emission from the lensed quasar, as well as extended X-ray emission from the cluster itself, which is consistent with the lensing mass measurement and the cluster velocity dispersion. | The rare chance alignment of a quasar behind a strong-lensing cluster provides unique opportunities for studies of different astrophysical objects. Through careful lens modeling, these systems can probe the mass distribution of the foreground lens; the high magnification enhances our ability to study the background quasar, and galaxies between us and the quasar can be seen in absorption along multiple lines of sight in the light of the background quasar. Lensing configurations that involve a quasar lensed by a single massive galaxy are more common; however, the lensing magnification of a single galaxy is typically significantly lower than in the galaxy cluster case. Unique to the cluster-lensed quasar configurations, the multiple images of the lensed quasar have large separations ($14\farcs6 - 22\farcs5$; Inada et al. 2003, 2006; Dahle et al. 2013) and high magnifications; the lensed active nucleus is point-like, providing accurate positional constraints, and is variable –- enabling measurements of the time delay between images of the same source. The high tangential magnification stretches the host galaxy of the quasar into a giant arc, thus resolving it from the light of the active nucleus, which usually dominates in a high-redshift quasar. To date, only three cases of high-redshift quasars strongly-lensed by a galaxy cluster are published: SDSS J1004+4112 (Inada et al. 2003), SDSS J1029+2623 (Inada et al. 2006), and \clustername\ (Dahle et al. 2013). \clustername\ was discovered as part of the Sloan Giant Arcs Survey (SGAS; Gladders et al. in prep, Bayliss et al. 2011a,b; Hennawi et al. 2008; Sharon et al. 2014). SGAS is a systematic survey of highly magnified lensed galaxies, also refered to as ``giant arcs,'' in the imaging data of the Sloan Digital Sky Survey (SDSS, York et al. 2000). The lensing identification process starts with optical selection of galaxy clusters from the SDSS photometry catalogs, using the cluster red sequence algorithm of Gladders \& Yee (2000). Sections of the imaging data around each cluster were then retrieved and processed to generate color images, with scaling parameters selected to optimize the visibility of possible lensing features. The images were visually inspected and ranked for lensing evidence by several observers in a process that enables a calculation of the selection statistics (the process will be described in full in Gladders et al., in preparation). All candidates were followed up for confirmation, and the survey purity and completeness were quantified. Bayliss et al. (2011a,b) give the results of the initial spectroscopic followup campaign, and measure the redshift distribution of the lensed galaxies. \clustername\ was detected in the SGAS search in SDSS Data Release 8 (Aihara et al. 2011) owing to a prominent giant arc that appears $8\farcs5$ south of the brightest cluster galaxy. A further investigation of the field revealed the multiply-imaged lensed quasar. The field was followed up by Dahle et al. (2013) using the Mosaic Camera (MOSCA) and the Andalucia Faint Object Spectrograph and Camera (ALFOSC) at the 2.56 m Nordic Optical Telescope (NOT). We have recently obtained \hst\ imaging data of this target (Figure~\ref{fig.hst}; Section~\ref{s.data}). As can be seen in Figure~\ref{fig.hst}, a background quasar is lensed by \clustername, forming six images around the core of a galaxy cluster at $z=$\zcluster. Three bright images appear north of the cluster core (labeled A, B, C; our labeling scheme follows Dahle et al. 2015), and three faint images (D, E, F) can be seen near the central cluster galaxies (G2, G3, G1, respectively). The cluster also lenses other background galaxies, the most prominent of which is seen as a blue arc south of the cluster core (labeled A1 in Figure~\ref{fig.hst}). Dahle et al. (2013) reported on the discovery of \clustername, confirmed the lensing interpretation, presented spectroscopic identification of the lensed quasar, spectroscopic confirmation of the six lensed images of the quasar, and measured its redshift to be $z=2.82$. In addition, we measured the spectroscopic redshifts of several cluster member galaxies, and of the lensed galaxy A1 at $z=$\zarcA. Stark et al. (2013) also measure the spectra of the quasar, $z=2.807$ and of galaxy A1. Interestingly, the spectrum of the quasar shows strong \Lya\ absorption at the redshift of the foreground lensed galaxy, as well as Si~II $\lambda$1526 and CIV~$\lambda$1549 (Stark et al. 2013), indicating the presence of neutral hydrogen and metals associated with gas surrounding the galaxy. Stark et al. (2013) estimated that the projected distance between the quasar image A and the interloper galaxy A1 is $\sim 50$ kpc. We refine this estimate in Section~\ref{s.absorber}. Following the discovery of \clustername, we have initiated an imaging monitoring program with the NOT to measure the time delays between the images of the quasar. The results from the first three years of ongoing photometric monitoring with the NOT and the first season of Gemini monitoring are presented in Dahle et al. (2015). The light curves of the brighter three images of \clustername\ are measured from an analysis of 42 distinct epochs, resulting in time delays of \tab=\timeABobs\ days, and \tac=\timeACobs\ days. A robust measurement of the time delays of images D, E, and F requires deeper observations; a monitoring campaign with Gemini was initiated in 2015 (GN-2016A-Q-28; PI: Gladders) for this purpose. This paper is structured as follows. In Section~\ref{s.data} we describe the \hst\ imaging data of \clustername, Gemini spectroscopy, and \Swift\ X-ray observations. We present a new strong lensing analysis based on the new data in Section~\ref{s.lensing}. In Section~\ref{s.results}, we present and discuss the predicted time delays, cluster mass, lensing magnification, source reconstruction, and absorbing systems. We conclude with future work in Section~\ref{s.conclusions}. Throughout this paper, we assume a flat cosmology with $\Omega_{\Lambda} = 0.7$, $\Omega_{m} =0.3$, and $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$. In this cosmology, $1\arcsec$ corresponds to 6.0384 kpc at the cluster redshift, $z=$\zcluster. Magnitudes are reported in the AB system. \begin{figure*} \centering \includegraphics[scale=0.5]{fig1.pdf} \caption{Color composite image of \clustername\ from our \hst\ program GO-13337 (PI: Sharon) in WFC3/IR F160W+F110W (red), ACS F814W (green), ACS F606W+F435W (blue). These data confirm the sixth quasar image (F) that was identified by \dahleA\ as tentative. The six images of the lensed quasar at $z=\zQSO$ and the previously identified giant arc A1 at $z=\zarcA$ are labeled in cyan; newly discovered secure multiply-imaged galaxies are labeled in white (B1, B2, B3 at $z=\zarcB$, see \S~\ref{s.spec}), yellow (D1, D2, D3) and magenta (C1, C2, C3). Other possible arc candidates are not labeled. Note that the point source that is seen embedded in the A1 arc is a foreground white dwarf star (Dahle et al. 2013). } \label{fig.hst} \end{figure*} | \label{s.results} \subsection{Time Delays}\label{s.posterior} Observational measurement of the time delays between the images of the quasar can provide valuable constraints on the lens model, as the arrival time is sensitive to the lensing potential. \clustername\ gives us a unique opportunity to obtain observational constraints on the time delay from six images of the same background quasar, three of which appear close to the cluster core, in close proximity to cluster galaxies. In Dahle et al. (2015) we report on the measurement of time delays between the three bright images of the quasar, \tab=\timeABobs\ days, and \tac=\timeACobs\ days (all the time delays are measured as excess arrival time relative to image A of the quasar). Our basic lensing analysis does not use the time delays as constraints, and is done strictly without any a priori knowledge of the time delays. We now confront the lens model with the time delay observations. The basic lensing analysis predicts that the arrival time is shortest for image C of the quasar, followed by images A,B,F,D,E. Quantitatively, we find \tab=\timeAB\ days, and \tac=\timeAC\ days, in good agreement with the observed measurements of \dahleB. Next, we use the observed time constraints and their 95\% confidence limits to further constrain the parameter space. We select the sets of parameters from the MCMC sampling that result in lens models with $\chi^2$ in the range [$\chi^2$,$\chi^2+4.5$]. We consider these models as producing reasonable scatter in the predicted vs. observed positions of images of the lensed galaxies, and their parameters are drawn from a range larger than the 1$\sigma$ confidence interval of the parameter space of well-constrained parameters, as sampled by the MCMC process. Models with larger $\chi^2$ were rejected. We then compute the Fermat potential for each one of these sets of parameters (Equation~\ref{eq.dt}), assuming that the quasar source position $\vec\beta$ is at the mean of the predicted source positions of the six quasar images. We compute the excess arrival time relative to image A of the quasar, i.e., the predicted time delay for each of the images. We identify the models that predict time delays \tab\ and \tac\ within the 95\% confidence limit of the observed values of \dahleB. In Figure~\ref{fig.parameters}, we plot the positional $\chi^2$ against each of the parameters of the main cluster halo, and color-code the models that predict either \tab\ in the range \timeABobs\ days, \tac\ in the range \timeACobs\ days, or both. As can be seen in Figure~\ref{fig.parameters}, while models that predict the observed \tab\ span the entire parameter space, the measured time delay between image C and A, \tac, has good constraining power over some of the parameters, mainly the overall mass of the main cluster halo ($\sigma_0$), and its ellipticity ($e$). \begin{figure*} \centering \includegraphics[scale=0.4]{fig8.pdf} \caption{ The goodness of fit plotted against the parameter value, for sets of parameters from the MCMC analysis. The goodness of fit is estimated via the $\chi^2$ value, computed by \lenstool\ as the scatter between observed and predicted image positions. Models that predict \tab\ in the range \timeABobs\ days are plotted in red squares; models that predict \tac\ in the range \timeACobs\ days are plotted in green circles; models that satisfy both criteria are circled in black. All other models are plotted in gray circles. The observed time delay \tac\ has constraining power over the parameters of the main cluster halo, mainly the normalization $\sigma_0$, which is correlated with the overall mass of the cluster, and the ellipticity $e$. } \label{fig.parameters} \end{figure*} \begin{figure*} \centering \includegraphics[scale=0.4]{fig9.pdf} \caption{Correlations between the predicted relative time delays of the six quasars. The top and bottom rows show \tab\ and \tad\, respectively, plotted against the other relative time delays. Colors and symbols are the same as in Figure~\ref{fig.parameters}. We see correlation between all the relative time delays except for \tab. Thus the observational measurements of \tab\ and \tac\ (Dahle et al. 2015) have strong constraining power over the model. Furthermore, observational measurement of either of the A-D,E,F time delays will narrow the uncertainty on the predicted time delays of the other images.} \label{fig.dtdt} \end{figure*} We find strong correlation between the predicted time delays $\tau_{\rm AD}$, $\tau_{\rm AE}$, and $\tau_{\rm AF}$, as can be seen in Figure~\ref{fig.dtdt}. Interestingly, $\tau_{\rm AD}$, $\tau_{\rm AE}$, and $\tau_{\rm AF}$ do not correlate with \tab, but they have strong correlation with \tac. This correlation places a tight constraint on the predicted time delays of the three central images. Moreover, the arrival times of images D, E, and F are strongly correlated, which means that a time delay measurement of one of them will provide an additional strict constraint on the time delays of the other images. The correlation of the time delays of the central images is not surprising. The arrival time lag of the central images is dominated by gravitational time delay as the light travels close to the center of mass, due to the deep potential well of the cluster; light will take longer to travel on this path, although this path is geometrically shorter (with smaller impact parameter and smaller deflection). Thus \tac\ is linked to \tad, \tae, \taf\ through its correlation with the overall normalization of the cluster halo, i.e., the effective velocity dispersion, $\sigma_0$. Applying the time delay observational cut on the parameter space, we are able to narrow down the uncertainties on the predicted time delays of the central images. Interestingly, we find that these time delays are short enough to be measured within the next few years: $\tau_{\rm AD}=$\timeADpred, $\tau_{\rm AE}=$\timeAEpred, and $\tau_{\rm AF}=$\timeAFpred\ days; Moreover, the arrival time of E and F relative to D is short -- of order 3-5 months: $\tau_{DE}=$\timeDEpred, $\tau_{DF}=$\timeDFpred\ days, thus measuring $\tau_{\rm DE}$ and $\tau_{\rm DF}$ can be achieved within a year or two of cadenced imaging with a large telescope (Section~\ref{s.conclusions}). In the following sections, the results of the lensing analysis take into account the constraints from the observed time delays, as measured by Dahle et al. (2015), and their 95\% confidence interval as described above. \subsection{Cluster Mass}\label{s.mass} We report the lensing-inferred total projected mass density of the lens (cylindrical mass) within projected radii of 100, 200, and 500 pc: \monehundred, \mtwohundred, and \mfivehundred, $\pm10\%$ systematic uncertainty. The statistical uncertainties are derived from the MCMC sampling of the parameter space, combined with the Dahle et al. (2015) 95\% confidence interval of the time delay measurements (see Section~\ref{s.posterior}). An additional 10\% systematic uncertainty should be applied, given the relatively small number of constraints and spectroscopic redshifts, that limit the accuracy of the lens model. Johnson \& Sharon (2016) found that while the enclosed mass is well constrained at the radius of the lensing evidence, its systematic uncertainty decreases with increasing number of lensing constraints and spectroscopic redshifts. The analysis in Johnson \& Sharon (2016) is tuned to the typical number of constraints in high cross-section lensing clusters such as the Frontier Fields (Lotz et al. 2016), and therefore they do not sample the affect on systematics in a case like \clustername, a much lower-mass cluster with four multiply-imaged lensed sources and three spectroscopic redshifts. We therefore conservatively adopt a $10\%$ systematic uncertainty on the enclosed mass, which is the typical uncertainty for a case of five sources and no spectroscopic redshifts. Interestingly, the observational measurement of the \tac\ time delay places a tight constraint on the total enclosed mass and is what drives the relatively small statistical uncertainty. Figure~\ref{fig.xray} shows the contours of the projected mass density distribution from the strong lens model, and the X-ray contours from \Swift\ observations (Section~\ref{s.swift}). We find that the X-ray emitting gas and the dark matter distribution are generally aligned, with no significant offset between their centroids. A more robust measurement of the X-ray distribution will be enabled with the superior resolution of Chandra observations. \subsection{Magnification}\label{s.time} The magnification map for a source at the quasar redshift, $z=\zQSO$, and the magnifications measured at the position of each image of the quasar, are shown in Figure~\ref{fig.magnific}b and Table~\ref{tab.magnific}, respectively. The uncertainties are estimated by computing magnification maps for a series of lens models sampled from steps the MCMC that correspond to $1\sigma$ in the parameter space, and the 95\% confidence interval of the time delay measuremens of Dahle et al. (2015). Since quasars are variable sources and are not standard candles, we cannot compare the absolute predicted lensing magnification with an observational measurement. Nevertheless, we can compare the predictions to the {\it relative} magnifications between images A, B, and C of the quasar, for which time delays have been measured. Dahle et al. (2015) find that the light curves of images A, B, and C, can be matched with time delays of \tab=\timeABobs\ and \tac=\timeACobs, and magnitude shift of $\Delta m_{\rm AB}=0.340\pm0.007$ mag and $\Delta m_{\rm AC}=0.483\pm0.012$ mag. We find that the model is in agreement with the observed relative magnification of image A and B. The model-predicted magnification of C is $\sim30\%$ too high to agree with the observed magnification ratio between A and C, indicating that the systematic uncertainties may be underestimated. We note that substructure in the cluster, as well as structure along the line of sight, may contribute to discrepancy between the measured and model-derived relative magnifications. Compared to the initial model in Dahle et al. (2013), which was based on ground-based observations, we find that the magnifications of A, B, and C are $\sim2.5\times$ higher than those derived in Dahle et al. (2013), but well within the large statistical uncertainties reported there. Moreover, systematic uncertainties, which are not taken into consideration, are large for lens models that are based on few lensed sources and few spectroscopic redshifts (Johnson \& Sharon 2016). We also note that the new constraints from the \hst\ data required a more massive component at the south of the cluster (G4) to explain the lensing evidence that was not identified from the ground. Compared to the magnifications in other wide-separation lensed quasars, we find that the magnifications of A, B, and C in \clustername\ are similar to the best-fit model-predicted magnifications of the three brightest images in SDSSJ1029, from Oguri et al. (2013), while in Oguri et al. (2010), the lens model of SDSSJ1004 predicts magnifications a factor $\sim2\times$ higher. \begin{deluxetable}{lclll} \tablecolumns{5} \tablecaption{Model-predicted magnifications and time delays \label{tab.magnific}} \tablehead{ \colhead{Image } & \colhead{F435W} & \colhead{Magnification} & \colhead{} & \colhead{Time delay} \\ \colhead{ } & \colhead{magnitude} & \colhead{$\mu$} & \colhead{} & \colhead{[days]} } \startdata A &21.861& \magA & \nodata & \nodata \\ B &22.261& \magB & $\tau_{\rm AB}$ & [~\timeABobs] \\ C &22.227& \magC & $\tau_{\rm AC}$ &[\timeACobs] \\ D &23.827& \magD & $\tau_{\rm AD}$ & \timeADpred \\ E &24.070& \magE & $\tau_{\rm AE}$ &\timeAEpred \\ F &24.909& \magF & $\tau_{\rm AF}$ & \timeAFpred \\ \enddata \tablecomments{Magnitudes in the ACS/F435W filter are measured within an aperture of radius $0\farcs56$ in an observation starting on JD 2456941.06751. Time delay is given in days, relative to image A. \tab\ and \tac\ are observational constraints from Dahle et al. (2015). The uncertainties represent the 95\% confidence level from the combined MCMC analysis and the observational time delay constraints.} \end{deluxetable} \begin{figure*} \centering \includegraphics[scale=0.4]{fig10.pdf} \caption{ Source plane reconstruction of galaxy A1 at $z=$\zarcA\ (left) and of the quasar host galaxy at $z=$\zsource\ (right). The quasar host is generated from image A of the quasar. The foreground white dwarf and the point-source emission from the quasar are masked to reveal the underlying information (see text). A horizontal bar indicates the scale in arcseconds and kpc at each source redshift. The yellow lines in the left panel are the locations of the source plane caustics, which map to the critical curves in the image plane. } \label{fig.arcs_src} \end{figure*} \subsection{Source Plane Reconstruction}\label{s.source} We reconstruct the source image of the lensed galaxy A1 at $z=$\zarcA, and the host galaxy of the quasar at $z=$\zsource, by ray-tracing the image-plane pixels through the lens equation, $\vec\beta = \vec\theta -\vec\alpha(\vec\theta)$, where $\vec\beta$ is the source position of each pixel, $\vec\theta$ is its observed position, and $\vec\alpha(\vec\theta )$ is the deflection matrix scaled by $d_{LS}/d_{S}$, the ratio between the distance from the lens to the source and from the observer to the source. The high lensing magnification resolves small substructure in these galaxies, which would otherwise be too small for \hst\ resolution. Galaxy A1 is highly distorted by the lensing potential due to its close proximity to the caustic. It is likely that a small region of this galaxy is multiply imaged within the giant arc. Prior to ray tracing the images, we subtract the light of the point source quasar light and the foreground white dwarf to reveal the underlying information. In each band we select a star in the field of view with similar brightness. We generate a second image by shifting the data so that the star is at the exact pixel position of the point source we wish to mask. We then scale the shifted image and subtract it from our data. Figure~\ref{fig.arcs_src} shows the reconstructed source-plane image of galaxy A, and of the quasar host galaxy. From the reconstructed source image, A1 measures $\sim13$~kpc in diameter and the quasar host is measured to be $\sim3$~kpc in diameter. A thorough investigation of the physical properties of these galaxies is left for future work. \subsection{Absorbing Systems}\label{s.absorber} Stark et al. (2013) find strong evidence for an absorption system at $z=$\zarcA\ in the spectrum of image A of the quasar, indicating that the extended gas halo of galaxy A1 has neutral hydrogen and metals, from absorption lines of \Lya, Si~II~$\lambda1526$ and CIV~$\lambda1549$. Stark et al. (2013) estimate the projected distance between A1 and image A of the quasar at $\sim50$ kpc. A proper estimate of the impact parameter takes into account the path of the light from the quasar source plane to each of its images, and where these paths traverse the source plane of A1, at $z=$\zarcA. In the left panel of Figure~\ref{fig.interloper} we show a reconstruction of the source plane at the redshift of galaxy A1. By ray-tracing the quasar images to the same redshift of A1, we find that the quasar light passes \qsoAdist\ kpc north of the center of A1. At this redshift, the quasar paths are separated by as much as 5 kpc. We are therefore presented with a unique opportunity to sample the uniformity of the gas halo on scales of a few kpc, with at least three bright lines of sight. Our GMOS multi-object spectroscopy masks targeted images A, B, C, and D of the quasar. Slits were placed on these sources on both nod and shuffle positions, and on both masks, resulting in a total of 2400 s on target for A, B, D, and 3600 s on target for C. The wavelength coverage allows the detection of FeII~2586,2600 and MgII~2796,2803 at the redshift of A1. The intervening absorption system is detected in the spectra of all three bright images of the quasar (A, B, C), at a redshift of $z = 2.2988 \pm 0.0002$. In the right panel of Figure~\ref{fig.interloper}, we plot the two strongest features of this system, the Mg~II~2796,2803~\AA\ doublet and the Fe~II~2600~\AA\ line. The spectra were continuum normalized, with the continuum calculated by smoothing the spectra with a 40~\AA\ boxcar. In Table~\ref{tab.obslog} we tabulate the equivalent width and redshift measurements for this absorption system in each quasar spectrum. Since the blue wing of the Mg~II~2976~\AA\ feature is affected by the [Ne IV] and Fe~III emission complex at 2423~\AA\ rest-frame \citep{VandenBerk:2001cd}, we do not try to fit this line, but instead consider the weaker transition Mg~II~2803. The intervening system is clearly detected in Mg~II in all three images of the quasar; the weaker Fe~II~2600 is detected in quasar images A and B. The equivalent widths are comparable given the uncertainties listed in Table~\ref{tab.obslog}. While Figure~\ref{fig.interloper} shows some variation in the absorption profiles from quasar image to image, particularly in the amount of redshifted absorption, these variations may not be significant given the signal-to-noise ratio of the data. Deeper spectra are required to measured differences in the absorption along these three lines of sight. We also detect FeII and MgII absorption from a second absorber at z=1.202 in the three spectra. The corresponding object is not currently identified in the imaging or spectroscopic data. The largest separation between the quasar lines of sight at this redshift is $\sim$40 kpc. Co-adding the spectra from the forthcoming spectroscopic followup campaign (Section~\ref{s.future}) will result in a deep spectrum of each of the quasar images, and high enough signal to noise to determine some of the physical properties of the gas halo in the absorbing systems. \begin{figure*} \centering \includegraphics[scale=0.5]{fig11a.pdf} \includegraphics[scale=0.53]{fig11b.pdf} \caption{{\it Left:} The $z=2.3$ plane, reconstructed from our lens model. At this redshift plane, the light from the $z=$\zsource\ quasar passes $\sim34$ kpc from galaxy A; the separation between the light rays of the quasar is a few kpc. The extended gas halo around galaxy A is seen in absorption in the spectra of the quasar, which permits a study of the spatial distribution of the physical properties of the gas. The grid is given in kpc, centered on the brightness peak of galaxy A. {\it Right: } Continuum-normalized spectra of the intervening absorption system at $z = 2.2988 \pm 0.0002$, for quasar spectra A (plotted in black), B (blue), and C (green). The Mg~II~2796, 2803 doublet and the Fe~II~2600 transition are shown at the top and bottom panels, respectively. The intervening absorber is clearly detected in Mg~II in all three quasar spectra, and in Fe~II in quasar images A and B. } \label{fig.interloper} \end{figure*} | 16 | 9 | 1609.08848 |
1609 | 1609.04061_arXiv.txt | Compositional convection is thought to be an important energy source for magnetic field generation within planetary interiors. The Prandtl number, $Pr$, characterizing compositional convection is significantly larger than unity, suggesting that the inertial force may not be important on the small scales of convection as long as the buoyancy force is not too strong. We develop asymptotic dynamo models for the case of small Rossby number and large Prandtl number in which inertia is absent on the convective scale. The relevant diffusivity parameter for this limit is the compositional Roberts number, $q = D/\eta$, which is the ratio of compositional and magnetic diffusivities. Dynamo models are developed for both order one $q$ and the more geophysically relevant low $q$ limit. For both cases the ratio of magnetic to kinetic energy densities, $M$, is asymptotically large and reflects the fact that Alfv\'en waves have been filtered from the dynamics. Along with previous investigations of asymptotic dynamo models for $Pr=O(1)$, our results show that the ratio $M$ is not a useful indicator of dominant force balances in the momentum equation since many different asymptotic limits of $M$ can be obtained without changing the leading order geostrophic balance. Furthermore, the present models show that inertia is not a requirement for driving low $q$, large-scale dynamos. | Planetary dynamos are thought to be generated by buoyancy-driven turbulence. Both thermal and compositional heterogeneities are possible buoyancy sources. The physical characteristics of thermal and compositional scalars are distinguished by their respective diffusion coefficients. Studies suggest that the thermal diffusivity in the Earth's outer core is larger than the compositional diffusivity by a factor of a thousand \citep{mP13}; i.e.~in non-dimensional terms, the Lewis number, $Le = \kappa/D = O(10^3)$, where $\kappa$ and $D$ are the thermal and compositional diffusivities, respectively. Denoting the fluid kinematic viscosity as $\nu$, the Prandtl numbers for the two scalars are therefore related via $Pr_C = Le Pr_T$, where $Pr_C = \nu/D = O(100)$ and $Pr_T = \nu / \kappa = O(10^{-1})$. Compositional convection has long been thought to play an important role in powering the geodynamo \citep{sB63,dS83,bB96b}. Recent studies have found that the thermal conductivity of the Earth's core may be larger than previously thought \citep{mP12,dK12}, suggesting that thermal convection may not play an important role in powering the geodynamo due to an associated increase in the conductive, adiabatic heat flux \citep{cD15}. Compositional convection has been suggested as the alternative power source for sustaining the geodynamo over the course of the Earth's evolution; this has stimulated interest in understanding the origin of such compositional heterogeneities \citep[e.g.][]{jO16}. There is clearly a need to better understand compositional convection and its influence on magnetic field generation. Provided the buoyancy forcing is not too strong, it is well known that the influence of inertia may be weak in large Prandtl number (thermal or compositional) convection \citep[e.g.][]{bS06}. The so-called large Prandtl number asymptotic limit can be used to remove inertia from the momentum equation; the resulting reduced system of equations are routinely employed for numerical investigations of convection in the Earth's subsolidus mantle where Prandtl numbers in excess of $O(10^{20})$ are typical \citep[e.g.][]{gS01}. In liquid metal planetary interiors the large Prandtl number limit is certainly invalid for the case of thermal convection, though it may be a useful approximation for describing compositional convection. We stress here, however, that the magnetic Prandtl number $Pm = \nu / \eta = O(10^{-5})$ in the Earth's core, where $\eta$ is the magnetic diffusivity. The large Prandtl number limit provides an interesting end-member case of convection in which all wave motion (inertial and magnetic) is absent; this approximation may then allow for assessing the influence that different waves might have on the dynamo when compared to cases that include inertia. Many previous dynamo investigations have neglected the influence of inertia \citep[e.g.][]{kZ90,gG95b,cJ00b,jR02,dH16}. One of the arguments given in the literature for this simplification is that such effects are likely small in the Earth's core because the Rossby number, $Ro = \mU/2\Omega L$, is small ($\mU$ is a characteristic flow speed, $\Omega$ is the rotation rate, and $L$ is a characteristic length scale). However, a flow that is characterized by a small Rossby number does not, by itself, imply that inertia is not important in the dynamics. It is well known that small Rossby number flows can be highly turbulent, i.e.~that the available potential energy in small departures from geostrophic balance is substantial \citep[e.g.][]{jP87}. It is therefore necessary to rely on employing either the large Prandtl number limit, or to arbitrarily restrict the flow to weakly supercritical states in order to justify rigorously the neglect of inertia; both of these approximations require the Reynolds number $Re = U L/\nu$ to be small. Despite the significant numerical advantages associated with eliminating inertia, the resulting set of equations still possess significant stiffness owing to the intrinsic separation in scales that occurs when the Rossby number is small. The mathematical result of this limit is that spatial derivatives perpendicular to the rotation axis become asymptotically larger than derivatives parallel to the rotation axis \citep{sC61} and that fast inertial waves become weakly damped. For this reason, the computational models that have employed the large Prandtl number approximation are still unable to reach the relevant geophysical limit of small Ekman number, $Ek = \nu / 2\Omega L^2 \ll 1$. It is therefore advantageous to pursue further reduction strategies that exploit the scale separation of rapidly rotating convection with the use of multiscale asymptotics \citep[e.g.][]{kJ07}. Such an approach has proven invaluable for investigating the onset of linear convection in spherical geometries \citep{pR68,fB70,cJ00,eD04}, weakly nonlinear dynamo action in the plane layer geometry \citep{sC72,aS74,yF82}, and for the development of a fully nonlinear reduced convection model in the plane layer \citep[][]{mS06,kJ12}. \cite{mC15b} recently extended the Childress-Soward weakly nonlinear dynamo model by developing a fully nonlinear reduced dynamo model that is capable of simulating dynamo action well above the onset of convection; we refer to this new reduced model as the quasi-geostrophic dynamo model (QGDM). In the present work we extend the QGDM to the large Prandtl number limit for the purpose of investigating the physical ramifications. | In the present work we have developed two self-consistent asymptotic dynamo models valid for inertia-less, large Prandtl number, rapidly rotating convection. In contrast to the previous work of \cite{mC15b}, it is the Roberts number $q$, rather than $Pm$, that becomes the important diffusivity parameter in the final model. However, we can relate these two parameters upon noting that $q = Pm/Pr$. We emphasize here that $q$ should be interpreted in terms of a compositional Roberts number. For $q = O(1)$, this implies that $Pm = O(Pr)$, i.e.~$Pm \gg 1$, which is clearly not relevant geophysically. For the $q = O(Ek^{1/2})$ model we have $Pm = O(Pr Ek^{1/2})$; to obtain $Pm \ll 1$ we require that $Pr < Ek^{-1/2}$, or in terms of $Ek_H$ this becomes $Pr < Ek_H^{-1/6}$. This latter requirement is fully consistent with the asymptotics, showing that it is possible to drive a low $Pm$, low $q$ dynamo with inertia-less convection in the limit of rapid rotation. Recalling that for the Earth's core $Ek_H = O(10^{-15})$, so that the large Prandtl number approximation is only valid provided $Pr \lesssim O(300)$; studies suggest that this inequality does hold for most chemical species \citep{mP13}. We note that both the small $q$ model of the present work and the small $Pm$ model of \cite{mC15b} are small magnetic Reynolds number models, $Rm = U\ell/\eta \ll 1$. In addition, the small $q$ model is characterized by a small convective-scale Reynolds number, $Re = U\ell/\nu \ll 1$; this follows from the fact that the P\'eclet number $Pe = U\ell/D = Pr Re = O(1)$ to allow convective heat transport. How, then, are such flows capable of driving small $Pm$ dynamos in light of the relation $Rm = Re Pm$? The answer lies in the fact that the large-scale magnetic Reynolds number, $Rm_H = Rm/Ek$, remains large since we have \be Rm_H = \frac{Rm}{Ek} = Ek_H^{-1/6} \gg 1 . \ee Thus, magnetic induction dominates ohmic diffusion on the large ($Z$) scale, consistent with the fact that both of the models presented here, and those developed in \cite{mC15b}, are large-scale dynamos. Indeed, setting the mean field to be identically zero in any of the variations of the QGDM will eliminate all dynamo action. We can also place a bound on the Reynolds number by utilizing the relation $Rm = Re Pm$. For the small $Rm$ limit this implies that $Re > Ek^{1/2}$ for the fluid to be considered low $Pm$ since $Re = Pe/Pr = O(1/Pr)$. The neglect of inertia is therefore limited to a finite range in $\Rat$ since $Re$ will grow as the buoyancy forcing is increased. Since $Re$ is an unknown function of $\Rat$, we require detailed numerical simulations of the QGDM to determine the appropriate range in $\Rat$ over which the large Prandtl number limit remains accurate. Table \ref{T:lims} shows the various asymptotic limits employed in the present work and the $Pr=O(1)$ cases of \cite{mC15b}. Taken together, these investigations show that the ratio of magnetic to kinetic energy densities, $M$, is not indicative of dominant balances in the governing equations since all of these models are geostrophically balanced. Rather, the size of $M$ is related to the presence of Alfv\'en waves in the reduced dynamics; order one values of $M$ indicate the presence of these waves, whereas such waves are filtered if $M \gg 1$. This filtering can be obtained by asymptotically eliminating the time derivative of either the velocity field in the fluctuating momentum equation or the time derivative of the magnetic field in the fluctuating induction equation. As shown in the present work, the elimination of both of these time derivatives yields a separation in the convective and Alfv\'en timescales that is larger than when only eliminating one of these time derivatives. Consistent with DNS, only for the $(Pr,Pm)=O(1)$ case do we find that $M=O(1)$. There is very little available data for small $q$ or small $Pm$ due to the computational cost of such studies, so no comparison can be made at present with the asymptotic relations for these limits. However, \cite{sS04} have reported values of $M$ (denoted as $E_{mag}/E_{kin}$ in their Table II) for all of their DNS cases with $q=1$. Of particular interest for comparison with the current work are their $Ek_H = 5 \times 10^{-6}$ cases for which they consider the three Prandtl numbers $Pr=1$, $10$, and $30$, where the Rayleigh number is the same for all three cases. In order of increasing Prandtl number they find that $M \approx 2, 11$, and $28$, respectively. These computed values compare well with the predictions $M=O(1)$, $O(10)$, and $O(30)$ based on the $Pr=O(1)$ model of \cite{mC15b} and the present $q=O(1)$ model. \begin{table} \begin{center} \begin{tabular}{lcccc} $Pr$ & $q$ & $Pm$ & $M$ \\ \hline \hline $\gg 1$ & $O(1)$ & $O(Pr)$ & $O(Pr)$ \\ $\gg 1$ & $O(Ek^{1/2})$ & $O(Pr Ek^{1/2})$ & $O(Pr Ek^{-1/2})$ \\ $O(1)$ & $O(1)$ & $O(1)$ & $O(1)$ \\ $O(1)$ & $O(\ep^{1/2})$ & $O(\ep^{1/2})$ & $O(\ep^{-1/2})$ \\ \end{tabular} \caption{Various asymptotic limits for the dynamo models developed in the present work ($Pr \gg 1$) and in \cite{mC15b} where $Pr=O(1)$. $Pr$ is the Prandtl number, $q$ is the Roberts number and $Pm = q Pr$ is the magnetic Prandtl number. The parameter $M$ is the ratio of magnetic to kinetic energy densities. The Rossby number is given by $\ep = Re Ek$, where $Re$ and $Ek$ are the Reynolds and Ekman numbers based on the small convective scale.} \label{T:lims} \end{center} \end{table} Although the present models are significantly oversimplified with respect to the Earth's outer core, it is nevertheless a useful exercise to determine if the asymptotic limits taken here are consistent with what is known about the geodynamo. Studies suggest a magnetic field strength in the core of $\mathcal{B} \sim 1 \, \text{mT}$ \citep[e.g.][]{nG10} and a flow speed of $\mathcal{U} \sim 10^{-4} \, \textnormal{m} \, \textnormal{s}^{-1}$ \citep[e.g.][]{cF11}, leading to an Alfv\'en speed of $\mathcal{U}_\mathcal{A} \sim 10^{-2} \, \textnormal{m} \, \textnormal{s}^{-1}$. These values lead to an estimate of $M \sim 10^4$ where, to remain consistent with the asymptotic derivation, we assume that the observed velocity scales diffusively as given by equation \eqref{E:Mdiff}. Since we know that $(q,Pm) \ll 1$ in the outer core, the small $q$ inertia-less model of the present work and the small $Pm$ model of \cite{mC15b} give values of $M \sim 3\times10^4$ and $M \sim 3\times10^2$, respectively, where we have assumed that $Ek_H = 10^{-15}$ and $Pr=100$; utilizing smaller values of $Pr$ will yield corresponding reductions in our estimate of $M$. Although the large $Pr$ asymptotic estimate of $M$ is in better agreement with our observational estimate of $M$, we emphasize that there is enough variation and uncertainty in the values of $\mathcal{U}$ and $\mathcal{U}_\mathcal{A}$ that may imply the $Pr=O(1)$ model is more relevant to core dynamics. Nevertheless, such a large value of $M$ for the core does suggest that Alfv\'en waves are not likely to be important for the geodynamo on the small scales of convection, though they can play an important role on the large scales of the core \citep[e.g.][]{nG10}. Simulations of the reduced dynamo models will better allow us to assess which approximation is more suitable for understanding the geodynamo. Recent kinematic investigations for the $Pr=O(1)$ QGDM are reported in \cite{mC16} and \cite{mC16b}. It is common in the literature to distinguish dynamo models based on the strength of the self-generated magnetic field. \cite{sC72} defined the field strength based on the asymptotic ordering of the Hartmann number $Ha_H = \mathcal{B} H/ \lb \rho \mu \eta \nu \rb^{1/2} = Ek_H^{-1/6} \lb q M / Pr \rb^{1/2}$, and the deviation of the Rayleigh number from its critical value. With regards to the \cite{sC72} terminology, the weak-field and intermediate-field regimes were investigated by \cite{aS74} and \cite{yF82}, and characterized by Hartmann numbers $Ha_H = O(1)$ and $Ha_H = O(Ek_H^{-1/6})$, respectively. All four of the QGDM variations identified thus far (listed in Table \ref{T:lims}) are characterized by a Hartmann number $Ha_H = O(Ek_H^{-1/3})$. However, we emphasize that both \cite{aS74} and \cite{yF82} were weakly nonlinear investigations in the sense that the mean temperature profile is always near that of the conductive state. As a result, the ratio of magnetic to kinetic energy $M$ is asymptotically small for both \cite{aS74} and \cite{yF82} \citep[e.g.~see][]{mC15b}. Therefore, both weak-field and intermediate-field dynamos have a magnetic energy that is asymptotically smaller than the kinetic energy in the flow. This characteristic is in stark contrast to the four variations of the QGDM, which describe \textit{saturated} dynamos that can have magnetic energy of the same order as, or asymptotically larger than, the kinetic energy. It is thus an oversimplification to categorize turbulent dynamos in rapidly rotating systems as either weak- or strong-field. Indeed, the Lorentz can have just as much influence on the quasi-geostrophic flow as does the buoyancy force in the QGDM, as demonstrated by equation \eqref{E:fmom}. An important question is whether there is a signature of the predominant forcing mechanism for the geodynamo in geomagnetic field observations. For instance, does the morphology of the geomagnetic field, or behavior of the secular variation depend upon whether thermal or compositional forcing is dominant in the core? Convection simulations find that the dynamics can indeed exhibit quite distinct behavior depending upon the value of the Prandtl number \citep[e.g.][]{mB10,mC12b}. The dynamics become further complicated when both forcing mechanisms are present \citep{tT12,fT14}. Dynamo simulations find that this Prandtl number dependence can result in significant changes in the structure of the resulting magnetic field \citep{fB06,bS06}. Moreover, `two-and-a-half' dimensional models find that inertia changes the temporal character the resulting dynamo \citep{dF01}, and mean-field models have found that the presence of inertia tends to facilitate dynamo action and therefore leads to stronger magnetic fields \citep{dF04}. The present model will allow for an extension of this previous work to the limit of rapid rotation and realistic fluid properties. | 16 | 9 | 1609.04061 |
1609 | 1609.03531_arXiv.txt | ASTROSAT is India's first satellite fully devoted to astronomical observations covering a wide spectral band from optical to hard X-rays by a complement of 4 co-aligned instruments and a Scanning Sky X-ray Monitor. One of the instruments is Large Area X-ray Proportional Counter with 3 identical detectors. In order to assess the performance of this instrument, a balloon experiment with two prototype Large Area X-ray Proportional Counters (LAXPC) was carried out on 2008 April 14. The design of these LAXPCs was similar to those on the ASTROSAT except that their field of view (FOV) was 3$^{\circ}$ $\times$ 3$^{\circ}$ versus FOV of 1$^{\circ}$ $\times$ 1$^{\circ}$ for the LAXPCs on the ASTROSAT. The LAXPCs are aimed at the timing and spectral studies of X-ray sources in 3-80 keV region. In the balloon experiment, the LAXPC, associated electronics and support systems were mounted on an oriented platform which could be pre-programmed to track any source in the sky. A brief description of the LAXPC design, laboratory tests, calibration and the detector characteristics is presented here. The details of the experiment and background counting rates of the 2 LAXPCs at the float altitude of about 41 km are presented in different energy bands. The bright black hole X-ray binary Cygnus X-1 (Cyg X-1) was observed in the experiment for $\sim$ 3 hours. Details of Cyg X-1 observations, count rates measured from it in different energy intervals and the intensity variations of Cyg X-1 detected during the observations are presented and briefly discussed. | \label{introduction} ASTROSAT is a multi-wavelength astronomy satellite which will carry four co aligned X-ray and UV instruments and a Scanning Sky Monitor (SSM), for wide spectral band studies of cosmic sources (\cite{Agrawal 2006a}, \cite{Agrawal 2006b}). Out of the 5 instruments aboard the ASTROSAT, four are sensitive in the X-ray band and one instrument namely Ultraviolet Imaging Telescope (UVIT) covers Visible (320-530 nm), Near UV (180-300 nm) and Far UV (130-180 nm) bands. The 4 X-ray instruments are : (1) The LAXPC instrument, with an effective area of $\ge$ 6000 cm$^{2}$ covers 3-80 keV spectral band for high time resolution studies and low spectral resolution studies of X-ray sources. (2) The Soft X-ray Telescope (SXT) uses conical foils based X-ray reflecting mirrors with an X-ray sensitive Charge Coupled Detector (CCD) at the focal plane to record low resolution ($\sim$ 3 arc minutes angular resolution) X-ray images and with moderate resolution X-ray spectra in the 0.5-8 keV spectral band. (3) The Cadmium - Zinc - Telluride Imager (CZTI) with geometrical area of 1000 cm$^{2}$ and energy resolution of about 10\% at 60 keV uses an array of CZT detectors and has a coded aperture mask placed above the detector plane for recording intensity and spectra of sources with angular resolution of $\sim$ 8 arc minute in the 10-100 keV hard X-ray band. (4) A Scanning Sky Monitor (SSM) consisting of 3 position sensitive proportional counters equipped with one dimensional coded aperture masks is aimed at studies of time variability of new transients and known X-ray sources in 2-10 keV interval. All the X-ray instruments except SSM and the UVIT are co-aligned and mounted on the top deck of the satellite for simultaneous observations of cosmic sources over the broad spectral band. For a more detailed description of ASTROSAT refer to \cite{Agrawal 2006a}. In the following section we present a detailed description of the design of the LAXPC instrument and its characteristics. | The main objective of LAXPC Balloon Flight Experiment was to validate the design of the LAXPC detectors by testing their performance in space-like environment. An additional aim was to measure the detector background rates and assess sensitivity for studying cosmic sources. The experiment also provided assessment of performance of the HV units, signal processing electronics of detectors and data modes of LAXPC instrument. The LAXPC detectors and associated electronic systems performed as expected during the balloon flight. The measured background rates are consistent with those estimated from an independent Monte Carlo simulation of the detector response (unpublished). A timing accuracy of 10 $\mu$s in the time tagging mode of data acquisition during the balloon-borne experiment was realized. This is necessary to study kHz QPO and other types of rapid variability. Results from the balloon experiment validate the design of the LAXPC instrument and it is hoped that this instrument will fulfil its science objectives. | 16 | 9 | 1609.03531 |
