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1808 | 1808.03036_arXiv.txt | We report on the detection of a pulsating Fe \Ka line in the High Mass X-ray Binary (HMXB) GX 301-2, from a 40-ks \cha observation near periastron. The pulsations in the Fe \Ka emission appeared only in the first 7 ks of the observation, with a period and phase profile similar to those of the continuum. The presence of pulsed fluorescent lines is an unusual property in HMXBs. After 7 ks, the continuum flux increased by a factor of three, the Fe \Ka flux increased only by about 10\%, and the pulsating signal in the line disappeared. Finally, in the second half of the observation, both the continuum and the line flux dropped by a similar factor of 2. We suggest that the pulsating component of the Fe \Ka line is coming from a transient non-isotropic distribution of dense gas around the neutron star, for example an accretion stream induced by periastron passage, or from the illuminated surface of the donor star. | The classical High Mass X-ray Binary (HMXB) system GX 301-2 consists of a pulsar immersed in a slow ($\approx$300 km\,s$^{-1}$) and dense stellar wind ($\dot{M}_{\rm w} \sim 10^{-5} M_\odot$\,yr$^{-1}$), arising from the B1 hypergiant Wray 977 \citep{Kap06}. The donor star has an estimated mass of $\approx$39--53 $M_\odot$, a radius of $\approx$62 $R_\odot$, and a distance from us of $\approx$3 kpc \citep{Kap06}. The orbit of the pulsar has a period $P_{\rm orb} = 41.506 \pm 0.003$ d, with reference phase 0 at the Modified Julian Date (MJD) $T_0 = 43,906.06 \pm 0.11$, and an orbital period derivative estimated as $\dot{P}_{\rm orb} = (-3.7 \pm 0.5) \times 10^{-6}$ s s$^{-1}$ \citep{Dor10}. The system has an eccentricity $e \approx 0.46$, with a periastron distance of $\approx$25 $R_{\odot}$ between the neutron star and the surface of the donor star \citep{Sat86, Koh97}. The inclination angle is moderately high, between 44$^{\circ}$--78$^{\circ}$ \citep{LK08,Kap06}. The pulsar has a long spin period $P_{\rm s} \approx 680$ s, which showed erratic changes over the last 30 years \citep[e.g.][]{Dor10,Eva10}. Such a long pulse period has been attributed to a strong magnetic field \citep[$\sim10^{14}$ G,][]{Dor10}. The field measured from the cyclotron line is much weaker ($\approx$5 $\times 10^{12}$ G), but the two values can be consistent with each other if the cyclotron line region comes from a tall accretion column of height $\approx$2.5--3 $R_{NS}$ \citep[][but also see \citet{IF12}]{Dor10}. \begin{figure*} \hspace{-1.3cm} \includegraphics[width=7.3in]{G301lcT68A.ps}\\[-20pt] \hspace{-1.2cm} \includegraphics[width=7.3in]{G301lcT68B.ps} \vspace{-1.3cm} \caption{Light curves of the continuum-subtracted Fe \Ka line (6.34--6.44 keV) and the continuum (6.5--10 keV), for the first 17 ks of the observation (top panel) and for the interval between 17--40 ks (bottom panel). The time resolution of the data points is 68 s = 0.1 times the pulse period. Letters indicate different sub-intervals of the observation, used for our timing and spectral analysis. The error bar plotted in the top panel at about $10^4$ s represents a typical 1-$\sigma$ error, for a count rate of 0.2 ct s$^{-1}$. We took half of the count rate within the 5.8--6 keV band as the continuum level for the Fe K$\alpha$ line. We did not use the 7.0--7.2 keV band for the continuum estimate because it is contaminated by the Fe \Kb line. } \end{figure*} The X-ray continuum flux of GX 301-2 shows regular orbital modulations and reaches a maximum about 1.4 days before periastron passage \citep{Sat86}. The flux increase has been explained by a gas stream induced by the neutron star near periastron \citep[][and references therein]{LK08, IP14}. GX 301-2 shows a strong 6.4-keV Fe \Ka fluorescent line, as well as many other fluorescent lines \citep{Sar96,Fur11,Suc12}. The Fe \Ka fluorescent line is produced when the X-ray emission illuminates the surrounding gas; its intensity and spectral profile are important tools for the study of HMXBs \citep[][and reference therein]{Tor10,Gim15}. For example, from the detection of a spectrally resolved Fe \Ka Compton shoulder at periastron, \citet{Wat03} inferred an absorption column density of $\approx$10$^{24}$ cm$^{-2}$ and an electron temperature of $\approx$0.5 eV. The location of the Fe \Ka fluorescent region for GX 301-2 