1609 | 1609.04916_arXiv.txt | New accurate and homogeneous optical {\it UBVRI} photometry has been obtained for variable stars in the Galactic globular $\omega$ Cen (\ngc{5139}). We secured 8202 CCD images covering a time interval of 24 years and a sky area of 84$\times$48 arcmin. The current data were complemented with data available in the literature and provided new, homogeneous pulsation parameters (mean magnitudes, luminosity amplitudes, periods) for 187 candidate $\omega$ Cen RR Lyrae (RRLs). Among them we have 101~RRc (first overtone), 85~RRab (fundamental) and a single candidate RRd (double-mode) variables. Candidate Blazhko RRLs show periods and colors that are intermediate between RRc and RRab variables, suggesting that they are transitional objects. The comparison of the period distribution and of the Bailey diagram indicates that RRLs in $\omega$ Cen show a long-period tail not present in typical Oosterhoff II (OoII) globulars. The RRLs in dwarf spheroidals and in ultra faint dwarfs have properties between Oosterhoff intermediate and OoII clusters. Metallicity plays a key role in shaping the above evidence. These findings do not support the hypothesis that $\omega$ Cen is the core remnant of a spoiled dwarf galaxy. Using optical Period-Wesenheit relations that are reddening-free and minimally dependent on metallicity we find a mean distance to $\omega$ Cen of 13.71$\pm$0.08$\pm$0.01 mag (semi-empirical and theoretical calibrations). Finally, we invert the {\it I\/}-band Period-Luminosity-Metallicity relation to estimate individual RRLs metal abundances. The metallicity distribution agrees quite well with spectroscopic and photometric metallicity estimates available in the literature. | \label{chapt_intro_omega} The Galactic stellar system $\omega$ Cen lies at the crossroads of several open astrophysical problems. It is the most massive Milky Way globular cluster ($4.05 \cdot 10^6 M_\odot [d/(5.5 \pm 0.2 \kpc)]^3$ where d is the distance, \citealp{dsouza_e_rix2013}) and was the first to show a clear and well defined spread in metal-abundance \citep{norrisdacosta1995,johnson_e_pilachowski2010} in $\alpha$ and in s- and r-process elements \citep{johnson2009}. On the basis of the above peculiarities it has also been suggested that $\omega$ Cen and a few other massive Galactic Globular Clusters (GGCs) might have been the cores of pristine dwarf galaxies \citep{dacosta_e_coleman2008,marconi2014}. The distance to $\omega$ Cen has been estimated using primary and geometrical distance indicators. The Tip of the Red Giant Branch (TRGB) was adopted by \citet{bellazzini2004,bono2008b} with distances ranging from 13.65 to 13.70 mag. The {\it K\/}-band Period-Luminosity (PL) relations of RR Lyrae stars (RRLs) have been adopted by \citet{longmore1990,sollima2006b,bono2008b}. The distance moduli they estimated range from 13.61 to 13.75 mag. On the other hand, $\omega$ Cen distance moduli based on the relations between luminosity and iron abundance for RRLs range from 13.62 to 13.72 mag \citep{delprincipe_etal2006}. The difference in distance between the different methods is mainly due to the intrinsic spread in the adopted diagnostics and in the reddening correction. Optical PL relations for SX Phoenicis stars were adopted by \citet{mcnamara2011} who found a distance of 13.62$\pm$0.05 mag. One eclipsing variable has been studied by \citet{kaluzny_etal2007}, and they found a distance modulus of 13.49$\pm$0.14 mag and 13.51$\pm$0.12 mag for the two components. The key advantage in dealing with eclipsing binaries is that they provide very accurate geometrical distances \citep{pietrzynski_etal2013}. Estimates based on cluster proper motions provide distance estimates that are systematically smaller than obtained from the other most popular distance indicators (13.27 mag, \citet{vanleeuwen_etal2000}; 13.31$\pm$0.04 mag, \citet{watkins_etal2013}). The reasons for this difference are not clear yet. The modest distance and the large mass of $\omega$ Cen make this stellar system a fundamental laboratory to constrain evolutionary and pulsation properties of old (t>10 Gyr) low-mass stars. The key advantage in dealing with stellar populations in this stellar system is that they cover a broad range in metallicity (--2.0$\lesssim$ [Fe/H] $\lesssim$--0.5, \citet{pancino2002}; --2.5$\lesssim$ [Fe/H] $\lesssim$+0.5, \citet{calamida_etal2009}; --2.2$\lesssim$ [Fe/H] $\lesssim$--0.6, \citet{johnson_e_pilachowski2010}) and they are located at the same distance \citep{castellani_etal2007}. Moreover, the high total stellar mass does provide the opportunity to trace fast evolutionary phases \citep{monelli_etal2005,calamida_etal2008} together with exotic \citep{randall_etal2011} and/or compact objects \citep{bono_etal2003}. Exactly for the same reasons mentioned before $\omega$ Cen was a crucial crossroads for RRLs. The first detailed investigation of RRLs was provided more than one century ago in a seminal investigation by \citet{bailey1902}. Using a large set of photographic plates he identified and characterized by eye 128 RRLs, providing periods, amplitudes and a detailed investigation of the shapes of the light curves. In particular, he suggested the presence of three different kind of pulsating variables (RRa, RRb, RRc) in which the luminosity variation amplitude steadily decreases and the shape of the light curve changes from sawtooth to sinusoidal. This investigation was supplemented more than thirty years later by \citet{martin1938} on the basis of more than 400 photographic plates collected by H. van Gent on a time interval of almost four years and measured with a microdensitomer. He provided homogeneous photometry and very accurate periods for 136 RRL variables. We needed to wait another half century to have a detailed and almost complete census of RRL in $\omega$ Cen based on CCD photometry, by the OGLE project \citep{kaluzny_etal1997,kaluzny_etal2004}. They collected a large number of CCD images in {\it V\/} and {\it B\/} covering a time interval of three years \citep{kaluzny_etal1997} and one and half years \citep{kaluzny_etal2004} and provided a detailed analysis of the occurrence of the Blazhko effect \citep[a modulation of the light amplitude on time scales from tens of days to years,][]{blazhko1907}. A similar analysis was also performed by \citet{weldrake_etal2007} using the observing facility and photometric system of the MACHO project. They collected 875 optical images covering a period of 25 days. A detailed near-infrared (NIR) analysis was performed by \citet{delprincipe_etal2006} using time series data collected with SOFI at NTT. They provided homogeneous {\it JK$_s$\/} photometry for 180 variables and provided a new estimate of the $\omega$ Cen distance modulus using the {\it K\/}-band PL relation (13.77$\pm$0.07 mag). A similar analysis was recently performed by \citet{navarrete_etal2015} based on a large set of images collected with the VISTA telescope. They provided homogeneous {\it JK$_s$\/} photometry for 189 probable member RRLs (101 RRc, 88 RRab) and discussed the pulsation properties of the entire sample in the NIR. In particular, they provided new NIR reference lines for Oosterhoff I (OoI) and Oosterhoff II (OoII) clusters. Moreover, they further supported the evidence that RRab in $\omega$ Cen display properties similar to OoII systems. These investigations have been complemented with a detailed optical investigation covering a sky area of more than 50 square degrees by \citet{fernandeztrincado2015}. They detected 48 RRLs and the bulk of them (38) are located outside the tidal radius. However, detailed simulations of the different Galactic components and radial velocities for a sub-sample of RRLs indicate a lack of tidal debris around the cluster. This is the fourth paper of a series focussed on homogeneous optical, near-infrared, and mid-infrared photometry of cluster RRLs. The structure of the paper is as follows. In \S~2 we present the optical multi-band {\it UBVRI\/} photometry that we collected for this experiment together with the approach adopted to perform the photometry on individual images and on the entire dataset. In subsection 3.1 we discuss in detail the identification of RRLs and the photometry we collected from the literature to provide homogeneous estimates of the RRL pulsation parameters. Subsection 3.2 deals with the period distribution, while subsection 3.3 discusses the light curves and the approach we adopted to estimate the mean magnitudes and the luminosity variation amplitudes. The Bailey diagram (luminosity variation amplitude vs period) is discussed in \S~3.4, while the amplitude ratios are considered in \S~3.5. Section 4 is focussed on the distribution of RRLs in the color-magnitude diagram (CMD) and on the topology of the instability strip. In \S~5 we perform a detailed comparison of the period distribution and the Bailey diagram of $\omega$ Cen RRLs with the similar distributions in nearby gas-poor systems (globulars, dwarf galaxies). Section 6 deals with RRL diagnostics, namely the PL and the Period-Wesenheit (PW) relation, while in \S~7 we discuss the new distance determinations to $\omega$ Cen based on optical PW relations. Section 8 deals with the metallicity distribution of the RRLs, based on the {\it I}-band PL relation, and the comparison with photometric and spectroscopic estimates available in the literature. Finally, \S~9 gives a summary of the current results together with a few remarks concerning the future of this project. | \label{chapt_final} We present new accurate and homogeneous optical, multi-band---{\it UBVRI\/}---photometry of the Galactic globular $\omega$ Cen. We collected 8202 CCD images that cover a time interval of 24 years and a sky area of 84$\times$48 arcmin across the cluster center. The bulk of these images were collected with the Danish telescope at ESO La Silla as time-series data in three main long runs (more than 4,500 images). The others were collected with several telescopes ranging from the 0.9m at CTIO to the VLT at ESO Cerro Paranal. The final photometric catalog includes more than 180,000 (Danish) and 665,000 (others) stars with at least one measurement in two different photometric bands. The above datasets were complemented with optical time series photometry for RRLs available in the literature. The global photometric catalog allowed us to accomplish the following scientific goals. {\em Homogeneity}---We provide new, homogeneous pulsation parameters for 187 candidate $\omega$ Cen RRLs. All in all the photometry we collected (proprietary$+$literature) covers a time interval of 36 years and the light curves of RRLs have a number of phase points per band that ranges from $\sim$10-40 ({\it U\/}), to$\sim$ 20-770 ({\it B\/}), to $\sim$20-2830 ({\it V\/}), to $\sim$10-280 ({\it R\/}) and to $\sim$10-445 ({\it I\/}). These numbers sum up to more than 300,000 multi-band phase points for RRLs, indicating that this is the largest optical photometric survey ever performed for cluster RRLs \citep{jurcsik2012,jurcsik2015}. The above data allowed us to provide new and accurate estimates of their pulsation parameters (mean magnitudes, luminosity variation amplitudes, epoch of maximum and epoch of mean magnitude). {\em Period distribution}---The key advantage in dealing with $\omega$ Cen is that its RRL sample is the 3$^{\mathrm{rd}}$ largest after M3 (237 RRLs) and M62 (217) among the globulars hosting RRLs. On the basis of the current analysis we ended up with a sample 187 candidate cluster RRLs. Among them 101 pulsate in the first overtone (RRc), 85 in the fundamental (RRab) mode and a single object is a candidate mixed-mode variable (RRd). We estimate the mean periods for RRab and RRc variables and we found that they are $<P_{ab}>= 0.668$ days, $<P_c>= 0.359$ days. The above mean periods and the population ratio, i.e., the ratio between the number of RRc and the total number of RRLs ($N_c/(N_{ab}+N_d+N_c)$) support previous findings suggesting that $\omega$ Cen is a Oosterhoff II cluster. {\em Bailey Diagram}---The luminosity variation amplitude vs period plane indicates a clear lack of HASP RRLs, i.e., RRab variables with P$\lesssim$0.48 days and $AV$ $>$ 0.75 mag \citep{fiorentino2015}. These objects become more popular in stellar systems more metal-rich than [Fe/H] $\approx$ --1.4, thus suggesting that RRL in $\omega$ Cen barely approach this metallicity range. The RRab variables that, from our investigation, appear to be more metal-rich than --1.4, have periods ranging from 0.49 to 0.72 days. Moreover, we also found evidence that RRc can be split into two different groups: a) short-period---with periods ranging from $\sim$0.30 to $\sim$0.36 days and visual amplitudes ranging from a few hundreths of a magnitude to a few tenths; b) long-period---with periods ranging from $\sim$0.36 to $\sim$0.45 days and amplitudes clustering around $AV\sim$0.5 mag. Theoretical and empirical arguments further support a well defined spread in iron abundance. {\em Amplitude ratios}---The well known spread in iron abundance of $\omega$ Cen stars makes its RRL sample a fundamental test-bench to characterize the possible dependence of amplitude ratios on metal content. We performed a detailed test and we found that both RRab and RRc attain similar ratios: $AB/AV$ = 1.26$\pm$0.01; $AR/AV$ = 0.78$\pm$0.01; $AI/AV$ = 0.63$\pm$0.01. Moreover, they do not display any clear trend with iron abundance. {\em Visual magnitude distribution}---We performed a detailed analysis of the visual magnitude distribution of RRLs and we found that they can be fit with four Gaussians. The two main peaks included a significant fraction of RRL ($\sim$76\%) and attain similar magnitudes ({\it V\/}$\sim$14.47, 14.56 mag). The fainter ({\it V\/}$\sim$14.71 mag) and the brighter ({\it V\/}$\sim$14.31 mag) peak include a minor fraction (11\%, 13\%) of the RRL sample. The above finding is suggestive of a spread in iron abundance of the order of 1.5 dex and paves the way for new solid estimates on the absolute age of the different stellar populations in $\omega$ Cen. {\em Blazhko RR Lyrae}---Empirical evidence based on the location of candidate Blazhko RRLs in the Bailey diagram and in the color-magnitude diagram clearly indicate that they are located between RRc and RRab variables. Indeed, we found that a significant fraction (79\%) of them (22 out of 28) have periods shorter than 0.6 days. Moreover, their location inside the instability strip indicates that a significant fraction (39\%) of them belongs to the fainter peak (V$\geq$14.6 mag), thus suggesting that this sub-sample is more associated with the more metal-rich stellar component. {\em Oosterhoff dilemma}---Dating back to the seminal investigation by \citet{oosterhoff1939} in which he recognized that cluster RRLs can be split, according, to their mean periods, into two different groups, the astronomical community undertook a paramount observational effort in order to constrain the physical mechanism(s) driving the empirical evidence. We performed a detailed comparison between the period distribution and the Bailey diagram of $\omega$ Cen RRLs with globulars hosting a sizable sample (>35) of RRLs and with RRLs in nearby dSphs and UFDs. We found, as expected, that the mean F and FO periods display a steady decrease when moving from the more metal-rich (Oosterhoff I) to the more metal-poor (Oosterhoff II) clusters. In this context dSphs and UFDs attain values that are intermediate between the OoInt and the OoII clusters, while $\omega$ Cen appears as the upper envelope of the distribution. On the other hand, the population ratio---$N_c/(N_{ab}+N_d+N_c)$---has a nonlinear trend, since it attains a well defined minimum for OoInt clusters. In spite of the possible differences, the iron abundance appears to be the key parameter in driving the transition from short mean periods to long mean periods stellar systems. The above results do not support the working hypothesis that $\omega$ Cen is the core remnant of dwarf galaxy \citep{bekkifreeman2003}. Moreover, there is mounting empirical evidence that cluster RRLs might not be the appropriate sample to address the Oosterhoff dichotomy, since they might be either biased by statistics or affected by environmental effects. {\em $\omega$ Cen disguised as a dwarf galaxy}---The number of globulars hosting long-period (0.82--0.85$\lesssim$P$\lesssim$1 days) RRLs is quite limited. Three honorable exceptions are $\omega$ Cen and the two metal-rich Bulge globulars hosting RRLs, namely \ngc{6388} and \ngc{6441}. The mean periods of the metal-rich clusters appear as an extreme case of OoI clusters. This is the reason why we suggest they should be classified as Oosterhoff type 0 instead of Oosterhoff type III. We note that the main common feature among these clusters is that the HB luminosity function shows a well developed blue tail. This indicates that the appearance of long-period RRLs is more nurture that nature. The environment, and in particular, the high central stellar density, might play a crucial role in the presence of a blue tail, and in turn of long-period RRLs. However, the observational scenario appears much more complex, since the RRLs in the metal-intermediate cluster \ngc{2808} hosts 11 RRab variables, but they have periods shorter than 0.62 days \citep{kunder2013a}. {\em Distance determination}---We take advantage of optical PW relations that are reddening independent by construction and minimally dependent on iron abundance to provide new estimates of the distance to $\omega$ Cen. We adopted both a semi-empirical and a theoretical calibration and we found a true distance modulus of 13.71$\pm$0.08$\pm$0.01 mag. They agree quite well with similar estimates available in the literature. In particular, we found that the agreement is within 1$\sigma$ with the geometrical distances based on eclipsing binaries \citep[13.49$\pm$0.14 / 13.51$\pm$0.12,][]{kaluzny_etal2007}. {\em Metallicity distribution}---We inverted the {\it I\/}-band PLZ relation for F and FO pulsators to provide individual metallicity estimates for 160 cluster RRLs. We found that the metallicity distribution agrees quite well with the metallicity distribution of RRLs based on spectroscopic measurements (74, S06) and on photometric indicators (131, R00). We also found evidence of a metal-poor tail that is not present in previous spectroscopic investigations of $\omega$ Cen RRLs. The current long-term photometric surveys are providing new and homogeneous measurements concerning field and cluster stars. The current status is going to experience a quantum jump as soon as the ongoing (Gaia) and near future ground-based experiments will release their data. The project we started more than 15 years ago on $\omega$ Cen, may be defined as a {\em local survey}. The number of optical images adopted to individuate main sequence and evolved variable stars have been discussed in detail on Section~\ref{chapt_obs_opt}. In dealing with the Danish dataset we analyzed 4539 images and we performed $\approx2.5\times10^8$ measurements, which means roughly 55,000 stars per image. In dealing with all the other optical datasets we analyzed 3663 images and performed $\approx1.1\times10^8$ measurements, which means roughly 27,000 stars per image. A similar number of measurements have been also performed in dealing with NIR images. The above numbers indicate once accounting for the preliminary steps in approaching the final photometric catalog, that we are dealing with an experiment that included more than one giga measurements. The results concerning variable and static stars will be addressed in a series of future papers in which we plan to use homogeneous multi-band optical, NIR and MIR photometry. | 16 | 9 | 1609.04916 |
1609 | 1609.00532_arXiv.txt | Modified gravity theories are a popular alternative to dark energy as a possible explanation for the observed accelerating cosmic expansion, and their cosmological tests are currently an active research field. Studies in recent years have been increasingly focused on testing these theories in the nonlinear regime, which is computationally demanding. Here we show that, under certain circumstances, a whole class of theories can be ruled out by using background cosmology alone. This is possible because certain classes of models (i) are fundamentally incapable of producing specific background expansion histories, and (ii) said histories are incompatible with local gravity tests. As an example, we demonstrate that a popular class of models, $f(R)$ gravity, would not be viable if observations suggest even a slight deviation of the background expansion history from that of the $\Lambda$CDM paradigm. | In the past decade or so, alternative theories of gravity as a possible explanation for the accelerating expansion of the Universe have received a great deal of attention \cite{cfps2011,joyce2015}. Such theories affect the dynamics of the expansion on cosmological scales, where General Relativity (GR) is usually assumed to break down, without invoking a mysterious new matter species commonly known as dark energy. Thanks to the development of linear and nonlinear computational tools in recent years, this area has advanced quickly, with the formation of large-scale structures in many of the new models being fairly well understood by now, and the study of baryonic and galaxy evolution in them already initiated by some groups \cite[e.g.,][]{aps2014, hlmw2015, hl2015}. There are, however, a few challenges hindering further development of the field. Many of the alternative theories, such as $f(R)$ gravity \cite{sf2010,dt2010}, indeed have GR as a limit, which means that there is some point (typically characterised by one or more model parameters) after which the theory is no longer distinguishable from GR in practice. In $f(R)$ gravity, for example, Ref.~\cite{jvj2013} shows that a model parameter, $|f_{R0}|$ (to be explained below), has to be smaller than $\sim10^{-7}$ for it to satisfy astrophysical constraints, thus making the cosmology of the model very similar to the general-relativistic prediction. We therefore face the situation that a cosmological model might never be ruled out by cosmological observations. Adding to this is the fact that studies of nonlinear structure formation in the remaining allowed parameter space are increasingly more challenging with ever higher resolution requirements, and systematics and uncertainties start to dominate over model differences from GR. Hence, it is beneficial to find other, hopefully cleaner, ways of testing the models using cosmology. One place we can look into, as we shall show below, is background cosmology. This may sound counter-intuitive: after all, given the purpose of modified gravity theories, fitting background cosmology seems to be the first test they need to pass. However, many of these alternative theories are known to have great flexibility -- for example, the fourth-order nature of the $f(R)$ gravity equations means that there is an infinite family of models which can exactly reproduce the background expansion history of the $\Lambda$CDM scenario \cite{shs2007}, thus giving us the freedom to simply adopt this standard background and focus on other effects (e.g., the fifth force) on cosmic structure formation. In this paper, we revisit the role of background cosmology in constraining modified gravity theories. With $f(R)$ gravity as a working example, we will demonstrate that the model is incapable of reproducing certain expansion histories. Furthermore, we exemplify the restrictions on the expansion history itself brought about by the findings of Ref.~\cite{bbds2008}, namely that for this model to be viable its background cosmology has to be very close to the $\Lambda$CDM prediction. This result is generic and model-independent, as it is not a direct constraint on $f(R)$ model parameters. Therefore, if future observations support a dark energy equation-of-state parameter $w$ that is different from $-1$ or evolves in time, the whole $f(R)$ class of theories as an explanation to the cosmic acceleration could be ruled out. This highlights the importance and potential benefits of employing future background cosmological observations in tests of gravity. This paper is organised as follows. In Section II, we give a brief overiew of the theory behind $f(R)$ gravity and the relevant field equations. In Section III, we explain how there are certain expansion histories that cannot be reproduced whatever the functional form of $f(R)$, in spite of the fourth-order nature of the theory. In Section IV, we show how deviations from the $\Lambda$CDM expansion history would require $f(R)$ to take on a form that makes it difficult to satisfy local gravity tests. We then give an example of the constraints that can be placed using these arguments in Section V, and present our conclussions in Section VI. | The results indicate a special property of $f(R)$ gravity, namely that, for it to be viable, the expansion history cannot be arbitrary but has to be very close to $\Lambda$CDM. More precisely, parameter values $(w_0,w_1)$ in the dark energy equation-of-state $w(a)$ shown in Eq.~(\ref{eq:eos}) {\it cannot} differ from $(-1,0)$ by more than $\mathcal{O}(10^{-5})$, if $f(R)$ gravity is the underlying gravity model. Likewise, any observational evidence for a significant time evolution of $w(a)$ would rule out $\Lambda$CDM and the entire $f(R)$ class of models simultaneously. \begin{figure} \includegraphics[scale=0.45]{countour.eps} \caption{\label{fig:countour}(Color online) {\it Main panel}: Region of the $w_0$-$w_1$ plane where $R(a)$ has no extrema at $a\leq1$, allowing the specified expansion history to be consistently reproduced by $f(R)$ gravity. {\it Insert}: Region where the reconstructed $f(R)$ model can pass solar system tests (see the main text for more details). Red crosses correspond to $\Lambda$CDM.} \end{figure} For this reason, it is crucial to further improve the observational constraints on $w(a)$. Most current limits do not rule out $w(a)=-1$, though in some cases it is not the best fit \cite{zcpz2012}. Future galaxy surveys, such as Euclid \cite{euclid} and {\sc desi} \cite{desi}, have the potential of reducing the uncertainty on $w_0$ and $w_1$ to $\Delta w_0\sim0.01$ and $\Delta w_1\sim0.05$. This means that small deviations -- if they exist -- from $w=-1$ can be measured, and in turn be used to rule out $f(R)$ gravity. One caveat, however, is that these estimates often stem from the synergy of different probes or even different surveys (e.g., including Planck): if the probes rely on the growth rate of matter perturbations, the derived constraints on $w_0, w_1$ depend on the gravity model (which in most forecasts is taken to be GR), and can't be used universally. Geometric measures, such as the baryon acoustic oscillation peak positions, could be used in a more model-independent way, but on their own the constraints would be weaker. Although most research efforts so far have focused on the goodness-of-fit of $w(a)$ parameterisations, the dark energy equation of state can indeed also be reconstructed nonparametrically with the fewest possible assumptions. Ref.~\cite{zcpz2012} contains such an example of reconstruction, where it is found that a time-varying $w(a)$ is slightly preferred over the $\Lambda$CDM case. Such work can prove invaluable in constraining theories like $f(R)$ gravity. From a more general point of view, $f(R)$ gravity is a subclass of the chameleon theory \cite{kw2004}, so its screening mechanism relies similarly on the scalar field staying small from very high to low curvatures, or equivalently from very early times to today \cite{bdl2012}. This means that the scalar field will barely evolve and its potential energy, which drives the cosmic acceleration, will stay nearly constant in time: once more, we have a background expansion history that has to be close to $\Lambda$CDM, as generically described by \cite{bbds2008}. To summarise, alternative theories of gravity have been extensively studied in the past few years due to their potential to explain the accelerated cosmic expansion. Although such theories have rich phenomenology in terms of structure formation, and can therefore be constrained using observations associated with the latter, we argue that in certain cases the expansion history itself may be used as a smoking gun to rule out classes of theories. This is based on two observations: first, a theory could be intrinsically incapable of producing certain background expansion histories, by analogy with the observation that quintessence models cannot produce a phantom ($w<-1$) background evolution; second, a theory capable of producing a certain expansion history might yield unexpected and unwanted phenomena at small scales. We show that both possibilities indeed happen in one of the most popular theories -- $f(R)$ gravity or the chameleon theory -- leading to the strong constraint that the expansion history must be very close to that of $\Lambda$CDM. Consequently, precise measurements of the dark energy equation of state, $w(a)$, can be useful in ruling out such theories in the future. In contrast, studies of structure formation, which have been the main focus in recent times, are more likely to result in reduced parameter spaces rather than the exclusion of whole classes of models. | 16 | 9 | 1609.00532 |
1609 | 1609.09763_arXiv.txt | We derive metallicities for 41 cataclysmic variables (CVs) from near-infrared spectroscopy. We use synthetic spectra that cover the 0.8 $\mu$m $\leq \lambda \leq$ 2.5 $\mu$m bandpass to ascertain the value of [Fe/H] for CVs with K-type donors, while also deriving abundances for other elements. Using calibrations for determining [Fe/H] from the $K$-band spectra of M-dwarfs, we derive more precise values for T$_{\rm eff}$ for the secondaries in the shortest period CVs, and examine whether they have carbon deficits. In general, the donor stars in CVs have sub-solar metallicities. We confirm carbon deficits for a large number of systems. CVs with orbital periods $>$ 5 hr are most likely to have unusual abundances. We identify four CVs with CO emission. We use phase-resolved spectra to ascertain the mass and radius of the donor in U Gem. The secondary star in U Gem appears to have a lower {\it apparent} gravity than a main sequence star of its spectral type. Applying this result to other CVs, we find that the later-than-expected spectral types observed for many CV donors is mostly an affect of inclination. All of the magnetic CVs, except the low accretion rate polar MQ Dra, have donors with subsolar metallicities. We find that two systems with unusual spectra, EI Psc and QZ Ser, that have large excesses of sodium, and extreme deficits of carbon. Synthetic spectra that have a reduced abundance of hydrogen are best able to explain the spectra of these two objects. | Cataclysmic variables (CVs) result when the cool, low mass secondary in a binary contacts its Roche lobe and transfers material to the white dwarf primary. To arrive at this stage, the standard evolutionary paradigm involves having the more massive component in the binary evolve off the main sequence and engulf its companion in a short-lived common envelope (CE) phase (e.g., Iben \& Livio 1993). During the CE phase, the outer atmosphere of the primary star is removed by interaction with the secondary star, and this interaction leads to either a stellar merger, or a more compact binary (Politano \& Weiler 2007). Angular momentum loss, such as through ``magnetic braking'' (Verbunt \& Zwaan 1981), eventually leads to a semi-contact binary and a CV is born. In most of the early population synthesis calculations (cf., Howell et al. 2001, and references therein) the vast majority of donor stars in CVs were expected to be unevolved. If true, the donor stars should generally appear to be similar to their main sequence counterparts, except for the effects introduced by the mass loss process. Thus, while they will have a slightly different mass-radius relationship due to the secondary being out of thermal equilibrium, their photospheres should not show any abundance peculiarities. The standard evolutionary paradigm though is slowly changing to reflect the fact that many long period CVs appear to have donors with unusual masses and/or radii. This suggests that the secondaries in some subset of CVs have undergone significant evolution prior to becoming interacting binaries. Podsiadlowski et al. (2003) constructed a population synthesis for CVs that allowed for evolved secondaries with initial masses of up to 1.4 M$_{\sun}$. Inclusion of these objects allowed them to explain the wide dispersion in the observed temperatures of donor stars in CVs with P$_{\rm orb}$ $>$ 5 hr. This result confirmed the conclusions of Baraffe \& Kolb (2000), that long period CVs are dominated by systems containing evolved secondaries. Recently, Goliasch \& Nelson (2015) have produced a new population synthesis calculation that includes detailed nuclear evolution. They confirm the results of Podsiadlowski et al., showing that there is a large range of masses and radii for donors in CVs with P$_{\rm orb}$ $>$ 5 hr when nuclear evolved secondary stars are considered. Two other interesting results from this study are that CVs with evolved secondaries do not have a period gap\footnote{The observed dearth of CVs with 2 hr $\leq$ P$_{\rm orb}$ $\leq$ 3 hr (see Kolb et al. 1998).}, and they can evolve to much shorter periods than predicted by models where evolved donors are not considered. Podsiadlowski et al. suggest that the secondaries that have undergone pre-CV evolution should show evidence for nuclear processing, in particular, products produced via the CNO cycle. Two alternative paths that might allow the donor stars in CVs to acquire peculiar abundance patterns are for them to accrete such material during the original common envelope phase, or during the briefer common envelope phase that results from a classical nova eruption. Marks \& Sarna (1998) have investigated a large number of scenarios for the composition of the surface layers of CV donor stars that have accreted both types of material, and find that significant changes are possible. A common theme amongst these models are large deficits of carbon. In a series of papers we have presented near-IR $K$-band spectra that demonstrated that the secondaries of many CVs appeared to have significant carbon deficits (Harrison et al. 2004a, 2005ab, 2009). Recently, Harrison \& Hamilton (2015, ``H\&H'') used synthetic spectra to quantify that SS Cyg, RU Peg, and GK Per had carbon abundances that were 10 to 30\% of the solar value. In contrast, however, they found that the metallicities ([Fe/H]) of the donor stars were solar, or only slightly subsolar. The study by H\&H was limited to the $K$-band and, as they describe, there are not many strong atomic absorption lines in that bandpass to derive more precise values of [Fe/H], nor enough spectral lines from any species other than carbon to allow derivation of their abundances. To remedy this issue, we have developed the capability to generate synthetic spectra that cover the $IJHK$ bandpasses. This allows us to investigate the abundances of other important elements for the subset of CVs we have observed using cross-dispersed spectrographs. From modeling these data we find a wide range of values for [Fe/H], and for [C/Fe]. In addition, however, we find a number of CVs that appear to have unusual abundances of aluminum, magnesium, and/or sodium. Due to the lack of an extensive line list (and other uncertainties discussed below) our abundance analysis has been limited to CVs whose secondaries have spectral types earlier than M0. Recently, relationships have been developed to extract metallicities for M dwarfs. Using those relationships it is possible to examine the cooler secondary stars found in short period CVs for abundance anomalies. We describe the observations in the next section, the generation of synthetic spectra in Section 3, the results on CVs with K-dwarf secondaries in Section 4, the process and data needed to generate metallicities for M dwarfs in Section 5, discuss our results in Section 6, and present our conclusions in Section 7. | We have ascertained the metallicity for the donor stars in a large number of CVs using NIR spectroscopy. For CVs with K-type donor stars, temperatures and abundances have been derived through the comparison of synthetic spectra to observations. For CVs with M-type stars, the values for [Fe/H] were derived from recently developed relationships for field M dwarfs that employ $K$-band spectroscopy. We find that the donor stars in CVs generally have subsolar metallicities. This is consistent with results for field stars in the solar neighborhood, where the local stellar population is shown to be a mix of thin and thick disk stars (Nordstr\"{o}m, et al. 2004). An analysis of the kinematics of CVs by Ak et al. (2015) derives a similar result. The donors in longer period CVs appear to have more dramatic abundances anomalies than do the short period systems. In agreement with predictions from population synthesis modeling. The presence of Mg deficits in several long period CVs argues that the donor stars in these systems had large initial masses. A caveat for the determination of [Fe/H] exists for CVs with M-type donors, as variations in the C/O ratio could lead to the derivation of artificially low values for [Fe/H] (Veyette et al. 2016). This concern is especially relevant given that many CVs have carbon deficits. Our comparison of the high S/N spectrum of U Gem to that of GJ 402 showed that the predicted metallicity offset implied by the observed carbon deficit produced a better agreement between the two objects. However, application of a similar offset to the derived [Fe/H] value for carbon-deficient YY/DO Dra made it {\it less} consistent with the best-fitting template. Due to the extreme carbon deficits seen in some donors, CVs could provide the data to better establish and/or calibrate the effect of the C/O ratio on the metallicity estimates for normal field stars. This could be easily accomplished with new multi-bandpass NIR spectroscopy of those CVs in Table 3 with proven carbon deficits. It is also now abundantly clear that while the donor stars in CVs may superficially resemble their main sequence cousins, many of them have undergone significant nuclear evolution prior to becoming a CV. The inclusion of such objects into recent population studies is a welcome change, and demonstrates that stellar evolution processes alone are able to explain many of our results. It would be extremely useful to have model atmospheres for cool stars that had both hydrogen deficiencies and enhanced levels of helium, so as to enable the generation of more realistic synthetic spectra. New models for the outbursts of dwarf novae that employ non-solar abundance patterns and/or hydrogen deficiencies should also be considered, so as to investigate whether composition has any affect on outburst properties. The same is true for models of classical novae eruptions, as the composition of the accreted material appears to alter the nature of these events (Shen \& Bildsten 2009). We close by noting that the total mass of the U Gem system is larger than the Chandrasekhar limit, and if the mass lost to CNe eruptions is not too great, the system could become a Type Ia supernova. Note that if the donor star in U Gem is even modestly hydrogen deficient, one of the main arguments against single-degenerate CVs as SNIa progenitors is removed: the lack of H I emission lines in their post-deflagration spectra. Pakmor et al. (2008) have found that models for the amount of matter stripped from a main sequence donor star in an SNIa explosion is roughly consistent with the limits implied by observations. Hydrogen deficits in CV donors would only aid this result. | 16 | 9 | 1609.09763 |
1609 | 1609.07062_arXiv.txt | {M-type subdwarfs are metal-poor low-mass stars and are probes for the old populations in our Galaxy. Accurate knowledge of their atmospheric parameters and especially their composition is essential for understanding the chemical history of our Galaxy. } {The purpose of this work is to perform a detailed study of M-subdwarf spectra covering the full wavelength range from the optical to the near-infrared. It allows us to perform a more detailed analysis of the atmospheric composition in order to determine the stellar parameters, and to constrain the atmospheric models. The study will allow us to further understand physical and chemical processes such as increasing condensation of gas into dust, to point out the missing continuum opacities, and to see how the main band features are reproduced by the models. The spectral resolution and the large wavelength coverage used is a unique combination that can constrain the processes that occur in a cool atmosphere.} {We obtained medium-resolution spectra (R = 5000-7000) over the wavelength range 0.3-2.5 $\mu$m of ten M-type subdwarfs with X-SHOOTER at VLT. These data constitute a unique atlas of M-subdwarfs from optical to near-infrared. We performed a spectral synthesis analysis using a full grid of synthetic spectra computed from BT-Settl models and obtained consistent stellar parameters such as effective temperature, surface gravity, and metallicity.} {We show that state-of the-art atmospheric models correctly represent the overall shape of their spectral energy distribution, as well as atomic and molecular line profiles both in the optical and near-infrared. We find that the actual fitted gravities of almost all our sample are consistent with old objects, except for LHS 73 where it is found to be surprisingly low. } {} | M-type stars are the most common stars in our Galaxy \citep[70$\%$ of the Galactic stellar population;][]{Bochanski2010}. They are an important probe for our Galaxy as they carry fundamental information regarding the composition history, a record of galactic structure and formation, and of its dynamics. Although they are not expected to be as numerous as compared to M-dwarfs \citep[0.25$\%$ of the Galactic stellar population;][] {Reid2005b}, the metal-poor M-type subdwarfs can serve as probes for the old populations in our Galaxy (old disc, thick disc and halo). The locus of M-subdwarfs (sdMs) in the Hertzsprung-Russel (HR) diagram deviates from most of the field stars owing to the metallicity difference that translates into different opacities from those of regular M-dwarfs. Since some subdwarfs lie close to the hydrogen burning limit, they can be used to probe the lower end of the stellar mass function, extending it towards the brown dwarf regime. Because of their intrinsic faintness and the difficulty in getting a homogeneous sample of unique age and metallicity, very little is known about them. As we go from earlier to later M-subtypes, more molecules form in their atmospheres, making the spectral continuum very hard or impossible to identify both in the optical and in the near-infrared (NIR). Furthermore they provide the right condition for studying molecules and dust formation in the low-metallicity environment as well as radiative transfer in cool, metal-poor atmospheres. With decreasing temperature, sdMs spectra show an increase in abundances of diatomic and triatomic molecules in the optical and in the near-infrared (e.g. SiH, CaH, CaOH, TiO, VO, CrH, FeH, OH, H$_2$O, CO). The molecules TiO and VO dominate the opacity sources in the optical and H$_2$O and CO at lower spectral resolution in the infrared, having more complex and extensive band structures which leave no window for the true continuum and create a pseudo-continuum that at low spectral resolution only shows the strongest, often resonant atomic lines \citep{Allard1990,Allard1995}. However, owing to their low metallicity, and hence low Ti and O abundances \citep{Savcheva2014}, the TiO bands are not as strong, and the pseudo-continuum is brighter as a result. This increases the contrast with the other opacities, such as hydride bands and atomic lines that feel the higher pressures of the deeper layers from where they emerge \citep{Allard1990,Allard1995}. Different classification schemes have been proposed to assign metallicity and spectral types of sdMs. \cite{Ryan1991a} and \cite{Ryan1991b} used metallic lines, such as the CaII K line, to determine the metallicity of subdwarfs. \cite{Baraffe1997} published the first evolution tracks and isochrones based of the NextGen model atmosphere \citep{Allard1997,Allard2000,Hauschildt1999} spanning the complete range of composition of (s)dMs stars. \cite{Gizis1997} proposed a first classification of sdMs and extreme subdwarfs (esdMs) based on the NextGen isochrone luminosity and TiO and CaH band strengths in low-resolution optical spectra. \cite{Lepine2007} has revised the adopted classification and proposed a new classification for the most metal poor, the ultra subdwarfs (usdMs). \cite{Jao2008} has compared his model grids with the optical spectra to characterise the spectral energy distribution of subdwarfs using three parameters-- temperature, gravity and metallicity-- and thus gave an alternative classification scheme of subdwarfs. The proper classification of these objects requires the grid of synthetic spectra to be compared with the observations, which then help to quantify their basic properties: elemental abundances, effective temperature, and surface gravity. At present, these physical properties are not yet particularly well determined for sdMs. Traditional techniques to estimate stellar effective temperature based on black-body approximations and broadband photometry are at best dangerous for cool M-dwarfs whose true continua are masked by extensive molecular absorption. Furthermore, the complexity of the stellar atmosphere increases significantly with decreasing effective temperature as dust cloud formation occurs \citep{Tsuji1996a,Tsuji1996b,Allard1998}. This is revealed by the weakening of condensible bearing opacities such as TiO, VO, CaH,and CaOH bands in the optical wavelengths by dust Rayleigh scattering, and a reddening of the infrared spectral energy distribution with weakening water bands due to dust backwarming or the greenhouse effect \citep{Allard2001}. Over the last decades tremendous development in the model atmospheres of cool low-mass stars, in particular M-dwarfs, has been achieved \citep{Brott2005,Helling2008a,Allard2012,Allard2013} that has boosted the number of studies deriving accurate physical parameters of these stars both in the optical and in the near-infrared \citep{Bayo2011,Bayo2012,Rajpurohit2012a,Rajpurohit2013,Neves2014}. \cite{Bayo2014} have shown the differences between estimating the parameters in the optical and in the near-infrared, with spectra and photometry. Thanks to the large improvement of atomic and molecular line opacities which dominate the optical and infrared spectral range of these objects and to the revision of the solar abundances by \cite{Asplund2009} and \cite{Caffau2011}, synthetic spectra such as the new BT-Settl \citep{Allard2013} has achieved major improvements in modelling these complex systems. These models now even include dust cloud formation which becomes important for cool M-dwarfs and subdwarfs and yield promising results in explaining the stellar-substellar transition \citep{Allard2013,Baraffe2015}. Metallicity effects on the physics of cool atmospheres of very low-mass stars have been studied theoretically \citep{Allard1990,Allard1997} and then compared with observations \citep{Leggett1996,Leggett1998,Leggett2000,Leggett2001,Burgasser2002}. Such comparisons have revealed possible inaccuracies and/or incompleteness of the opacities used in the model at the time. The spectra of M-subdwarfs show the strengthening of hydride bands (OH, FeH) and pressure induced absorptions by H$_2$ around 2 $\mu$m relative to double metal bands (TiO,VO), and the broadening of atomic lines due to the effect of gravity confirming the predictions of \cite{Allard1990} and \cite{Allard1995}. Therefore we see these molecular bands in more detail in M-subdwarfs than in M-dwarfs and under more extreme gas pressure conditions. This can help reveal the remaining inaccuracies and/or incompleteness of the opacities used in the model. In this paper, we present a homogeneous sample of M-subdwarf X-SHOOTER spectra covering the full temperature sequence from the optical to the near-infrared, and perform a spectral synthesis analysis using a full grid of BT-Settl synthetic spectra\footnote{https://phoenix.ens-lyon.fr/Grids/BT-Settl/CIFIST2011bc}. This is necessary to determine how well the models reproduce the overall observed spectral properties. We therefore determine the atmospheric parameters from the overall spectral energy distribution (SED), i.e. the simultaneous optical and near-infrared measurements, so that variability does not include further uncertainty in our estimations. We assemble ten optical to near-infrared (0.3-2.5 \,$\mu$m) spectra of sdMs, as described in \S~\ref{obs}. Section 3 present the most recent atmosphere models that we used for comparison using a method described in \S~\ref{comp}. Section 5 gives the results of the comparison and the conclusion is given in \S~\ref{ccl}. | \label{ccl} In this paper, we obtained 0.3-2.5 $\mu$m simultaneous X-SHOOTER spectra for ten subdwarfs including extreme and ultra-subdwarfs, and described the results of their comparison to the BT-Settl synthetic spectra. We found that the parameters of an observed star are best determined by using a combination of both low-resolution and high-resolution spectra in order to disentangle the atmospheric parameters $\teff$, log\,$g$, and [Fe/H] which can have compensating effects. Models at the low metallicity used in this analysis have not been extensively tested before over full wavelength coverage for M-subdwarfs. This model grid constitutes a significant improvement over the earlier models by \cite{Allard1990,Allard1998} and \cite{Allard2000} owing to more complete and accurate line lists for TiO and H$_2$O, to the current computing possibility to include all the lines in the radiative transfer, and to the revised solar abundances. The synthetic spectra demonstrate the influence of the metallicity and gravity. We found that gravity has a relatively small influence on the spectra and the overall energy distribution, as was also found with earlier models \citep{Leggett1998,Leggett2000}, but it has a significant effect on high-resolution line profiles and details of the band systems. However, the effects of gravity become stronger with lower effective temperatures. The metallicity has, on the other hand, a large effect on the spectra. We determined stellar parameters and found a good agreement with those derived from higher spectral resolution observations in the optical wavelength \citep{Rajpurohit2014}. The difference between the parameters determined from UVES spectra and from X-SHOOTER spectra agree within the error bars defined by the grid spacing. The disentangling between the effects of temperature, metallicity and gravity is possible because the models show a large variation in their spectral features due to either molecular or atomic opacities, and the parameters are obtained with a high confidence level despite the existing grid spacing. Features formed by the CIA bands are reproduced well enough through the NIR continuum. The overall agreement of models with the observations is much improved when compared to the earlier work. \cite{Gizis1997} obtained values for $\teff$ and [Fe/H] of sdMs and esdMs by comparing optical spectra in the region from 6200-7400 $\AA$ (resolution = 3-4 $\AA$). \cite{Gizis1997} used the extended model grid \citep{Allard1995} which assumed local thermodynamic equilibrium (LTE) for their comparison with observed spectra. \cite{Leggett1996,Leggett1998,Leggett2000} compared spectra in the range from 0.6 - 2.5 $\mu$m with the NextGen \citep{Hauschildt1999} model grid. The BT-Settl models used in this paper also assume LTE and resolve the historical discrepancy in the stellar parameters obtained from the infrared and optical spectral regions also seen by \cite{Viti1997}. Even though the BT-Settl collection does a better job in reproducing simultaneously the optical and near-infrared features of these cool metal-poor sources, there is still room for improvement since there are regions where the fit is not optimal, in particular below 0.45 $\mu$m and in the H and K bands. This can be due to missing CaOH bands in the V bandpass in the models. Also, a complete FeH line list is currently missing in the H bandpass. An accurate and complete TiO line list is currently being developed by the ExoMol group. Upgrading these opacities is the next step in improving these models in the near future, before computing detailed model atmosphere grids and interior and evolution models at finer steps in the atmospheric parameters. We also note that at low metallicity, the effects of temperature inhomogeneities in the atmosphere begin to have greater impact on the spectrum formation, which can only be accurately modelled with 3D RHD simulations and ultimately 3D radiative transfer. | 16 | 9 | 1609.07062 |
1609 | 1609.00197_arXiv.txt | {Theoretically, rotation-induced chemical mixing in massive stars has far reaching evolutionary consequences, affecting the sequence of morphological phases, lifetimes, nucleosynthesis, and supernova characteristics. } {Using a sample of 72 presumably single O-type giants to supergiants observed in the context of the VLT-FLAMES Tarantula Survey (VFTS), we aim to investigate rotational mixing in evolved core-hydrogen burning stars initially more massive than 15\,\msun\ by analysing their surface nitrogen abundances.} {Using stellar and wind properties derived in a previous VFTS study we computed synthetic spectra for a set of up to 21 N\,{\sc ii-v} lines in the optical spectral range, using the non-LTE atmosphere code {\sc fastwind}. We constrained the nitrogen abundance by fitting the equivalent widths of relatively strong lines that are sensitive to changes in the abundance of this element. Given the quality of the data, we constrained the nitrogen abundance in 38 cases; for 34 stars only upper limits could be derived, which includes almost all stars rotating at $\vrot\ >200\,\kms$. } {We analysed the nitrogen abundance as a function of projected rotation rate \vrot\ and confronted it with predictions of rotational mixing. We found a group of N-enhanced slowly-spinning stars that is not in accordance with predictions of rotational mixing in single stars. Among O-type stars with (rotation-corrected) gravities less than $\log\,g_c = 3.75$ this group constitutes 30$-$40 percent of the population. We found a correlation between nitrogen and helium abundance which is consistent with expectations, suggesting that, whatever the mechanism that brings N to the surface, it displays CNO-processed material. For the rapidly-spinning O-type stars we can only provide upper limits on the nitrogen abundance, which are not in violation with theoretical expectations. Hence, the data cannot be used to test the physics of rotation induced mixing in the regime of high spin rates.} {While the surface abundances of 60-70 percent of presumed single O-type giants to supergiants behave in conformity with expectations, at least 30-40 percent of our sample can not be understood in the current framework of rotational mixing for single stars. Even though we have excluded stars showing radial velocity variations, of our sample may have remained contaminated by post-interaction binary products. Hence, it is plausible that effects of binary interaction need to be considered to understand their surface properties. Alternatively, or in conjunction, the effects of magnetic fields or alternative mass-loss recipes may need to be invoked.} | \label{sec:intro} Despite the importance of massive stars for Galactic and extragalactic astrophysics, many of the physical processes that control the evolution of these objects are still not well understood \citep[see e.g.,][]{langer2012}. As one of the key agents of massive star evolution, the effects of rotation are manifold. The internal structure of spinning stars becomes latitude dependent \citep{zeipel1924} and centrifugal forces resulting from rotation may lead to deviation % from a spherical shape \citep[e.g.,][]{collins1963,townsend2004}. The von Zeipel effect results in such stars showing relatively hot and bright polar regions and relatively cool and dim equatorial zones -- effects that are actually observed \citep{domicianodesouza2003,domicianodesouza2005}. Spinning stars have longer main-sequence lifetimes \citep{brott2011a,ekstrom2012,kohler2015} and may follow different paths in the Hertzsprung-Russell diagram (HRD), as the centrifugal force % reduces the effective gravity and because rotation may also impact mass-loss and angular momentum loss at the surface \citep[see][for an extensive discussion]{maeder2009,langer2012}. Rotation induced instabilities trigger internal mixing, transporting material from deep layers to the surface. In extreme cases, this mixing may be so efficient that the stars remain chemically homogeneous throughout their lives and -- because of this -- avoid envelope expansion \citep{maeder1987}. In special cases, this type of evolution has been proposed to lead to long-duration gamma ray bursts \citep[e.g.,][]{yoon2005,woosley2006}. Chemical homogeneity is further a pivotal ingredient for certain formation channels of close-binary black holes, that may over time merge and emit a gravitational wave signal \citep{demink2009,mandel2016,marchant2016}. Helium may, in principle, serve as a tracer of the efficiency of rotationally induced mixing processes. However, as it is produced on the nuclear timescale, the anticipated surface enrichment of He % is relatively minor. In this paper we focus on the surface abundance of nitrogen, which is a much more sensitive agent. In the CN (and CNO) cycle, carbon (and oxygen) are converted into nitrogen. Processed material that can escape from the core before a chemical gradient is established at the core boundary, can be mixed into the envelope \citep{maedermeynet1997}. This material will, over time, reach the stellar surface. The faster the star is spinning, the more quickly the material will surface and a greater surface N abundance will be reached. \citet{hunter2008} searched for a correlation between surface nitrogen abundance and projected spin velocity of a sample of evolved main-sequence early B-type stars, observed in the context of the VLT-FLAMES Survey of Massive Stars \citep{evans2006}. A subset of their sample displayed such a correlation and was used by \citet{brott2011b} to calibrate the efficiency of rotational mixing at the metallicity of the Large Magellanic Cloud (LMC). However, two groups of stars did not concord with the expectations of rotational mixing. We discuss this further in Sect.~\ref{sec:otherNstudies}, comparing their findings for B stars with ours for O stars. Here, we determine the nitrogen abundance of LMC O-type giants and supergiants that have been observed in the context of the VLT-FLAMES Tarantula Survey \citep[VFTS;][]{evans2011}. The main questions we want to address are: what is the behaviour of these O-type stars in terms of N-abundance as a function of projected spin velocity? Does our sample also contain groups of stars that are not in agreement with the predictions of rotational mixing in single stars? If so, can these be identified as the counterparts of the peculiar groups identified among the B dwarfs by \citet{hunter2008}? The paper is organised as follows. In Sect.~\ref{sec:sample} we briefly introduce the sample of O giants and supergiants. The method of nitrogen abundance analysis is explained in Sect.~\ref{sec:method}. The results are presented in Sect.~\ref{sec:results}, compared to population synthesis predictions in Sect.~\ref{sec:popsyn}, and discussed in Sect.~\ref{sec:discussion}. Finally, we summarise our findings in Sect.~\ref{sec:conclusions}. | \label{sec:discussion} In this section we discuss the nature of the stars in Box\,2 that do not concur with predictions of rotational mixing in massive stars. We also compare our finding to previous studies. We start by examining the current mass of the stars in our sample. \subsection{Current evolutionary masses} \label{sec:massdependence} Current evolutionary masses of the stars in our sample have been determined by Ram\'{i}rez-Agudelo et al. (subm.) using {\sc Bonnsai}\footnote{The {\sc Bonnsai} web-service is available at \newline \url{www.astro.uni-bonn.de/stars/bonnsai}.}, a Bayesian method to constrain the evolutionary state of stars \citep{schneider2014}. As independent prior functions these authors adopt a \citet{salpeter1955} initial mass function, an initial rotational velocity distribution as given by \citet{ramirezagudelo}, a random orientiation of spin axes, and a uniform age distribution. The observables on which the mass estimates are based are $L$, \Teff\ and \vrot. Figure~\ref{fig:massplot} presents a summary of the mass properties of the low gravity stars in Fig.~\ref{fig:hunterplotgroups}, comparing the stars in Box\,2 to the rest of the sample. Two types of statistics are provided, one based on median statistics (Tuley's five-number summary) and one based on the mean of the sub-samples. Each of these shows that the objects in Box\,2 are appreciably more massive than the remainder of the stars: the median evolutionary mass of stars in Box\,2 is 34.8\,\msun, while that of the remaining stars is 25.8\,\msun. \citet{brott2011b} performed a similar analysis as done here for a sample of 107 main-sequence B-type stars in the LMC with projected spin rates up to $\sim$300 km\,s$^{-1}$ discussed by \citet{hunter2008,hunter2009}. They too note that the nitrogen enhanced slowly spinning stars appear to have higher masses. Though the bulk of their sample has a mass $\leq 12$\,\msun\, the stars that populate Box\,2 almost all have masses in between 12$-$20\,\msun. Further on in this section we discuss scenarios regarding the nature of the Box\,2 stars that may explain this behaviour. \begin{figure}[t!] \begin{center} \resizebox{0.8\hsize}{!}{\includegraphics{\figpath massplot.png}} \end{center} \caption{Distribution of evolutionary masses for the stars in Fig.~\ref{fig:hunterplotgroups} that occupy Box~2 (left) and the remainder of stars (right). Black crosses indicate data points. The middle column in each half of the diagram shows the median of the sample in red, where the boxplot indicates the 0.25 and 0.75 percentiles with whiskers extending out no further than 1.5 times the inter-quartile range. The right column indicates the weighted average and standard deviation (in parentheses) of the sample in blue. Evolutionary masses are calculated using the {\sc Bonnsai} tool (Sect.~\ref{sec:massdependence}). } \label{fig:massplot} \end{figure} \subsection{Current surface nitrogen abundances as predicted by single star evolutionary tracks} \label{sec:bonnsaicomp} In addition to the posterior probability distribution of the current mass, {\sc Bonnsai} also provides the distribution for the current surface nitrogen abundance. For all but one of the stars for which our analysis resulted in upper limits on \nabun\ only, the upper limit is in agreement with the \nabun\ predicted by {\sc Bonnsai}. Out of the 38 stars for which we have a direct \nabun\ measurement, only 7 agree within 1$\sigma$ with the nitrogen abundance expected from their evolutionary state. Interestingly, among these are three binaries (namely VFTS\,333, 399 and 440). In the other 31 cases our nitrogen measurement is higher than that predicted by {\sc Bonnsai}. Since all direct measurements have low \vrot, this is an alternative way of expressing that the bulk of our nitrogen rich, slowly rotating sources are not expected from the evolutionary sequences by \citet{brott2011a}. We repeated the {\sc Bonnsai} analysis for the 9 sources in Fig.~\ref{fig:hunterplotgroups} that have measured abundances and occupy Box~2. In this run, we added the measured surface abundances of helium and nitrogen to the observables \Teff, $L$, and \vrot. While in the analysis without these extra two observables {\sc Bonnsai} returned in all cases a modestly rotating star with most probable inclination $\sin\,i \gtrsim 0.8$, the new run produced fast rotators seen almost pole-on ($\sin i \lesssim 0.5$) in 8 out of 9 cases. This reflects that in order to reproduce the high He- and N-abundances in the tracks of \citet{brott2011a} more mixing is required, demanding more rapid rotation. Assuming inclinations randomly distributed in 3D space, the probability that out of the 27 stars with $\logg_c \leq 3.75$, 8 have $\sin i \leq 0.5$ is 0.015.\footnote{Assuming randomly distributed spin axes in space, the probability for an individual system having $\sin\,i \leq x$ is given by $P(x)=1-\sqrt{1-x^2}$. As such, the probability for a population of $n$ stars containing $k$ stars with $\sin\,i\leq x$ is given by $\frac{n!}{k! (n-k)!} P(x)^k \left(1-P(x)\right)^{n-k}$.} This too demonstrates that the stars in Box\,2 cannot be straightforwardly understood in the context of current single-star evolutionary tracks. \subsection{Correlation between nitrogen and helium abundance} \label{sec:NvsHe} In Fig.~\ref{fig:hunterplotNvsHe}, we plot the nitrogen surface abundance derived here versus the helium surface fraction (by mass) as derived by Ram\'{i}rez-Agudelo et al. (subm.). Shown in the background is the same population synthesis calculation as discussed in Sect.~\ref{sec:popsyn_lowg} and shown in Fig.~\ref{fig:hunterplotgroups}, i.e., for O-type sources that have $\log g \leq 3.75$. The results are grouped in bins of $0.01\,(Y) \times 0.05\,(\nabun)$ dex. Simulated stars that are helium- as well as nitrogen enriched are all rotating initially at $\gtrsim 350\,\kms$. Though the scatter is sizable, the N- and He-abundances of the stars in our sample appear correlated. Such a correlation is also found by \citet{riverogonzalez2012a}. Moreover, the trend agrees quite well with the models, even though the number of predicted and observed stars with He, as well as N, enrichment do not agree. This suggests that though the mechanism that brings N and He to the surface may still be unidentified, it is one that brings a mixture of gases to (or deposits a mixture of gases on) the surface that has the He/N abundance ratio that is expected from the CNO process. % If the exposed material is CNO-processed, an enhancement in N is expected to be accompanied by a depletion of C. An independent verification would thus be to simultaneously measure the surface C abundances of these stars. It should also be noted that several other surface-enrichment mechanisms (e.g., those discussed in Sect.~\ref{sec:alternativescenarios}) are expected to result in the display of CNO-processed material at the surface. \begin{figure}[t!] \resizebox{1.05\hsize}{!}{\includegraphics{\figpath propertyplot-XaxisModeY.pdf}} \caption{Nitrogen abundance versus surface helium mass fraction. In the background the outcome of population synthesis is projected, as explained in Sect.~\ref{sec:popsyn} and shown in Fig.~\ref{fig:hunterplotgroups}} \label{fig:hunterplotNvsHe} \end{figure} \subsection{Comparison to other \nabun-\vrot\ studies} \label{sec:otherNstudies} We limit a comparison to other studies to main sequence OB stars in the LMC, notably to the work on early-B stars by \citet{hunter2008,hunter2009} and \citet{brott2011b}, and on O-type stars by \citet{riverogonzalez2012a}. For an analysis of LMC B-type supergiants see \citet{hunter2008} and \citet{mcevoy2015}; of Galactic and SMC B-type stars see \citet{hunter2009}; for that of Galactic O-type stars, see \citet{martins2015,martins2015b,bouret2012}; of SMC O-type dwarfs, see \citet{bouret2013}. \citet{riverogonzalez2012a} determined the nitrogen abundance of the LMC stars studied by \citet{mokiem2007} in the context of the VLT-FLAMES Survey of Massive Stars \citep{evans2006}. This sample of O and early-B stars of all luminosity classes comprises sources in the central LH9 and LH10 associations in the young star-forming region N11, augmented by LMC field stars. Though the authors refrain from a statistical analysis due to the modest sample size, they do compare with theoretical predictions. They find an even larger fraction of stars with N-enrichments that appear too large for their (current) rotation rate. In their work, the correlation between He- and N-abundances (their Fig.~8, our Fig.~\ref{fig:hunterplotNvsHe}) also appears present. It is interesting that the eight O-type objects in their sample that have $\log g_c \leq 3.75$ span essentially the same range in nitrogen abundance found here and all would appear in Box~2. The mean \vrot\ (and associated standard deviation) of this sample of eight is $59 \pm 20$\,\kms, which is lower than the 98 $\pm$ 12\,\kms\ of our sample in Box~2. The main reason for this is that \citeauthor{riverogonzalez2012a}\ accounted for macro-turbulent broadening of their lines, which in the \vrot\ estimates used here was ignored. In the analysis of the full presumed-single VFTS O star sample, \citet{ramirezagudelo} employed similar techniques to correct for the effect of macro-turbulent broadening. Using these results, the mean projected spin velocity of our stars in Box~2 would be 78 $\pm$ 19\,\kms, i.e., within uncertainties consistent with the findings of \citet{riverogonzalez2012a}. We conclude that the population of N-enriched slow rotators presented in \citet{riverogonzalez2012a} is similar in characteristics to the Box~2 population found here. Its presence (in the sample studied by \citeauthor{riverogonzalez2012a}) supports our statement that the N-enriched slowly spinning sources are not in agreement with predictions of rotational mixing in single stars. \citet{hunter2008}, in their analysis of over 100 B-type main-sequence stars in the VLT-FLAMES Survey of Massive Stars, identified a population of N-enriched intrinsically slow rotators ($\sim$20\% of their sample) and a population of relatively unenriched fast rotators (a further $\sim$20\% of their sample) that both challenge the concept of rotational mixing. The stars in the group of slowly rotating stars preferentially have masses in the range of 12\,\msun\ to 20\,\msun\ and reach nitrogen abundances \nabun\ $\sim$ 8, so somewhat lower than the $\sim$8.7 that is reached by the O stars. % The definition of the box of N-enriched low \vrot\ sources in \citeauthor{hunter2008} is somewhat more stringent than in our case, in that all have a projected spin velocity $\vrot < 50$\,\kms. \citet{brott2011b} compared the incidence of these stars to the same models being tested here and found that this region of the \nabun\ versus \vrot\ diagram contained 15 times as many stars as predicted. Whether the fact that the nitrogen-enriched B dwarf population spins somewhat slower than the nitrogen-enriched slowly rotating O-type stars and reach somewhat lower N-enrichments points to a fundamental difference in the nature of this group remains to be investigated. However, given that the global similarities between the two groups are substantial a common nature seems plausible. For the group of relatively unenriched fast rotators identified by \citet{hunter2008} mostly upper limits to the nitrogen abundance could be constrained. By far the bulk of this population seems compatible with an abundance that is consistent with the LMC baseline value. The upper limits of our O stars sample are less constraining and it is therefore too early to tell whether the fast rotators in both samples behave alike. \subsection{Alternative scenarios} \label{sec:alternativescenarios} In this subsection we discuss scenarios that may yield slowly rotating, nitrogen enriched sources (i.e. Box\,2). It may well be possible that several processes are working in parallel, such that sources may occupy Box~2 for independent, or a combination of, reasons. \subsubsection{Efficiency of rotational mixing} The efficiency of rotational mixing in the single-star evolutionary models of \citet{brott2011a} is calibrated using efficiency factors. Rotationally induced instabilities contribute in full to the turbulent viscosity \citep{heger2000a}, while their contribution to the total diffusion coefficient (the way in which all mixing processes are treated in the models of \citeauthor{brott2011a}) is reduced by a factor $f_{\rm c}$. \citet{brott2011a} calibrated this parameter using the B dwarfs discussed in Sect.~\ref{sec:otherNstudies}. The inhibiting effect of chemical gradients on the efficiency of rotational mixing processes is regulated by the efficiency parameter $f_{\mu} = 0.1$, calibrated following \cite{yoon2006}. The steepness of the correlations shown in Figs.~\ref{fig:hunterplot} and~\ref{fig:hunterplotgroups} is essentially the result of the $f_{\rm c}$ calibration for the quoted $f_{\mu}$. This calibration ignored the nitrogen enriched slow rotators (the Box~2 sources) and the relatively unenriched rapid rotators. Assuming the sources in Box~2 reflect a much larger mixing efficiency, e.g. a much increased $f_{\rm c}$,% an explanation would be required as to why two vastly different regimes of mixing efficiencies would occur in massive stars (i.e., one for Box~2 and one for Box~3 stars). \subsubsection{Envelope stripping through stellar winds}% \label{sec:windstripping} The finding that the mass of the Box\,2 stars is on average higher (Sect.~\ref{sec:massdependence}) could mean that these may have endured a more severe stripping of their envelopes -- revealing chemically enriched layers -- as stars of higher mass (or luminosity) are in general expected to have a stronger mass loss. Mass loss can also drain angular momentum from the surface, resulting in spin-down. This scenario would be in line with the correlation pattern of the nitrogen and helium abundance as discussed in Sect.~\ref{sec:NvsHe}. \citet{bestenlehner2015} investigated the stars on the upper main sequence in the VFTS sample. The authors find a clear correlation between helium mass fraction $Y$ and $\log\,M/\dot{M}$ at $\log\,M/\dot{M} \gtrsim -6.5$ (their Fig.~14), consistent with the idea that severe mass loss in very massive stars indeed reveals the deeper layers. Ram\'{i}rez-Agudelo et al. (subm.) added the sources studied here to this diagram, populating the regime at $\log\,M/\dot{M} \lesssim -6.5$. They find that, save for a few outliers, a monotonic trend of $Y$ versus $\log\,M/\dot{M}$ is missing in this regime\footnote{For completeness we note that one of our sources, VFTS\,180, populates the regime studied by \citet{bestenlehner2015}. This source does seem to follow the trend reported by these authors.} (their Fig.~7). The latter is also expected from the models of \citet{brott2011a}. For example, Fig.~4 of \citeauthor{brott2011a} shows the expected surface enrichment at the TAMS, as a function of initial mass and rotation rate. It can be seen that, at LMC metallicity and typical rotation rates, significant surface enrichment is expected only for stars in excess of $60\,\msun$. Note also that these are values as predicted for the TAMS, while most of our observed stars are still undergoing main-sequence evolution. Furthermore, even though the average mass of Box~2 sources compared to the remaining stars is higher, most stars have evolutionary masses on the order of $\sim30\,\msun$. The implicit assumption in the above discussion is that the mass-loss rates that are adopted in the evolutionary tracks being scrutinised here are correct. The $\dot{M}$ recipe used in these models are from \citet{vink2000,vink2001}. These are consistent with the assumption of moderately clumped winds (\citeauthor{mokiem2007} \citeyear{mokiem2007}; Ram\'{i}rez-Agudelo et al. subm.). Hence significantly higher mass-loss rates would require the outflows to be close to homogeneous, which is not expected. In the context of envelope stripping it would be hard to reconcile a mass dependence of Box\,2 versus non-Box\,2 sources in both the B stars sample of \citet{hunter2008} and our O star sample, the Box\,2 B-type stars having a lower average mass than the non-Box\,2 O-type stars. Moreover, at typical rotation rates wind stripping is thought to be important only for the most massive stars (as explained above), such that the anomalous B stars of \citeauthor{hunter2008} are unlikely to have resulted from this process. We conclude that, although for some stars, envelope stripping could be an important aspect, for the bulk of the Box~2 objects it is unlikely to be the cause of the observed N-enrichment. \subsubsection{Binary evolution} \label{sec:binaryevolution} Binary interaction in general might yield stars of higher mass (as the mass gainer may dominate % the light of the system, or when two stars yield a more massive single star after a merger event), in accordance with the observations described in Sect.~\ref{sec:massdependence}. It is not improbable that our sample of presumed single stars is polluted with binary products. \citet{demink2014} showed that using radial velocities variations to select a sample against binaries may inadvertently {\em increase} the fraction of binary interaction products. This is a result of two considerations. First, merger products are essentially single stars. Second, in a typical post mass transfer system the light is dominated by the mass gainer, which in many cases displays only modest radial velocity variations (less than 20\,\kms). Such a population of post-interaction binaries can explain the presence of a high-velocity tail ($\vrot > 300\,\kms$) found in the spin distribution of the VFTS single O star sample by \citet{ramirezagudelo}, which was proposed by \citet{demink2013}. Interestingly, this tail appears absent in the spin distribution of relatively wide pre-interacting binaries, i.e., sources that are not affected by tidal effects \citep{ramirezagudelo2015}, supporting the idea that the high-velocity tail of the spin distribution is the result of binary interaction. So indeed, a sizable fraction of our presumed single-star sample may in fact be a product of binary evolution. Several scenarios involving binary interaction might explain the stars in Box~2 (and stars in Box~3 if their nitrogen abundances turn out to be too low to concord with rotational mixing). \citet{glebbeek2013}, for instance, investigate the evolution of merger remnants. They find that mixing of the envelopes, for example through thermohaline circulation, can result in enhanced nitrogen abundances. These enhancements are stronger for more massive stars. The possibility of Box\,2 stars being binary products has also been brought forward by \citet{brott2011b} to explain the B-dwarfs analysed by \citet{hunter2008}, to which we refer for a more extended discussion. Note, however, that binary interaction through mass transfer is more likely to produce fast rotators, be they enriched or not. Channels reproducing the slow rotators in Box~2, for example through spin-down by resonant locking, are relatively less likely with respect to channels reproducing rapid rotators. A quantitative test of the binary hypothesis requires binary population synthesis models of rotationally mixed stars, which are not yet available \citep[but see][for a discussion of first steps in this direction]{langer2012}. An empirical approach to probing a post-interaction nature of the stars in Box~2 is to compare the nitrogen abundances of the sample studied here to those of the binaries in the VFTS. It should be noted that if binarity is at play, one must consider the contribution of a possibly undetected companion to the total observed continuum light. This contribution may dilute the strength of the nitrogen lines, implying the derived N-abundances reported here would formally represent lower limits in those cases. \subsubsection{Magnetic fields} Magnetic fields have been suggested to play a role in explaining the N-enhanced, slowly rotating early-B dwarfs that challenge rotational mixing \citep{morel2008,morel2012,przybilla2011}. Such a field may be either of fossil origin or due to a rotationally driven dynamo operating in the radiative zone of the star \citep{spruit1999,maeder2004}. \citet{meynet2011} proposed that magnetic braking \citep[see][]{udDoula2002,udDoula2008,udDoula2009} during the main-sequence phase leads to slow rotation. In their model, \citeauthor{meynet2011} assume the absence of a magnetic field in the stellar interior. Consequently, they find that the slowing down of the surface layers results in a strong differential rotation, resulting in a nitrogen surface enhancement. A situation where the observable magnetosphere is the external part of a strong toroidal-poloidal field inside the star \citep{braithwaite2004} might induce considerably less mixing \citep[see also][]{morel2012}. \citet{potter2012} discussed the $\alpha-\Omega$ dynamo as a mechanism for driving the generation of large-scale magnetic flux and found a strong mass dependence for the dynamo-driven field. For stars with initial masses greater than about 15 \msun\ (i.e., those born as O stars), they find that the dynamo cannot be sustained. Initially lower mass stars (i.e., those born as B or later type stars) with sufficiently high rotation rates are found to develop an active dynamo and so exhibit strong magnetic fields. They are spun down quickly by magnetic braking and magneto-rotational turbulence \citep{spruit2002} causing changes to the surface composition. Our finding that also the O stars \citep[in addition to the B stars; see][]{hunter2008} populate Box\,2 might thus indicate that if magnetic fields play a role in explaining the nature of these N-enriched slowly spinning stars their fields are of fossil origin. If not, two different dynamo models may need to be invoked to explain the Box\,2 objects, one for O-type stars and another for B-type stars. Alternatively, magnetic fields may play a role in conjunction with binary evolution. If the mass gainer is spun up by angular momentum transfer, efficient rotational mixing leads to the surfacing of nitrogen. It may generate magnetic fields that, after the surface is enhanced in nitrogen, spin down the star as a result of magnetic braking. As an extreme example of this mechanism one may envision the actual merger of the two stars. Such merger products may represent of the order of 10\% of the O star field population \citep{demink2014}, a similar percentage as the incidence rate of magnetic O-type stars \citep{grunhut2012,wade2014,fossati2015}. This is, however, less than the percentage of stars found in Box~2. | 16 | 9 | 1609.00197 |
1609 | 1609.05856_arXiv.txt | It has been suggested that the internal dynamics of dwarf spheroidal galaxies (dSphs) can be used to test whether or not ultralight axions with $m_a\sim 10^{-22}\text{ eV}$ are a preferred dark matter candidate. However, comparisons to theoretical predictions tend to be inconclusive for the simple reason that while most cosmological models consider only dark matter, one observes only baryons. Here we use realistic kinematic mock data catalogs of Milky Way dSph's to show that the ``mass-anisotropy degeneracy'' in the Jeans equations leads to biased bounds on the axion mass in galaxies with unknown dark matter halo profiles. In galaxies with multiple chemodynamical components this bias can be partly removed by modelling the mass enclosed within each subpopulation. However, analysis of the mock data reveals that the least-biased constraints on the axion mass result from fitting the luminosity-averaged velocity dispersion of the individual chemodynamical components directly. Applying our analysis to two dSph's with reported stellar subcomponents, Fornax and Sculptor, and assuming that the halo profile has not been acted on by baryons, yields core radii $r_{c}>1.5$ kpc and $r_c> 1.2$ kpc respectively, and $m_a<0.4\times 10^{-22}\text{ eV}$ at 97.5\% confidence. These bounds are in tension with the number of observed satellites derived from simple (but conservative) estimates of the subhalo mass function in Milky Way-like galaxies. We discuss how baryonic feedback might affect our results, and the impact of such a small axion mass on the growth of structures in the Universe. | axion dark matter is described by a classical scalar field, and differs from Cold Dark Matter (CDM, which is described by collisionless particles) on scales below the de Broglie wavelength due to the presence of gradient energy \citep[see][for a review]{2017PhRvD..95d3541H,2016PhR...643....1M,2014ASSP...38..107S}. For ultralight axions (ULAs) with $m_a/10^{-22}\text{ eV}\equiv m_{22}\approx 1$ this scale is large enough to be of relevance for the cusp-core problem in dSphs, as well as alleviating various other small scale issues with CDM~\citep{2000PhRvL..85.1158H,2014MNRAS.437.2652M,2014NatPh..10..496S,Matos:2000ss,Matos:2000ng,Sahni:1999qe,2000PhRvL..85.1158H}.\footnote{Note that a mass scale of order of $10^{-22} \text{eV}$ has for a long time been a recurring result in the studies of axion or scalar field models for galaxy halos and small scale structure, see for instance\citep{1990PhRvL..64.1084P,Sin:1992bg,Sahni:1999qe,Arbey:2001qi,Matos:2000ss,2000PhRvL..85.1158H}.} On non-linear scales, axion DM forms a class of pseudo-soliton known as an oscillaton, or ``axion star''~\citep{1969PhRv..187.1767R,Seidel:1991zh,UrenaLopez:2001tw,Guzman:2004wj}. The soliton is supported against gravitational collapse by gradient energy, and is expected to form in the centres of ULA halos. On larger scales, since structure formation proceeds just as for CDM, ULA halos should resemble the NFW profile \citep{1997ApJ...490..493N}. Indeed, the NFW profile is found from collisionless $N$-body simulations, which are operationally equivalent to the axion model on scales above the de Broglie wavelength~\citep[e.g.][]{Widrow&Kaiser1993,Uhlemann:2014npa}. High-resolution cosmological simulations and other numerical experiments~ \citep{2014NatPh..10..496S,2014PhRvL.113z1302S,2016arXiv160800802V,2016PhRvD..94d3513S} reveal just this: ULA/scalar field DM halos comprise a central soliton core transitioning to an NFW-like profile at large radii. The size of the core depends on the axion mass and local density, with larger cores occurring for smaller particle masses and lower densities. Standard CDM halos are well described by the NFW profile at all radii and display a central cusp. For almost a decade now, it has been suggested in that the cusp-core problem in dwarfs, as well as other ``small-scale crises'' \citep{2015PNAS..11212249W}, may be evidence for DM physics beyond CDM\citep[e.g.][]{2001ApJ...556...93B,2013PhRvD..87k5007T,2014MNRAS.437.2652M,2016MNRAS.460.1399V,2016PDU....12...56B}. It is not necessary that a DM model solve all of the apparent small-scale crises at once (a catch-all solution), but proposed solutions to any given problem must, of course, be consistent with cosmology and structure formation. The stellar dynamics of dwarf spheroidal (dSph) galaxies in the Milky Way (MW) can be used to study the distribution of DM in these systems \citep[see e.g.][for a review]{2013pss5.book.1039W}. dSphs are DM dominated at all radii, and so the stars can be seen as test particles orbiting in the DM halo. In particular, Fornax and Sculptor galaxies have two distinct stellar sub-populations of different metallicty. \citep{2011ApJ...742...20W} (henceforth, WP11) used the virial quantity $\langle \sigma^2_{\rm los}\rangle$ to measure the DM density profile slope, and showed a preference for cores ($ \rho\propto r^0$) over cusps ($\rho\propto r^{-1}$). Different particle physics models for DM predict different halo profiles; therefore the dSph measurements can be used to test the consistency of these models, or even to give evidence for one model over another \citep[e.g.][]{2007ApJ...657L...1S}. First attempts to use Stellar dynamics of dSphs to constrain axion and scalar field DM models are discussed in e.g. \citep{2014PhRvD..90d3517D,2015MNRAS.451.2479M,2016arXiv160609030C} In this work we address how stellar velocity measurements in dSph's can be used to place \emph{unbiased} constraints to the dark matter particle mass for an axion DM halo model. We investigate this using a series of $N$-body mocks for stars as test particles orbiting in static DM halos. We identify the now-familiar $\beta$-degeneracy, which introduces significant bias in the extraction of halo parameters using Jeans analysis when the stellar velocity anisotropy, $\beta$, is unknown. We then show how certain parameters can be extracted in an unbiased way using virial (integrated) quantities, where dependence on the anisotropy is reduced. Fig.~\ref{fig:mass-lik} represents our main findings concerning the axion mass and the MW dSphs. A joint Jeans analysis of the velocity dispersion profile of the eight classical MW dSphs \citep[using the data from]{2010ApJ...710..886W} selects a particular axion mass, $m_a=2.44 ^{+1.3}_{-0.6}\times 10^{-22}\text{ eV}$.\footnote{While the present work was in preparation, a similar Jeans analysis was performed by \citep{2016arXiv160609030C}, whose results are broadly consistent with ours. We comment on their analysis later.} However, our analysis of mocks leads us to conclude that the Jeans analysis has an unknown bias in the recovered axion mass, caused by the $\beta$-degeneracy. Notice that the Jeans analysis is also in some tension with the constraints of \citep{2015MNRAS.451.2479M} (hereafter MP15) based on the mass profile slope and virial mass estimator of WP11, which limits $m_a< 1.1 \times 10^{-22}\text{ eV}$ at a 95\% confidence level (C.L.). We revise this upper limit in a new analysis, proved to work extremely well in mock data, finding $m_a< 0.4 \times 10^{-22}\text{ eV}$ at 97.5\% C.L., using $\langle \sigma^2_{\rm los}\rangle$ from direct integration of Jeans equation, which we dub the $\langle\sigma_{\rm los}^2\rangle$-fit. In the rest of this paper we carefully examine the source of the discrepancy in these bounds on the ULA mass from dSphs, and argue that our revised bound is unbiased. We then discuss possible implications from a cosmological perspective. This paper is organized as follows. We begin in Section~\ref{sec:cusp_core} by reviewing the status of the cusp-core problem in dSphs. In Section~\ref{sec:model} we describe the ULA halo density profile, the model for stellar kinematics, and the set of synthetic observations we use to test our methodology. We perform N-body simulations of stars in the DM potential to generate mock data. In Section \ref{sec:results} we present the results of a Markov Chain Monte Carlo (MCMC) analysis over synthetic observations where we fit the parameters of our model using: ({\it i}) the full velocity dispersion profile i.e. Jeans analysis; ({\it ii}) the averaged velocity dispersion of two stellar populations using a mass-velocity dispersion estimator, as it was first proposed in \citep{2011ApJ...742...20W} and referred here as the slopes method; and ({\it iii}) a variation of the slopes method, the aforementioned $\langle\sigma_{\rm los}^2\rangle$-fit, where we propose to use an unbiased estimator based on a further integration of the Jeans equation. Then the same three analyses are performed on the real data. We shall see that $\langle\sigma_{\rm los}^2\rangle$-fit is the only one provinding unbiased estimates on the axion mass from the kinematics of dSph's. Section~\ref{sec:cosmology} discusses the cosmological (in)consistency of the axion cores, and compares our constraints to other studies. We conclude in Section~\ref{sec:conclusions}. The Appendix presents some additional details on Jeans analysis, and some discussion of constraints from cosmological reionization. \begin{figure} \vspace{-0.2em}\includegraphics[width=\columnwidth]{mass_lik_new.pdf} \vspace{-2.5em}\caption{Marginalized constraint on axion mass from dSph stellar dynamics using three methodologies. Using mock data, we demonstrate that only the $\langle\sigma_{\rm los} ^2\rangle$-fit returns unbiased results. The constraint from this method, $m_a<0.4\times 10^{-22}\text{ eV}$, produces too few subhalos and is inconsistent with a conservative bound of $m_a>1\times 10^{-22}\text{ eV}$ from cosmology. Jeans analysis returns the most biased results, due to the anisotropy degeneracy, while the ``virial estimator'' of WP11 used by MP15 (see text) has a slight bias to larger axion masses.} \label{fig:mass-lik} \end{figure} | \label{sec:conclusions} We used mocks of dSphs embedded in an axion DM halo to test for the presence of bias in constraints to the particle mass in axion DM models. The main points to conclude are: \begin{itemize} \item Using Jeans analysis with constant unknown anisotropy to fit the line of sight velocity dispersion leads to biased constraints on the axion mass, to the point where one can conclude that galaxies are well fitted by the axion halo model (cored) when in reality the underlying model is a "cuspy" one. \item We also found that using the M-estimator to fit the slope defined by the mean velocity dispersion from two different stellar populations in the same galaxy also leads to biased constraints, though to a lesser extent. As expected, the bias is worse when the axion mass is smaller since this case corresponds to very large cores. \item An intermediate approach where we compute the mean velocity dispersion from direct integration of the Jeans equation, Eq.~\eqref{eq:meanslos_2}, and fit the luminosity averaged velocity dispersion of two stellar subpopulations seems to provide unbiased constraints on the halo parameters. \end{itemize} Fitting $\langle\sigma_{\rm los}^2\rangle$ rather than using the M-estimator (used by MP15) in the joint analysis for data from Fornax and Sculptor galaxies taken from WP11 leads to a tighter limit to the axion mass, $m_a<0.4\times 10^{-22}\text{ eV}$. We have corroborated that our new approach works very well independently of the true anisotropy profile, by applying it to isotropic and anisotropic mock data. This is clearly in tension with a ``blind'' Jeans analysis, performed using a joint likelihood of the eight classical dSph's, with free constant anisotropy parameter $ \beta$, which gives $m_a=2.4^{+1.3}_{-0.6}\times 10^{-22}\text{ eV}$. However, our analysis of mocks leads us to believe that the estimate of $m_a$ from Jeans analysis suffers from significant bias, see Fig.~\ref{fig:mass-lik} for a comparison. Without proper knowledge of the dSph velocity anisotropy it is not possible to extract unbiased constraints on DM models. Here are two important points to remark. First, we are not attempting to extract more information from the other dSphs because we cannot be certain that those galaxies actually have a cored density profile. Second, we have verified that constraints set individually from Fornax and Sculptor are compatible, otherwise stating a joint likelihood would not be correct. From the $\langle\sigma_{\rm los}^2\rangle$ individual fit we obtained $m_{22}<0.48$ and$m_{22}<0.79$ for Fornax and Sculptor respectively. The tight bound on $m_a<0.4\times 10^{-22}\text{ eV}$ required if ULAs are to provide kpc-scale cores to Fornax and Sculptor (WP11 cores) runs into several problems when faced with cosmology. Firstly, we performed a simple estimate of the subhalo mass function, and demonstrated that such a ULA cannot provide enough dwarf satellites of the MW to be consistent with observations. This suggest that ULAs, like WDM, may suffer from a \emph{Catch 22} in that ``if you want large cores, you don't get enough dwarfs; if you want enough dwarfs, you don't get big enough cores''. Existing cosmological bounds also appear to rule out a ULA origin for the WP11 cores: high-$z$ galaxies rule out $m_a\lesssim 1\times 10^{-22}\text{ eV}$~\citep{2015MNRAS.450..209B,2016ApJ...818...89S}. Constraints from reionization, such as the CMB optical depth $\tau$, are less conclusive, and may even favour low axion mass. If ULAs satisfying our bound are indeed the DM, they can be definitively ruled out either by proper analysis of existing Lyman-$\alpha$ forest power spectrum data~\citep{2017arXiv170309126A}, or by upcoming CMB Stage-IV lensing power spectrum measurements~\citep{2016arXiv160708208H}. However, we would like to emphasis the need of detailed studies of the interplay between the physics of the axion DM and the physics of baryons, in order to make more consistent comparisons. Improvements on DM constraints from stellar kinematics may be obtained from future measurements of proper motions \citep[e.g.][]{2007ApJ...657L...1S}. Extending our present analysis (chemodynamical modelling of multiple stellar populations combined with the $\langle\sigma_{\rm los}^2\rangle$-fit) to other dwarfs beyond Fornax and Sculptor can be achieved by multi-object spectrographs attached to several-metre telescopes. On a short time scale we have WEAVE\citep{2014SPIE.9147E..0LD}, MOONS \citep{2012SPIE.8446E..0SC} and 4MOST \citep{2012SPIE.8446E..0TD}. Using these data and our methodology could significantly improve our ability to test models of DM using stellar dynamics. | 16 | 9 | 1609.05856 |