is still actively debated. Using spectral data from the {\it Advanced Satellite for Cosmology and Astrophysics} ({\it ASCA}), \citet{End02} measured a Gaussian $\sigma$ of $\approx$40--80 eV for Fe \Ka and inferred an emission region within $\approx$10$^{10}$ cm (0.3 lt-s) of the neutron star. An {\it XMM-Newton} study \citep{Fur11} did not reveal any significant time delays between the observed continuum and the reprocessed line emission, which implies a distance smaller than $\approx$6 $\times 10^{10}$ cm between the two sources of emission. Instead, using {\it Suzaku} observations, \citet{Suc12} inferred a distance greater than $\approx$2 $\times 10^{13}$ cm (700 lt-s), based on the relatively flat phase profile of the line compared with that of the continuum. Despite this discrepancy in the proposed size of the line emission region, one thing that appeared well ascertained was that the Fe \Ka line was less pulsed compared to the continuum \citep[e.g.][and reference therein]{Fur11}. In this paper, based on {\it Chandra X-ray Observatory} observations, we show that this is not always true: the Fe \Ka line pulsates at times as strongly as the continuum. | The key new result of this paper is that we detected a pulsating Fe \Ka line with the same period as the continuum, during the first 7 ks (interval A) of a \cha observation of GX 301-2. The phase profile of the Fe \Ka line was also similar to that of the continuum. The pulsation signal of the Fe \Ka line then disappeared; at that time, the continuum flux increased by a factor of 3 and kept its pulsation profile; the line flux increased only by about 10\%. X-ray fluorescence lines are one of the hallmarks of neutron star HMXBs \citep[{\it e.g.},][]{LC93,Bas78,HM77}; they may be formed in many locations: in the accretion column, on the surface of the donor star or of an accretion disk, in the stellar wind, or in an accretion stream, shell or wake. Usually, the line emission is not pulsed \citep[{\it e.g.},][]{Ber15,Suc12,Wil11,Lei09,Pau02}. A few well-studied exceptions, where fluorescent line emission from iron, oxygen and/or neon is pulsed, are 4U\,1626$-$67 \citep{Ber15,Sch01,Ang95}, Hercules X-1 \citep{Vas13,Zan04,Cho94}, and GX 1+4 \citep{Yos17a,Yos17b,Kot99}. GX 301-2 was one of the X-ray pulsars in which the pulsed fraction of the Fe \Ka line was always absent or much smaller than the pulsed signal in other energy bands; this was explained as the effect of smearing of the pulsed signal in an isotropic line-emitting region around the neutron star \citep{Tas91,End02,Fur11}. Instead, we have now shown that at some epochs, the pulsed fraction of the line is as high as that of the continuum. Thus, the simplest qualitative explanation of this finding is that during those epochs, the pulsating Fe \Ka line originates in an anisotropic gas structure, illuminated by the pulsed emission of the neutron star. More remarkably, we have shown that the line pulsation disappears at a phase $\phi = 0.000 \pm 0.003$ (extrapolating from the ephemeris and period derivative solution of \citealt{Dor10}). The inferred column density of the scattering medium responsible for the \Ka fluorescence lines (inferred from the relative strength of the \Ka Compton shoulder), and for the scattered continuum (inferred from broad-band spectral modelling) are similar, that is $\approx1.5\times10^{24}$ cm$^{-2}$. Such a high column density cannot be produced by the quasi-isotropic stellar wind of a super-giant donor, but can be reached in the transient accretion stream formed between the donor star and the neutron star near periastron \citep{Ste88}. Evidence for the existence of this dense stream flow comes from studies of optical lines \citep{Kap06}, and the recurrent pre-periastron X-ray flare \citep[][and references therein]{LK08}. When the neutron star enters and leaves the accretion stream, an anisotropic configuration will be formed. Specifically, the {\it Chandra} observation studied in this work was taken shortly after the flare; in the accretion stream scenario, it might correspond to the phase when the neutron star leaves the accretion stream. The illumination of the anisotropic stream distribution would produce the pulsating Fe \Ka line. When the neutron star moves further away from the stream, the contribution of the pulsating Fe \Ka line becomes less important. Such