1609 | 1609.02186_arXiv.txt | We outline a new method to compute the Bayes Factor for model selection which bypasses the Bayesian Evidence. Our method combines multiple models into a single, nested, Supermodel using one or more hyperparameters. Since the models are now nested the Bayes Factors between the models can be efficiently computed using the Savage-Dickey Density Ratio (SDDR). In this way model selection becomes a problem of parameter estimation. We consider two ways of constructing the supermodel in detail: one based on combined models, and a second based on combined likelihoods. We report on these two approaches for a Gaussian linear model for which the Bayesian evidence can be calculated analytically and a toy nonlinear problem. Unlike the combined model approach, where a standard Monte Carlo Markov Chain (MCMC) struggles, the combined-likelihood approach fares much better in providing a reliable estimate of the log-Bayes Factor. This scheme potentially opens the way to computationally efficient ways to compute Bayes Factors in high dimensions that exploit the good scaling properties of MCMC, as compared to methods such as nested sampling that fail for high dimensions. | One of the key questions underlying science is that of model selection: how do we select between competing theories which purport to explain observed data? The great paradigm shifts in science fall squarely into this domain. In the context of astronomy - as with most areas of science - the next two decades will see a massive increase in data volume through large surveys such as the Square Kilometre Array (SKA) \citep{hollitt2016overview} and LSST \citep{becla2006designing}. Robust statistical analysis to perform model selection at scale will be a critical factor in the success of such future surveys. The basic problem of model selection is easy to state. As one considers models with more and more free parameters, one must expect that such models will fit any dataset better and better, irrespective of whether they have anything to do with reality. This problem of overfitting has led to many proposed methods to deal with this kind of situation: that is, finding a way to suitably penalise extra parameters. One method is LASSO (Least Absolute Shrinkage and Selection Operator) \citep{hastie2005elements}. Other methods such as Akaike Information Criterion (AIC) \citep{akaike1974new} and Bayesian Information Criteria (BIC) \citep{schwarz1978estimating} penalise the best fit likelihood based on the number of free parameters \citep{gelman2014understanding}. From a Bayesian point of view, model selection is not viewed as a question to be answered looking only at a single point in the parameter spaces, e.g. the point of maximum likelihood of the models in question, but rather should also depend on the full posterior distribution over the parameters. Hence selection is performed by choosing the model with the maximum model probability $\PP(\MM|\DD)$, derived from the Bayesian Evidence (or marginal likelihood) $\PP(\DD|\MM)$. This automatically expresses Occam's razor, thus penalising extra parameters which are not warranted by the data. Here and throughout this paper we will use $\DD$ to denote data and $\MM$ for a model. Given two competing models, one would typically compute the Bayesian Evidence for each model and hence the Bayes Factor, which is the ratio of the evidences. There are a number of issues with the Bayesian evidence. It is very sensitive to priors and, of key interest to us, since it involves integrals over the full parameter spaces of each model, is hard to compute efficiently. Techniques such as nested sampling (\citep{skilling2004nested}) scale exponentially with the number of parameters and cannot be used for high-dimensionality problems. However, if one model is nested within the other (i.e. all the parameters of one model are shared by another), we can use the Savage-Dickey Density Ratio (SDDR) (\citet{dickey1971weighted} and \citet{verdinelli1995computing}) to directly calculate the Bayes Factor. As an example, consider the case where the parameters in model $\MM_{1}$ are $\phi$ and $\theta$ while the parameter in model $\MM_{2}$ is $\theta$. Then, $\MM_{2}$ is nested in $\MM_{1}$ at some value of $\phi$ which we can take to be $\phi=0$. The Bayes Factor is then given directly by \begin{equation} \textrm{B}_{21}=\left.\dfrac{\mathcal{P}\left(\phi\left|\DD,\,\MM_{1}\right.\right)}{\mathcal{P}\left(\phi\left|\MM_{1}\right.\right)}\right|_{\phi=0} \end{equation} where $\mathcal{P}\left(\phi\left|\DD,\,\MM_{1}\right.\right)$ is simply the normalised posterior probability distribution of $\phi$ in the extended model, that is: \[ \mathcal{P}\left(\phi\left|\DD,\,\MM_{1}\right.\right)=\int\mathcal{P}\left(\theta,\,\phi\left|\DD,\,\MM_{1}\right.\right)\,d\theta \] The core of this paper is the idea that it is possible to embed any two models into a {\em Supermodel} such that each model is nested within the supermodel. Related ideas can be found in \citep{hee2016bayesian, hlozek2012photometric, kamary2014testing}. In the next sections, we shall illustrate this in detail. The paper is organised as follows: in \S\ref{sec:methods}, we describe our idea in the general context. In \S\ref{sec:application_linear_model} and \S\ref{sec:application_non_linear}, we test our approach using both the linear and non-linear models while in \S\ref{sec:reparam_exploit}, we also consider one example of reparameterization of $\alpha$, the hyperparameter with respect to which the models are nested. We conclude in \S\ref{sec:conclusion}. | \label{sec:conclusion} In this work we have used the Savage-Dickey Density Ratio (SDDR) to show that we can calculate the Bayes Factor of two non-nested models by introducing a new hyperparameter that combines the models into a single supermodel. This Savage-Dickey Supermodel (SDSM) method does not need the Bayesian evidence (Marginal Likelihood) to be computed. The core supermodel embedding can be done either at the level of the model (eq. (\ref{combmod})) or at the level of the likelihood (eq. (\ref{comblike})) and effectively makes the the models nested and hence amenable to the SDDR approach to computing the Bayes Factors. In the context of Gaussian linear models we show that the SDDR both analytically and numerically reproduces the Bayes Factors computed analytically. We then consider a nonlinear example and show that our supermodel approach agrees well with that from nested sampling. Though we have a clever way of avoiding multidimensional integrals to calculate the Bayesian Evidence, this new method requires very efficient sampling and for a small number of dimensions is not faster than individual nested sampling runs. The major reason for this is that we require independent samples for $\alpha$ and one way to ensure we are doing so is to have a short autocorrelation length. Hence the thinning factor for the MCMC chain needs to be adjusted as well as the number of the steps, especially for large log-Bayes Factor. However, generically the scaling of MCMC methods with the number of dimensions is much more benign than the scaling of nested sampling methods. The approach presented here is thus expected to work also for very high numbers of dimensions where nested sampling fails. Additionally, if we only keep in a MCMC chain the elements for which $\alpha=1$ or $\alpha=0$ then we obtain a model-averaged posterior. For this application we do not need a very high number of samples, so that the method is competitive with nested sampling for model averaged posteriors also at a smaller number of dimensions. For future work we note that other, nonlinear, combinations of models/likelihoods are also possible. For example, consider product combined model and likelihood $\mathcal{M}_{3}=\mathcal{M}_{1}^{\alpha}\mathcal{M}_{2}^{\left(1-\alpha\right)}$ and $\mathcal{L}_{3}=\mathcal{L}_{1}^{\alpha}\mathcal{L}_{2}^{\left(1-\alpha\right)}$ in which case, the general condition (\ref{eq:general_condition}) still holds for $\alpha\in\left[0,\,1\right]$. Such nonlinear supermodels, choices of reparametrisation function $f\left(\alpha\right)$ or other innovations (such as using simulated annealing) may greatly simplify some aspects of the sampling and provide a clever way of not only obtaining the log-Bayes Factor, which helps us to understand the relative strength of the models but also to have model averaged posteriors of all the parameters in both models. Study of these generalisations is left to future work. | 16 | 9 | 1609.02186 |
1609 | 1609.02379_arXiv.txt | We present spectroscopic and interferometric measurements for a sample of nine K giant stars. These targets are of particular interest because they are slated for stellar oscillation observations. Our improved parameters will directly translate into reduced errors in the final masses for these stars when interferometric radii and asteroseismic densities are combined. Here we determine each star's limb-darkened angular diameter, physical radius, luminosity, bolometric flux, effective temperature, surface gravity, metallicity, and mass. When we compare our interferometric and spectroscopic results, we find no systematic offsets in the diameters and the values generally agree within the errors. Our interferometric temperatures for seven of the nine stars are hotter than those determined from spectroscopy with an average difference of about 380 K. | Giant stars are excellent candidates for both interferometric and asteroseismic observations. Interferometers have been used for many years to measure the angular diameters of giants, from the Mark III Interferometer \citep[e.g.,][]{2003AJ....126.2502M} to the Palomar Testbed Interferometer \citep[e.g.,][]{1999AJ....117..521V} to the Navy Precision Optical Interferometer \citep[e.g.,][]{1999AJ....118.3032N}. More recently, a sample of 25 K giant stars was measured by our team using the Center for High Angular Resolution Astronomy (CHARA) Array \citep{2010ApJ...710.1365B}. The other technique under consideration for this sample is asteroseismology, the study of stellar oscillations. It is a unique tool to infer the structure of stellar interiors with very little model dependence \citep[see, e.g.,][]{1994ARAandA..32...37B, 2004SoPh..220..137C}. Photometric space missions focusing on asteroseismology, i.e., \emph{MOST} \citep[\emph{Microvariability and Oscillations of STars,}][]{2003PASP..115.1023W}, \emph{CoRoT} \citep[\emph{Convection, Rotation, and planetary Transits,}][]{2006ESASP.624E..34B, 2009AandA...506..411A}, and \emph{Kepler} \citep{2013ApJ...765L..41S, 2010Sci...327..977B, 2010ApJ...713L..79K}, have dramatically increased both the number of stars with oscillation measurements as well as the quality of the data. These are critical measurements because the frequencies observed are dependent on the sound speed inside the star, which in turn depends on interior properties such as density, temperature, and gas motion \citep{2010AandA...509A..73C}. The stellar parameters resulting from these observations are key for testing stellar interior and evolutionary models \citep[see, e.g.,][]{2011Sci...332..213C}. Most giant stars, if not all of them, display measurable oscillations \citep[e.g.,][]{2013ApJ...765L..41S, 2006AandA...448..689D, 2002AandA...394L...5F, 1994ApJ...422..366H}, which makes them an ideal class of objects for deriving fundamental stellar parameters such as mass, radius, and temperature. They are bright, abundant, large enough to measure easily with interferometry, and exhibit radial velocity amplitudes from a few to several tens of m s$^{-1}$. The observed oscillation frequencies put constraints on the star's internal structure \citep{2006ApJ...647..558B}, namely the mean density of the star, while interferometry measures the star's size. The combination leads to the masses for these single stars. The defining characteristic of a star is its mass but for giant stars, determining this quantity is indirect and heavily model dependent. Often spectroscopic observations are used to measure a star's surface gravity (log~$g$), effective temperature ($T_{\rm eff}$), and iron abundance ([Fe/H]). The radius and mass are then determined by fitting evolutionary tracks to the star's position on the H-R diagram. This is an tricky process because the evolutionary tracks of stars with a large range of masses converge on the H-R diagram in the same region, and different evolutionary track models produce different masses for a given set of inputs. Without good calibrating objects, no set of tracks can be proven to be the best. Once we can test them by comparing theoretically determined mass and radius to measured values, we can have faith in applying the tracks to stars for which direct measurements are not possible. Several of the stars in our sample are ear-marked for asteroseismic studies using precise stellar radial velocity (PRV) measurements. It is difficult to obtain sufficient data in order to detect all pulsation modes using ground-based facilities. This requires a large amount of observing time often using multi-site campaigns. However, it is still possible to derive the stellar mass using a modest amount of ground-based data even taken at one site {\it if} one knows the stellar radius. This was done with some success for $\beta$ Gem \citep{2012AandA...543A..98H} and $\iota$ Dra \citep{2008AandA...491..531Z, 2011ApJ...743..130B}. PRV measurements will be made using the Thuringia State Observatory's 2 m telescope and McDonald Observatory's 2.7 m telescope, and results will be presented in a forthcoming paper. In the near future, network telescopes such as the Stellar Oscillations Network Group \citep[SONG,][]{2013POBeo..92...39G} should be able to investigate better the pulsations in these stars using PRVs. It is important, however, to first obtain stellar radii measurements, which is the goal of this paper. The measured angular diameters, when combined with other measurements from the literature, ultimately lead to radii ($R$) and $T_{\rm eff}$ for the giant stars. These are important properties that characterize the star as well as the environment in which any possible exoplanets reside. Section 2 discusses the spectroscopic measurements of $T_{\rm eff}$, log $g$, and [Fe/H], section 3 describes the interferometric observations and calibrator star selection, section 4 outlines how we measure the angular diameter and calculate the $R$, luminosity, and $T_{\rm eff}$ for our sample, section 5 explores the physical implications of our measurements and plans for oscillation measurements, and section 6 summarizes our findings. | \subsection{Comparing Spectroscopic and Interferometric Diameters} We compared the angular diameters predicted using the Girardi tracks using spectroscopically determined $T_{\rm eff}$ and [Fe/H] against the interferometric measurements in Figure \ref{angdiam}. For the most part, the diameters agree within the errors and there is no clear bias. The error bars on the interferometric measurements are substantially smaller than those on the Girardi diameters, between 3$\times$ and 19$\times$ smaller: the errors for $\theta_{\rm interf}$ are on the order of 1 to 3$\%$ with just one at 6$\%$, while the errors for $\theta_{\rm Girardi}$ range from 11$\%$ to 18$\%$. The largest outliers in Figure \ref{angdiam} are HD 31579 and HD 157681. The latter was observed as part of \citet{2010ApJ...710.1365B}, and its interferometric diameter of 1.664$\pm$0.010 mas was larger than the diameter predicted by spectroscopy (1.27$\pm$0.24 mas). Baines et al. concluded it was likely due to the calibrator star used (HD 158460) so we observed it again using two different calibrators. Our new diameter of 1.901$\pm$0.013 mas is even larger than the previous measurement. However, when the data are analyzed using each calibrator star separately, the resulting angular diameters are remarkably consistent with a mere 0.003 mas difference. When the calibrators are used to calibrate each other, no systematic offsets are present. We also used the relationship described in \citet{1999PASP..111.1515V} between the angular diameter and the ($V-K$) color to estimate HD 157681's diameter and obtained 2.05$^{+0.45}_{-0.82}$ mas, which agrees with our new interferometric measurement to within the errors. As for HD 31579, the spectroscopically determined angular diameter (0.91$\pm$0.36 mas) is the outlier when considered against the those determined using the SED fit (1.60$\pm$0.12 mas), the ($V-K$) color (1.67$\pm$0.27 mas), and the interferometric measurement (1.593$\pm$0.008 mas). All the diameters are consistent and agree to within the errors except for the spectroscopic calculation. \subsection{Comparing Spectroscopic and Interferometric Temperatures} We plotted the spectroscopically determined $T_{\rm eff}$ versus our interferometric results in Figure \ref{tcompare}. There is some scatter off the 1:1 line with the spectroscopic values tending to be cooler than the interferometric ones by an average of $\sim$380 K. The discrepancy may be due to the atmospheric models of K giant stars in the near-ultraviolet lacking a source of thermal extinction, which could affect the $T_{\rm eff}$ measurements \citep{2009ApJ...691.1634S}. Another cause may lie in the methods used to determine $T_{\rm eff}$: interferometry measures the overall $T_{\rm eff}$ of the star while spectroscopic values rely on Fe\,I and Fe\,II lines and measure the $T_{\rm eff}$ in the thin layers of the atmosphere where those lines are formed. In dwarf stars, local thermodynamic equilibrium is a reasonable assumption and the $T_{\rm eff}$ determined using the iron lines is the same as the $T_{\rm eff}$ of the atmosphere overall. For giant stars, the atmosphere is more extended and the models may not be correct due to factors such as convection. Another consideration may be that the 1-D models do not include geometrical surface cooling and the 2-D models may not be as extended as real stars, so do not perfectly describe the atmospheres. HD 157681 is again an object of interest when it comes to determining its $T_{\rm eff}$. In order to match $\theta_{\rm spec}$, $T_{\rm eff}$ would have to drop from 4400 K to 3844 K, which is much closer to the 3900 K predicted by the ($B-V$) color. This has the effect of moving the star from below the 1:1 line in Figure \ref{tcompare} to above it, which is consistent with the rest of the stars except for HD 31579. HD 157681 is the coolest giant in the sample, which is expected because it is a K5 star while the others are K0 to K3. As an independent check on $T_{\rm eff}$, we used the equations from \citet{2010MNRAS.403.1592B} that relate ($B-V$) color, bolometric correction (BC$_V$), and $T_{\rm eff}$ for stars between 3300 and 5000 K. The results are listed in Table \ref{tcompare_table}. Color $T_{\rm eff}$ are even cooler than the spectroscopic $T_{\rm eff}$, except for HD 216174 where they are equal. On average the spectroscopic $T_{\rm eff}$ are hotter than the color $T_{\rm eff}$ by $\sim$320 K, while the interferometric $T_{\rm eff}$ are hotter on average by $\sim$580 K. We also did a search in the literature using the VizieR service and averaged all available $T_{\rm eff}$ values, and these are included in Table \ref{tcompare_table}. As a final check, we calculated both $\theta_{\rm LD}$ and $T_{\rm eff}$ using the relations between them and the surface brightness and ($V-K$) color, respectively, described in \citet{2003AJ....126.2502M}. Table \ref{tcompare_table} lists the resulting values, which are also plotted in Figure \ref{diams_teff}. The diameters show a scatter around the 1:1 ratio but are within the errors, and we see a similar offset in $T_{\rm eff}$, where our new measurements are hotter than those predicted using Mozurkewich et al.'s equations for seven of the nine stars. When we compare the temperatures determined spectroscopically, interferometrically, and using the ($V-K$) colors, four of the nine stars have $T_{\rm inf}$ that fall in between the $T_{\rm spec}$ and $T_{(V-K)}$. \subsection{Future Oscillation Studies} The velocity amplitude of the K giant stars' p-mode oscillations range from a few to tens of m\,s$^{-1}$, depending on the evolutionary state of the star \citep{1995AandA...293...87K}. The mode periods range from several hours to days. These amplitudes and periods are measureable with 2-3 m class telescopes using precise stellar RV measurements, which typically reach a precision of $\sim$1 m\,s$^{-1}$. We intend to use the Coud\'e echelle spectrograph of the 2-m Alfred Jensch Telescope of the Thuringia State Observatory to detect the stellar oscillations in those stars for which we have interferometrically measured $R$. An iodine absorption cell will be used to provide the wavelength calibration for the RV measurement. This instrument is able to achieve and RV precision of $\sim$2 m\,s$^{-1}$ on bright K giant stars \citep{2012AandA...543A..98H}. Fundamental stellar parameters of K giant stars are important for exoplanet studies because of their masses, which can be 1.5--3 $M_\odot$. Main-sequence stars of this mass range are ill-suited for RV measurements due to a paucity of stellar lines and high stellar rotation rates. Thus K giants offer us a means to study planet formation around stars more massive than the Sun. \subsection{Stellar Masses} Determining $M$ for these giant stars is key to understanding whether or not planet populations orbiting massive stars are different than planets found orbiting solar-type stars. Some scientists argue that more massive stars host more massive planets, and that A stars are at least five times more likely to host a giant planet than an M dwarf \citep{2010ApJ...709..396B, 2010PASP..122..149J, 2010PASP..122..701J, 2012AandA...544A...9V}. There are models that support this theory: e.g., \citet{2008ApJ...673..502K, 2013ApJ...778...78H}. However, \citet{2011ApJ...739L..49L, 2013ApJ...774L...2L} disagrees, claiming the masses determined for the exoplanet host stars are in error due to the convergence and crossing of evolutionary tracks from stellar models. This leads to degeneracies, and Lloyd believes the masses of the evolved stars are not as high as those claimed by previous studies. Our ultimate contribution to this controversy will be the direct determination of $M$ for our sample of giant stars by combining our interferometric $R$ with the asteroseismic density measurements. We will then be able to determine if the models are indeed correct, and test if the idea that more massive stars host more massive planets is valid. | 16 | 9 | 1609.02379 |
1609 | 1609.09631_arXiv.txt | {Explaining the accelerated expansion of the Universe is one of the fundamental challenges in physics today. Cosmography provides information about the evolution of the universe derived from measured distances, assuming only that the space time geometry is described by the Friedman-Lemaitre-Robertson-Walker metric, and adopting an approach that effectively uses only Taylor expansions of basic observables.} {We perform a high-redshift analysis to constrain the cosmographic expansion up to the fifth order. It is based on the Union2 type Ia supernovae data set, the gamma-ray burst Hubble diagram, a data set of 28 independent measurements of the Hubble parameter, baryon acoustic oscillations measurements from galaxy clustering and the Lyman-$\alpha$ forest in the SDSS-III Baryon Oscillation Spectroscopic Survey (BOSS), and some Gaussian priors on $h$ and $\Omega_M$.} {We performed a statistical analysis and explored the probability distributions of the cosmographic parameters. By building up their regions of confidence, we maximized our likelihood function using the Markov chain Monte Carlo method.} {Our high-redshift analysis confirms that the expansion of the Universe currently accelerates; the estimation of the jerk parameter indicates a possible deviation from the standard $\Lambda$CDM cosmological model. Moreover, we investigate implications of our results for the reconstruction of the dark energy equation of state (EOS) by comparing the standard technique of cosmography with an alternative approach based on generalized Pad\'e approximations of the same observables. Because these expansions converge better, is possible to improve the constraints on the cosmographic parameters and also on the dark matter EOS.} {The estimation of the jerk and the DE parameters indicates at $ 1 \sigma $ a possible deviation from the $\Lambda$CDM cosmological model.} | In the past dozen years a huge and diverse set of observational data revealed that the Universe is now expanding at an accelerated rate, see for instance \cite{Riess07}, \cite{SNLS}, \cite{Riess11}, \cite{Union2.1}, \cite{PlanckXXVI}, and \cite{PlanckXIII}. It is usually assumed that this accelerated expansion is caused by the so-called dark energy, a cosmic medium with unusual properties. The pressure of dark energy $p_{de}$ is negative, and it is related to the positive energy density of dark energy $\epsilon_{de}$ by the equation of state (EOS), $p_{de}=w\epsilon_{de}$, where the proportionality coefficient is $ < -1/3$. According to current estimates, about 75\% of the matter energy in the Universe is in the form of dark energy, so that today the dark energy is the dominant component in the Universe. The nature of dark energy is not known. Models of dark energy proposed so far include at least a non-zero cosmological constant (in this case $w=-1$), a potential energy of some scalar field, effects connected with inhomogeneous distribution of matter and averaging procedures, and extended theories of gravity (an accelerated expansion can be obtained by generalizing the Einstein theory of gravity to some theory derived from a modified action with respect to the Hilbert-Einstein action: the simplest extension of General Relativity is achieved assuming that the gravitational Lagrangian is an arbitrary continuous function $f(R)$ of the Ricci scalar $R$ . In this case, in general, $w\not= -1$ and it is not constant and depends on the redshift $z$. Extracting the information on the EOS of dark energy from observational data is then at the same time a fundamental problem and a challenging task. To probe the dynamical evolution of dark energy in these circumstances, we can parameterize $w$ empirically, usually using two or more free parameters. Of all the parametrization forms of the dark energy EOS, the Chevallier-Polarski-Linder (CPL) model \cite{cpl1}, \cite{cpl2} is probably the most widely used, since it presents a smooth and bounded behavior for high redshifts and a manageable two-dimensional parameter space and also provides a simple and effective instrument of computations. However, it would result a in physically incomplete parametrization of dark energy if we were to take into account the inhomogeneities of the late-time Universe. Linear parametrizations of the dark energy EOS (the CPL EOS is linear in the scale factor $a$) are not compatible with the theory of scalar perturbations in the late Universe. Therefore these EOS are not the fundamental and can only be used to approximate the real EOS \cite{akarsu15}. In our approach, this model is only used to investigate whether the EOS is constant, independently of any assumption on the nature of the DE: according to this point of view, even the small number of parameters of the CPL model is not as important as this independence (in some scalar field models of dark energy, the so-called quintessence, first introduced in \cite{Peebles88a}, \cite{Peebles88b}, the scalar field has one free parameter less than CPL). Moreover, it is worth noting that even neglecting the inhomogeneities, several dark energy models considered so far agree reasonably well with the observational data, so that, unless higher precision probes of the expansion rate and the growth of structure are developed, these different approaches cannot be distinguished. This degeneration suggests a kinematical approach to the problem of cosmic acceleration, relying on quantities that are not model dependent. The cosmographic approach is related to the derivatives of the scale factor and enables fitting the data on the distance - redshift relation without any a priori assumption on the underlying cosmological model. It is based on the sole assumption that the Universe is spatially homogeneous and isotropic, and that it can be described by the Friedman-Lemaitre-Robertson-Walker (FLRW) metric. In our high-redshift investigation, extended behind the supernova type Ia (SNIa) Hubble diagram, we require at least a fifth-order Taylor expansion of the scale factor to obtain a reliable approximation of the distance - redshift relation. As a consequence, it is in principle possible to estimate up to five cosmographic parameters, $(h, q_0, j_0, s_0, l_0)$, although the available high-redshift data sets are still too small and do not allow us to obtain a precise and realistic determination of all of them, see \cite{salzcosmo}. When these quantities have been determined, we can use them to set constraints on the dark energy models.To constrain the cosmographic parameters, we use the Union2 SNIa data set, the gamma-ray burst (GRB) Hubble diagram, constructed by calibrating the correlation between the peak photon energy, $E_{\mathrm{p, i}}$, and the isotropic equivalent radiated energy, $ E_{\mathrm{iso}}$ (see Paper I), a sample of 28 measurements of the Hubble parameter, compiled in \cite{farooqb}, Gaussian priors on the distance from the baryon acoustic oscillations (BAO), and the Hubble constant $h$ (these priors have been included to help break the degeneracies of the model parameters). Our statistical analysis is based on Monte Carlo Markov Chain (MCMC) simulations to simultaneously compute the full probability density functions (PDFs) of all the parameters of interest. The structure of the paper is as follows. In Sect. 2 we describe the basic elements of the cosmographic approach and explicitly derive series expansions of the scale factor and other relevant parameters. In Sect. 3 we describe the observational data sets that are used in our analysis. In Sect. 4 we describe some details of our statistical analysis and present results on cosmographic parameters obtained from three sets of data. In Sect. 5 we present constraints on dark energy models that can be derived from our analysis. General discussion of our results and conclusions are presented in Sect. 6. | We investigated the dynamics of the Universe by using a cosmographic approach: we performed a high-redshift analysis that allowed us to set constraints on the cosmographic expansion up to the fifth order, based on the Union2 SNIa data set, the GRB Hubble diagram, constructed by calibrating the correlation between the peak photon energy, $E_{\mathrm{p, i}}$, and the isotropic equivalent radiated energy, $ E_{\mathrm{iso}}$, and Gaussian priors on the distance from the BAO, and the Hubble constant $h$ (these priors were included to help break the degeneracies among model parameters). Our statistical analysis was based on MCMC simulations to simultaneously compute the regions of confidence of all the parameters of interest. Since methods like the MCMC are based on an algorithm that randomly probes the parameter space, to improve the convergence, we imposed some constraints on the series expansions of $H(z)$ and $d_{L}(z)$, requiring that in each step of our calculations $d_{L}(z) > 0$ \,, and $H(z) > 0$. We performed the same MCMC calculations, first considering the SNIa Hubble diagram and the BAO data sets or the GRBs Hubble diagram, and the BAO data sets separately (Cosmography Ia and Ib, respectively), and then constructing an overall data set joining them together (Cosmography II). Our MCMC method allowed us to obtain constraints on the parameter estimation, in particular for higher order cosmographic parameters (the jerk and the snap). The deceleration parameter confirms the current acceleration phase; the estimation of the jerk reflects at $1 \sigma$ the possibility of a deviation from the $\Lambda$CDM cosmological model. Moreover, we investigated implications of our results for the reconstruction of the dark energy EOS by comparing the standard technique of cosmography with an alternative approach based on generalized Pad\'e approximations of the same observables. Owing to the better convergence properties of these expansions, it is possible to improve the constraints on the cosmographic parameters and also on the dark matter EOS: our analysis indicates that at the $1 \sigma$ level the dark energy EOS is evolving. \subsection* | 16 | 9 | 1609.09631 |
1609 | 1609.01716_arXiv.txt | Photometric measurements are prone to systematic errors presenting a challenge to low-amplitude variability detection. In search for a general-purpose variability detection technique able to recover a broad range of variability types including currently unknown ones, we test 18 statistical characteristics quantifying scatter and/or correlation between brightness measurements. We compare their performance in identifying variable objects in seven time series data sets obtained with telescopes ranging in size from a telephoto lens to 1\,m-class and probing variability on time-scales from minutes to decades. The test data sets together include lightcurves of 127539 objects, among them 1251 variable stars of various types and represent a range of observing conditions often found in ground-based variability surveys. The real data are complemented by simulations. We propose a combination of two indices that together recover a broad range of variability types from photometric data characterized by a wide variety of sampling patterns, photometric accuracies and percentages of outlier measurements. The first index is the interquartile range (IQR) of magnitude measurements, sensitive to variability irrespective of a time-scale and resistant to outliers. It can be complemented by the ratio of the lightcurve variance to the mean square successive difference, $1/\eta$, which is efficient in detecting variability on time-scales longer than the typical time interval between observations. Variable objects have larger $1/\eta$ and/or IQR values than non-variable objects of similar brightness. Another approach to variability detection is to combine many variability indices using principal component analysis. We present 124 previously unknown variable stars found in the test data. | A variety of phenomena manifest themselves as changes in apparent brightness of astronomical objects. The amplitudes and time-scales of these changes vary from tens of magnitudes and weeks for supernova explosions to a fraction of a magnitude and minutes for stellar pulsations. With the notable exceptions of light echoes \citep[e.g.][]{2003Natur.422..405B}, variable reflecting nebulae \citep[e.g.][]{1997ApJ...489..210C} and the M87 jet \citep[e.g.][]{2011ApJ...743..119P} variable objects are unresolved by single-dish telescopes\footnote{The light travel time argument implies that an object varying on a time-scale $t$ cannot be larger than $ct$ light seconds, otherwise its variability would be smeared.}. Variable point-like objects are often embedded in light of a resolved non-variable source (active nucleus or a supernova in a galaxy; young stellar object embedded in a nebula) that complicate measurements of the variable object's brightness. The variations may be associated with a single catastrophic event (supernova), may be approximately (dwarf novae) or strictly periodic (eclipsing binaries) or aperiodic (active galactic nuclei) in nature. Our understanding of these events depends on the efficient and reliable detection of brightness variations. Photometric measurements are prone to systematic errors that are difficult to characterize. This makes it challenging to distinguish true low-amplitude variability from the apparent one caused by systematic effects and measurement errors. Imaging artefacts such as cosmetic defects of a CCD, diffraction spikes from bright objects and cosmic ray hits as well as blending between images of nearby objects can corrupt photometry and mimic high-amplitude variability. Three different lines of attack on the problem of variable object detection are described in the literature: direct image comparison, (``transient detection''), lightcurve analysis using variability indices and periodicity search. Transient detection techniques seek to identify changes between two sets of sky images taken at different times (epochs). The changes may be found by subtracting the images pixel-by-pixel after resampling them to a common coordinate grid and accounting for seeing changes (difference image analysis -- DIA; \citealt{1996AJ....112.2872T}, \citealt{1998ApJ...503..325A,2000A&AS..144..363A}, \citealt{2008MNRAS.386L..77B,2012MNRAS.425.1341B}, \citealt{2015arXiv151206872Z}, \citealt{2016MNRAS.457..542B}; applications of the method include \citealt{2003AJ....126..175B,2003AstL...29..599Z,2013AcA....63..429A,2014PASA...31...12S,2015RAA....15..215Z}). Large surveys such as OGLE \citep{2015AcA....65....1U}, PTF~\citep{2009PASP..121.1395L}, Pan-STARRS \citep{2014ApJ...795...44R} and DES~\citep{2015AJ....150..172K} implement the image subtraction technique. Alternatively, one may extract astronomical objects (sources) from each image independently and compare the resulting source lists (\citealt{2014MNRAS.439.1829C}, CRTS -- \citealt{2009ApJ...696..870D}). The second-epoch images are often taken in pairs, triplets or even longer series with dithering to eliminate image artefacts that are usually associated with a given position on the image detector, not in the sky. More sophisticated detection strategies may be applied if measurements are obtained at more than two epochs. Their obvious advantage over the simple two-epoch data comparison is the potential to average-out individual measurement errors and thus detect variability with a lower amplitude. One class of methods employs various ``variability indices'' characterizing the overall scatter of measurements in a lightcurve and/or degree of correlation between consecutive flux measurements \citep[some recent examples:][see the detailed discussion in Section~\ref{sec:indices}]{2014JAD....20....4M,2015MNRAS.447.3973J,2015AJ....150..107Y}. The other class of methods % search for significant periodicity in flux variations \citep[e.g.][]{2014MNRAS.438.3383M,2014ApJS..213....9D,2014AcA....64..309K,2015ApJ...802L...4C,2015MNRAS.447.3536N,2016MNRAS.455.2337N,2015AcA....65..297S}. While many types of variable stars show periodic or semi-periodic light variations, flux measurement errors are expected to be aperiodic, or associated with a known periodic process inherent to the observations (diurnal or seasonal cycle, synodic month, periodic guiding errors, orbital period of a spaceborne telescope, etc.). If a search is aimed at a specific variability type for which a lightcurve shape is generally known in advance (e.g. exoplanet transits or eclipsing binaries in general, Cepheids, RR~Lyrae stars, novae), template fitting (e.g. \citealt{1996Icar..119..244J}, \citealt{1999ApJ...521..155M}, \citealt{2011AJ....141...83P,2013AJ....146...21S,2014A&A...567A.100A}, \citealt{2016arXiv160708658H}) with various trial periods/flare development time-scales can be performed. Simple cuts on lightcurve parameters \citep{2008A&A...477...67H,2010AcA....60..109G} as well as advanced machine learning techniques \citep{2005AJ....130...84F} can be used to select lightcurves of a known shape from a large photometric data set. A pre-selection based on colour can be used to reduce the number of candidates when searching for variables of a specific type \citep[][]{2006AJ....132.1202K,2013A&A...551A..77T,2014ApJ...781...22Z,2016MNRAS.455.2163O,2016MNRAS.459.1687M}. Since period search and template-fitting algorithms are computationally expensive, a two-step approach can be applied. Candidate variable stars are pre-selected using a fast-to-compute variability index (and/or colour) and only the lightcurves that passed this selection are subjected to period search \citep[e.g.][]{2000AJ....119.1901A,2013ApJ...763...32D,2014A&A...562A.125K,2015A&A...573A.100F,2015A&A...574A..15F,2015A&A...575A.114G,2016AJ....151..118V} or template fitting \citep[e.g.][]{2011ApJ...733..124S,2015AJ....149..183H}. If the total number of observed objects is low, both period search and lightcurve scatter-based selection criteria are applied independently of each other to conduct exhaustive search for both periodic and non-periodic variables \citep{2014AcA....64..115S,2015AJ....150..175R}. Selection based on period search may be followed by even more computationally intensive steps like binary system modelling \citep{2005ApJ...628..411D}. \cite{2014A&A...566A..43K} used the period along with other variability features as an input for the random forest algorithm to select periodic variable star candidates in the EROS-2 database and simultaneously classify them. The methods described above may efficiently select variable object candidates from a large set of photometric data. However, the final decision to designate an object as ``variable star'' rather than a ``candidate'' is usually made only after visual inspection of the object's lightcurve by a human expert (e.g. \citealt{2005AcA....55..275P,2011AcA....61..103G,2013AcA....63..115P,2013AcA....63..323P,2013AJ....146..101P}, \citealt{2013ApJ...779....7C}, \citealt{2016AJ....151..110K}, \citealt{2016arXiv160604792S}). If the number of observations is small, the original images are checked for the presence of obvious problems [image artefacts, cosmic ray hits, point spread function (PSF) wings of a bright nearby object] affecting photometry of a candidate variable \citep[e.g.][]{2003AJ....125.1261D,2010ApJ...712.1259B,2011AstL...37...91D,2014MNRAS.437..132R}. While advanced image artifact rejection procedures exist \citep{2002PASP..114..144F,2016A&C....16...67D}, visual image inspection remains an important data quality control tool as it may uncover unexpected problems \citep{2016A&C....16...99M}. Variable star detection may be considered in the framework of classical hypothesis testing \citep[e.g.][]{2003psa..book.....W}: to establish that an object is variable, one needs to rule~out the null hypothesis that it is constant given the observations \citep{2006ASPC..349...15E}. One may compare a value of variability-detection statistic (Section~\ref{sec:indices}) derived from the lightcurve to the distribution of this value expected for non-variable objects. The problem is that objects with corrupted measurements produce long tails in the aforementioned distributions. In the presence of badly measured objects one is forced to set a low threshold for accepting candidate variables (Section~\ref{sec:cutoffvalue}) and rely on additional information not captured by the variability-detection statistic to distinguish true variables from badly measured objects in the distribution tail. Alternatively one may view the search for variable stars as a classification problem that may be approached with machine learning techniques. The task is to classify a set of objects characterized by their lightcurves, images associated with each lightcurve point and possibly additional pieces of information associated with each brightness measurement (object's position on the CCD frame, airmass, seeing, temperature, etc.). One needs to distinguish various classes of variable stars from the class of well-measured constant stars and classes of stars affected by various types of measurement errors (bad pixels, diffraction spikes, blending). Objects that do not belong to one of the known classes should also be identified. While considerable progress has been made in lightcurve-based automated classification of stars already known to be variable \citep{2007A&A...475.1159D,2014AJ....148...31P,2016A&A...587A..18K}, an automated system that could reliably identify variable stars among non-variables remains to be developed. In practice, the following approach to variable star detection is often adopted. {(i)\,}Objects affected by blending and image artefacts are flagged at source extraction stage. {(ii)\,}The lightcurves of the detected objects are constructed and may be refined using the available additional information (Section~\ref{sec:filtering}). {(iii)\,}The techniques described in the previous paragraphs are used to select promising variable star candidates based on their lightcurves. {(iv)\,}The list of candidates is examined by a human expert who performs the final classification and removes false variables from the list. In this work we explore the limits of the traditional approach outlined above and identify the best ways to select candidate variables. We compare the performance of popular variability detection techniques on various real and simulated photometric data sets. We refer to any value that quantifies `how variable' a given object is as a `variability index'. The discussion is limited to variability indices based on lightcurve scatter (Sections~\ref{sec:chi2}--\ref{sec:peaktopeak}) and correlation (Sections~\ref{sec:lag1autocorr}--\ref{sec:sb}) while the period-search based techniques will be discussed elsewhere. We attempt to find a general-purpose variability detection technique able to recover a broad range of variability types including currently unknown ones \citep{2009MNRAS.400.1897S}. Such a technique would also be useful for solving the opposite problem: reliable selection of non-variable objects that can be used as photometric standards \citep[e.g.][]{2012PASP..124..854O} or targets for searches of variations not intrinsic or not typical to the objects such as microlensing events \citep{1994AcA....44..227U}, occultations of stars by distant Solar system bodies \citep{2013AJ....146...14Z}, tidal disruption events in nuclei of non-active galaxies \citep{2011ApJ...741...73V} and failed supernovae \citep{2008ApJ...684.1336K}. Publications focused on comparing performance of variability search techniques include \cite{2012A&A...548A..48E} who compared planetary-transit detection algorithms, while \cite{2010AJ....139.1269D} and \cite{2010ApJ...723..737V} discussed a number of variability detection tests in the context of active galactic nuclei. \cite{2016A&A...586A..36F} compared performance of some multi-band correlation-based variability indices. \cite{2003MNRAS.345.1271V} and \cite{2013ApJ...771....9A} discussed properties of `excess variance' (Section~\ref{sec:sigmaxs}) and `fractional variability amplitude', the variability measures often used in X-ray astronomy. \cite{2013MNRAS.434.3423G} compared the accuracy and performance of period finding algorithms. \cite{2015ApJ...798...89F} compare various methods of extracting a characteristic time-scale from an irregular lightcurve. \cite{2015arXiv150600010N} provide an extensive list of features useful for lightcurve characterization and classification. Preliminary results of our variability index comparison based solely on photographic lightcurves are presented by \cite{2016arXiv160503571S}. This paper is structured as follows. Section~\ref{sec:indices} defines the variability indices we investigate. Section~\ref{sec:testdata} describes the test data. Section~\ref{sec:comparison_tech} presents the technique for comparison of effectiveness of variability indices in selecting variable objects. Section~\ref{sec:resdisc} discusses the results of the comparison and Section~\ref{sec:conclusions} summarizes our findings. | \label{sec:conclusions} We compare 18 variability indices quantifying the overall scatter and/or degree of correlations between consecutive measurements in a lightcurve. The ability of these indices to distinguish variable stars from non-variable ones is tested on seven data sets collected with various ground-based telescopes and on simulated data incorporating actual lightcurves of non-variable objects as realistic models of photometric noise. We apply the PCA in search for an optimal combination of multiple variability indices. We find that correlation-based indices are more efficient in selecting variable objects than the scatter-based indices for data sets containing hundreds of measurement epochs or more. The indices $1/\eta$, $L$, ${\rm MAD}$ and ${\rm IQR}$ perform better than others in selecting candidate variables from data sets affected by outliers. We suggest using the $1/\eta$ index together with the ${\rm IQR}$ as the pair of indices applicable to a wide variety of survey strategies and variability types. The indices $1/\eta$ and ${\rm IQR}$ provide stable high performance, albeit not always the highest one for each of the investigated data sets. However, the overall quality of a photometric data set including the percentage of outlier measurements and number of badly measured objects has a higher impact on the efficiency of variability search than the choice of a specific (set of) variability index(es). Another efficient approach to variability detection is to compute many scatter- and correlation-based variability indices and perform the PCA over them. The admixture coefficient of the first principal component can be used as the composite index most suitable for the particular data set under investigation. This `composite index' will perform on par with the best individual variability indices in this data set, but it requires no a~priori knowledge of which indices are the best for the data set under investigation. We also find that in practice, all the discussed variability indices as well as their combinations are not sufficient on their own to automatically select variable stars from a large set of lightcurves. The reason is that both variable and non-variable stars are diverse groups: variables may have various lightcurve shapes, while non-variable stars include both the majority of objects displaying just noise and objects with photometry corrupted by nearby objects, cosmetic defects of a CCD, etc. The investigated indices cannot distinguish the badly measured objects from real variables because the corrupted measurements not only increase the lightcurve scatter (compared to a non-variable object of similar brightness), but may also mimic correlated variability (due to night-to-night seeing variations, drift of the object's image across a cosmetic defect and so on). If all causes of measurement corruption in a particular data set can be identified and all such cases flagged at the source extraction stage, the discussed variability indices may efficiently distinguish variable objects standing out among the majority of non-variable stars. At the same time, the variability indices are perfectly suitable to solve the inverse problem: identify well-measured constant stars in a photometric data set. The list of well-measured non-variable stars may be useful as photometric standards for calibration or targets for a search of variations not intrinsic to these objects such as microlensing events, occultations of stars by distant Solar system objects, etc. The data sets used to test the variability indices were searched for variable objects previously. Despite that, we were able to identify 124 new variable stars during the tests. This highlights the fact that variability search techniques originally used to investigate the data sets can be improved by the application of the multiple variability indices tested here. The information about the new variables is summarized in Table~\ref{tab:newvar} and their lightcurves are presented in Fig.~\ref{fig:newvarlightcurves}. The variability types are assigned according to the GCVS system\footnote{\url{http://www.sai.msu.su/gcvs/gcvs/iii/vartype.txt}} \citep{2009yCat....102025S} and high-amplitude $\delta$~Scuti/SX~Phoenicis stars are indicated as HADS. | 16 | 9 | 1609.01716 |
1609 | 1609.08771_arXiv.txt | We present the first VLBI detection of HCN molecular absorption in the nearby active galactic nucleus NGC~1052. Utilizing the 1 mas resolution achieved by the Korean VLBI Network, we have spatially resolved the HCN absorption against a double-sided nuclear jet structure. Two velocity features of HCN absorption are detected significantly at the radial velocity of 1656 and 1719 km~s$^{-1}$, redshifted by 149 and 212 km~s$^{-1}$ with respect to the systemic velocity of the galaxy. The column density of the HCN molecule is estimated to be $10^{15}$---$10^{16}$ cm$^{-2}$, assuming an excitation temperature of 100---230 K. The absorption features show high optical depth localized on the receding jet side, where the free-free absorption occurred due to the circumnuclear torus. The size of the foreground absorbing molecular gas is estimated to be on approximately one-parsec scales, which agrees well with the approximate size of the circumnuclear torus. HCN absorbing gas is likely to be several clumps smaller than 0.1 pc inside the circumnuclear torus. The redshifted velocities of the HCN absorption features imply that HCN absorbing gas traces ongoing infall motion inside the circumnuclear torus onto the central engine. | According to the most common active galactic nuclei (AGNs) model, i.e. the Unified Scheme of AGNs \citep{antonucci93}, AGNs consist of a supermassive black hole (SMBH) or a central engine with a rotating dense gas disk or torus surrounding the SMBH. It is generally accepted that pronounced activity in an AGN is generated by accreting gas from its torus onto a SMBH. The distribution and kinematics of the circumnuclear gas are keys for understanding the fueling of the AGN. High angular resolution studies of molecular gas in the center of the external galaxies (radius of $\sim$ 1 kpc) have been obtained with millimeter interferometers. The size of the circumnuclear torus, however, is smaller than 10 pc \citep[e.g.][]{garcia16}, and a milliarcsecond (mas) resolution is required to study its internal structure in nearby AGNs. While conventional millimeter interferometers, even ALMA, did not achieve such a high angular resolution, VLBI observations have revealed the parsec- or subparsec-scale morphology of nearby AGNs. Generally, thermal emission lines from molecular gas are not luminous enough to detect with the VLBI. VLBI maps, however, can display thermal absorption lines of the gas in silhouette against a bright background synchrotron radiation source with a mas resolution. NGC 1052 is a nearby elliptical galaxy with a systemic velocity ($V_{\rm sys} = cz$) of 1507 km~s$^{-1}$ \citep{jensen03}, implying a distance of 20~Mpc assuming $H_0= 75$ km~s$^{-1}$~Mpc$^{-1}$ and $q_0 = 0.5$. Its nuclear activity is classified as occurring within a low-ionization nuclear emission-line region \citep[LINER; e.g.][]{gabel00}. NGC 1052 has a nearly symmetric double-sided radio jet structure along the east--west direction from kiloparsec to parsec scales \citep[e.g.][]{jones84, wrobel84, kellermann98}. Past VLBI studies constrained the jet inclination angle to lie $\geq 57^{\circ}$ using the optically thin radio flux density of the twin jet \citep{kadler04, sss08, baczko16}. The parsec-scale jet structure shows a prominent gap between the eastern (approaching) and western (receding) jets at various centimeter wavelengths \citep[e.g.][]{claussen98, vermeulen03}. VLBI images at 43~GHz, however, show an innermost component in the gap \citep{kadler04, sss08}. \cite{kameno01} proposed the existence of a parsec-scale circumnuclear torus consisting of cold dense plasma from the convex radio continuum spectra due to the free-free absorption (FFA) by the foreground plasma. Besides the ionized gas, several atomic and molecular gases are also found toward the center of NGC 1052. An H$_2$O megamaser emission is detected at velocities by redshifted 50---350 km~s$^{-1}$ with respect to $V_{\rm sys}$ \citep{braatz94, braatz03, kameno05}, and the H$_2$O maser gas lies on the inner components in the center of NGC~1052, where the plasma torus is obscured \citep{claussen98, sss08}. Neutral atomic hydrogen (\ion{H}{1}), OH, HCO$^+$, HCN, CO transitions are found toward the center of NGC~1052 as absorptions \citep{vangorkom86, omar02, liszt04, impellizzeri08evn}. The velocity range over which those absorption lines are observed spans from 1500 to 1800 km~s$^{-1}$, redshifted with respect to $V_{\rm sys}$-like H$_2$O maser emission. These facts imply that ionized, neutral, and molecular gases coexist in the vicinity of the central engine. To investigate the geometry and physical properties of the molecular gas in the circumnuclear region of NGC~1052, we have performed high-resolution observations toward HCN (1--0) absorption lines with the Korean VLBI Network (KVN). Here we show the first parsec-scale HCN (1--0) maps of the center of NGC~1052. One milliarcsecond corresponds to 0.095 pc in the galaxy. | \begin{table*}[t] \begin{center} \caption{Column density of HCN (1--0) absorption.\label{tb:hcn}} \begin{tabular}{lcccccc} \tableline \tableline Label & $V_{\rm p}$ & $V_{\rm p} - V_{\rm sys}$ & $\Delta v$ & $\int \tau dv$ & $N_{\rm HCN,100}$ & $N_{\rm HCN,230}$ \\ & [km~s$^{-1}$] & [km~s$^{-1}$] & [km~s$^{-1}$] & [km~s$^{-1}$] & [10$^{14}$ cm$^{-2}$] & [10$^{14}$ cm$^{-2}$] \\ (1) & (2) & (3) & (4) & (5) & (6) & (7) \\ \tableline I & 1656 & 149 & 31.7 & $9.3\pm0.6$ & $9.5\pm0.6$ & $50\pm3$ \\ II & 1719 & 212 & 52.9 & $19\pm1$ & $20\pm1$ & $101\pm5$ \\ \tableline \end{tabular} \tablecomments{(1) Absorption feature ID, shown in Figure~\ref{spc}. (2) Peak velocity of the absorption feature. (3) Velocity with respect to $V_{\rm sys}$. (4) Line width. (5) Velocity-integrated optical depth. (6) HCN column density with $T_{\rm ex}=100$ K. (7) HCN column density with $T_{\rm ex}=230$ K.} \end{center} \end{table*} \subsection{Absorbing Molecular Gas Properties} The HCN opacity obtained from our observations was beyond 0.1 in the circumnuclear region of NGC~1052, and one order higher than that from the PdBI observations \citep{liszt04}. This indicates that the HCN covering factor is much smaller on the scale of a few hundreds parsecs. We estimated the mean opacity over the whole parsec-scale source structure $\left< \tau \right>$, using \begin{equation} \left< \tau \right> = \frac{\displaystyle \int \!\! \int \tau(x,y) I(x,y) \,dx \,dy}{\displaystyle \int \!\! \int I(x,y) \,dx \,dy} \end{equation} where $\tau(x,y)$ is the HCN opacity distribution shown in the color scale of Figure~\ref{hcnmap}, and $I(x,y)$ is the 89 GHz continuum image. The estimated $\left< \tau \right>$ is 0.027 and 0.028 for the absorption features I and II, respectively. They are on the same order of HCN opacity yielded from the normalized flux of 0.98 and 0.94 for the absorption features I and II observed with the PdBI \citep{liszt04}. Thus, the HCN covering factor is likely to be $\sim1$ on parsec scales. Under the assumption of the local thermodynamic equilibrium, and a covering factor of 1, the total molecular column density of the J=1--0 absorption line is derived as \begin{equation} N_{\rm tot} = \frac{3 k}{8 \pi^3 \mu^2 B} ~ \frac{(T_{\rm ex} + hB /3k)}{\bigl[1-\exp{\bigl( - h\nu/ kT_{\rm ex}\bigr) \bigr]}} \int \tau dv \label{eq:nhcn} \end{equation} where $k$ is the Boltzmann constant, $h$ is the Planck constant, $\mu$ is the permanent dipole moment of the molecule, $B$ is the rotational constant, $T_{\rm ex}$ is the excitation temperature, and $\int \tau dv$ is the velocity-integrated optical depth of the absorption feature. Using the equation~(\ref{eq:nhcn}), $\mu=2.98$ Debye, $B= 44315$ MHz for HCN, and assuming $\int \tau dv = \tau_{\rm max} \Delta v$, where $\tau_{\rm max}$ is the maximum of optical depth, we give the total column density of HCN for each absorption feature in Table~\ref{tb:hcn}. To determine the column density of molecular hydrogen (H$_2$), we assume a HCN-to-H$_2$ abundance ratio of 10$^{-9}$, from the HCN-to-H$_2$ abundance ratio estimates of (0.2---4.1)$\times10^{-9}$ of several HCN clumps in the Galactic Circumnuclear Disk \citep[CND;][]{smith14}. Here we take $T_{\rm ex}=100$ K and 230 K, as recent millimeter and submillimeter interferometric observations have detected the vibrationally excited HCN emission lines in the dust-enshrouded AGN of the luminous infrared galaxies \citep{sakamoto10, imanishi13}. The derived H$_2$ column density ($N_{\rm H_2}$) ranges $10^{24}$---$10^{25}$ cm$^{-2}$, using the values of $N_{\rm HCN}$ in Table~\ref{tb:hcn}. Hence, equivalent $N_{\rm H}$ would be of the same order as $N_{\rm H_2}$. It is one or two orders higher than the estimate of the $N_{\rm H} \sim 10^{23}$ cm$^{-2}$ from modeling of X-ray spectra \citep{guainazzi99, weaver99, guainazzi00, brenneman09}, and the electron column density of $\sim10^{23}$ cm$^{-2}$ estimated from the FFA opacity of the dense plasma \citep{sss09}. \subsection{Circumnuclear torus model} \begin{figure*}[th] \epsscale{0.9} \plotone{./fig3.eps} \caption{(a) Possible model of the oriented double-sided jet and the circumnuclear torus. The near side of the torus covers the receding jet side. (b) Schematic diagram for the intersection of the circumnuclear torus in NGC~1052. The torus has two-phase layers on the inner surface, the X-ray heated plasma region and the X-ray dissociation region (XDR). HCN absorbing gas is located in the cooler molecule region next to XDR. HCN absorbing gas could be clumpy, and the line of sight passes through several clumpy clouds with a different radial velocity. \label{model}} \end{figure*} \cite{kameno05} and \cite{sss08} proposed that the circumnuclear torus consists of several phase layers. On the inner surface of the torus, a hot ($\sim8000$ K) plasma layer is formed by X-ray emission from the central engine and it causes FFA. A heated (above $\sim400$ K) molecular layer or X-ray dissociation region \citep[XDR;][]{neufeld94, neufeld95} lies next to the plasma layer, and the H$_2$O megamaser emission arises from here. As well as FFA opacity distribution, the HCN opacity shows high values on the western receding side of the jet. Thus, HCN absorbing gas could be associated with the torus, like the dense plasma. If $T_{\rm ex}$ of HCN is $\sim230$ K, HCN molecules should lie in the cooler molecule layer next to XDR. The redshifted HCN absorption against the continuum emission is likely indicative of ongoing gas infall. The velocities of the HCN absorption features are close to the centroid velocity of the broad H$_2$O maser emission ($\sim1700$ km~s$^{-1}$). Therefore, H$_2$O and HCN could trace the same infall motion inside the torus. The HCN absorption spectrum consists of at least two narrow absorption features, and HCN absorbing gas is more likely to be several small ($\le$ 0.1 pc) clumps or layers with a different velocity, rather than a large homogeneous structure with a single velocity. We have seen the opacity contrast between the absorption features I and II, and the contrast also suggests the inhomogeneity inside the molecular torus. In Figure~\ref{model}, we present a possible model for the geometry of the circumnuclear torus and the jet in NGC 1052. Since the jet axis is inclined from the sky plane, the near side of the torus should cover the receding jet side. The line of sight intersects at least two HCN gas clumps inside the torus. As we have mentioned in Section \ref{sec:results}, there is no significant detection of absorption features around 1500---1600 km~s$^{-1}$, close to $V_{\rm sys}$. A possible explanation is that the missing absorption features could arise not from the compact molecular clumps in the circumnuclear region, but from the foreground diffuse interstellar medium in the host galaxy. If so, their covering factor would not vary on scales between a few parsecs and a few hundreds parsecs, and thus the depth of the absorption features would be $\sim -3\%$, the same as the PdBI results. It is comparable to the rms noise level of our spectral profile with the KVN, and no significant detection comes as a consequence. Furthermore, when their background continuum sources have a faint and extended structure, the background continuum emissions should be fully resolved-out. Therefore, the absorption features against the resolved-out background sources would be invisible in our KVN data. We can estimate an approximate size of the foreground absorbing molecular gas using the relation $N_{\rm H_2} = f_v n_{\rm H_2} L $, where $f_v$ is the volume filling factor, $n_{\rm H_2}$ is the molecular hydrogen volume density, and $L$ is the size of the molecular gas. Here we simply assume $f_v=$ 0.01. Adopting $N_{\rm H_2}$ of $10^{24}$--$10^{25}$ cm$^{-2}$ from our results and $n_{\rm H_2}$=(0.1---2)$\times10^6$ cm$^{-3}$ in the Galactic CND \citep{smith14}, the resulting $L$ is approximately on 1-pc scales. It is consistent with the approximate size of the circumnuclear torus of NGC~1052 \citep{kameno01}. On the other hand, the size of the innermost receding jet component can be estimated as 0.2---0.4 mas from the VLBA images at 43 GHz \citep{kadler04, sss08}, and it corresponds to 0.02---0.04 pc in linear scale. Based on our measurements of the radial velocities and the distribution of HCN absorption features, the mass infall rate of the gas accretion toward the central engine can be estimated by $\dot{M} = f_v R_{\rm in}^2 \rho V_{\rm in} \Omega$, where $f_v$ is the volume filling factor, $R_{\rm in}$ is an infall radius, $\rho$ is the mass density of the infalling materials, $V_{\rm in}$ is the infall velocity, and $\Omega$ is the solid angle of the torus from the center. Here we assume $\rho = N_{\rm H} m_{\rm H} / R_{\rm in}$, where $N_{\rm H}$ is the column density of a hydrogen atom and $m_{\rm H}$ is the mass of a hydrogen atom. Giving $R_{\rm in}=$ 1 pc, $N_{\rm H}=$ $10^{24}$---$10^{25}$ cm$^{-2}$, and $V_{\rm in}=$ 200 km~s$^{-1}$ as the approximate velocity of the HCN absorption feature with respect to $V_{\rm sys}$ ($V_{\rm p}-V_{\rm sys}$ in Table~\ref{tb:hcn}), the derived mass infall rate ranges $\dot{M} =$ (47---470) $f_v (\Omega / 4\pi) M_{\sun}$ yr$^{-1}$. If ($\Omega/ 4 \pi$) takes a few 0.1 and $f_v$ is 0.01, it would be comparable to the calculation of the accretion rate of $10^{-1.39} M_{\sun}$ yr$^{-1}$ from the hard X-ray luminosities by \cite{wu06}. We note that the X-ray luminosity indicates an instantaneous accretion rate. However, our estimation of the accretion rate in the molecular torus gives a long-term ($R_{\rm in}/V_{\rm in} \sim 5000$ year) activity. The coincidence of the two accretion rate values suggests the continuity of AGN activity in NGC~1052. | 16 | 9 | 1609.08771 |
1609 | 1609.07145.txt | %@abs As part of our program to build a complete radio and X-ray database of all the 3CR extragalactic radio sources, we present an analysis of 93 sources for which \chn\ archival data are available. Most of these sources have been already published. Here we provide a uniform re-analysis and present nuclear X-ray fluxes and X-ray emission associated with radio jet knots and hotspots using both publicly available radio images and new radio images that have been constructed from data available in the VLA archive. For about 1/3 of the sources in the selected sample a comparison between the \chn\ and the radio observations was not reported in the literature: we find X-ray detections of 2 new radio jet knots and 17 hotspots. We also report the X-ray detection of extended emission from the intergalactic medium of 15 galaxy clusters, two of which were most likely unknown previously. | %@intro \label{sec:intro} The first release of the Third Cambridge catalog (3C), performed at 159 MHz, was published in 1959 \citep{edge59}. In 1962 Bennett et al. revised the whole 3C catalog using observations at 178 MHz and this revised version (3CR) was considered a definitive list of the brightest radio sources in the Northern Hemisphere for many years. The flux limit of the 3CR catalog is set to 9 Jy at 178 MHz and it covers the whole Northern Hemisphere above -5\degr\ in Declination. Then, in 1985, Spinrad, Djorgovski, Marr and Aguilar presented the last revised version of the Third Cambridge catalog (3CR) \citep{bennett62} listing 298 extragalactic radio sources \citep[see also][]{edge59,mackay71,smith76,smith80} including new revised positions, redshifts and magnitudes having 91\% of the sources out of the Galactic plane (i.e., Galactic latitude $|b|>$10\degr). Since then several photometric and spectroscopic surveys have been carried out to obtain multifrequency coverage of the 3CR catalog. All the 3CR sources at redshift $z<$0.3 have been already observed with the Hubble Space Telescope (HST) \citep[e.g.,][]{chiaberge00,tremblay09} while a near infrared, optical and ultraviolet survey for higher redshift sources is still ongoing. A large fraction of the 3CR radio sources were also targets of the spectroscopic survey carried out with the Telescopio Nazionale Galileo \citep[TNG; e.g.,][]{buttiglione09}. Radio images with arcsecond resolution for the majority of the 3CR sources are available from the NRAO Very Large Array (VLA) Archive Survey (NVAS)\footnote{\underline{http://archive.nrao.edu/nvas/}} and from the MERLIN archive\footnote{\underline{http://www.jb.man.ac.uk/cgi-bin/merlin\_retrieve.pl}}. As a radio low frequency catalog, the selection criteria for the 3CR are unbiased with respect to X--rays. Since it spans a wide range of redshift and radio power and has a vast multifrequency database of ground and spaced based observations for comparison, it is an ideal sample to investigate properties of active galaxies. Motivated by the large number of multifrequency observations already available for the 3CR sources, we have undertaken a project to ensure that each 3CR extragalactic source has at least an exploratory/snapshot \chn\ observation. We have chosen to achieve this goal in a step wise strategy, working out in redshift with modest proposals each cycle to minimize the impact on the \chn\ schedule. A description of our progress in this endeavor is given in the following sections. In this paper we present the X-ray analyses of most of the 3CR sources present in the \chn\ archive which have not already been published by us with our standard procedures: i.e. the snapshot surveys \citep{massaro10,massaro12,massaro13,massaro15} and the 3CR sources in the XJET sample \citep{massaro11}. Our main aim is to provide a uniform analysis for all the archival observations. X-ray flux maps were constructed and compared with radio images to search for any X-ray emission associated with radio jet knots, hotspots and lobes. In some cases new radio images have been been constructed from archival VLA data for comparison with the X-ray images. We report the measurements of the X-ray nuclear emission for all sources in our sample, but we did not perform a detailed spectral analysis because most of them (i.e., $>$70\%) have already been reported in the literature \citep[see e.g.,][]{hardcastle09,balmaverde12,wilkes13,kuraszkiewicz15}. The paper is organized as follows. A brief historical overview of the \chn\ observations of the 3CR sources is provided in \S~\ref{sec:history} while the description of the selected sample is presented in \S~\ref{sec:sample}. Data reduction procedures are given in \S~\ref{sec:obs} while results are discussed in \S~\ref{sec:results}. Then, \S~\ref{sec:summary} is devoted to our summary and conclusions. Finally, in the Appendix, we show the X-ray images with radio contours superposed for all the sources analyzed (\S~\ref{sec:images}) and a summary of the \chn\ observations for the entire sample of 3CR extragalactic sources (\S~\ref{sec:state}). For numerical results, cgs units are used unless stated otherwise and a flat cosmology was assumed with $H_0=72$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{M}=0.27$ and $\Omega_{\Lambda}=0.73$ \citep{dunkley09}, to be consistent with our previous analyses \citep[e.g.,][]{massaro10,massaro12,massaro13}. Spectral indices, $\alpha$, are defined by flux density, S$_{\nu}\propto\nu^{-\alpha}$. | %@summary \label{sec:summary} We have described the combined radio-X-ray analyses of 93 3CR radio sources for which \chn\ observations requested by others for many different reasons, were already present in the archive. The main objectives of the present analysis are: (1) to present a uniform X-ray and radio database for the 3CR catalog, (2) to search for possible detections of X-ray emission from radio jet knots, hotspots and lobes and (3) to look for new galaxy cluster detections surrounding the 3CR radio sources. In order to perform the radio--X-ray comparison we reduced archival radio observations for 6 sources. We focused on the comparison between the radio and the X-ray emission from extended components such as radio jet knots, hotspots, and lobes. We discovered 2 new radio jet knots and 17 hotspots emitting in the X-rays. Flux maps for all the X-ray observations were constructed and we provided photometric results for all the extended components detected. All the radio knots and hotspots have been classified on the basis of the radio morphology of their parent source, adopting the definition suggested by Leahy et al. (1997) for the hotspots, i.e., brightness peaks which are neither the core nor a part of the jet, usually lying where the jet terminates, and considering all other discrete brightness enhancements as jet knots. The following conventions for labeling the extended structures detected in the X-rays was adopted. We indicated with the letter `k' the jet knots and with `h' the hotspots; then the name of each component is a combination of one letter (indicating the cardinal direction of the radio feature with respect to the nucleus) and one number (indicating the distance from the core in arcsec) as described in Massaro et al. (2011). We also reported the presence of 15 X-ray galaxy clusters associated with the selected 3CR source, 13 already known in the X-rays and 2, namely 3CR\,427.1 and 3CR\,449, reported here for the first time to the best of our knowledge. In the Appendix we present X-ray images with radio contours for all the 93 sources analyzed in this paper (\S~\ref{sec:images}) and in (\S~\ref{sec:state}) we give the \chn\ status of the observations for all extragalactic 3CR sources. | 16 | 9 | 1609.07145 |
1609 | 1609.09361_arXiv.txt | We calculate the structure and evolution of a gamma ray burst central engine where an accreting torus has formed around the newly born black hole. We study the general relativistic, MHD models and we self-consistently incorporate the nuclear equation of state. The latter accounts for the degeneracy of relativistic electrons, protons, and neutrons, and is used in the dynamical simulation, instead of a standard polytropic $\gamma$-law. The EOS provides the conditions for the nuclear pressure in the function of density and temperature, which evolve with time according to the conservative MHD scheme. We analyze the structure of the torus and outflowing winds, and compute the neutrino flux emitted through the nuclear reactions balance in the dense and hot matter. We also estimate the rate of transfer of the black hole rotational energy to the bipolar jets. Finally, we elaborate on the nucleosynthesis of heavy elements in the accretion flow and the wind, through computations of the thermonuclear reaction network. We discuss the possible signatures of the radioactive elements decay in the accretion flow. We suggest that further detailed modeling of the accretion flow in GRB engine, together with its microphysics, may be a valuable tool to constrain the black hole mass and spin. It can be complementary to the gravitational wave analysis, if the waves are detected with an electromagnetic counterpart. | Gamma Ray Bursts (GRB), are extremely energetic transient events, visible from the most distant parts of the Universe \citep{Piran04,Kumar15}. When a newly born black hole forms, either via the core collapse supernova, or after the merger of two compact stars, a large amount of matter is accreted onto it within the timescale between tens of milliseconds, to thousands of seconds, with hyper-Eddington rates \citep{woosley93,paczynski98}. In such conditions, the nuclear reactions inside the plasma lead to creation of neutrinos, mainly via the electron-positron capture on nucleons, which acts in the so-called NDAF models \citep{Narayan00,Lei09,Cao14}. The annihilation of neutrino-antineutrino pairs is a source of power to the relativistic jets launched along the axis of the disk-black hole system. This can be comparable to the power extracted on the cost of the black hole spin, with the disk mediated by the magnetic fields. Both neutrinos and magnetic fields in the accreting matter can drive the wind outflow from the disk surface, and act as a collimation mechanism for the relativistic jets (see, e.g., \citealt{LeeRamirez2007} and references therein). The latter, are then the sites of the high-energy radiation produced at large distances from the engine, on the cost of the magnetic reconnection and conversion of the jet electromagnetic energy flux into heat (e.g., \citealt{Bromberg2016}). Numerical computations of the structure and evolution of the accretion flow in the gamma ray bursts engine begun with steady-state, and time-dependent models, which were axially and vertically averaged, and based on the classical prescription for viscosity with the so-called $\alpha$-parameter \citep{popham1999, dimatteo2002, kohri2002, kohri2005, reynoso2006, chen, janiuk2004, janiuk2007, janiuk2010}. More recently, these accretion flows are being described by fully relativistic, MHD computations \citep{Shibata2007, nagataki09, barkovkomis2010, barkov2011, rezzola2011, janiuk13, kyuchi2015}. In the outer parts of the disc, the heavy elements can be formed, and then the isotopes are prone to the radioactive decay \citep{Fujimoto2004}. This may in the future bring observable effects on the measured hard X-ray spectra of GRBs \citep{aj2014} as well as the additional peaks in the blue and infrared bands a few days after the prompt GRB emission \citep{martin2015}. In addition, the gamma ray bursts contribute in this way to the galactic chemical evolution \citep{surman2014}. The physical conditions in this engine, its density, temperature, and electron fraction, depend on the global parameters of the flow, such as the amount of matter and accretion rate, black hole mass and its spin, and the magnetic field. These conditions in turn will affect the total abundances of heavier elements synthesized within the flow. Therefore, a possible detection of the signatures of the radioactive decay of these species, and quantitative modeling of nucleosynthesis together with the central engine structure and evolution, and the jets gamma ray fluence, could give a constraint on the black hole parameters. These, in turn, can now be independently estimated, if the black hole is born via a merger event, detectable in the gravitational wave interferometer. In this work, we study the central engine of a gamma ray burst, which is composed of a stellar mass, rotating black hole and accreting torus that has formed from the remnant matter at the base of the GRB jet. We compute the time-dependent, two dimensional model of the rotationally supported, magnetized disc, rapidly accreting onto the center. The simulation is based on the HARM (High Accuracy Magnetohydrodynamics) scheme, which works on the stationary metric around black hole. The code integrates the total energy equation and updates the set of 'conserved' variables, i.e. comoving density, energy-momentum, and magnetic field \citep{gammie}. The basic version of the HARM scheme was designed to model the accretion flows in the centers of galaxies, where the gas equation of state is with a good accuracy given by that of an ideal gas, and can be described using an adiabatic relation (e.g., \citealt{mmosc2013}). In our current dynamical model, we use the non-ideal equation of state, computed on the basis of the $\beta$-equilibrium \citep{janiuk2007}. We allow for the partial degeneracy of relativistic nucleons, electrons and positrons in the accreting plasma, and we compute the neutrino cooling rate due to the weak interaction processes, electron-positron pair anihillation, as well as bremsstrahlung and plasmon decay. Here, we for the first time incorporate this detailed microphysics into the heart of the GR MHD scheme. The neutrino cooling was computed already from the nuclear reaction balance in the context of gamma ray bursts in our previous work \citep{janiuk13}. However, in that work, still the simple, adiabatic relation for the pressure ($p=(\gamma-1)u$, with $\gamma=4/3$) was used, and only the internal energy of the gas was updated during the simulation. In the current model, we mainly address the question of microphysics in the accretion flow, and we are now also able to self-consistently estimate the efficiency of the disk neutrino cooling. We check if it is capable of powering the relativistic jets in the gamma ray bursts, now in the case of a rather weak magnetic field. We also check the effect of the rotation speed of the black hole. In addition, we analyze the structure, velocity and amount of mass loss through the uncollimated winds launched from the accretion disc. Finally, we examine the resulting nucleosynthesis in the body of the disk, which can occur already at small distances from the black hole, as well as in the winds, and we show that the accretion flow can be the site of significant production of heavy element. The article is organized as follows. In \S~\ref{sec:model}, we describe our MHD model of the GRB central engine. In \S~\ref{sec:results}, we present the results, describing the structure of the torus and the effects of adopted microphysics treatment. We also compute the power transferred to the GRB jets via the neutrino anihillation and compare it to the energy extracted by the magnetic fields dragged through the BH horizon. In \S~\ref{sec:nucleo} we describe the process of heavy elements formation in the outskirts of the torus in the GRB engine. We present the resulting abundances of elements computed under the assumption of nuclear statistical equilibrium. Finally, we compare our results with the previous simulations. We discuss the results in \S~\ref{sec:diss}. | \label{sec:diss} We calculated the structure and short-term evolution of a gamma ray burst central engine in the form of a magnetized torus accreting onto a black hole. Our computations were performed in a fixed background metric around a spinning black hole, and the physical conditions in the plasma driven by the hyper-Eddington accretion rate were taken into account by the adopted microphysics. We describe the flow properties using the numerical equation of state of dense and hot matter, where the free nucleons and electron-positron pairs are relativistic and partially degenerate. We estimate the resulting neutrino luminosity, which cools the flow via the nuclear reactions in which neutrinos of the three flavors are produced. We conclude that the total integrated luminosity is larger than the power possibly extracted from the rotating black hole via the Blandford-Znajek process. Taking into account the uncertainty of the efficiency of the neutrino anihillation in the bi-polar jets \citep{zalamea2011}, we still find that the power supply to the jets by neutrinos is plausible and energetically competitive to the magnetic extraction of the black hole rotation energy. The discussion of interplay between the Blandford-Znajek mechanism versus neutrino power supply to the GRB jets has been proposed in a number of works based on the so-called NDAF disk, e.g., \citet{liu2015}. These conclusions are however far from robust, regarding the very simplified treatment of accretion flows in these models, which is described by a classical 1-dimensional steady state accretion flow with $\alpha$ viscosity, and pseudo-Newtonian gravity. Our computations realistically treat the nuclear equation of state, together with other relevant physical ingredients, like the general relativity and magnetic fields. The latter is still lacking in some of the otherwise advanced, three-dimensional dynamical simulations of neutrino-cooled accretion flows in the short GRB engines \citep{Dessart2009, Perego2014}. The essentially inviscid nature of such models might lead to their somewhat underestimated results for the accretion speed, dissipation and cooling rates. On the other hand, the relativistic models which are based on the ideal gas EOS (e.g., \citet{rezzola2011} use the adiabatic EOS with index of 2), can predict overestimated profiles of the disk scaleheight. Also, as noted by \citet{kyuchi2015}, who use the gamma index of 1.8 in their accretion flow computations, the value of this index will affect the amount of matter ejected in the torus wind. By accounting for the proper microphysics, instead of the simplistic EOS, we can get therefore most reliable description of the conditions in neutrino driven winds. The exact determination of the accretion disk thickness is crucial with respect to the possible gravitational or thermal instabilities \citep{Perna2006, janiuk2007}, which can lead to the episodic accretion events and hence observable flares. As it was shown by \citet{BegelmanPringle2007}, the strong magnetic fields can be important for the stabilisation of the accretion flow against self-gravity. The ordered magnetic field configuration, which is mostly toroidal in the accretion torus with but predominantly poloidal and jet-like along the BH spin axis, is produced in simulation of \citet{rezzola2011} at time about 6 ms after the black hole formation. In our simulations the poloidal field lines appear after about 15 ms for the black hole of a 3 $M_{\odot}$, so the results are consistent within an order of magnitude. The difference in the timescale may be connected to some extent with the interplay between magnetic field dynamics and gas properties, but also affected by the 2-dimensional setup which is used in our model. We plan to expand the current computations into 3-dimensions in the future work, and then test quantitatively the resulting models against the observables of GRBs, such as their time variability, duration, and energetics. The ultimate impact of such simulations can also lead to the constraints on the magnetic fields strengths and configuration in the cosmic environment related to GRB progenitors. Finally, based on the computed structure of the flow, its density, temperature, and electron fraction, we are able to estimate the integrated abundances of heavy elements which are synthetized in the flow via the subsequent nuclear reaction chain. These elements, if decaying radioactively at some distance from the central engine, will then contribute to the heating of the circum-burst medium, and afterglow emission. As shown by \citet{sekiguchi2015}, the dynamical ejecta including neutrino transport may develop the shock waves, and related heating is responsible for enhanced reaction rates. The electron fraction is then increased, and their nuclear reaction networks are able to produce the r-process elements up to A=195. This was however not encountered in the Newtonian simulations, which produce elements up to A=130 \citep{goriely2011}. Neutron star mergers are therefore expected to produce significant quantities of neutron-rich radioactive species, whose decay should result in a faint transient emission. The latter is called a `kilonova', and should appear in the days following the burst. The term refers to the luminosity peak in the near infrared band, rather than UV or Optical \citep{lipaczynski1998}, and the emission produced in the regions of the opacities much higher than in the case of of supernovae. An observational support for such event has been reported e.g. by \citet{tanvir2013} for a short duration GRB 130603B. Recently, re-examination of the observed long burst emission in GRB 060614 has shown that it is consistent with the 'kilonova' scenario, rather than an undelying supernova \citep{Yang2015}. The gamma ray bursts can play a unique role in the metal enrichment of the Universe, as they are the brightest sources at all redshifts, and, for long GRBs, they occur in the star forming regions. Because the nucleosynthesis yields from the explosions of stars in different populations could differ by orders of magnitude, the metal abundance patterns, imprinted in the interstellar medium, should depend on the initial mass function \citep{heger}. The abundance patterns measured with distant GRBs can be used to determine the typical masses of the early stars (Pop III), and disentangle them from the younger Pop II SNe. Such studies are planned e.g. with the newly developing instruments, e.g., on-board the Athena X-ray satellite, which is aiming to probe the GRB afterglow spectra, in combination with the studies of the quasars on the line of sight \citep{jonker2013}. The GRBs can be observed up to the highest redshifts (e.g., z=9.4, for GRB 090429, \citet{Cucchiara2011}) and hence the old stellar populations can be traced in their X-ray afterglow spectra, e.g., with absorption lines of Fe, Mg, Si, S, Ar, from the ionized gas in the GRB environment. Nevertheless, these data can be contaminated by the effects of the intergalatic medium and hence the observed absorption column is a combined effect of the warm-hot intergalactic medium, Lyman alpha clouds, as well as the circumburst medium intrinsic to the host galaxy \citep{Starling2013}. From the point of view of stellar evolution models, \citet{heger} computed the nucleosynthesis yields of elements produced in the massive metal-free stars, but they ignored both the neutrino-powered winds and the gamma-ray bursts accretion disks. The latter, as studied by \citet{Surman2006}, have large overproduction factors in their outflows for the nuclei like $^{44}Ti$, $^{45}Sc$, and $^{64}Zn$. Also, the light p-nuclei, such as $^{92}Mo$, and $^{94}Mo$, are produced in the rapidly accreting disks, as found by \citet{Fujimoto2003}). The neutrino-driven winds, on the other hand, were modeled recently by \citet{Perego2014} in the context of binary neutron star mergers. They found that the wind can successfully contribute to the weak r-process in the range of atomic masses from 70 to 110. Obviously, the long GRBs are only a fraction of the massive star explosions and luminous supernovae \citep{Guetta2007, HjorthBloom2012}, and their impact for the majority of elements may not be crucial. However, verifying the signatures of these elements which are produced mainly in the GRB engines, may in the future occur to be a prospective tool to study the chemical evolution in the Universe due to the collapsing massive stars, via the available GRB afterglow X-ray spectra. Also, some of the bursts may be connected with peculiar type of events, such as the burst GRB 111209A/SN 2011kl \citep{Kann2016}. In such case, the conditions in GRB central engine, determined during the explosion, can also lead to peculiar distributions of electron fraction, and affect the synthesis process for neutron-rich isotopes. \begin{table*} \begin{center} \caption[]{Summary of the models with a massive black hole, $M_{\rm BH}=62 M_{\odot}$, as estimated for the first LIGO-detected black hole. The mass of the torus is given in $M_{\odot}$, accretion rate in $M_{\odot}$s$^{-1}$, the time is in seconds, and luminosity in erg s$^{-1}$. \label{table:models_lbh} } \begin{tabular}{cccccccccc} \hline Model & $a$ & $R_{\rm max}$ & $t_{\rm e}$ & $M_{\rm t}(t_{0})$ & $M_{\rm t}(t_{\rm e})$ & $<\dot M>$ & $\dot M(t_{\rm e})$ & $L^{\rm tot}_{\nu}(t_{\rm e})$ & $L_{\rm BZ}(t_{\rm e})$ \\ \hline lbhAc & 0.6 & 9.1 & 0.91 & 16.8 & 14.36 & 5.01 & 4.06 & $4.80 \times 10^{53}$ & 0 \\%361_cool_a06 lbhBc & 0.7 & 8.4 & 0.91 & 13.0 & 11.17 & 4.64 & 1.99 & $6.18 \times 10^{53}$ & 0 \\%361_cool_a07 lbhCc & 0.8& 7.5 & 0.92 & 10.6 & 9.02 & 4.07 & 3.59 & $6.60 \times 10^{53}$ & $ 1.93 \times 10^{51}$ \\%376 lbhDc & 0.9& 7.0 & 0.91 & 10.4 & 8.89 & 3.98 & 2.55 & $9.72 \times 10^{53}$ & $ 3.30 \times 10^{51}$ \\%376 lbhEc & 0.98& 6.6 & 0.91 & 16.4 & 15.23 & 2.93 & 1.85 & $5.64 \times 10^{54}$ & $3.31 \times 10^{52}$ \\%376 \hline \end{tabular} \end{center} \end{table*} The synthesized heavy elements can also be the source of a faint emission of a kind of an 'orphan' afterglow, if the main GRB jet is directed off-axis, or it is intrinsically very weak in its gamma ray fluence. Still, the detection of such an afterglow emission will be a hint of the existence of a black hole in the engine site. Measurements of nucleosynthesis signatures may therefore help constrain the model parameters for such engine, namely the black hole spin and its mass. The recent detections of the gravitational wave signal \citep{Abbott2016} confirmed independently the existence of black holes in the Universe. The analysis of the gravitational wave signal allows to constrain the black hole mass and spin (with some uncertainty). In the case of the source GW150914, it was found that the signal is related to the merger of a binary black hole. So far, the electromagnetic counterpart for this first event was tentatively reported only by Fermi satellite, who gave a limit for detection of a weak gamma ray signal \citep{Fermi2016}. No infrared emission was reported, and also the signal was not confirmed by other gamma-ray missions for the case of this burst \citep{greiner}. Nevertheless, the detection of the electromagnetic counterpart of a gravitational wave source in the future events is appealing. The related estimates for the heavy elements synthesized in the ejecta could help constrain the parameters for the GRB engine, if the latter happens to coincide with the gravitational wave source. In particular, if the binary black holes merge within a remnant circumbinary disk, which is a leftover after the past supernova explosion \citep{perna2016}, then the nucleosynthesis of elements within the disk engine and ejected winds can also be responsible for an infrared counterpart. Direct measurements of the escaping gamma-rays, which will be produced at larger distance in the r-process nucleosynthesis, are probably difficult with current X-ray and gamma-ray missions \citep{Hotokezaka2016}. The black hole, which was the product of the merger of two black holes of about 29 and 36 Solar mass black holes, was also a moderately spinning one. The constraints for its dimensionless spin parameter obtained from the amplitude and phase evolution of the observed gravitational waveform gave a value of $0.67^{+0.05}_{-0.07}$, while the final black hole mass was equal to $62^{+4}_{-4}\ M_\odot$. As can be shown, such a binary black hole merger, if accompanied by the matter accretion onto a rotating black hole, which happens either before, or after the merger, can lead to an unusual gamma ray burst \citep{janiuk2013b,janiuk2016}. In the case of the relatively small rate of the black hole rotation, such as in GW150914, the accretion will result in a very small electromagnetic flux dragged through the black hole horizon and hence a negligibly small power of the Blandford-Znajek process. We computed an additional set of the gamma ray burst engine models, with parameters taken to be representative for the above scenario (see Table \ref{table:models_lbh}). \begin{figure} \includegraphics[width=8cm]{Fig_abund_lbh_time2000.ps} \caption{Heavy elements synthesized in the torus and its wind, for the model lbhBc and lbhDc in the Table 2. The black hole mass is $M_{\rm BH} = 62 M_{\odot}$, and its spin is a=0.9 (blue), or a=0.7 (red). The torus mass is about 15 $M_{\odot}$ and results were taken at time $t=2000 M$ in the dynamical simulation.} \label{fig:nse2Dm62} \end{figure} We have also computed the mass fraction distribution of heavy elements produced in this putative engine (Figure \ref{fig:nse2Dm62}). We found, that contrary to the presented in the previous Section 'fiducial' models, with a standard (small) black hole mass and accretion torus, here the nuclear conditions preferably lead to production of mainly Iron peak elements. In the model \textit{lbhD} (see Table \ref{table:models_lbh}), the mass fraction of Iron 56 reached almost 1.0, while the isotopes from the group of Selenium and Manganese were $10^{3}$ times less abundant. This result suggests a possible additional test for the conditions present in the GRB engine, e.g., if the event happened to be coincident with a merger of the two $\sim 30 M_{\odot}$ black holes with masses of about $\sim 30 M_{\odot}$ each. It is also worth to notice that the supernova explosion which might have left a 30 Solar mass black hole, should have originated from at least 80 Solar mass star on zero-age main sequence, in a low-metallicity environment \citep{Spera2015}. Even if a significant amount of the star's mass was ejected during the evolution and explosion, some remnant disk is plausible, but detailed modeling of its structure should be also confronted with the supernova theory. More generally, the gravitational waves discovery, GW150914, may reveal the new population of weak gamma ray bursts originating from the low Lorentz factor jets. In this case, the prompt electromagnetic emission may be difficult to discover, but the observation of a radioactive decay of elements produced in the engine of the GRB may be confronted with the engine simulation predictions. The black hole parameters, namely its spin and mass, determine the shape of the gravitational wave signal and hence can be independently constrained. As shown in Figure \ref{fig:nse2Dm62}, the results for nucleosynthesis yields somewhat depend on the black hole spin value. The abundance of isotopes produced in the accreting torus around a faster spinning black hole, is about 10 times higher at about $A=80$ (Selenium group) than for a slowly spinning black hole. Thus, the predictive power of our models is worth exploring in the future work, if only the accreting black hole parameters are determined and the nucleosynthesis yields constrained with some future observations, i.e. X-ray probes. | 16 | 9 | 1609.09361 |