a scenario also predicts a pulsating Fe \Ka line when the neutron star enters the accretion stream, i.e., before the X-ray flare. Moreover, the pulsating behaviour should repeat at similar phases every orbital cycle. Such a scenario can be tested with future observations. Another possible anisotropic structure is the surface of the donor star. At periastron, the donor star subtends a solid angle $\Omega_{\rm p} \approx 1.88$ sr, that is, it is seen by the neutron star with an apex half-angle $\theta\approx 45^{\circ}$. Thus, for a moderate misalignment between the beam of X-ray emission and the normal to the orbital plane, the beamed X-ray emission of the neutron star could directly illuminate the surface of the donor star, and produce the pulsating Fe \Ka line. As the neutron star moves away from periastron, the solid angle subtended by the donor star decreases. At apastron, the subtended solid angle is only $\Omega_{\rm a} \approx 0.22$ sr, corresponding to a half-angle $\theta \approx 15^{\circ}$. In this scenario, we suggest that the disappearance of the pulsating line signal corresponds to the X-ray beam missing the stellar surface. Such scenario predicts that the pulsating phenomenon repeats periodically at periastron, and can also be tested with future observations. The phase profile of the Fe \Ka line of GX 301-2 during interval A is quite different from those reported in the literature for other sources with pulsating line emission. For example, the phase profile of the Fe \Ka line from Her X-1 shows a sharp minimum around the continuum peak, which might indicate a hollow cone of the accretion structure \citep{Vas13}. The \OVII\ line fluxes of 4U 1626$-$67 show a variation stronger than that of the continuum, which could be due to variable illumination of the warped region of an accretion disk \citep{Ber15}. In contrast, the phase profile of the Fe \Ka line of GX 301-2 during interval A is similar to that of the continuum. It indicates that a major part of the observed continuum might come from the Compton-thick medium producing the Fe \Ka line, as evidenced by the spectral modelling in \S 3.2. In principle, even a homogeneous, spherically symmetric scattering medium illuminated by a central X-ray pulsar can produce a fluorescent signal that appears pulsed to the distant observer, because of the finite-light-speed effect \citep{Yos17b}. However, the characteristic size of the medium for this effect to be significant is of order of the pulsar spin period times the speed of light. In the case of GX 301-2, that would require a scattering medium with a size $\sim$10$^{13}$ cm, implausibly larger than the binary system itself. The possibility that the observed continuum could have a significant contribution from reprocessed emission may change the interpretation of the time delays between the line and continuum. Using \xmm data, \cite{Fur11} found no time delays between Fe \Ka line and continuum in the time range of 2--5,000 s, and they inferred a fluorescent region smaller than 2 lt-s ($6\times 10^{10}$ cm). However, if a major part of observed continuum come from scattering, far from the neutron star surface, the lack of relative time delays would no longer constrain the size of the region. A reanalysis of \xmm data is needed to check this possibility. Another way to constrain the region of the fluorescent lines is through their velocity broadening. Thanks to the \cha HEG resolution, we have shown that typical velocity broadening of the Fe \Ka line is $\approx$200--400 km s$^{-1}$, a few times lower than previously inferred. This is comparable to the typical terminal velocity of the stellar wind and the ballistic motion of an accretion stream. If the broadening is due to gas in virial motion, such velocities would correspond to distances $\sim5\times10^{10}$--$2\times10^{11}$ cm for a neutron star of 1.5 $M_\odot$, similar to the accretion radius of GX 301-2 ($\sim10^{11}$ cm). Such a relatively large scale is consistent with fluorescence lines from cool and nearly neutral Fe ions. We note that it is much larger than the magnetosphere of a neutron star even with a field of $10^{14}$ G ($4\times10^{9}$ cm). On the other hand, the fluorescent region can not be larger than the light travel distance of a significant fraction of the pulse period (e.g., 100 ls), otherwise the pulsation signal