1609 | 1609.01300_arXiv.txt | One aim of cosmic ray measurements is the search for possible signatures of annihilating or decaying dark matter. The so-called positron excess has attracted a lot of attention in this context. On the other hand it has been proposed that the data might challenge the established diffusion model for cosmic ray propagation. We investigate an inhomogeneous diffusion model by solving the corresponding equations analytically. Depending on the propagation parameters we find that the spectral features of the positron spectrum are affected significantly. We also discuss the influence of the inhomogeneity on hadronic spectra. | \label{chap:introduction} In many particle physics models, dark matter can annihilate or decay into Standard Model particles, which would then propagate in the galaxy and enhance the amount of observed cosmic rays. A major issue for testing this hypothesis is a precise determination of the astrophysical background. Since antiparticles are mainly produced as secondary particles by the spallation of cosmic rays on the interstellar medium (ISM), an estimate of the background contribution is possible. Before reaching the Earth, the particles propagate in the Milky Way scattering off the magnetic field, eventually suffering from energy losses, annihilation or escape from the galactic halo. It turns out that the cosmic ray spectra are strongly affected by these processes. An appropriate description for cosmic ray transport is provided by diffusion models \cite{Ginzburg:1990sk,Strong:1998pw,Moskalenko:1997gh,Maurin:2001sj}. However, cosmic ray propagation is far from being completely understood. It gives rise to the main uncertainties in the derivation of the background and thus demands for more realistic and complex models. In this light the measured positron flux received significant attention as it exceeds the theoretical predictions for the astrophysical background. This observation was first made by the PAMELA collaboration in 2008 \cite{Adriani:2008zr}, which discovered a rise in the positron fraction $e^+/(e^-+e^+)$. Interpretations of this excess consider additional primary sources like dark matter \cite{Bergstrom:2008gr,Cirelli:2008jk}, pulsars \cite{Hooper:2008kg, Linden:2013mqa} or supernova remnants \cite{Fujita:2009wk, Kohri:2015mga}, but the processes responsible for a possible production of positrons are speculative and poorly understood. At the experimental side the situation has further improved with the recent AMS-02 data \cite{Aguilar:2014mma} which resulted in further dark matter studies \cite{Bergstrom:2013jra,Kopp:2013eka,Ibarra:2013zia} and background parameterizations including pulsars \cite{Gaggero:2013nfa,DiMauro:2014iia,DiMauro:2015jxa}. See also \cite{Serpico:2011wg} for a review about the positron excess. An alternative idea was proposed in \cite{Katz:2009yd, Blum:2013zsa, Israel:2014mja, Cowsik:2013woa, Blasi:2009hv, Tomassetti:2015mha, Dado:2015eda, Ahlen:2014ica, Lipari:2016vqk}, stating that a modified diffusion model could explain the observed positron signal from a purely secondary origin. An interesting inspiration for this assumption is presented in \cite{Katz:2009yd, Lipari:2016vqk}. Studying the spectrum of cosmic antiprotons which are expected to have (mainly) a secondary origin \cite{Giesen:2015ufa,Evoli:2015vaa,Kappl:2015bqa} an upper bound on the secondary positron spectrum is estimated which agrees with the data. This hypothesis is encouraged by recent AMS-02 data on the antiproton to positron ratio \cite{Aguilar:2016kjl}. In a model where energy losses are less substantial, for example if the propagation time is reduced, the theoretical prediction for the spectrum would be flatter. This condition can be realized if the diffusion coefficient is taken to increase with galactic height. This approach is motivated by the spatial distribution of the galactic magnetic field which is responsible for the diffusion process. From observations of radio data it is expected to decrease exponentially with distance to the galactic disk (see e.g. \cite{Jansson:2009ip}). As a particle propagates, it experiences fewer disturbances for weaker mag\-netic field stren\-gths and prop\-agates freely in the limit of large galactic height $z$. In contrast, close to the galactic disk the field lines may capture a charged particle for a while as the direction is frequently changed. Hence a decreasing magnetic field corresponds to an increasing diffusion coefficient \cite{Evoli:2008dv,Gebauer:2009hk,Grajek:2010bz}. A consequence of the spatial dependence is that particles leaving the vicinity of the disk have a lower probability to come back but will rather drift away. This can be understood recalling the escape time which is anti-proportional to the diffusion coefficient \cite{Ginzburg:1990sk}. With increasing galactic height $z$ the escape time gets smaller, signifying the decreasing probability of a particle to return to the disk. In opposition to a homogeneous and isotropic diffusion coefficient this implies that particles detected in the disk are less likely to originate from a source which is far away. In other words the radius from where particles reach the Earth is reduced and energy losses become less important. This does not spoil the constraints from secondary to primary ratios of cosmic nuclei as their spectrum is nearly insensitive to the propagation time since energy losses have a much smaller impact on hadronic spectra \cite{TalkBlum}. Some special cases of inhomogeneous diffusion can be found in the literature. In the \textsc{Dragon}-package \cite{Evoli:2008dv} a vertically, exponentially increasing diffusion coefficient is implemented. An updated version also takes the spiral structure of the source distribution into account \cite{Evoli:2016xgn}. Another discussion is presented in \cite{Gebauer:2009hk} where a diffusion coefficient increasing linearly in vertical direction is implemented into the \textsc{Galprop}-code. A similar setting is investigated in \cite{Grajek:2010bz} studying the influence of the same modification on the antiproton spectrum. In these works it is pointed out that a diffusion model with vertically increasing diffusion coefficient is compatible with B/C and $^{10}\text{Be}/^{9}\text{Be}$ measurements. All these approaches focus on a numerical treatment, whereas we present analytic solutions for different kinds of vertical inhomogeneous diffusion models for leptons and hadrons\footnote{Some analytical results for special cases of inhomogeneous diffusion are discussed in the literature (see e.g. \cite{1975ICRC....2..706B,1980ApJ...239.1089L}).}. Concretely we consider a diffusion coefficient with an arbitrary power-law and exponential dependence in the vertical coordinate $z$. The paper is organized as follows. In chapter \ref{chap:isotropicDiffusion} we review the isotropic two-zone diffusion model. In chapter \ref{chap:spaceDepDiffusion} we present an analytic solution to the inhomogeneous diffusion equation for positrons and hadrons. We show that the high energy part of the positron spectrum and the B/C ratio can be reproduced in this framework. After a short discussion on the propagation parameters we conclude in section \ref{chap:conclusion}. All details of the calculation are sketched in appendix \ref{app:solve}. | \label{chap:conclusion} The observation of cosmic rays offers the opportunity to study dark matter via indirect detection. A good signal to background ratio is expected for antiparticles. In this light we studied the origin of the observed positron excess by modifying the established two-zone diffusion model, giving rise to a significant change for the spectral features of the background prediction. We presented an inhomogeneous propagation model with a diffusion coefficient that increases with galactic height. Our approach is well motivated from the spatial distribution of the magnetic field which again is responsible for the diffusion process. For the first time we derived an analytic solution for the inhomgeneous model. Considering two different cases we explored a diffusion coefficient depending on the galactic height first as a power-law and second as an exponential function. The modification reduces the propagation time scale such that energy losses are less relevant. In this model the high energy part of the measured positron spectrum can be reproduced with secondary positrons only. We also pointed out that this is a particular feature of the inhomogeneity. As a drawback, the model in such a setup is in contrast to the isotropic model no longer able to describe the low energy part of the observed positron spectrum. As a conclusion we claim that the inhomogeneity strongly affects the spectral shape of leptonic spectra and that these effects have to be taken into account for constraining additional primary sources. Our extended model can be tested experimentally. One feature of the inhomogeneous model is the prediction of a spectral softening in the positron spectrum. This offers a possibility to discriminate its spectral features from a pulsar contribution, where an exponential cutoff is expected. Furthermore the modification implies a reduction of the propagation time, such that the model can be constrained by analyzing the spectra of long lived radioactive particle species, e.g.\! $^{10}\text{Be}/^{9}\text{Be}$. Lastly an analysis of synchrotron radiation could quantify the spatial particle density distribution which provides another chance to discriminate between the different models. In order to check consistency of the inhomogeneous model with other particle species we investigated an analytic derivation of the hadronic fluxes. We found that these spectra are not significantly affected by the modification. This can be understood as energy losses play a minor role for the propagation of nuclei. Further investigation is required to examine whether the high energy part of the positron spectrum and the B/C spectrum can be described with the same set of propagation parameters. However, in the literature the method of using the same propagation parameters for leptons and hadrons is under debate. If the model can further be extended to describe the positron spectrum at lower energies is beyond the scope of this work. Finally we want to stress how many exciting possibilities are currently offered by indirect detection for the exploration of dark matter. If we want to benefit from the precision data which is expected to be available within the next few years, it therefore becomes a major issue to understand the astrophysical backgrounds better. Going beyond the ordinary two-zone diffusion model a more realistic description of cosmic ray propagation has been achieved in our work. The resulting analytic expressions can easily be applied to future studies. | 16 | 9 | 1609.01300 |
1609 | 1609.06716_arXiv.txt | Standard accretion disk theory \citep{ss1973} predicts that the total pressure in disks at typical (sub-)Eddington accretion rates becomes radiation pressure dominated. However, radiation pressure dominated disks are thermally unstable. Since these disks are observed in approximate steady state over the instability time-scale, our accretion models in the radiation pressure dominated regime (i.e. inner disk) need to be modified. Here, we present a modification to the Shakura \& Sunyaev model, where radiation pressure is in equipartition with gas pressure in the inner region. We call these flows Accretion in Radiative Equipartition (AiRE) Disks. We introduce the basic features of AiRE disks and show how they modify disk properties such as the Toomre parameter and central temperature. We then show that the accretion rate of AiRE disks is limited from above {\rm and} below, by Toomre and nodal sonic point instabilities, respectively. The former leads to a strict upper limit on the mass of supermassive black holes as a function of cosmic time (and spin), while the latter could explain the transition between hard and soft states of X-ray binaries. | Since the early years after \citet{ss1973} introduced thin disks, they have been known \citep{ss1976, piran1978} to be fraught with thermal instabilities. The instability occurs in the inner regions of the disk, where the pressure is dominated by radiation pressure and the opacity is mostly due to Thomson scattering, resulting in a much stronger temperature dependence in the heating of the disk compared to its cooling. \\ Decades later, the resolution of thermal instabilities still remains one of the major outstanding problems in understanding thin and slim disks. One of the major uncertainties in the Shakura \& Sunyaev (SS) $\alpha$ disk model is the assumption that the viscous stress is proportional to the total pressure. Early attempts to model thermal (and viscous) stability, such as the works of \citet{sakimoto, stella, merloni}, explored the possibility that the viscous stress might instead be only proportional to the gas pressure. These disks were called $\beta$ disks, where $t_{r\phi}=\beta\, p_{gas}$. Recent numerical simulations \citep{jiang2013,bran1995,stone1996, mishra2016, 2016MNRAS.459.4397S, 2016MNRAS.462..960S} see the presence of thermal instabilities, where the onset of thermal instability causes the disk to expand or collapse at the time scale of only a few orbits. These local simulations do not see evidence for such $\beta$ disks. \\ Radiative MHD simulations such as those in \citep{2016MNRAS.459.4397S, 2016MNRAS.462..960S} find stable radiatively efficient and strongly radiation pressure dominated disks, in the presence of strong magnetic fields. This has led to the claim that strong magnetic fields could stabilize disks against thermal instabilities \citep{2016MNRAS.459.4397S, 2016MNRAS.462..960S, begelman07, oda09}. If there is not enough magnetic flux, however, the instability once again sets in and a different means for stabilization must be sought. An iron opacity bump has also been suggested by \citet{jiang2016} as a means to postpone (but not avoid) thermal instabilities. In this paper, we present Accretion in Radiative Equipartition (AiRE) disks as an alternative solution to the thermal instability problem in thin and slim disks. | To summarize our main results, we have introduced AiRE disks as a solution to thermal instability in thin disks. The key feature of AiRE disks is that the radiation pressure is in equipartition with the gas pressure in the inner region. We have presented some features of these flows such as their central temperature and Toomre parameter profiles. We have derived upper limits for the mass of supermassive black holes due to the gravitational Toomre instability in AiRE disks. We have also found a transition from saddle to nodal type of the sonic points in AiRE disks and used nodal point instability to place a lower limit on the mass accretion rate as a function of viscosity parameter $\alpha$ and black hole spin. We conjecture that this transition might be responsible for the observed lower limits on the Eddington ratio of the soft state in X-ray binaries. \\ While we introduced AiRE disks to provide a thermally stable description of thin accretion flows, they may also significantly refine our understanding of other disk properties. With new observations from advanced Laser Interferometer Gravitational-wave Observatory (aLIGO) and the Event Horizon Telescope (EHT), we are at the advent of a new era of black hole physics. Disk Theory may play a major role in explaining some of their future findings. \\ In upcoming work \citep{ykyna}, we study the spectrum of AiRE disks and its properties. Furthermore, the onset of Toomre instability in the inner regions of AiRE disks around active galactic nuclei can lead to formation and merger of binary black hole systems, such as the ones recently detected by LIGO \citep{ligo}, and lead to characteristic detectable electromagnetic signatures \citep[e.g.,][]{bartos}.\\ Another important future direction is the study of the AiRE disk regime in full MHD radiative transfer simulations, and whether enhanced cooling leading to pressure equipartition can indeed arise in a realistic setting. | 16 | 9 | 1609.06716 |
1609 | 1609.02571.txt | { %\textbf{ We in this paper investigate the formation and evolution of primordial black holes (PBHs) in nonsingular bouncing cosmologies. We discuss the formation of PBH in the contracting phase and calculate the PBH abundance as a function of the sound speed and Hubble parameter. Afterwards, by taking into account the subsequent PBH evolution during the bouncing phase, we derive the density of PBHs and their Hawking radiation. Our analysis shows that nonsingular bounce models can be constrained from the backreaction of PBHs. %} %We in the present paper investigate the formation and evolution of primordial black holes (PBHs) within the scenario of nonsingular bouncing cosmology. We first analyze the process of PBH formation in a contracting background, and evaluate the PBH abundance at the end of matter contracting phase. %Our result shows that the abundance is generally small unless the energy scale associated with the bouncing phase is as high as the Planck scale, i.e., $|H_{-}|\gtrsim M_p$, or the sound speed of cosmological perturbations is sufficiently small, which implies, $c_s \ll 1$. %Our results imply that PBHs can form only when the background fluid has a small sound speed $c_s \leq 0.39$, %due to the fact that black holes ought to be within the cosmic apparent horizon, %and the PBH abundance, if $c_s \leq 0.39$, depends on the Hubble parameter $H(t)$ of the moment when PBHs were formed. %Afterwards, the subsequent evolution of the PBHs during the bouncing phase are studied. If the evaporation of PBHs due to Hawking radiation is not considered, a relation upon the model parameters $\Upsilon \geq c_s^2 \pi^2 H^2_{-}$ is expected to be satisfied, otherwise the PBH mass would grow to infinity before the bouncing point. We also calculate the back-reaction of PBHs \textbf{for $c_s \leq 0.39$} to constrain the bouncing models, considering both the PBH accretion and the evaporation due to Hawking radiation. The constraint is in accordance to the aforementioned relation $\Upsilon \geq c_s^2 \pi^2 H^2_{-}$ on the relatively low energy scale $H_{-}^2 \ll 10^9 c_s^5 M_\text p^2$. } %\pacs{98.80.Cq} | The matter bounce scenario \cite{Brandenberger:2012zb, Wands:1998yp, Finelli:2001sr} is one type of nonsingular bounce cosmology \cite{Brandenberger:2009jq, Brandenberger:2016vhg, Novello:2008ra, Lehners:2008vx, Cai:2014bea, Battefeld:2014uga}, which is often viewed as an important alternative to the standard inflationary paradigm \cite{Guth:1980zm, Starobinsky:1980te, Sato:1980yn, Fang:1980wi}. By suggesting that the universe was initially in a contracting phase dominated by dust-like fluid (with a vanishing equation-of-state parameter $w=0$), then experienced a phase of nonsingular bounce, and afterwards entered a regular phase of thermal expansion. The matter bounce cosmology can solve the horizon problem as successfully as inflation and match with the observed hot big bang history smoothly. Based on primordial fluctuations generated during matter contracting and their evolution through the nonsingular bounce, one can obtain a scale invariant power spectrum of cosmological perturbations. Unlike inflation, the matter bounce scenario does not need a strong constrain on the flatness of the potential of the primordial scalar field that drives the evolution of the background spacetime \cite{Adams:1990pn, Copeland:1994vg}. Also, this scenario can avoid the initial singularity problem and the trans-Planckian problem, which exists in inflationary and hot big bang cosmologies \cite{Borde:1993xh, martin2001trans}. The aforementioned scenario has been extensively studied in the literature, such as the quintom bounce~\cite{Cai:2007qw, Cai:2007zv}, the Lee-Wick bounce \cite{Cai:2008qw}, the Horava-Lifshitz gravity bounce \cite{Brandenberger:2009yt, Cai:2009in, Gao:2009wn}, the $f(T)$ teleparallel bounce \cite{Cai:2011tc, deHaro:2012zt, Cai:2015emx}, the ghost condensate bounce \cite{Lin:2010pf}, the Galileon bounce \cite{Qiu:2011cy, Easson:2011zy}, the matter-ekpyrotic bounce \cite{Cai:2012va, Cai:2013kja, Cai:2014zga}, the fermionic bounce \cite{Alexander:2014eva, Alexander:2014uaa}, etc.~(see, e.g. Refs. \cite{Brandenberger:2010dk, Brandenberger:2012zb} for recent reviews). In general, it was demonstrated that on length scales larger than the time scale of the bouncing phase, both the amplitude and the shape of the power spectrum of primordial curvature perturbations can remain unchanged through the bouncing point due to a no-go theorem \cite{Quintin:2015rta, Battarra:2014tga}. A challenge that the matter bounce cosmology has to address is how to obtain a slightly red tilt on the nearly scale invariant primordial power spectrum. To address this issue, a generalized matter bounce scenario, which is dubbed as the $\Lambda$-Cold-Dark-Matter ($\Lambda$CDM) bounce, was proposed in \cite{Cai:2014jla} and predicted an observational signature of a positive running of the scalar spectral index \cite{Cai:2015vzv, Cai:2016hea}. As a candidate describing the very early universe, the matter bounce scenario is expected to be consistent with current cosmological observations and to be distinguishable from the experimental predictions of cosmic inflation as well as other paradigms \cite{constraint, Cai:2014bea}. Meanwhile, a possible probe of primordial black holes (PBHs) may offer a promising observational approach to distinguish various paradigms of the very early universe \cite{Carr:1975qj, Carr:2009jm}. PBHs could form at very early times of the universe, where a large amplitude of density perturbations would have obtained. Correspondingly, the formation process and the abundance of PBHs strongly depend on those early universe models, in which fluctuations of matter fields are responsible for such large amplitudes of density perturbations \cite{Carr:1974nx}. In the literature, most of attentions were paid on the computation of PBH predictions from the inflationary paradigm (for instance see \cite{sub-Hubble,sub-Hubble2, saperate, threshold, Josan}), while so far, only a few works addressed the PBH formation in a bouncing scenario\cite{carr2011persistence, carr2016primordial}. Furthermore, those studies of PBHs in a bouncing scenario have not yet been discussed in detail, for specific cosmological paradigms or been applied to falsify various early universe cosmologies, especially the matter bounce scenario. In the context of matter bounce cosmology, there are several differences on the computation of the PBH abundance comparing with that in an expanding universe. First, comparing with inflation where the primordial fluctuations become frozen at the moment of the Hubble exit, those primordial fluctuations on matter fields in bounce cosmology would continue to increase after the Hubble exit during the contracting phase until the universe arrive at the bouncing phase \cite{Cai:2008qw, Brandenberger:2016vhg}, and the contracting phase would yield a different initial condition for the PBH formation and evolution. Second, once these PBHs have formed, the contraction of spacetime could also compress and enlarge the primordial matter density, thus change the PBH horizon radius which then can lead to effects on their evolution. In this paper, we perform a detailed survey on the PBH formation and evolution in the background of the matter bounce cosmology. In Section \ref{model}, we briefly introduce the matter bounce scenario and describe the formation of the power spectrum of primordial curvature perturbation in an almost model-independent framework. In Section \ref{formation}, a physical picture of the PBH formation in the contracting background is presented. After a process of detailed calculations, the threshold for forming PBHs and the corresponding mass fraction are provided. In Section \ref{evolve}, we discuss the evolution of PBHs in the bouncing phase by taking into account the effects arisen from the contraction of the background and the Hawking radiation. In Section \ref{conclusion}, we summarize our results and discuss on some outlook of the PBH physics within the nonsingular bouncing cosmology. | In this paper, we have investigated the formation and evolution of PBHs in the matter bounce cosmology. Firstly, we described the general matter bounce models by some parameters like $c_s$, $H_{-}$ and $\Upsilon$. The comoving curvature perturbation $\mathcal R_k$ is also calculated during the matter contracting phase, which seeds the PBH formation. Then we had a discussion about the condition of PBH formation in the contracting background, which is different from that in the expanding universe. By taking a simple collapsing model, the threshold of the density fluctuation for forming a PBH is derived. Furthermore, in the comoving gauge, the density fluctuation and its threshold are all related to the curvature perturbation $\mathcal R_k$. Therefore, we can calculate the mass fraction of PBHs $\beta$ in the contracting phase from the Press-Schechter theory, and constrain the bouncing models from PBHs. PBH formation depends on the model parameters $c_s$ and $H_{-}$. For instance, PBHs can form in the contracting phase only if $c_s\leq 0.39$, since the BHs are not allowed to be larger than the cosmic apparent horizon. When $c_s\leq 0.39$, $\beta$ is small and the model is safe for $|H_{-}|\ll M_\text p$ or $|H_{-}|\gg M_\text p$, but only the low energy regime is favored from the model construction. The subsequent evolution of PBHs in the bouncing phase is also investigated, with the PBH accretion and Hawking radiation considered. The growth behavior of PBH yields a constraint to the model $\Upsilon \geq c_s^2 \pi^2 H^2_{-}$, in case that the mass of PBHs grow to infinity before the bouncing point. Moreover, the back reaction of PBH and its Hawking radiation is calculated to constrain models, in order not to neutralize the negative pressure $p_{\rm bg}=-2M_\text p^2 \Upsilon$, since in such case the universe cannot expand again. The constraint reduces to $\Upsilon \geq c_s^2 \pi^2 H^2_{-}$ at the low energy scale $H^2_{-}\ll 10^9 c_s^5 M_\text p^2$. Afterwards, a rough constraint of bouncing model though the PBH observations is given, and the constraint is stringent only when $\Upsilon$ is slightly larger than $c_s^2 \pi^2 H^2_{-}$. A more precise analysis taken into account both the PBH growth and the Hawking evaporation as well as the detailed constraints from cosmological observations will be addressed in our follow-up works. We note that while our paper was being prepared, an independent work was being carried out by another group \cite{Quintin:2016qro}, which explores similar features of BH formation in bouncing cosmology. | 16 | 9 | 1609.02571 |
1609 | 1609.06650_arXiv.txt | The Massive Young Star Forming Complex in Infrared and X-ray (MYStIX) project provides a new census on stellar members of massive star forming regions within 4~kpc. Here the MYStIX Infrared Excess catalog (MIRES) and $Chandra$-based X-ray photometric catalogs are mined to obtain high-quality samples of Class~I protostars using criteria designed to reduce extragalactic and Galactic field star contamination. A total of 1,109 MYStIX Candidate Protostars (MCPs) are found in 14 star forming regions. Most are selected from protoplanetary disk infrared excess emission, but 20\% are found from their ultrahard X-ray spectra from heavily absorbed magnetospheric flare emission. Two-thirds of the MCP sample is newly reported here. The resulting samples are strongly spatially associated with molecular cores and filaments on \textit{Herschel} far-infrared maps. This spatial agreement and other evidence indicate that the MCP sample has high reliability with relatively few 'false positives' from contaminating populations. But the limited sensitivity and sparse overlap among the infrared and X-ray subsamples indicate that the sample is very incomplete with many 'false negatives'. Maps, tables, and source descriptions are provided to guide further study of star formation in these regions. In particular, the nature of ultrahard X-ray protostellar candidates without known infrared counterparts needs to be elucidated. | In the densest cores of molecular clouds, gravitational collapse overcomes thermal pressure, turbulence, and rotational and magnetic energies to form protostars. The protostellar phases of pre-main sequence evolution, lasting an estimated 0.5~Myr, are characterized by infall of material from the circumstellar envelope, formation of a protoplanetary disk, and bipolar outflows collimated by magnetic fields \citep{Dunham14}. Protostars can be identified in three primary ways: \begin{enumerate} \item A strong photometric infrared excess due to thermal emission from the envelope and disk dust grains discriminates protostars from less interesting foreground and background stars. The infrared spectral energy distribution (SED) provides an evolutionary classification classification \citep{Adams87, Andre00}. Class~0 protostars appear in the far-infrared, too obscured to be detected in the near- and mid-infrared bands. Class~I protostars have a positive spectral index over the $2-25 \mu m$ band. A flat spectral index is sometimes classified as a Class I/II transition, Class~II stars have a mildly negative spectral index, and stars without disks exhibiting the Rayleigh-Jeans slope are classified as Class III \citep{White07}. The SED may depend on the chance orientation of the disk; for example, a Class II T Tauri star with disk seen edge-on can mimic Class I protostellar SED \citep{Whitney03}. Furthermore, protostellar infrared colors can be similar to contaminating cosmic populations: dust-enshrouded AGB post-main sequence stars, starburst galaxies, Type 2 obscured active galactic nuclei, and shocked interstellar gas emission. Cuts in color-color and color magnitude diagrams have been developed help discriminate protostars from these contaminants \citep{Gutermuth09}. \item Protostars are distinguished by strong outflows that are directly traced by molecular maps showing redshifted and blueshifted lobes. When molecular maps are not available, infrared line emission from outflow shocks can sometimes be detected photometrically. The {\it Spitzer Space Telescope} 4.5 $\mu m$ band contains excited H$_2$ lines and CO band heads that can produce a distinctive photometric [4.5] excess \citep{Cyganowski08}. \item Throughout the Class I-II-III stages of pre-main sequence evolution, low mass stars show very high levels of surface magnetic activity that are most easily detected by X-ray flaring \citep{Feigelson99}. X-ray flares from protostars and T Tauri stars are similar to, but much stronger and more frequent than, magnetic reconnection events flares above the surface of the contemporary Sun. The high extinction of deeply embedded protostars absorbs softer energy photons, leaving only harder X-ray photons; spectral hardness is thus a signature of Class~I stars \citep{Imanishi01, Ozawa05}. The median energy $ME$ (in keV) of detected photons from a young star is a reliable measure of soft X-ray absorption \citep{Getman10}. Class~0 stars have never been detected in X-rays, but the limits are consistent with Class~I levels subject to heavy absorption \citep{Giardino07}. \end{enumerate} Infrared excess protostellar samples are readily obtained from imaging surveys obtained with the \textit{Spitzer Space Telescope} and \textit{Herschel Space Observatory} in nearby smaller star forming clouds \citep{Evans03, Andre10} and the nearest Orion giant molecular cloud \citep{Megeath12, Stutz13}. Less complete samples of Class~I protostars are obtained from X-ray surveys with the {\it Chandra X-ray Observatory} and {\it X-ray Multimirror Mission} in nearby clouds such as $\rho$ Ophiuchi \citep{Casanova95, Imanishi01}, Corona Australis \citep{Koyama96, Hamaguchi05}, and Serpens \citep{Preibisch04, Winston07}, as well as the Orion giant molecular cloud \citep{Grosso05, Prisinzano08, Pillitteri13}. However, it is difficult to obtain unbiased flux-limited protostellar samples in more distant and massive Galactic star forming clouds for three reasons. First, the protostellar infrared sources are fainter and more crowded, difficult to resolve from neighboring protostars or the more common Class~II stars in embedded clusters. Second, more distant star forming regions are concentrated in the Galactic Plane where contamination by dusty post-main sequence stars can be severe. The surface density of field stars in $Spitzer$ images at low Galactic latitudes and longitudes can reach the confusion limit fainter than $\sim 12$~mag in short exposures due to red giant field stars \citep{Ramirez08}. Third, in the vicinity of H\textsc{ii} regions ionized by recently formed OB stars, unresolved protostellar sources can be difficult to detect within the spatially complex diffuse emission from heated dust, especially in infrared bands with strong polycyclic aromatic hydrocarbon emission. In some cases, reasonable protostellar samples can be obtained despite diffuse emission, as in the Rosette Molecular Cloud \citep{Hennemann10}, but in other cases only luminous high-mass protostars can be identified, as in the NGC~6334 complex \citep{Brogan09}. A new opportunity for the discovery of protostars in giant molecular clouds emerges from the Massive Young Star-forming complex in Infrared and X-ray (MYStIX) project. MYStIX combines infrared and X-ray photometric surveys of 20 OB-dominated young star forming complexes at distances from 0.5 to 4 kpc. MYStIX uses spectral-imaging data from NASA's \textit{Chandra X-ray Observatory}, $1-2$~$\mu$m images from the UK InfraRed Telescope (UKIRT) wide-field camera, and $3-8$~$\mu$m images from NASA's \textit{Spitzer Space Telescope} \citep{Feigelson13}. This survey produced a catalog of 31,784 MYStIX Probable Complex Members (MPCM) in the 20 regions \citep{Broos13}. Most MPCM stars are Class~II and Class III pre-main sequence stars, but a small fraction should be Class~I protostars. The much larger catalogs of unmatched MYStIX Infrared-Excess Source \citep[MIRES;][]{Povich13} catalog and X-ray sources \citep{Kuhn13a, Townsley14} produced as intermediate products during the construction of the MPCM catalog may contain additional protostars, although these single-waveband catalogs suffer strong contamination by Galactic field stars and/or extragalactic sources. The present study seeks protostars from these MYStIX survey catalogs. We combine objects with protostellar infrared SEDs and 4.5~$\mu$m excesses with X-ray sources exhibiting ultra-hard spectra denoting very heavy obscuration. These criteria filter away nearly all of the older Class II-III stars and contaminant populations, but give very incomplete samples. The result is a list of 1,109 protostellar candidates in fourteen star forming regions (in order of increasing right ascension): W~3, Flame Nebula, Rosette Nebula, NGC~2264, RCW~38, RCW~36, NGC~6334, NGC~6357, Trifid Nebula, Lagoon Nebula, M~17, Eagle Nebula, W~40, and DR~21. A fifteenth MYStIX region with the 5 Myr cluster NGC~2362, was chosen as a control field as no protostars are expected in this region without molecular material. The reliability of the catalog is strengthened because a large majority (86\%) are found to be associated with dense cores seen in \textit{Herschel} 500~$\mu$m maps that trace cold dust emission. However, the candidate list requires more detailed study for confirmation, and cannot be viewed as an unbiased view of star formation in the clouds. Section 2 details the methods of selection for the candidate protostars, sections 3-4 present results for individual star forming regions, and section 5 discusses the overall findings. | \subsection{Summary} \label{summary.sec} The results of this search for protostars in the MYStIX joint infrared/X-ray survey of 14 OB-dominated star forming regions are summarized in Table~1. We identify 1,109 MYStIX Candidate Protostars (MCPs) of which about 700 are newly reported here. The great majority (955 candidates or 86\% of the total) are closely associated with dense cloud cores or filaments seen on $Herschel$ satellite far-infrared maps. Except for the tightening of X-ray selection criteria to reduce candidates lying far from cloud cores (\S\ref{ultrahardX.sec}), the MCP selection criteria were based on infrared and X-ray photometric properties with no reference to location in the field. About 1/5 of the MCPs were identified by virtue of their X-ray photometry (\S\ref{ultrahardX.sec}) while most were identified using infrared photometry. X-ray selection is particularly needed in locations where bright PAH nebulosity impedes detection of protostellar disks in the infrared. We discuss the novel selection of protostars by X-ray spectral hardness in \S\ref{Xray.sec} below. Our infrared selection procedures differ from those used by other groups, using combinations of SED modeling, location in the IRAC color-color diagram, and [4.5] band excess emission (\S\ref{irexcess.sec}). In addition, photometric values are derived from MYStIX pipelines designed to treat crowded and nebulous regions in UKIDSS/UKIRT near-infrared images and $Spitzer$ mid-infrared images \citep{King13, Kuhn13b}. We therefore do not expect that MYStIX protostellar lists to be identical to other infrared photometric studies. One useful result of the maps of MCPs in Figure~1 is the identification of new small star forming clouds in the vicinity of well-known young stellar clusters. In the Rosette Nebula, for example,we see a previously unrecognized star forming cloudlet $\sim 10$~pc southwest of the main molecular cloud, and several cloudlets with active star formation are seen around RCW~38. In other cases, candidate protostars are found in areas where only older pre-main sequence stars were known, such as around S~Mon in NGC~2264 or in the IRAS~09002-4732 cloud near RCW~38. The reliability of the candidate list, discussed in \S\ref{reliability.sec}, appears to be high in the sense that the fraction of non-protostellar interlopers (False Positives) is low. The main sources of contamination are expected to be extragalactic active galactic nuclei in the X-ray images, and Galactic AGB red giants and extragalactic starburst galaxies in the infrared images. All of these populations should give spatial distributions that either avoid the molecular clouds or are randomly distributed across the fields. But Figure~1 shows strongly clumped distributions of MCP stars along cloud cores and filaments. This is clear evidence for the high reliability of our sample. Interpretation of the scattered candidates is unclear; they may be residual contaminant populations, or indications of distributed star formation in small cloudlets. The completeness of the candidate list, discussed in \S\ref{incompleteness.sec}, is low. The high fraction of MCPs that are very bright in the IRAC images $-$ such as those with [8.0]$ \leq 8$~mag in the Rosette, NGC~6334, NGC~6357, Lagoon and Eagle Nebulae $-$ clearly suggests that a large population of lower brightness protostars is missing from the sample. We attempt to estimate the incompleteness \S\ref{incompleteness.sec} and infer that the full population of protostars is probably an order of magnitude larger than our sample; that is, many thousands of protostars are probably present compared to our sample of $\sim 1,000$ MCPs in the 14 MYStIX regions under study. Due to this large and uncertain incompleteness, we emphasize that the MCP list should not be used for astrophysical calculations that assume complete samples; for example, for estimation of star formation rates and efficiencies in molecular clouds. \subsection{Reliability of results} \label{reliability.sec} Several lines of evidence give confidence that most of these candidates are real protostars: \begin{enumerate} \item About half of the MCPs satisfied a different set of constraints to be listed MYStIX Probable Complex Member (MPCM) catalog of \citet{Broos13} that consists mainly of Class~II and III pre-main sequence stars. The MPCM catalog is specifically designed to reduce contamination by Galactic field stars and extragalactic sources. But it was not designed to capture protostars efficiently, particularly in the X-ray band (\S\ref{Xsel.sec}). \item As mentioned above, 86\% of the MCPs are spatially associated with dense cores or filaments in $Herschel$ 500$\mu$m maps. In NGC~6334, for example, several dozen MCPs are closely aligned with the principal molecular filament or outlying cores. In NGC~2264, NGC~6334, Eagle Nebula, Trifid Nebula, W~40 and M~17, most candidates are aligned with the molecular structure. In the Rosette Nebula, NGC~6357, Trifid Nebula, M~17, Eagle Nebula, W~40 and elsewhere the candidate protostars are spatially distinct from the older cluster rich in Class~II and III stars ionizing the giant H\textsc{ii} region that was historically the first indication of star formation. \item X-ray selected MCPs are sometimes more dispersed; some coincide with cores and filaments, but others are widely dispersed in the $Chandra$ fields. Some of the dispersed X-ray selected MCPs are probably the expected contaminants from Type 2 active galactic nuclei (\S\ref{AGN.contam.sec}). \item It is possible that a few of the MCPs are Class~II stars with nearly edge-on protoplanetary disks that can temporarily exhibit the high absorption expected in Class~I stars. The best studied example of such a system is AA~Tau, and \citet{Morales11} estimate that $\gtrsim 5$\% of Class~II stars in the Orion Nebula Cluster exhibit brief photometric dips from orbiting disk material in the $Spitzer$ IRAC bands. Two cases of X-ray absorption from an edge-on disk producing an ultrahard spectrum are also known in the Orion Nebula Cluster \citep[COUP~241 and COUP~419;][]{Kastner05}. \item The MCP sample recovers $\sim 300$ protostars found by past studies. The level of association depends on the historical coverage of each region in the infrared. For example, many MYStIX candidates are associated with known protostars in the M~17 region where past research has been extensive, but few are previously known in NGC~6357 where past study is weak. As our listing is very incomplete (\S\ref{incompleteness.sec}), we do not expect to recover all previously known protostars. \end{enumerate} Taking these issues into consideration, we estimate that $80-90$\% of the 1,109 MCPs are likely to be true protostars. The selection criteria derived in \S\ref{selection.sec} was sufficiently conservative to remove the large contaminating populations of dusty Galactic field stars and obscured extragalactic objects from our sample. \subsection{Incompleteness of the protostar survey} \label{incompleteness.sec} The protostar selection procedures described in \S\ref{selection.sec} are very conservative in the sense that criteria were chosen at levels to increase reliability of the sample (i.e., set to reduce false positives) and not chosen to capture a large fraction of the true protostellar population (i.e., set to reduce false negatives). These constraints can be summarized as follows: \begin{enumerate} \item Our criterion of $ME>4.5$~keV requires extremely high absorption to the X-ray source, and excludes the majority of protostars that have $3.0 < ME < 4.5$~keV ($22.5 <\log N_H < 23.3$~cm$^{-2}$) range. See Table~1 of Imanishi et al.\ 2001, Tables~2 and 4 of Grosso et al.\ 2005 and Table~6 of Prisinzano et al.\ 2008 for X-ray hardness distributions of protostars in the $\rho$ Ophiuchi and Orion Nebula regions. A few of these less-absorbed X-ray emitting protostars are found through our infrared excess criteria; for example, MCP Flame 7 has 36 net counts with $ME = 3.6$~keV. \item The X-ray flux limits for heavily absorbed stars in the MYStIX fields severely truncates our sample. For a hard intrinsic spectrum (say, thermal bremsstrahlung with $kT = 5$~keV or $T \simeq 100$~MK), less than 5\% of the photons that would be detected by the $Chandra$ ACIS detector from a hypothetical unabsorbed protostar are in fact detected when $ME > 4.5$~keV. Thus our requirement $>10$ net counts is equivalent to a requirement of $>200$ net counts if absorption were not present. This essentially confines the ultrahard X-ray sources to the upper tail of the protostellar X-ray luminosity function with $\log L_X \simeq 31-32$~erg s$^{-1}$. Protostars with X-ray flares at this level have been found \citep{Hamaguchi05, Pizzocaro16} but are rare \citep{Imanishi01, McCleary11, Prisinzano08}. \item The infrared color-color criteria should detect many Class~I protostars, limited by the impacts of crowding and spatially varying PAH nebular emission. But the criterion of [4.5]$\mu$m excess emission excludes the large majority of infrared sources with protostar-like SEDs (see Figures~3 and 6 of Povich et al.\ 2013). This photometric characteristic is sufficient indicator of protostars but is not at all a necessary indicator. \item Most Class~0 protostars, and any Class~I stars with extremely heavy absorption ($A_V \geq 200$~mag), are missed by our procedures described in \S\ref{selection.sec}. Possible Class~0 sources captured in the MCP catalog include SerpS-MM24 (MCP W40 63), W40 MM5 (MCP W40 79), and members of the Class~0 microcluster in NGC~2264 (MCP NGC2264 27, 30-32). For extremely heavily absorbed stars, photometric identification requires longer wavelengths than available from $Spitzer$'s IRAC detector and higher sensitivity than available from $Chandra$'s X-ray telescope with typical exposure times. For example, a deep $Chandra$ observation of the Serpens cloud detected 9 Class~I and 9 Class~I/II (flat spectrum) protostars, but none of the 6 Class~0 or Class~0/I stars \citep{Giardino07}. \end{enumerate} Combining these effects is not simple. The survey depths in the X-ray and infrared bands differ considerably among regions, so quantitative conclusions about the full MYStIX survey are not reliable. The relationship between sensitivity and protostellar mass is not clear for $any$ protostellar survey; completeness with respect to the underlying Initial Mass Function is thus always difficult to establish. Our rough qualitative estimate based on comparison with previous protostellar surveys that our selection procedure in \S\ref{selection.sec} misses at least half, and probably most, of the true protostellar population. The full underlying protostellar population would then be several our observed 1,109 MCPs in all 14 MYStIX regions examined. Another rough estimate of the incompleteness can be made with the mathematics of `capture-recapture' modeling used for over a century in ecological studies \citep{Amstrup05}. Here we estimate the total population of a class based on the overlap when two independent but incomplete samples are obtained assuming equal probability of capture in each experiment. The population is assumed to be 'closed' with no additions or losses among the experiments. Consider the first 'capture' experiment as the detection of candidate protostars based on infrared excess. These are then `returned' to the environment and a second `capture' experiment detects candidate protostars based on X-ray properties. We omit X-ray observations in regions where strong PAH nebulosity inhibits infrared detection, and we omit infrared observations in regions without $Chandra$ detection. These regions have only one effective `capture' experiment. The fact that very few protostars are simultaneously detected with X-ray and infrared selection techniques is, in the context of capture-recapture theory, indicative that the total population is much larger than the combined samples. We treat the infrared detection of $\sim 700$ MCPs in $Chandra$ fields to be the first `capture' experiment, and the X-ray detection of $\simeq 100$ candidate protostars in regions where PAH contamination is not severe to be the second `recapture' experiment. In our MCP survey, only five sources are recovered by both X-ray and infrared methods: MCP Flame 14, MCP NGC2264 36, MCP NGC2264 54, MCP DR21 32, and MCP DR21 66. However, about 70 candidates would have been recovered by both surveys is the source ha redder infrared colors, a few more X-ray photons, or higher X-ray median energy. Standard mathematical procedures from capture-recapture analysis are used: the Petersen estimator of the total protostellar population with Chapman's bias correction and Seber's standard error estimate \citep{Amstrup05}. This computation was made with function $closedp.bc$ in CRAN package $Rcapture$ \citep{Rivest14} within the public domain R statistical software environment \citep{RCoreTeam15}. For 5 recovered stars from experiments with 700 and 100 captures, the estimated total population of $27 \pm 11$ thousand protostars. If we more optimistically consider 70 stars recovered using both methods by relaxing X-ray and infrared selection rules, the estimated population would be $3.1 \pm 0.3$ thousand protostars. We conclude that there are $\sim 3,000 - 30,000$ protostars in these 14 MYStIX fields, so that the MCP identification rate is between 3\% and 30\% of the full protostellar population. We conclude that the MCP sample provides a very incomplete complete survey of the full protostellar population in the 14 examined MYStIX fields. An order of magnitude more protostars are probably present that we do not identify, and quantitative estimates of incompletness are very inaccurate. {\it As a consequence of this large and poorly quantified incompleteness factor, the MYStIX Protostellar Candidates sample should not be used for any statistical or astrophysical purpose, such as estimation of the current star formation rate or efficiency in these molecular clouds.} \subsection{The nature of X-ray ultrahard sources} \label{Xray.sec} Perhaps the most unusual element of our search for MYStIX protostars is the inclusion of ultrahard X-ray sources that do not have known infrared counterparts. This may seem risky, as infrared excess with an ascending SED has historically been the sole criterion for identifying Class~I protostars. But the risk may not be so high. First, the MYStIX fields are inhospitable for self-consistent and sensitive identification of protostars due to sensitivity reduction from the nebulosity of bright H\textsc{ii} regions, contamination by dusty red giants, and crowding. So the absence of catalogued infrared sources from low-resolution telescopes like $Spitzer$ associated with ultrahard X-ray sources does not clearly indicate that a protostar is not present. Second, the combination of X-ray selection criteria of $ME > 4.5$~keV and Net counts $>10$ is so severe that, except for a handful of randomly located Type~2 quasars, the resulting sample should be confined to X-ray emitting Class~I stars. Third, most of the the ultrahard sources are clustered around $Herschel$ cloud cores and, in some cases, recover previously known protostars. We therefore encourage followup of the X-ray ultrahard MCP sources using infrared telescopes. It is possible that they represent a subclass of protostars with different properties than those located with traditional infrared color criteria. The most valuable instrumental characteristic for followup would be imaging and photometry at high spatial resolution at $2-10 \mu$m wavelengths, needed to reduce the surface brightness of nebulosity. Large ground-based telescopes with active optics in the infrared band, and NASA's forthcoming {\it James Webb Space Telescope}, would be particularly effective for finding the expected infrared counterparts. | 16 | 9 | 1609.06650 |
1609 | 1609.01593_arXiv.txt | {\small Further investigation of data on quasars, especially in the ultraviolet band, yields an amazingly coherent narrative which we present in this paper. Quasars are characterised by strong continuum emission and redshifted emission and absorption lines which includes the famous Lyman $\alpha$ forest. We present irrefutable evidence in support of (1) the entire line spectrum arising in matter located inside the quasar system, (2) the range of redshifts shown by the lines being due to the variable contribution of the gravitational redshift in the observed line velocity, (3) existence of rotating black holes and of matter inside its ergosphere, (4) quasars located within cosmological redshifts $\sim 3$, (5) $\gamma$ ray bursts being explosive events in a quasar. These results are significant and a game-changer when we realise that the absorbing gas has been postulated to exist along the line-of-sight to the quasar and observations have accordingly been interpreted. In light of these definitive results which uniquely constrain the quasar structure, we need to drastically revise our understanding of the universe built on the faulty assumptions of observed redshifts of quasars having an entirely cosmological origin and the absorption lines arising in the intervening medium. } | Quasars are a well-observed and well-studied subgroup of objects generally known as active nuclei - we refer to quasi-stellar objects as quasars in this paper. However, they famously remain one of the least understood objects. The study which identified quasars \citep{1963Natur.197.1040S} reported wide emission lines ($\sim 50$A) on a strong blue continuum in 3C 273. The lines could only be idenfied if a redshifted velocity component was included although the object had a star-like appearance. Quasars have spurned extensive research but continue to intrigue even after 50+ years of discovery. Quasars show rich observational signatures which includes strong thermal emission in the ultraviolet (and optical) and a power law continuum emission which dominates the optical to infrared bands and continuing in the radio when detected. Their line spectra show the presence of wide emission lines and numerous absorption features spanning a range of velocities and widths. Some of the observed properties of active nuclei, in particular quasars, can be summarised to be: \begin{itemize} \item Flat or power law continuum emission from radio to ultraviolet wavelengths. Enhanced ultraviolet emission referred to as the `blue bump' or 'uv upturn'. Many active nuclei also detected in X-rays. \item Quasar spectra are characterised by broad emission lines and narrow/broad absorption lines. BL Lac objects show a featureless continuum with only a few showing spectral lines. Seyfert 1 galaxies show broad emission lines while Seyfert 2 galaxies show narrow emission lines. \item High ionization lines such as doublets of C IV (1548.188A 1550.762A), S IV (1393.755A,1402.770A), N V (1238.808A, 1242.796A), O VI (1031.928A, 1037.619A) are detected from many active nuclei especially quasars. Low ionization lines such as C II (1335A), Fe II (2383A,2586A), Si II (1260A), Mg II (2795.528A, 2802.704A) are also detected in the spectra of many active nuclei. These lines are detected either in emission and/or absorption in the quasar spectrum. \item Quasar spectra show a host of redshifts with the emission line redshifts being the largest. Absorption lines span a range of redshifts. \item Many active nuclei show variability especially quasars, blazars and Seyfert 1 galaxies. \end{itemize} While there is general agreement that a supermassive black hole is the central object in all active nuclei, rest of the details remain perplexing at best. It is instructive to glimpse the exciting research that quasars sparked due to their exotic nature as captured by their observations. Soon after quasars were identified \citep{1963Natur.197.1040S}, the debate on whether these objects were extragalactic and very distant or whether these were Galactic or in the neighbourhood has been going on. \citet{1963Natur.197.1040S} and \citet{1964ApJ...140....1G} concluded that quasars were distant extragalactic objects. The controversy arose since the large observed redshifted velocities of the spectral lines, if interpreted to indicate Hubble expansion, would make quasars very distant objects. This, then, led to the observed magnitudes translating to very high luminosities for the quasars which had, hitherto, not been observed in any extragalactic object. However, there was a group of astrophysicists who were convinced observations indicated that quasars were local and the redshifts were intrinsic. \citet{1967ApJ...148..321A} suggested that radio sources were associated with nearby peculiar galaxies \citep{1966ApJS...14....1A} or bright galaxies \citep[e.g.][]{1974IAUS...58..199A}. In fact a few such pairs were also found to be physically connected by a bridge \citep[e.g. NGC 4319 and Mrk 205;][]{1971ApL.....9....1A}. However the local origin did not find favour with most astronomers. One major problem with the local origin and association of quasars with nearby galaxies were the distinct redshifts noted for the quasar (high) and the nearby galaxy (low). Since one of the explanations was that quasars are ejected from nuclei of galaxies \citep[e.g.][]{1967ApJ...148..321A}, the quasars can show large redshifts. However in this scenario, the quasars would show both redshifted and blueshifted velocities wrt to the nearby galaxy whereas the quasars always showed a redshift wrt to the nearby galaxy. This essentially ruled out the ejection origin for the redshift and an hitherto unknown non-velocity intrinsic origin for the redshifts of quasars was postulated. \citet{1974IAUS...58..199A}, \citet{2007ARA&A..45....1B} and others continued to advocate the scenario of quasars being local objects and the observed redshifts having an intrinsic origin. In this paper, we revisit the quasar redshifts and present evidence for a sizeable intrinsic redshift component in quasar spectra. Another perplexing observational result noted around the same time was the arrangement of galaxies in the Coma cluster along bands in the redshift-magnitude diagram which was especially significant when the nuclear redshifts and magnitudes were plotted \citep{1972ApJ...175..613T,1973ApJ...179...29T}. This result, in addition to a periodicity observed in the velocity differences in pairs of galaxies \citep{1980ApJ...236...70T} suggested that the velocities were quantised in factors or multiples of $\sim 72$ kms$^{-1}$ \citep[e.g.][]{1980ApJ...236...70T}. Although these results are not yet understood, its interesting that this value is close to the currently accepted value of the Hubble constant. Since the Hubble constant gives the velocity difference between two galaxies separated by 1 Mpc, the Hubble law can be understood as quantifying the radial velocity distribution of galaxies in space or in other words `redshift quantisation'. In fact, it is interesting that Tifft had estimated a value for the Hubble constant without realising it. Obviously, observational astronomy was throwing up several puzzling results which were difficult to understand. We present a possible explanation for the redshift-magnitude bands in the paper. Most of the astronomical community has currently accepted the cosmological origin of redshifts of quasars and observations have been examined with this implicit assumption. Observations which could not be explained like Tifft's bands and Arp's quasar/galaxy associations were considered to be faulty or spurious which is alarming since even if one did not agree with their interpretation, these were observational results by solid astronomers and needed to be scientifically investigated. A careful study clearly brings out the fantastic nature of quasars and one senses that the unique properties of quasars are likely extreme signatures of an active nucleus. In this paper, we present our study which provides strong evidence to a sizeable internal contribution to the observed redshifts of quasars. We show that one of the physical processes considered for the origin of the observed high redshifts by \citet{1963Natur.197.1040S,1964ApJ...140....1G} is indeed the most plausible explanation. We start with a study of the origin of the large observed velocity shifts of the lines in a quasar spectrum followed by the ultraviolet continuum emission. Then we suggest a structure for the quasar which can explain most ultraviolet observations within the framework of known physics, discuss other active nuclei and variability and end with some concluding remarks. | We have examined the origin of the continuum emission and the large redshifts shown by the emission and absorption lines in a quasar spectrum in the ultraviolet. The study has resulted in several important watertight inferences, all fitting together like a jigsaw puzzle and uniquely constraining the quasar structure. The important results/inferences can be summarised to be: \begin{itemize} \item All the emission and absorption lines detected in a quasar spectrum arise from matter inside the quasar system. The observed redshift ($z=z_{em}$) of a quasar is large due to contribution from both the cosmological ($z_c$) and intrinsic ($z_{in}$) redshifts. $z_{in}$ comprises of a Doppler component ($z_D$) and a gravitational redshift ($z_g$). We show that $z_g$ going upto values of one comprises the dominant component of $z_{in}$. \item For quasars, we find $z_c<z_{em}$, $z_{in} < 1.25$ and $z_c\le3$. \item The observations, particularly in the ultraviolet band, determine a unique structure for a quasar. The proposed quasar structure consists of a supermassive black hole surrounded by degenerate matter shells (where the neutron and electron degeneracy pressure balance the black hole gravity) emitting thermal continuum peaking in the ultraviolet which heats and ionizes matter around it giving rise to a Stromgren shell where emission lines arise. Around the emitting shells are dense shells from where the absorption lines arise. The entire structure is arranged just outside the event horizon of a black hole in quasars (see Figure \ref{fig12}). The spectral lines are shifted by different $z_g$ depending on the separation of the line forming region from the black hole i.e. gravitational potential. This structure explains the enhanced ultraviolet continuum of quasars and the multiple redshifted spectral lines. \item We marvel at and believe that the existence of maximally gravitationally redshifted lines in quasar spectra are one of the most compelling proofs of Einstein's general theory of relativity \citep{1915SPAW.......844E, 1915SPAW...47..831E, 1915SPAW.......778E, 1915SPAW.......799E}. There is no doubt that there exist black holes with properties determined by the exact analytical solutions derived for non-rotating black holes by \citet{1916SPAW.......189S} and for rotating black holes by \citet{1963PhRvL..11..237K}. Quasars strongly support $z_g$ upto 1 as expected for maximally rotating black holes and rule out $z_g >1$. \item We suggest a method to separate $z_c$ and $z_{in}$ in a quasar spectrum. We assume that the lowest detected redshift of the Mg II absorption line $z_{MgII}$ in a quasar spectrum includes no contribution from $z_{in}$ and hence $z_c=z_{MgII}$. Using this, we find the remarkable result that the difference between the highest ($z_{em}$) and lowest ($z_{MgII}$) redshifts deduced in a quasar spectrum are $ < 1.25$. This proves the existence of a non-trivial intrinsic redshift contribution to the velocity of the lines. \item Spectral lines from quasars show a $z_g$ upto one. $z_g\sim0.5$ can be shown by lines arising in shells close to the Schwarzchild radius while $z_g\sim1$ is possible only from matter inside the ergosphere of a rotating black hole. This is the first time, to the best of our knowledge, that $z_g$ between 0.5 and 1 has been inferred in astronomical spectral lines providing irrefutable evidence to the existence of rotating black holes and matter inside its ergosphere. \item The gravitational redshift component can trivially explain several intriguing observational results on quasars. \item We suggest that the peculiar morphology of the ergosphere is taken up by matter within as it arranges itself in equipotential shells around the black hole. The matter inside the ergosphere will resemble a thick accretion disk which is often postulated to exist around active nuclei. \item The emission and absorption lines detected in a quasar spectrum are broadened due to the varying gravitational potential in the line forming region and hence cannot be used to estimate the mass of the black hole. \item We find that $z_g <0.5 $ for most BL Lac objects whereas $z_g$ upto one is shown by spectral lines from FSRQs - both comprising blazars. Thus blazars are indeed quasars with the structure shown in Figure \ref{fig12}. \item We show that GRBs are transient events in quasars consisting of $\gamma$ ray photons generated in an explosive thermonuclear reaction on the hot degenerate matter surface of the quasar. This explosion illuminates the degenerate matter surface giving rise to the multiband afterglow emission and a glimpse of the line spectrum of the quasar. It needs to be investigated whether this formation scenario explains all observed GRBs. \item We suggest the observed variability in quasars, especially in the ultraviolet continuum emission, can be unequivocally associated with energetic events on the degenerate matter surface - GRBs being the most energetic. Thus, the variability timescales alongwith the decoded shell structure from the observed gravitational redshifted lines should be able to uniquely constrain the quasar - black hole mass and the physical properties of the surrounding degenerate matter and line formingshells. \item The quasar model is applicable to other active nuclei with the variables being the separation between the black hole, degenerate matter surface and the line forming zones. \item We explain the band structure in the nuclear redshift-magnitude diagram shown by galaxies in the Coma cluster as being due to the effect of a gravitational redshift component and obscuration. This result strongly supports the presence of a supermassive black hole at the centres of all galaxies. \item Now that we have undeniable proof of the existence of matter inside the ergosphere of rotating black holes in quasars, it should be possible to further investigate how the energy of the black hole is tapped \citep[e.g.][]{1969NCimR...1..252P,1977MNRAS.179..433B}. \item We now revisit the four questions posed in Section 2 and answer them based on the results: (1) the spectral lines arise inside the quasar; (2) quasars exist in the same volume as other active galaxies; (3) the ultraviolet continuum arises on the degenerate matter surface and a gravitational instability on the same leads to the some of the observed variability; (4) quasars are isolated black holes surrounded by matter with comparable masses but smaller physical sizes than galaxies. \item In light of the conclusive results on the presence of an intrinsic component in the quasar redshifts presented here, the suggestion that nearby galaxies and quasars/active nuclei are related \citep{1967ApJ...148..321A,1974IAUS...58..199A} needs to be revisited. \item Now that we know that observed redshifts of quasars contain a sizeable redshift of non-cosmological origin, the existence and explanation of superluminal motions need to be revisited. For example, \citet{1971ApJ...170..207C} inferred that while the quasars 3C273 and 3C279 showed superluminal expansion, the jets in the Seyfert galaxy NGC 1275 show non-superluminal expansion and in M87 show no expansion. This kind of gradation, if found to be widespread, could be due to the wrong cosmological redshifts (and hence distance) which have been used for quasars (and to a lesser extent other active nuclei) and needs to be examined. \item The existence and implications of `dark' quasars should be examined. \item Possibility of a gravitational redshift contribution to the Hubble constant should be examined. \item This research then gives us new questions to ponder on - for example: How are metals synthesized in quasars ? What determines the separation between the event horizon and the line forming zones ? Where does the synchrotron emission arise and is it due to a shock set up by the instability on the degenerate matter surface ? What part of the structure of a quasar is common to the structure of any accreting black hole? Can we observe gravitationally redshifted lines from stellar mass black holes? \end{itemize} | 16 | 9 | 1609.01593 |
1609 | 1609.01841_arXiv.txt | \noindent The emergence of turbulence in shear flows is a well-investigated field. Yet, there are some lingering issues that have not been sufficiently resolved. One of them is the apparent contradiction between the results of linear stability analysis quoting a flow to be stable and experiments and simulations proving it to be otherwise. There is some success, in particular in astrophysical systems, based on Magneto-Rotational Instability (MRI), revealing turbulence. However, MRI requires the system to be weakly magnetized. Such instability is neither a feature of general magnetohydrodynamic (MHD) flows nor revealed in purely hydrodynamic flows. Nevertheless, linear perturbations of such flows are nonnormal in nature which argues for a possible origin of nonlinearity therein. The concept behind this is that nonnormal perturbations could produce huge transient energy growth (TEG), which may lead to non-linearity and further turbulence. However, so far, nonnormal effects in shear flows have not been explored much in the presence of magnetic fields. In this spirit, here we consider the perturbed visco-resistive MHD shear flows with rotation in general. Basically we recast the magnetized momentum balance and associated equations into the magnetized version of Orr-Sommerfeld and Squire equations and their magnetic analogues. We also assume the flow to be incompressible and in the presence of Coriolis effect solve the equations using a pseudospectral eigenvalue approach. We investigate the possible emergence of instability and large TEG in three different types of flows, namely, the Keplerian flow, the Taylor-Couette (or constant angular momentum) flow and plane Couette flow. We show that, above a certain value of magnetic field, instability and TEG both stop occurring. We also show that TEG is maximum in the vicinity of regions of instability in the wave number space for a given magnetic field and Reynolds number, leading to nonlinearity and plausible turbulence. Rotating shear flows are ubiquitous in astrophysics, especially accretion disks, where molecular viscosity is too low to account for observed data. The primary accepted cause of energy-momentum transport therein is turbulent viscosity. Hence, these results would have important implications in astrophysics. | The origin of linear instability and turbulence, and subsequent angular momentum transport in various classes of shear flows, specifically in astrophysical accretion disks, which are rotating shear flows, has not been explained completely yet. However, it is understood from observed data that, to explain the accretion in astrophysical disks, some sort of viscosity is required. In the absence of adequate molecular viscosity \cite{pringle1981}, turbulent viscosity was argued to play the main role in the accretion process by Shakura and Sunyaev \cite{shasun1973}. Nevertheless, a Keplerian accretion disk is linearly stable, thus proving it difficult to explain the origin of turbulence in the absence of any unstable linear perturbation. Similar problem exists in some laboratory flows. For example, plane Poiseuille flow becomes turbulent in the laboratory at a Reynolds number $Re\sim 1000$, whereas linear theory predicts it to be stable up to $Re=5772$. An even more severe discrepancy, which has a direct interest to astrophysics, occurs in the case of plane Couette flow, which is shown to be turbulent for $Re$ as small as $350$ in laboratory experiments and numerical simulations. However theoretical analysis shows it to be linearly stable for all $Re$ up to infinity. Subsequently, Balbus and Hawley applied the idea of Magneto-Rotational Instability (MRI) \cite{balbushawley1991}, established originally by Velikhov \cite{velikhov1959} and Chandrasekhar \cite{chandrasekhar1960}, to resolve the issue of instability and turbulence in magnetized flows and, hence, in some kinds of accretion disks. But the puzzle remains in laboratory flows which are colder and, hence, MRI would not work there. Moreover, to work MRI successfully, magnetic field strength has to be weak (weaker than a critical value depending on $Re$ \cite{nath2015}). Hence, for global purposes, a full scale exploration of magnetohydrodynamic (MHD) flows is needed. Exploration of MHD instabilities in various fluid systems is nothing new. The comprehensive descriptions of their various properties including eigenspectra of perturbation and stability are given in \cite{mhdbook1,fusion} in the limit of ideal MHD and in \cite{mhdbook2} in the presence of visco-resistive effects. The properties of eigenspectra and instability have also been explored to a great degree, even in two dimensions, in the context of tokamak fusion physics (see, e.g., \cite{mhdbook2}). Moreover, ideal MHD spectra for cylindrical plasma column were explored in order to investigate that how the local criteria govern the existence of the accumulating eigenmodes \cite{jpp,pp}. In the context of astrophysical accretion disks and other transonic flows, full scale MHD instability was found in radially stratified flows \cite{apjl} as well as in axisymmetric plasmas having poloidal flow speed exceeding critical slow magnetosonic speed \cite{ppastro}. In a completely different approach, MHD instability and plausible turbulence were also argued in accretion disks and other magnetized flows by computing various types of correlation of perturbations \cite{sujitamit1,sujitamit2}. Generally, below a certain critical value of $Re$ ($Re_c$), the linear stability analysis would predict a flow to be stable, but sometimes the most minutely controlled experiments would result in turbulence below $Re_c$ set by the theory. That exactly is being observed in laboratory experiments and numerical simulations of plane Poiseuille flow mentioned above, when its $Re_c=5772$ \cite{orszag1971}. Such a discrepancy would lead one to believe that simple linear stability analysis is probably not the best tool to enlighten the onset of turbulence. In a related field, Trefethen, Embree, Schmid and Henningson \cite{trefethenembree,schmidhenningson, orszag1972} explored the idea of nonnormality. Under this idea, it is shown that even in the complete absence of a linearly unstable mode, perturbations could exhibit `Transient Energy Growth (TEG)' \cite{reddyhenningson1993}. This happens when the eigenfunctions of a linear system are not completely orthogonal in nature and, because of that, certain combinations of the eigenfunctions and initial conditions, may develop a significant amplitude of (transient) energy growth, despite being stable overall. This form of growth, as the name suggests, occurs only for a short period of time, but its magnitude could be sufficient (depending upon the parameters of flows) to cause nonlinearity and plausible turbulence in fluid flows. In this work, we consider the visco-resistive (including fluid viscosity and magnetic diffusivity) MHD equations for three cases of flows: with and without the presence of Coriolis (rotational) effects, to explore their linear stability and TEG analyses. We precisely consider a small section of \begin{itemize} \item plane Couette flow, \item Keplerian flow, \item constant angular momentum flow or classic Taylor-Couette flow. \end{itemize} The second class of flow often mimics a small section of an astrophysical accretion disk and, hence, our results, as will be shown, have important implications in astrophysics. The present work is the sequel of the work \cite{nath2015} by the present group towards the application of nonnormality to MHD shear flows, including astrophysical flows. In the latest work, the authors approached the problem in the Lagrangian formulation. While that is an elegant way of approaching it, in particular for the purposes of that work, to uncover certain other physics, Eulerian approach is more useful. Hence, in the present work, we undertake the Eulerian approach to fulfill the underlying physics. Overall, the latest work \cite{nath2015} and the present one complement to each other, to understand the full picture of the problem. We begin with the description of model with basic equations in \S \ref{basiceq}, followed by perturbed fluid equations in \S \ref{perturbeq}. We then describe these equations in an eigenvalue formulation in \S \ref{eigform}, introduce and apply the concept of TEG to them in \S \ref{transen} and discuss the numerical considerations used to solve the problem (for eigenvalue formulation and TEG) in \S \ref{numcond}. Subsequently, we explore a simpler analytical scheme in \S \ref{anal}, which is useful to interpret and understand the numerical results presented in \S \ref{numrel}. Finally we end with a summary and conclusion in \S \ref{summa}. | \label{summa} We have explored and compared how linear instability and TEG may arise in MHD flows for Keplerian disk, constant angular momentum flow and plane Couette flow in terms of an eigenvalue formulation of the shearing box model. The system considered is the incompressible visco-resistive MHD flow following the Orr-Sommerfeld and Squire operator formulation, supplemented by the Coriolis effects and magnetic fields. In terms of spectral decomposition, such a system, by design, does not exhibit any fast magneto-acoustic modes and the underlying slow modes are degenerate with the Alfv\'{e}n modes which, in the presence of rotation, may also exhibit MRI. The basic trends in the system can be understood by a simple plane wave perturbation analysis. The incorporation of visco-resistive effects (by having a finite $Re$ and/or $Rm$) and non-axisymmetry (non-zero azimuthal wavenumber $k_y$) result in a variety of modes (modifications of the basic Alfv\'{e}n modes), which manifest physically as lesser number of unstable modes as well as lowered TEG. It seems that, in the case of stable systems, the amplitude of TEG is directly correlated with the number of slowly-decaying low-frequency modes, since these are the modes that may allow an optimal linear combination over sufficient timescales to exhibit substantial TEG. Perturbations with non-axisymmetric component (non-zero $k_y$) tend to get sheared by the flow, resulting in high-frequency Doppler-shifted oscillatory modes. These modes have a lower possibility of optimal linear combinations and, hence, do not show significant TEG. Since the flows under consideration have plane (and linear) shear, axisymmetric perturbations therein remain unaffected and exhibit instability or substantial TEG. We posit that high frequency oscillations do not allow the modes to have optimal linear combination over a given time range, which is necessary for large TEG. Perturbations with vertical component are affected by both rotation (leading to non-zero vorticity; essential for turbulence) and the background magnetic field strength. A certain finite range of these perturbations allows for instability and significant TEG. This range (which is also dependent on the value of background magnetic field) specifies where MRI and where significant TEG can occur. Beyond this range (above certain value of vertical wavenumber $k_z$ for a given magnetic field), the magnetic field stabilizes the flow and in the absence of vertical perturbation the flow becomes irrotational. In either of the cases, instability is reduced (having less number of unstable modes with lower growth rates) as well as TEG is decreased. Last, strong background magnetic fields tend to have a stabilizing effect on the perturbations, which can be understood by invoking the ``rod''-like nature of these fields, compared to the ``spring"-like nature of weak fields governing MRI. What is more interesting to note is that these strong fields also kill off TEG. The type of modes that we consider is limited by the assumption of incompressibility, the shearing box model (which ignores the effects of curvature) and plane wave perturbation in the azimuthal and vertical directions. A lot of work, which includes some of these consideration, but limited to the scope of linear stability analysis, is already present in literature \cite{jpp,apjl}. A more complete picture of MHD TEG may emerge with the study of compressible flows in cylindrical coordinates with more generalised perturbations. \\ \\ T.S.B. would like to thank the Department of Physics, Indian Institute of Science, Bangalore, for providing support to pursue his Master's thesis, where this research was conducted. A partial financial support from the project with research Grant No. ISTC/PPH/BMP/0362 is acknowledged. Finally thanks are due to the referees for their valuable critique and suggestions. | 16 | 9 | 1609.01841 |
1609 | 1609.02358_arXiv.txt | We have obtained new H~{\sc i} observations with the 100 m Green Bank Telescope (GBT) for a sample of 29 extremely metal-deficient star-forming Blue Compact Dwarf (BCD) galaxies, selected from the Sloan Digital Sky Survey spectral data base to be extremely metal-deficient (12 + log O/H $\leq$ 7.6). Neutral hydrogen was detected in 28 galaxies, a 97\% detection rate. Combining the H~{\sc i} data with SDSS optical spectra for the BCD sample and adding complementary galaxy samples from the literature to extend the metallicity and mass ranges, we have studied how the H~{\sc i} content of a galaxy varies with various global galaxian properties. There is a clear trend of increasing gas mass fraction with decreasing metallicity, mass and luminosity. We obtain the relation $M$(H~{\sc i})/$L_g$ $\propto$ $L_g^{-0.3}$, in agreement with previous studies based on samples with a smaller luminosity range. The median gas mass fraction $f_{gas}$ for the GBT sample is equal to 0.94 while the mean gas mass fraction is 0.90$\pm$0.15, with a lower limit of $\sim$0.65. The H~{\sc i} depletion time is independent of metallicity, with a large scatter around the median value of 3.4 Gyr. The ratio of the baryonic mass to the dynamical mass of the metal-deficient BCDs varies from 0.05 to 0.80, with a median value of $\sim$0.2. About 65\% of the BCDs in our sample have an effective yield larger than the true yield, implying that the neutral gas envelope in BCDs is more metal-deficient by a factor of 1.5--20, as compared to the ionized gas. | The formation and evolution of the first galaxies in the universe remains a key issue in cosmology. It is now thought that large massive star-forming galaxies form in a hierarchical manner from the assembly of smaller dwarf systems, through accretion and merger processes. These galaxy interactions trigger the formation of stars which enrich the interstellar gas in metals by stellar winds and supernovae. In this scenario, the number of extremely metal-deficient (XMD) dwarf galaxies should be high in the early universe but considerably smaller at the present epoch \citep{M12}. However, while much progress has been made in finding large populations of galaxies at high ($z\geq$3) redshifts \citep[e.g. ][]{Ste03,Ade05}, truly chemically unevolved galaxies remain elusive in the high-$z$ universe. The spectra of distant galaxies generally indicate the presence of a substantial amount of heavy elements, implying previous star formation and metal enrichment. The discovery of XMD galaxies at high $z$ may have to wait until the advent of the {\sl JWST} and 30 m-class ground-based telescopes. We adopt here a different approach. Instead of searching for high-$z$ metal-deficient objects, we focus our attention on XMD star-forming dwarf galaxies in the local universe. They are the most promising local proxies of chemically unevolved galaxies in the early universe, and are usually found among a class of dwarf galaxies undergoing intense bursts of star formation called Blue Compact Dwarf (BCD) galaxies \citep{Thu81}. The optical spectra of the BCDs are characterized by a blue continuum on which are superimposed strong narrow emission lines. XMD BCDs are very rare \citep*{Izo12}. For more than three decades, one of the first BCD discovered, I~Zw~18 \citep{Sar70}, held the record as the most metal-deficient emission-line galaxy known, with an oxygen abundance [O/H]= 12 + log O/H = 7.17$\pm$0.01 in its northwestern component and 7.22$\pm$0.02 in its southeastern component \citep{TI05} \citep[$\sim$ 3\% solar, adopting the solar abundance 12+logO/H = 8.76 of ][]{Ste15}. Only in 2005 has I~Zw~18 been superceded in its rank by SBS~0335--052W with a metallicity [O/H] = 7.12 \citep*{ITG05}. Because of the scarcity of XMD emission-line galaxies, we stand a much better chance of discovering them in very large spectroscopic surveys such as the Sloan Digital Sky Survey \citep[SDSS, ][]{Y00}. We have carried out a systematic search for such objects with [O/H] in the SDSS spectroscopic data release 7 (DR 7) \citep{Aba09}. Imposing cut-offs in metallicity and redshift results in a total sample of 29 XMD BCDs. Similar searches for XMD galaxies in the SDSS have been carried out by \citet{Mor11} and \citet{San16}. We found a total of 10 galaxies in common between the present sample and that of \citet{San16} so that the two samples are likely to have similar properties, and the H~{\sc i} characteristics discussed here probably apply to the \citet{San16} objects as well. The focus of this paper is the study of the neutral hydrogen content of these XMD objects. There have been previous H~{\sc i} studies of this type of extremely metal-poor galaxies. Thus, \citet{Fil13} have carried out a single-dish H~{\sc i} study with the Effelsberg radio telescope of a subsample of 29 XMD galaxies selected from the \citet{Mor11} list. We will use part of the data of those authors to supplement our own and will compare with our results with theirs when warranted. A handful of interferometric H~{\sc i} maps of other XMD galaxies have also been obtained by the Pune group with the Giant Metrewave radio telescope \citep*[see ][ and references therein]{Ekt09,Ekt10}. In Section 2, we define the XMD BCD sample and describe observations of their neutral hydrogen content with the Robert C. Byrd Green Bank Telescope (GBT) at the National Radio Astronomy Observatory~\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.}. This sample will be referred to hereafter as the GBT sample. Section 3 describes the H {\sc i} data along with derived ancillary data such as metallicities, star formation rates, stellar masses needed to study trends of the neutral gas content with other properties of the BCDs. In Section 4, we discuss several comparison samples compiled from the literature, useful for studying the H~{\sc i} content of star-forming galaxies over a wider range of heavy element abundances and stellar masses. Our total galaxy sample, composed of the GBT and three comparison samples, includes 151 objects and covers the extensive metallicity range 7.20 $\leq$ [O/H] $\leq$ 8.76. We study in Section 5 the correlations of various quantities. In particular, we analyze the dependence of the neutral gas mass to light ratio on metal abundance, and that of the gas mass fraction on stellar mass. We also discuss the chemical evolution of XMD BCDs. We summarize our conclusions in Section 6. Throughout this paper, we adopt the cosmological model caracterized by a Hubble constant $H_{0}$ = 73 km s$^{-1}$ Mpc$^{-1}$, a matter density parameter $\Omega_M$ = 0.27 and a dark energy density parameter $\Omega_\Lambda$ = 0.73 \citep{R11}. | New H {\sc i} observations with the Green Bank Telescope (GBT) are presented for a sample of 29 extremely metal-deficient star-forming Blue Compact Dwarf (BCD) galaxies. The BCDs were selected from the spectal database of Data Release 7 of the Sloan Digital Sky Survey (SDSS) to have a well detected [O{\sc iii}]$\lambda$4363 line (for direct abundance determination) and an oxygen abundance 12 + log O/H $\leq$ 7.6. Neutral hydrogen was detected in 28 galaxies, a 97\% detection rate. For each galaxy, we have derived ancillary data from the SDSS optical spectrum such as oxygen abundance, star formation rate, and stellar mass. Because of the narrow metallicity range of the GBT sample (the lower limit of 12 + log O/H is 7.35), we have also added published H {\sc i} and optical data for three complementary galaxy samples to extend the metallicity and mass range and study statistically how the H {\sc i} content of a galaxy varies with various global galaxian properties. We have found the following: 1) The lowest-luminosity lowest-metallicity galaxies have the largest neutral hydrogen mass to light ratios, following the relation $M$(H {\sc i})/$L_g$ $\propto$ $L_g^{-0.3}$, in good agreement with the dependence found in previous studies of galaxy samples with a smaller luminosity range. Our derived mass-metallicity relation is also in good agreement with those derived by other authors. 2) Metal-deficient low-mass dwarf galaxies are gas-rich. The median gas mass fraction of the GBT sample is 0.94, while its mean gas mass fraction is 0.90$\pm$0.15. The vast majority of the GBT galaxies have more than 65\% of their baryonic mass in gaseous form. The existence of a lower limit, also found for larger dwarf samples, puts stringent constraints on feedback mechanisms in low-mass galaxies as they should not remove all of the galaxy's atomic gas. 3) The H {\sc i} depletion time is independent of metallicity or stellar mass. Although there is a large scatter about the median value of 3.4 Gyr, the relative constancy of the gas depletion time implies that external processes or feedback mechanisms that control the gas supply are important for regulating star formation in both low- and high-mass galaxies. 4) The ratio of the baryonic mass to the dynamical mass varies over a wide range, from 0.05 to 0.80, with a median value of $\sim$0.2 and no dependence on the dynamical mass. 5) About 35\% of the GBT galaxies have an effective yield less than the true yield, which can be understood as the result of the loss of metals due to supernova-driven outflows, and/or the consequence of dilution by inflows of metal-poor gas. However, the remaining 65\% show an effective yield larger than the true yield. This can be understood if the metallicity of the neutral gas is lower than the metallicity of the ionized gas by a factor $\sim$1.5--20, as UV absorption studies of BCDs also show. | 16 | 9 | 1609.02358 |
1609 | 1609.01737.txt | Using spatially resolved spectroscopy from SDSS-IV MaNGA we have demonstrated that low ionisation emission line regions (LIERs) in local galaxies result from photoionisation by hot evolved stars, not active galactic nuclei, hence tracing galactic region hosting old stellar population where, despite the presence of ionised gas, star formation is no longer occurring. LIERs are ubiquitous in both quiescent galaxies and in the central regions of galaxies where star formation takes place at larger radii. We refer to these two classes of galaxies as extended LIER (eLIER) and central LIER (cLIER) galaxies respectively. cLIERs are late type galaxies primarily spread across the green valley, in the transition region between the star formation main sequence and quiescent galaxies. These galaxies display regular disc rotation in both stars and gas, although featuring a higher central stellar velocity dispersion than star forming galaxies of the same mass. cLIERs are consistent with being slowly quenched inside-out; the transformation is associated with massive bulges, pointing towards the importance of bulge growth via secular evolution. eLIERs are morphologically early types and are indistinguishable from passive galaxies devoid of line emission in terms of their stellar populations, morphology and central stellar velocity dispersion. Ionised gas in eLIERs shows both disturbed and disc-like kinematics. When a large-scale flow/rotation is observed in the gas, it is often misaligned relative to the stellar component. These features indicate that eLIERs are passive galaxies harbouring a residual cold gas component, acquired mostly via external accretion. Importantly, quiescent galaxies devoid of line emission reside in denser environments and have significantly higher satellite fraction than eLIERs. Environmental effects thus represent the likely cause for the existence of line-less galaxies on the red sequence. | \label{intro} Large spectroscopic galaxy surveys (SDSS, \citealt{York2000}, GAMA, \citealt{Driver2011}, zCOSMOS, \citealt{Lilly2007}) have demonstrated that galaxies are strongly bimodal in several of their fundamental properties, including colours, star formation rates, stellar ages, morphology and gas content \citep{Strateva2001, Blanton2003, Baldry2004, Baldry2006, Wyder2007, Schawinski2007, Blanton2009, Peng2010, Thomas2010}. These observations support a scenario where some physical processes lead galaxies to transition from the `star forming' (SF) blue cloud to the passive `red sequence', causing the eventual shut-down (or `quenching') of star formation on a galaxy-wide scale. The physics driving this transformation is currently hotly debated and has not yet been unambiguously identified, although it is likely that more than one mechanism is at play \citep{Schawinski2007, Peng2010, Schawinski2014, Mendel2013, Peng2015, Smethurst2015}. A growing body of evidence points towards at least two quenching pathways: an \textit{environment-dependent} channel, acting differentially on central and satellite galaxies, and a \textit{mass-dependent} channel, which may be associated with internal processes \citep{Silk1977, Rees1977,Kauffmann2006}. This inference is supported by the separability of the effects of mass and environment on the passive fraction \citep{Baldry2006, Peng2010, Thomas2010}. These two quenching pathways have each been associated with several physical processes: ram pressure stripping \citep{Larson1980}, harassment and strangulation \citep{Gunn1972, Peng2015} for the environment-related channel and radio mode feedback from active galactic nuclei (AGN, \citealt{Dekel2006, Croton2006, Cattaneo2009}), halo mass shock heating \citep{Birnboim2003, Keres2005, Dekel2006, Keres2009} or morphological quenching \citep{Martig2009, Genzel2014, Forbes2014} for the mass-dependant channel. It has also been suggested that both mass and environmental quenching may be manifestations of the same quenching channel \citep{Knobel2015, Carollo2016}. Observationally much work has been dedicated to the study of environmental effects, demonstrating that local environment correlates well with morphology \citep[e.g][]{Dressler1980, Bamford2009}, colour \citep{Blanton2005a, Baldry2006, Peng2010} or gas content \citep{Cortese2011, Boselli2014}. The direct effect of the hot cluster atmosphere on infalling galaxies (ram pressure stripping) is further observed both via trails of H\textsc{i} \citep[e.g][]{Chung2009} and ionised gas \citep[e.g][]{Fumagalli2014, Poggianti2016}. Regarding the mass quenching regime, the observational picture is much more complex. Several studies agree on the fact that the presence of a massive bulge is the single best observational parameter correlating with the probability of a galaxy to be passive \citep{Pasquali2012, Cheung2012, Wake2012, Bluck2014}. Since black hole mass is correlated to the bulge mass \citep{Marconi2003, Haring2004, McConnell2013}, it can be expected that galaxies with more massive bulges experience a larger amount of energetic feedback from their black holes. The details of the coupling between the black hole and its host galaxy are, however, not currently understood. As recently reviewed by \cite{Heckman2014}, AGN appear as a bimodal population that can be broadly divided into radiative-mode (including Seyfert galaxies and quasars) and radio-mode AGN. In the low redshift Universe, no direct correlation is observed between radiative-mode AGN and the shutdown of star formation, as radiative-mode AGN are often associated with the presence of star formation in the central regions of galaxies (inner few kpc). While radiative-mode AGN are observed to power large scale outflows \citep{Maiolino2012, Cicone2014, Carniani2015}, the duty cycle and time-scales of these phenomena is uncertain, thus making their impact on the exhaustion of the cold gas reservoir in galaxies questionable. Radio mode AGN, on the other hand, are observed to inflate large bubbles of ionised gas in massive haloes, thus playing an important contribution to offsetting the cooling of the gas in the halo \citep{Fabian2012}. The extension of this feedback mechanisms to lower masses (halo mass $< 10^{12}~M_\odot$) is observationally uncertain. Both environmental and mass quenching (through radio-mode AGN heating) channels are included in recent semi-analytical models \citep{Croton2006, Bower2006} and in cosmological hydrodynamical simulations via sub-grid prescriptions \citep{Gabor2012, Vogelsberger2014, Schaye2015} in order to match the stellar mass function and the bimodality of the observed galaxy population. In recent years, models of galaxy formation and evolution have become able to qualitatively reproduce the main features of the bimodal galaxy population \citep[e.g.][]{Sales2015, Trayford2016}. Namely, in simulations environmental quenching appears to be the main cause of growth of the red sequence at low masses, while AGN radio-mode heating is responsible for the presence of high-mass central passive galaxies. However, while qualitative agreement with observations is promising, this success does not imply that the sub-grid recipes implemented for AGN feedback are physically correct. Observationally, further insight into the quenching process could be obtained by a statistical study of how it proceeds inside galaxies: `inside out' versus `outside in' scenarios \citep[e.g][]{Tacchella2015, Li2015}, and what role the different morphological components (disc, bar, bulge) have in the process \citep{Martig2009, Masters2012}. Current integral field spectroscopy (IFS) galaxy surveys (CALIFA, \citealt{Sanchez2012a}; SAMI, \citealt{Allen2015}; MaNGA, \citealt{Bundy2015}) promise to revolutionise the observational picture, by enabling the study of the relation between star formation and quiescence on resolved scales. Nebular emission lines, originating from SF regions, are the ideal tool to trace current ($< $100 Myr old) star formation on the kpc scales \citep{Kennicutt1998} probed by a survey like MaNGA (mean redshift $<z> \sim 0.03$), provided other contributions to line emission (AGN, old hot stars etc) can be disentangled. In this paper we make use of a large sample of $\sim$ 600 galaxies observed as part of the Sloan digital sky survey IV (Blanton et al., submitted) Mapping nearby galaxies at Apache Point Observatory (MaNGA) survey, and build on the analysis presented in \cite{Belfiore2016a} (henceforth \textit{Paper I}) to define spatially resolved `quiescence' and make an unbiased census of star forming and quiescent regions within the z$\sim$0 galaxy population. This work is structured as follows. In Sec. \ref{sec2} we briefly recap the relevant properties of the MaNGA data and our spectral fits. In Sec. \ref{sec3} we summarise the main findings of \textit{Paper I}, including our new emission-line based galaxy classification, the importance of low ionisation emission-line regions (LIERs) in the context of understanding quiescence in galaxies and the stellar populations properties of the different galaxy classes. In Sec. \ref{sec4} we explore the role of LIER galaxies within the bimodal galaxy population, while in Sec. \ref{sec5} we gather evidence regarding the origin of the ionised gas in LIERs. Finally in Sec. \ref{sec6} we discuss possible evolutionary scenarios for LIER galaxies. We conclude in Sec. \ref{concl}. Throughout this work redshifts, optical and UV photometry and stellar masses are taken from the MaNGA targeting catalogue, based on an extended version of the NASA Sloan Atlas (NASA-Sloan catalogue (NSA {\tt v1\_0\_1}\footnote{http://www.sdss.org/dr13/manga/manga-target-selection/nsa/}, \citealt{Blanton2011}). Stellar masses are derived using a \cite{Chabrier2003} initial mass function and using the {\tt kcorrect} software package (version {\tt v4\_2}, \citealt{Blanton2007}) with \cite{Bruzual2003} simple stellar population models. All quoted magnitudes are in the AB system and k-corrected to rest frame after correction for galactic extinction. Effective radii are measured from the SDSS photometry by performing a S\'ersic fit in the r-band. Single fibre spectroscopic data from the SDSS Legacy Survey data release 7, \citep{Abazajian2009} for the main galaxy sample targets \citep{Strauss2002} are referred to as legacy SDSS. When quoting luminosities, masses and distances we make use of a $\Lambda$CDM cosmology with $\Omega_m=0.3$, $\Omega_\Lambda=0.7$ and $\rm H_0=70 \mathrm{ \ km^{-1} s^{-1}Mpc^{-1}}$. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % SECTION 2 %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%% FIGURE 2 \begin{figure*} \includegraphics[width=1.0\textwidth, trim=0 0 0 0, clip]{display4.pdf} \caption{Example of the different properties of SF (top three), cLIER (middle three) and eLIER (bottom three) galaxies. Each row represents a different galaxy (labelled with its MaNGA-ID on the left). The columns represent respectively (1) SDSS g-r-i composite image with MaNGA bundle superimposed, (2) MaNGA H$\alpha$ map, (3) MaNGA EW(H$\alpha$) map, (4) MaNGA map of $\rm D_N(4000)$, (5) MaNGA map of $\rm H\delta_A$, (6) BPT map of the MaNGA galaxies based on the [SII] BPT diagram. SF regions are colour-coded in blue, LIER in orange, Seyfert-like in red, (7) velocity field of the stars, (8) velocity field of the gas. For all maps, the colour bar increases from blue to red and the values corresponding to the minimum and maximum of the colour-bars are annotated on the top right corner of each map. The average PSF for the MaNGA data is shown in the bottom right corner of the maps.} \label{fig2} \end{figure*} %%%%%%% | \label{concl} In this paper we study the properties of galaxies based on their spatially resolved ionised gas emission. We make use of spatially resolved spectroscopy for a sample of 586 galaxies from SDSS-IV MaNGA, and classify galaxies according to the classification scheme introduced in \cite{Belfiore2016a}, which properly accounts for the ubiquitous presence of low ionisation emission line regions (LIERs). Within this framework non-active galaxies are classified in four classes: (1) star forming galaxies (SF), (2) galaxies characterised by LIER emission at small galactocentric radii and star formation at larger radii (cLIERs), (3) galaxies characterised by LIER emission at all radii where line emission is detected (eLIERs), (4) Line-less galaxies, operationally defined to have average EW(H$\alpha$) $<$ 1 \AA\ within 1.0 $\rm R_e$. In this work we place these galaxy classes within the context of the evolving galaxy population by studying their integrated colours and SFR, morphology and kinematics. We complement the MaNGA data with environmental information for the much larger MaNGA parent sample ($\sim$ 30~000 galaxies with legacy SDSS spectroscopy) to be able to disentangle the effect of mass and environment on line emission on the red sequence. In the following we summarise the fundamental properties of the two new galaxy classes we have focused on. \paragraph*{Central LIER galaxies (cLIERs)} \begin{enumerate} \item{Stellar populations: cLIERs span a wide range of stellar ages, with old central LIER regions and outer regions characterised by young stellar populations.} \item{Colours and SFR: cLIERs lie preferentially in and around the green valley in UV-optical colours, although they appear increasingly consistent with the red sequence using optical colours alone ($u-r$ or $g-r$). In the SFR-$\rm M_\star$ diagram cLIERs lie 0.8 dex below the main sequence of star forming galaxies, albeit with a large scatter.} \item{Morphology: cLIER galaxies are visually classified as disc galaxies, but have larger bulge fraction than star forming galaxies of the same mass, as traced via photometric indices and spectroscopic stellar velocity dispersion measurements.} \item{Kinematics: cLIERs display regular disc-like kinematics in both gas and stars, which are mutually aligned.} \item{Environment: cLIERs live in slightly denser environments than SF galaxies for $\rm log(M_\star/M_\odot) = 10.5- 11.5$, but have the same satellite fraction.} \end{enumerate} cLIERs are consistent with originating from SF galaxies that are slowly quenching. The inside-out quenching pattern is consistent with both gas exhaustion and/or a decrease in the star formation efficiency. In either case, the presence of a bulge is closely connected to the lack of central star formation in these galaxies. The star forming disc is the likely source of the LIER emitting gas present in the central regions. \paragraph*{Extended LIER galaxies (eLIERs)} \begin{enumerate} \item{Stellar populations: eLIERs have old stellar populations, indistinguishable from those of line-less galaxies.} \item{Colours and SFR: eLIERs lie on the red sequence in colour-magnitude diagrams, with only a small contamination to the green valley.} \item{Morphology: eLIER galaxies are visually classified as early type galaxies and have similar bulges to those of line-less galaxies of the same stellar mass.} \item{Kinematics: eLIERs display a variety of gas kinematics, but generally present well-defined rotation/flows on kpc scales. The velocity field of gas and stars are often observed to be misaligned ($65 \pm 7 \%$ of eLIERs are misaligned by more than 30$^\circ$).} \item{Environment: eLIERs live in denser environments than SF and cLIERs at high masses ($\rm M_\star > 10^{10.5} M_\odot$), as expected from the well-known correlation between colour and passive fraction. More importantly eLIERs tend to live in less dense environments than line-less galaxies of the same mass. At the high mass end, line-less galaxies are mostly centrals in high mass haloes, while at the low mass end line-less galaxies are mostly satellites.} \end{enumerate} eLIERs are red sequence galaxies with residual cold gas, acquired mostly via external sources. While this gas does not form stars, it is illuminated by the diffuse ionisation field from hot old stars and shines with LIER line ratios. LIER emission is suppressed in high density environments, likely because of lack of gas, which has either been consumed due to the shut down of inflows or otherwise heated and/or tidally stripped by the hot halo gas. % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % ACKNOWLEDGEMENTS %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 16 | 9 | 1609.01737 |
1609 | 1609.05578_arXiv.txt | We present the Large Area Radio Galaxy Evolution Spectroscopic Survey (LARGESS), a spectroscopic catalogue of radio sources designed to include the full range of radio AGN populations out to redshift $z\sim0.8$. The catalogue covers $\sim800$\,deg$^2$ of sky, and provides optical identifications for 19,179 radio sources from the 1.4\,GHz Faint Images of the Radio Sky at Twenty-cm (FIRST) survey down to an optical magnitude limit of $i_{\rm mod} < 20.5$ in Sloan Digital Sky Survey (SDSS) images. Both galaxies and point-like objects are included, and no colour cuts are applied. In collaboration with the WiggleZ and Galaxy And Mass Assembly (GAMA) spectroscopic survey teams, we have obtained new spectra for over 5,000 objects in the LARGESS sample. Combining these new spectra with data from earlier surveys provides spectroscopic data for 12,329 radio sources in the survey area, of which 10,856 have reliable redshifts. 85\% of the LARGESS spectroscopic sample are radio AGN (median redshift $z=0.44$), and 15\% are nearby star-forming galaxies (median $z=0.08$). Low-excitation radio galaxies (LERGs) comprise the majority (83\%) of LARGESS radio AGN at $z<0.8$, with 12\% being high-excitation radio galaxies (HERGs) and 5\% radio-loud QSOs. Unlike the more homogeneous LERG and QSO sub-populations, HERGs are a heterogeneous class of objects with relatively blue optical colours and a wide dispersion in mid-infrared colours. This is consistent with a picture in which most HERGs are hosted by galaxies with recent or ongoing star formation as well as a classical accretion disk. | Over the past fifteen years, large surveys at optical, infrared and radio wavelengths have allowed us to make significant progress in understanding the typical radio properties of galaxies in the local and distant Universe. Two large-area radio surveys carried out by the Very Large Array (VLA) operated by the National Radio Astronomy Observatory (NRAO), the Faint Images of the Radio Sky at Twenty-cm \citep[FIRST;][]{becker95} and the NRAO VLA Sky Survey \citep[NVSS;][]{condon98} have been particularly influential. Both are 1.4 GHz continuum radio surveys covering a large fraction of the sky down to milli-Jansky flux densities. The high resolution and positional accuracy of the FIRST survey is complemented by the lower resolution of NVSS, which has better surface brightness sensitivity. Several studies \citep[e.g.][]{sadler02,hopkins03,best05,mauch07,best12} have matched NVSS and FIRST radio sources to counterparts in the optical or infrared. These optical/infrared identifications, combined with spectroscopic information such as redshifts, emission line and absorption line measurements, have advanced our understanding of the physical processes responsible for radio emission from nearby galaxies. For extragalactic radio sources, the radio continuum emission may arise from either an active galactic nucleus (AGN) or processes related to star formation. In star-forming galaxies, the observed radio emission is usually dominated by synchrotron emission from relativistic electrons accelerated by supernova remnants in \textsc{Hii} regions, with a smaller contribution from thermal free-free emission \citep{condon92}. The short-lived massive stars in the \textsc{Hii} regions of star-forming galaxies photoionize the surrounding gas and produce a characteristic pattern of emission lines in the observed spectrum. Spectroscopic studies of radio AGN reveal two main populations: those with prominent optical emission lines, and those with weak or no emission lines \citep{longair79,laing94}. We follow current practice and refer to the first (strong emission-line) population as high-excitation radio galaxies (HERGs) and the second as low-excitation radio galaxies (LERGs). The difference between these two populations is thought to reflect differences in the accretion efficiency of gas onto the central black hole \citep{hardcastle07}. A comprehensive review of the properties of the two classes of radio AGN is given by \cite{heckman14}. In the current paradigm, the HERGs undergo {\em cold-mode} (also known as {\it radiative mode}) accretion, characterised by a high accretion efficiency such that gas is accreted rapidly onto the galaxy's central black hole. This allows the formation of a radiatively-efficient accretion disk that photoionizes the surrounding gas to produce the observed high-excitation emission lines. The term cold-mode refers to the past temperature of the gas, which in this case has never reached the virial temperature of the halo \citep{keres09}. The LERGs on the other hand undergo {\em hot-mode} (also known as {\it jet-mode}) accretion, where the gas has at least reached the virial temperature in the past and is generally cooling from a surrounding hot X-ray corona. This is an inefficient accretion process without a radiatively efficient accretion disk, so the optical spectra of LERGs show weak or no emission lines. Hot-mode accretion is expected to occur in high halo-mass systems \citep[$>2-3 \times 10^{11} M_{\sun}$;][]{keres09}, particularly at low redshift, and cold-mode accretion in lower-mass systems over a wider range in redshift \citep{hardcastle07,van-de-voort11}. Recent observational studies of the properties \citep{best05a,smolcic09a,janssen12}, environments \citep{best04,bardelli10,gendre13,sabater13} and evolution \citep{smolcic09,best12} of HERGs and LERGs appear to confirm this picture, showing that HERGs are typically found in lower-mass galaxies with younger stellar populations, and in poorer environments than the LERGs, which are typically in the most massive galaxies, with an old stellar population, and found in rich environments. \ctable[ notespar, star, cap = {Spectroscopic survey regions}, caption={Regions covered by optical spectroscopic surveys. The coverage and overlap was calculated using the Virtual Observatory footprint service \citep[][http://www.voservices.net/footprint]{budavari07}.}, label={tab:regions}, ]{l r rr rr r rc cc}% { \tnote[a]{The official WiggleZ limit for this region has a maximum $\delta=8.0$; the additional 0.1 was mistakenly put into the original search. However, since the WiggleZ pointings included this extra small area, we include these objects as well.}\\ \tnote[b]{The actual 2SLAQ regions are several small strips along the equatorial region, but for simplicity we adopt the two large pseudo-2SLAQ strips shown here.} } { \FL & \multicolumn{1}{c}{Field} & \multicolumn{2}{c}{R.A. (deg)} & \multicolumn{2}{c}{$\delta$ (deg)} & Total Area & \multicolumn{2}{c}{FIRST-SDSS overlap} &\multicolumn{2}{c}{Spectral completeness to $i=20.5$}\\ \cline{8-9} \cline{10-11} Survey & \multicolumn{1}{c}{ID} & min & max & min & max & (deg$^{2}$) & (deg$^{2}$) & Fraction & Spectrum & Redshift \\ &&&&&&&&& observed & success rate \\ \hline\hline WiggleZ & & \\ & 0h & 350.1 & 359.1 & $-$13.4 & $+$1.8 & 136.0 & 44.7 & 0.33 & 55\% & 89\% \\ & 1h & 7.5 & 20.6 & $-$3.7 & $+$5.3 & 118.3 & 32.7 & 0.28 & 64\% & 95\% \\ & 3h & 43.0 & 52.2 & $-$18.6 & $-$5.7 & 116.0 & 7.2 & 0.06 & 55\% & 86\% \\ & 9h & 133.7 & 148.8 & $-$1.0 & $+$8.1\rlap{$^{a}$} & 137.8 & 136.7 & 0.99 & 71\% & 84\% \\ & 11h & 153.0 & 172.0 & $-$1.0 & $+$8.0 & 172.1 & 172.1 & 1.00 & 66\% & 84\% \\ & 15h & 210.0 & 230.0 & $-$3.0 & $+$7.0 & 201.7 & 200.0 & 0.99 & 63\% & 88\% \\ & 22h & 320.4 & 330.2 & $-$5.0 & $+$4.8 & 96.2 & 24.5 & 0.25 & 73\% & 88\% \\ GAMA & & \\ & 9h & 129.0 & 141.0 & $-$1.0 & $+$3.0 & 48.2 & 48.2 & 1.00 & 91\% & 94\% \\ & 12h & 174.0 & 186.0 & $-$2.0 & $+$2.0 & 48.2 & 48.2 & 1.00 & 86\% & 88\% \\ & 15h & 211.5 & 223.5 & $-$2.0 & $+$2.0 & 48.2 & 48.2 & 1.00 & 86\% & 89\% \\ 2SLAQ\tmark[b] & & \\ & - & 123.0 & 230.0 & $-$1.259 & $+$0.840 & 325.0 & 301.6 & 0.93 & 65\% & 91\% \\ & - & 309.0 & 59.70 & $-$1.259 & $+$0.840 & 347.9 & 224.3 & 0.64 & 57\% & 93\% \LL } The most powerful radio sources are known to undergo strong cosmic evolution, with their volume density at redshift $z\sim2$\ being up to a thousand times higher than it is today \citep[e.g.][]{longair66,dunlop90}. The cosmic evolution of lower-power radio AGN appears to be much less rapid \citep{sadler07,donoso09,simpson12}, but is only just starting to be mapped out separately for the HERG and LERG sub-populations beyond the local Universe \citep{best14}. Our aim in undertaking the work described in this paper was to produce a new, large and complete spectroscopic radio-source catalogue that would allow us to track the HERG and LERG populations in detail over a wide range in radio luminosity back to at least redshift $z\sim0.8$ (i.e. a lookback time almost half the age of the Universe) as well as studying the radio galaxy and radio-loud QSO populations across a common range in redshift. There is growing evidence that the redshift evolution of the HERG and LERG populations is very different \citep[e.g.][]{best12,simpson12,best14}, and that observed luminosity-dependent cosmic evolution of the radio luminosity function is driven mainly by the different cosmic evolution of these two populations \citep{heckman14}. One key motivation for this new study arose from earlier work on the evolving radio AGN luminosity function carried out by \cite{sadler07} and \cite{donoso09}. These authors used relatively large samples of radio-detected AGN (391 objects in the \cite{sadler07} spectroscopic sample; 14,453 objects with photometric redshifts in the larger-area \cite{donoso09} sample) to measure radio luminosity functions in the redshift range $0.4<z<0.7$ with unprecedented accuracy. Both samples were photometrically selected to target luminous red galaxies \citep[LRGs; ][]{eisenstein01} but exclude blue galaxies with ongoing star-formation. \cite{sadler07} explicitly noted that the rate of cosmic evolution measured for low-power radio galaxies in their study was only a lower limit, since the LRG sample they used had a strict colour cut-off, whereas no such colour restriction was applied to the $z\sim0$\ radio galaxy sample used as the local benchmark. By compiling a new sample of distant radio AGN without any pre-selection on colour, we wanted to find out whether there is indeed a significant population of `blue' radio galaxies in the distant Universe and (if so) how their properties compare with the better-studied population of `red' radio galaxies. The data catalogue presented in this paper includes over 10,000 spectroscopically-observed radio sources, with a median redshift of $z\sim0.44$ for the radio AGN which make up $\sim85$\%\ of the sample. Our sample of 2281 radio-source spectra at $0.5<z<1$ represents an order-of-magnitude increase over previous spectroscopic samples in this redshift range. For example, the recent \cite{best14}\ measurement of the radio luminosity function out to $z=1$ used a catalogue of 211 radio-loud AGN at $0.5<z<1.0$, while the \cite{simpson12} measurement used $\sim$100 spectroscopically-observed objects in the same redshift range (supplemented by a similar number of photometric redshift estimates). A companion paper by \cite{pracy06} uses the dataset presented here to make new measurements of the evolving radio luminosity functions of HERGs and LERGs out to redshift $z\sim0.8$. We describe the optical and radio catalogues used to compile our sample in \S\ref{sec:catalogues} and the radio-optical matching process in \S\ref{sec:sdss_first_matching} and \S\ref{sec:nvss_matching}. The spectroscopic follow-up program is discussed in \S\ref{sec:spec_data}, and the related completeness analysis presented in \S\ref{sec:spec_completeness}. \S\ref{sec:spec_class} describes the identification of star-forming galaxies and the classification of high- and low- excitation radio galaxies, and \S8 presents the full data catalogue. The full sample and some sub-samples are characterised in \S\ref{sec:sample_character}, while \S10 compares the properties of a matched sample of HERG and LERG host galaxies. Finally, we present a summary of the LARGESS sample properties in \S\ref{sec:summary1}. Throughout this paper we adopt the cosmological parameters $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\Lambda}=0.7$ and $\Omega_{m}=0.3$. All optical magnitudes are corrected for Galactic dust extinction and {\em k}-corrected using the \textsc{kcorrect} code \citep{blanton07}. An analysis of a subset of the LARGESS-GAMA (Galaxy And Mass Assembly) sample \citep{hardcastle13} found a mean radio spectral index (between 325 MHz and 1.4 GHz) of $\alpha = -0.7$ (where $S_{\nu} \propto \nu^{\alpha}$), and we adopt this value to calculate a radio {\em k}-correction. \begin{figure*} \centering \includegraphics[width=0.85\textwidth]{./figures/sample_phot2_small.eps} \includegraphics[width=0.85\textwidth]{./figures/sample_spec_all2.eps} \caption[ The LARGESS sky coverage]{Sky area covered by the LARGESS radio sample: (top) sky distribution of the full photometric catalogue of 19,179 radio-sources matched to optical objects with $i<20.5$\,mag, (bottom) distribution of the 10,764 objects in the final spectroscopic catalogue that currently have good-quality optical spectra and redshifts. Points are colour-coded according to the source of the redshift measurement: GAMA (dark blue), WiggleZ (magenta), SDSS (cyan), 2SLAQ (red).} \label{fig:coverage:specradio} \end{figure*} | \label{sec:summary1} We have compiled the Large Area Radio Galaxy Evolution Spectroscopic Survey (LARGESS) catalogue, a new, large dataset designed to study the evolution and environments of high- and low- excitation radio galaxies (HERGs and LERGs) out to $z\sim0.8$. The catalogue contains 19,179 radio sources detected by the FIRST survey that have an SDSS optical counterpart with $i_{\rm mod} < 20.5$. We have obtained new good-quality optical spectra for over 5,000 LARGESS objects, and adding these to lower-redshift archival data means that 10,856 LARGESS objects currently have a reliable optical redshift. The spectroscopic completeness varies across the survey area, and is highest ($\sim90$\%) in the three equatorial GAMA fields. The median redshift of radio AGN in the LARGESS catalogue is $z\sim0.44$, and the bright end of the radio luminosity function (P$_{1.4}> 10^{24}$\,W\,Hz$^{-1}$) is well-sampled for all the main radio AGN populations over the redshift range $0<z<0.8$. The catalogue also contains a minority population of nearby ($z<0.3$) star-forming galaxies. The great majority of LARGESS radio sources are compact (only 10\% have complex or multi-component radio structures on scales larger than a few arcsec), which makes the optical matching process straightforward in most cases. We estimate that our final matched sample is 95\% complete and has 94\% reliability. The LARGESS sources typically have between 5\% and 25\% of their 1.4\,GHz radio emission in a diffuse component that is seen by the NVSS survey but missed by the higher-resolution FIRST images. This needs to be taken into account in any subsequent analysis. In particular, as discussed in \S4.3, any analysis requiring accurate flux densities for extended radio sources should impose a $S^{\rm FIRST}_{\rm tot} \geq 3.5$ mJy limit for the LARGESS catalogue. For objects with a reliable redshift from the SDSS, GAMA and WiggleZ surveys, we made both a visual and an automated classification of the spectrum. These two classification schemes were combined to provide a single classification, and allow us to distinguish between (i) Galactic stars, (ii) radio galaxies with weak or no emission lines (LERGs), (iii) galaxies with emission lines or radio emission generated from star formation (SF), and (iv) radio galaxies where both emission line and radio emission arise from the AGN (broad emission lines, AeB and narrow emission lines, HERGs). While the optical counterparts of most radio AGN are massive red galaxies at all redshifts out to $z=0.8$, we find that at least 15\% of radio AGN at $0.4<z<0.8$ are hosted by blue galaxies which appear likely to have ongoing star formation. It is notable that at least half of the HERG population in this redshift range are `blue radio galaxies' rather than passively-evolving red galaxies. Roughly 8\% of the radio AGN population in the same redshift range are radio-loud QSOs with broad Balmer emission lines in their optical spectra. We find that LERGs, HERGs and AeB (QSO) objects have very different broad-band spectral-energy distributions (SEDs). The optical and mid-IR properties of LERGs are generally similar to those of luminous red galaxies dominated by an old stellar population, while the AeB (QSO) objects are consistent with a power-law SED. In contrast, the HERGs are a much more heterogeneous population in terms of both their optical and mid-IR colours. We used a matched sample of HERGs and LERGs to compare their physical properties, and found that that HERGs have bluer colours (probably indicative of a younger stellar population) compared to LERGs of similar mass, redshift, and radio luminosity. The concentration index is also higher for LERGs than HERGs, suggesting that LERGs are hosted by a higher fraction of elliptical galaxies than the HERGs. Both of these results are independent of redshift. The matched sample of LERGs and HERGs show no significant difference in the linear sizes in any redshift bin. These results confirm results of previous studies, but also extend them into a higher redshift range. By comparing the fraction with a reliable detection in various WISE mid-infrared bands, we found that HERGs are significantly more likely to be detected in the W3 and W4 bands than a matched sample of LERGs, both locally and at higher redshifts.. From this, we inferred the presence of a radiatively efficient accretion disk and dusty torus in the HERGs, which is absent in the LERGs. This is consistent with theories that suggest a dichotomy in the accretion mode of HERGs and LERGs. \vspace{1cm} \noindent{\bf ACKNOWLEDGEMENTS} \noindent JHYC would like to acknowledge the funding provided by the University of Sydney via an Australian Postgraduate Award, the Postgraduate Research Scholarship Scheme and the William \& Catherine McIlrath Scholarship. JHYC would also like to acknowledge the support provided by the Astronomical Society of Australia. EMS and SMC acknowledge the financial support of the Australian Research Council through Discovery Project grants DP1093086 and DP 130103198. SMC acknowledges the support of an Australian Research Council Future Fellowship (FT100100457). We thank Philip Best and Dick Hunstead for helpful comments. GAMA is a joint European-Australasian project based around a spectroscopic campaign using the Anglo-Australian Telescope. The GAMA input catalogue is based on data taken from the Sloan Digital Sky Survey and the UKIRT Infrared Deep Sky Survey. Complementary imaging of the GAMA regions is being obtained by a number of independent survey programs including GALEX MIS, VST KIDS, VISTA VIKING, WISE, Herschel-ATLAS, GMRT and ASKAP providing UV to radio coverage. GAMA is funded by the STFC (UK), the ARC (Australia), the AAO, and the participating institutions. The GAMA website is http://www.gama-survey.org/. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max- Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington. This publication makes use of data products from the Wide-field Infrared Survey Explorer, which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics and Space Administration. This research has made use of data obtained from or software provided by the US National Virtual Observatory, which is sponsored by the National Science Foundation. | 16 | 9 | 1609.05578 |
1609 | 1609.00731_arXiv.txt | We present medium-resolution spectroscopy and multicolor photometry for the optical transient PSN J09093496+3307204 (SN~2015bh) in the galaxy NGC~2770, which has transferred into the supernova phase. The observations were carried out between February 2015 to May 2016. Both at the phase of the SN impostor (2015a) and at the supernova phase (2015b), besides Balmer emissions, the strong Fe~II emissions are seen in the spectrum; so, these spectra resemble those of Williams Fe~II type classical novae. The star is located near the edge of a dark nebula and notably absorbed (A$_V$ = 1\fm14 $\pm$0\fm15). Taking into account this absorption, we determined maximum absolute magnitudes of $M_V$ = --15\fm0 $\pm$0\fm3 at the 2015a phase and of $M_V$ = --18\fm14 $\pm$0\fm30 at the 2015b phase. The light curve at the 2015b phase is similar to those of SN IIL. The supernova progenitor is a luminous blue variable (LBV) star with the powerful H$_\alpha$ emission. We considered several hypotheses of supernovae explosions following optical transients related with LBV. The hypothesis of core collapse of an evolved massive star interrupting the process of its merging with massive companion in a binary system (a failed luminous red nova) was chosen as the preferable one for this event. | The optical transient PSN J09093496+3307204 was discovered in the galaxy NGC 2770 on February 7, 2015 in the Catalina Real-Time Transient Survey, and on February 8, 2015 by S. Hoverton in the process of observations under the program SNhunt\footnote{http://www.rochesterastronomy.org/sn2015/snhunt275.html}. The transient is also known as SN 2015bh and SNhunt275. In the spectrum obtained in Asiago Observatory on February 9, a broad emission H$_\alpha$ (FWHM $\sim$ 6800 \kms) with a narrow line with P Cygni profile on top was visible \citep{ERea15}. The absorption component of the profile showed the expansion velocity $\sim$950 \kms. In 2008 -- 2009 HST images, a faint star was detected coinciding in coordinates with the transient and variable within 21\fm5 -- 22\fm8 in the red filter F606W. Its spectrum and chaotic variability resembled supernova impostors 2000ch and 2009ip, before the explosion of latter as a supernova in June 2012 \citep{ERea15}. The definition of "supernova impostor" was introduced by \citet{VDPKea00} to identify the explosions of LBV stars in which stars survived and being observed several years after outbursts in contrast to SN~IIn, which were destroyed at the core collapse. The eruption of $\eta$ Car in 1844 -- 1850 can be an analog of a SN impostor. So, the transient in NGC 2770 was classified as a SN impostor or an LBV explosion. Photometry and spectroscopy with the 10.4 m GTC telescope at Canarias within the time interval between March 27 and April 14, 2015 revealed a considerable increase of star brightness by 1\fm2, and the object reached the absolute magnitude --14\fm2 in the SDSS $r$ band \citep{UPLea15}. In the spectrum taken on April 14, the H$_\alpha$ emission, other Balmer lines, the distinct emissions Fe~II in the range $\lambda$ 4950--5400 \AA\ and the Na~I/He~I blend were predominating. The H$_\alpha$ line profile was asymmetric with a red-side wing extending up to $\sim$5000 \kms, and a steeper blue-side decline. But the P~Cyg type profile in H$_\alpha$ line was already not observable. On May 16, 2016, \citet{UPTea15} informed about a sharp increase of the transient brightness by 2$^m$. The absolute magnitude in the $R$ band reached \hbox{--16\fm4.} In new spectra of May 16 from the Keck~I telescope, the star displayed a hot continuum with Balmer, He~II and He~I emission lines \citep{DBLea15}. H$_\alpha$ remained asymmetric with the half-width FWHM $\sim$1200 \kms. It became clear that it was a new explosion of the impostor, and the events were developing according to the scenario of SN 2009ip \citep{MSF13,MMSea14,GSVea14,PCIea13}. The observations continued by \citet{CTLea15,VSV15} and \citet{RA15} confirmed the assumptions that the transient passed to the SN~II phase. By analogy with other papers, we designate the first outburst or the impostor phase as 2015a and the second brightening or the SN phase as 2015b. The host galaxy of SN~2015bh, NGC~2770 is located at a distance of d = 29.70 $\pm$3.4 (distance module $M-m$ = 32\fm35 $\pm$0.24); its radial velocity is $v_r$ = 1947 \kms\ (redshift z = 0.006494); the Galaxy extinction in this direction is $A_V$ = 0\fm062 (NED). Before the 2015 event, three SNe of the rare Ib type exploded in the galaxy NGC~2770 during a short period of 10 years, that is why the galaxy was called "a factory of type Ib supernovae"\ \citep{TMLea09}. These are SNe 1999eh, 2007uy and 2008D. The latter one is interesting by the fact that it was related to an X-ray transient. \citet{TMLea09} show in their Fig.~1 the galaxy image assembled from three VLT frames of March 16, 2008, including one with the filter centered on 6604~\AA\ at $z$ = 0.007 and the width 64~\AA, and another one in the filter H$_\alpha$ at $z$ = 0 with the center at 6563~\AA\ and the width 61~\AA. Besides the two SNe of 2007 and 2008, one can see there a SN~2015bh progenitor, which is distinguished by its excess of emission in the H$_\alpha$ line. This is one of the brightest stars of the galaxy. | Spectra of the optical transient 2015bh/NGC~2770 in the phases of SN impostor and real SN obtained with BTA/SCORPIO are unusual, and they are similar to spectra of Williams Fe~II type classical novae. From spectral data it was established that in the impostor phase (2015a), an optically thick expanding shell formed which was passed by a shock wave at the SN phase (2015b). Eventually, the shell was accelerated and transferred to the optically thin phase. Besides, we have revealed the absorption components in spectral line profiles related to the extended circumstellar medium which partly accelerated by radiation pressure of SN explosion. Near brightness maximum, we have detected the emission component of high-speed, 18000 \kms\ ejecta in the H$_\beta$ line profile along with slower moving, up to 4000 \kms\ absorption components, what suggests the asymmetric eruption. The light curve resembles the SN IIL ones with the rate of the secondary decline corresponding to the rates caused by the radioactive decay of $^{56}$Co isotope. In maximum, the SN reached the absolute magnitude $M_V$ = --18\fm14 $\pm$0\fm30. The SN~2015bh progenitor is a luminous blue variable star (LBV) with the strong emission H$\alpha$, one of the brightest stars of the galaxy NGC~2770. We consider that the most probable hypothesis to explain the SN~2015bh event is a core collapse of a more massive far-evolved star in a binary system breaking its merging with a less massive companion (a failed luminous red nova in a massive binary system). Arguments for and against other hypotheses were also considered. The core collapse of a single massive star induced by the PPI event is also possible.\\ | 16 | 9 | 1609.00731 |
1609 | 1609.06321_arXiv.txt | \noindent The absence of large cooling flows in cool core clusters appears to require self-regulated energy feedback by active galactic nuclei but the exact heating mechanism has not yet been identified. Here, we analyse whether a combination of cosmic ray (CR) heating and thermal conduction can offset radiative cooling. To this end, we compile a large sample of 39 cool core clusters and determine steady state solutions of the hydrodynamic equations that are coupled to the CR energy equation. We \C{find solutions} that match the observed density and temperature profiles for all our clusters well. Radiative cooling is balanced by CR heating in the cluster centres and by thermal conduction on larger scales, thus demonstrating the relevance of both heating mechanisms. Our mass deposition rates vary by three orders of magnitude and are linearly correlated to the observed star formation rates. Clusters with large mass deposition rates show larger cooling radii and require a larger radial extent of the CR injection function. Interestingly, our sample shows a continuous sequence in cooling properties: clusters hosting radio mini halos are characterized by the largest cooling radii, star formation and mass deposition rates in our sample and thus signal the presence of a higher cooling activity. The steady state solutions support the structural differences between clusters hosting a radio mini halo and those that do not. | The population of galaxy clusters can be divided into cool core (CC) and non-CC clusters. CC clusters are characterized by low entropies and short cooling times in the centre \citep{Cavagnolo2009, Hudson2010}. Unimpeded radiative cooling results in cooling flows with mass deposition rates of $1000\,\rmn{M_{\sun}\,yr^{-1}}$ \cite[see][for a review]{Peterson2006}. In contrast, only a moderate amount of cold gas and star formation is observed, which can be up to two orders of magnitude smaller than the predictions \citep{Peterson2006}. In order to solve the emerging cooling flow problem an additional heating mechanism is required that balances radiative cooling. In the centres of CC clusters, the temperature increases with radius such that the gas at the temperature peak functions as a heat reservoir. The transport of heat to the centres of clusters by means of thermal conduction has been studied in great detail \citep[e.g.][]{Bertschinger1986, Bregman1988, Zakamska2003, Guo2008globalstability}. Although it is possible to construct solutions in which thermal conduction balances radiative cooling, the required conductivity has to be fine-tuned \citep{Guo2008}. Moreover, in some clusters such a thermal balance requires a conductivity that exceeds the theoretical maximum, i.e.{\ }the \textit{Spitzer} value \citep{Zakamska2003}. In addition the solutions are not locally stable on scales larger than the Field length \citep{Kim2003, Soker2003}. Hence, thermal conduction cannot be the sole solution to the cooling flow problem. Nevertheless, it might still play an important role beyond the central region at intermediate cluster radii \citep{Voit2011}. Another source of energy that is in principle powerful enough to balance cooling is the feedback from the active galactic nucleus (AGN) of the brightest cluster galaxy \citep[see e.g.][for reviews]{McNamara2007, McNamara2012}. Here, the critical question is how to efficiently couple this energy to the intra-cluster medium (ICM). Various processes have been explored, including mechanical heating by hot bubbles visible in X-ray observations \citep[e.g.][]{Brueggen2002, Gaspari2012} as well as viscous dissipation of sound waves \citep{Ruszkowski2004}. Additionally, the rising AGN bubbles excite gravity modes that decay and thereby generate turbulence. Hence, dissipation of turbulent motions is another possibility for heating the cluster gas \citep[e.g.,][]{Zhuravleva2014}. However, recent X-ray data find a low ratio of turbulent-to-thermal pressure in the Perseus cluster at 4 per cent, thus challenging this scenario since low-velocity turbulence cannot spread far without being regenerated \citep{Hitomi2016}. This result is in line with idealized hydrodynamical simulations, which demonstrate that the conversion of gravity modes into turbulence is very inefficient and transfers less than 1 per cent of the injected AGN energy to turbulence \C{\citep{Reynolds2015, Yang2016b}}. Moreover, all these mechanisms can only make use of one quarter of the available enthalpy provided that the bubbles are disrupted by Kelvin-Helmholtz instabilities within a few exponential pressure scale heights \citep{Pfrommer2013}. The remaining enthalpy is most likely contained as internal energy of relativistic particles and magnetic fields inside the lobes\C{, which also modifies the interplay between jets and the cluster medium \citep{Sijacki2008, Guo2011}}. If the CRs are able to escape the bubbles and fill the ICM, they can heat the cluster gas through streaming. Streaming CRs excite Alfv\'en waves via the streaming instability (\citeauthor{Kulsrud1969} \citeyear{Kulsrud1969}; \citeauthor{Skilling1971} \citeyear{Skilling1971}; see also \citealt{Zweibel2013}, for a review). The CRs then scatter on these self-excited waves, which limits the macroscopic CR velocity in the rest frame of the gas to approximately the Alfv{\'e}n speed \citep[][ assuming pressure carrying CRs at GeV energies]{Wiener2013}. This self-confinement can be very efficient since it operates on time-scales of the order of $30\,\rmn{yr}$, which is much shorter than all other time-scales in the cluster \citep{Wiener2013, Zweibel2013}. The wave growth is counteracted by damping mechanisms such as non-linear Landau (NNL) and turbulent damping \citep{Farmer2004, Wiener2013}, which leads to an energy transfer from the CRs to the cluster gas \citep{Wentzel1971, Guo2008}. Importantly, CR heating allows for a self-regulated feedback loop. The CRs that are injected by the central AGN stream outwards and heat the cluster gas. Thereby, the CRs lose energy and become more and more dilute such that radiative cooling eventually starts to predominate. Cooling gas can then fuel the AGN, which launches relativistic jets that accelerate CRs. Once those escape into the ICM, they stream again outwards and provide a source of heat. An important aspect are the involved time-scales: if CR heating was much slower than the involved dynamical processes, it would not be able to efficiently heat the gas. The free fall time-scale for a typical total density of $\rho = 9 \times 10^{-25}\rmn{~g\,cm^{-3}}$ is $\tau_\rmn{ff} = \sqrt{3 \upi/(32 G \rho)} \approx 7 \times 10^{7}\rmn{~yr}$ \citep{KrumholtzNotes}. We compare this value to the Alfv\'en time since CR heating is mediated by Alfv\'en waves. If we approximate the Alfv\'en time-scale as $\tau_\rmn{A} = L/\vv_\rmn{A}$ and use a typical CR pressure scale height of $L =30\rmn{~kpc}$ and a characteristic Alfv\'en velocity of $\vv_\rmn{A} = 200\rmn{~km~s^{-1}}$ (corresponding to a magnetic field of 10~$\umu$G and $n_{\rmn{e}} = 0.01~\rmn{cm^{-3}}$), this yields $\tau_\rmn{A} \approx1.5 \times 10^{8}\rmn{~yr}$. Hence, the Alfv\'en time-scale is of the same order as the free fall time-scale, which demonstrates that CR heating is sufficiently fast to have an impact on dynamical processes. Moreover, these time-scales are in the range of typical AGN duty cycles of a few times $10^7~\rmn{yr}$ to a few times $10^8~\rmn{yr}$ \citep{1987MNRAS.225....1A, 2005Natur.433...45M, 2005ApJ...628..629N, 2008MNRAS.388..625S}, which is a necessary condition for sufficient replenishment of CRs. For these reasons, CR heating has the potential to play a significant role in solving the cooling flow problem \citep{Loewenstein1991, Guo2008, Ensslin2011, Fujita2011, Fujita2013b, Pfrommer2013}. In particular, there exists a steady state for spherically symmetric models, in which radiative cooling is balanced by CR heating in the central regions and by thermal conduction further out \citep{Guo2008}. Unlike thermal conduction, CR heating is locally stable to thermal fluctuations at $kT \sim 1\;\rmn{keV}$, coincident with the observed temperature floor in some CC clusters \citep{Pfrommer2013}. Moreover, detailed gamma-ray and radio observations of the Virgo cluster allow for a CR population that prevents cooling in this particular cluster \citep{Pfrommer2013}. Steady state solutions are a necessary condition for the viability of a mechanism to prevent cooling flows. There are various steady state solutions for the ICM that include different physical processes in the literature \citep{Zakamska2003, Guo2008globalstability, Fujita2013b}. If only the effects of thermal conduction are considered, steady state solutions exist but the required conductivity needs to be fine-tuned \citep{Zakamska2003}. This situation can be improved by including AGN feedback that is also able to reduce the conductivity to physical values \citep{Guo2008globalstability}. However, \citet{Guo2008globalstability} use the ``effervescent heating'' model by \citet{Begelman2001}, which describes AGN feedback by buoyantly rising bubbles. Motivated by the results of \citet{Guo2008} and \citet{Pfrommer2013}, we explore steady state solutions that simultaneously take into account thermal conduction and CR heating and discuss common characteristics of the solutions. In our companion paper \citep{Jacob2016b}, we assess the viability of our solutions by comparing the resulting non-thermal radio and gamma-ray emission to observational data. This enables us to put forward an observationally supported scenario for self-regulated feedback heating, in which an individual cluster can either be stably heated, is predominantly cooling or is transitioning from one state to the other. Previous works considered at most very small cluster samples. This precludes a sound statistical statement about the viability and applicability of the solution to the entire CC population. Hence, we extend our analysis to a considerably larger cluster sample. Here, we are especially interested in clusters in which CRs have already been observed, e.g. in the form of extended radio emission. In a subsample of CC clusters, such emission occurs as radio mini halos (RMHs) with typical radii of a only a few hundred kpc in contrast to the $\sim\rmn{~Mpc}$ radio halos of non-CC clusters \citep[see e.g.][for a review]{Feretti2012}. Thus, we also include those clusters in the sample selection. This paper is structured as follows. In Section~\ref{sec:sample}, we introduce our cluster sample and determine required properties from observations. The governing equations of our model and our parameter choices are described in Section~\ref{sec:model}. We discuss our results in Section~\ref{sec:discussion} and conclude in Section~\ref{sec:conclusions}. Throughout the paper, we assume a cosmology with $h= 0.7$, $\Omega_{\rmn{m}} = 0.3$ and $\Omega_{\Lambda} = 0.7$. | \label{sec:conclusions} The cooling flow problem in CC clusters remains one of the most interesting open questions in galaxy clusters. While the paradigm of self-regulated AGN feedback is very attractive, the physical heating mechanism that balances radiative cooling has not yet been identified. In this work, we have analysed whether a combination of CR heating and thermal conduction is able to provide the required heating. To this end, we have compiled one of the largest samples of CC clusters ever used for a theoretical investigation of the cooling flow problem. Here, we have focused on clusters for which non-thermal activity has either already been observed or which are predicted to be suitable targets for non-thermal emission. In particular, this includes all clusters that host a radio mini halo, i.e., an extended radio emission in the centres of the clusters. Clusters with an RMH are typically at slightly higher redshifts than clusters without RMHs, but the virial masses of most clusters are comparable with some outliers that we treat separately. We find that the observed infra-red SFR and the cooling radius, which we define as the radius where the cooling time equals $1~\rmn{Gyr}$, are correlated. Moreover, clusters with an RMH have larger SFRs and cooling radii than clusters without RMHs. For all clusters, we found steady state solutions to the system of hydrodynamic equations coupled to the CR energy equation. The thermal energy equation accounts for thermal conduction as well as Alfv\'en wave heating excited by streaming CRs. We choose the parameters of the gravitational potential, CR streaming and injection to obtain physical solutions and ask for maximum CR heating solutions. In consequence, we find solutions that match the observed density and temperature profiles well\C{, however requiring a somewhat high conductivity for some systems}. Radiative cooling is typically balanced by CR heating in the cluster centres and by thermal conduction in the intermediate cluster parts, closer to the peak in temperature. The combination of these two heating mechanisms has several advantages over models that include only one of the two processes. CR heating is locally stable \C{at temperature values corresponding to islands of stability that form at locations of cooling line complexes in the cooling function \citep{Pfrommer2013}} and it allows for self-regulated AGN feedback, in contrast to thermal conduction, which appears to be nonetheless required to balance cooling at large scales and to allow for mass deposition rates that are in agreement with observational findings. Our solutions predict modest mass deposition rates; consistent with the low star formation rates and the observed reservoirs of cold gas in the centres of those systems. The cooling gas can escape the detection of soft X-rays ($kT\lesssim0.5$~keV) by absorption in the filaments with a sufficiently high integrated hydrogen column density and/or by mixing the cooling gas with colder gas, thereby lowering its temperature non-radiatively \citep{Werner2013, Werner2014}. Furthermore, we used our comparably large cluster sample to analyse the parameters of these steady state solutions. We found weak correlations between the observed infra-red SFR and the mass deposition rate in our solutions as well as between the cooling radius and the radial extent of the CR injection. Particularly, clusters with and without RMHs occupy different parts of these relations. Clusters that are hosting an RMH have higher star formation and mass accretion rates in comparison to clusters without an RMH. In addition, the cooling and CR injection radii are typically larger in clusters with an RMH. Hence, the existence of an RMH delineates a homogeneous subclass within the population of CC clusters. In this work, we present steady state solutions that match X-ray observations well. However, these solutions predict a CR population that interacts hadronically with the ambient medium. As a result, pions are produced which decay into electrons and photons that can be observed in the radio and gamma-ray regime, respectively. The crucial question whether the CR populations of our solutions are in agreement with current observations and upper limits of this non-thermal emission will be addressed in our companion paper \citep{Jacob2016b}. | 16 | 9 | 1609.06321 |
Subsets and Splits
No community queries yet
The top public SQL queries from the community will appear here once available.