would be smeared out. In principle, any temporary anisotropic gas structure on scales of $\sim10^{11}$ cm with enough covering factor, like a torus or a warped accretion disk, could be invoked to explain the pulsating Fe \Ka line. Such gas structures are likely caused by the dynamical interaction between the donor star and the neutron star near periastron. Detailed hydrodynamical simulations are needed to test the existence of such structures. The increase in the observed continuum after the first 7 ks of the observation, accompanied by the non-increase of the line emission and the disappearance of the pulsed line signal remain puzzling, and can be due to a combination of increased intrinsic emission, decreased absorption column density along our line of sight, and decreased column density of the scattering medium. Simultaneous observations in wide X-ray bands are needed to disentangle these effects. Nonetheless, using only the \cha HEG data, we have already shown a decreasing trend in the column density of the scattering medium in the second half of the observation. This is seen both from our broad-band continuum modelling, and from the progressively lower flux ratio between the Compton shoulder and the Fe \Ka line. | 18 | 8 | 1808.03036 |
1808 | 1808.04420_arXiv.txt | We present a characterization of the new Volans-Carina Association (VCA) of stars near the Galactic plane ($b$\,$\simeq$\,-10\textdegree) at a distance of $\simeq$\,75--100\,pc, previously identified as group~30 by \cite{2017AJ....153..257O}. We compile a list of \likelymembers\ likely members from \emph{Gaia}~DR2 with spectral types B8--M2, and \candidatemembers\ additional candidate members from \emph{Gaia}~DR2, 2MASS and AllWISE with spectral types A0--M9 that require further follow-up for confirmation. We find an isochronal age of \age\,Myr based on MIST isochrones calibrated with Pleiades members. This new association of stars is slightly younger than the Pleiades, with less members but located at a closer distance, making its members $\simeq$\,3 times as bright than those of the Pleiades on average. It is located further than members of the AB~Doradus moving group which have a similar age, but it is more compact on the sky which makes it less prone to contamination from random field interlopers. Its members will be useful benchmarks to understand the fundamental properties of stars, brown dwarfs and exoplanets at $\simeq$\,90\,Myr. We also provide an updated version of the BANYAN~$\Sigma$ Bayesian classification tool that includes the Volans-Carina association. | 18 | 8 | 1808.04420 |
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1808 | 1808.01802_arXiv.txt | We clarify the line-of-sight structure of the Taurus Molecular Cloud 1 (TMC-1) on the basis of the CCS($J_N=4_3-3_2$) and HC$_3$N($J=5-4$) spectral data observed at a very high velocity resolution and sensitivity of $\Delta V \simeq 0.0004$ km s$^{-1}$ ($=61$ Hz) and $\Delta T_{\rm mb} \simeq 40$ mK. The data were obtained toward the cyanopolyyne peak with $\sim$30 hours integration using the Z45 receiver and the PolariS spectrometer installed in the Nobeyama 45m telescope. Analyses of the optically thin $F=4-4$ and $5-5$ hyperfine lines of the HC$_3$N emission show that the spectra consist of four distinct velocity components with a small line width ($\lesssim 0.1$ km s$^{-1}$) at $V_{\rm LSR}=$5.727, 5.901, 6.064, and 6.160 km s$^{-1}$, which we call A, B, C, and D, respectively, in the order of increasing LSR velocities. Utilizing the velocity information of the four velocity components, we further analyzed the optically thicker CCS spectrum and the other hyperfine lines of the HC$_3$N emission by solving the radiative transfer to investigate how the four velocity components overlap along the line of sight. Results indicate that they are located in the order of A, B, C, and D from far side to near side to the observer, indicating that TMC-1 is shrinking, moving inward as a whole. \\ | \label{sec:intro} Molecular emission lines provide us with essential information on the kinematics of molecular clouds. Especially, emission lines from low rotational transitions of carbon monoxide \citep[such as the $J=1-0$ transition of CO, e.g., ][]{Dame2001} have been widely used in various studies of molecular clouds and star formation. Combination of the optically thick $^{12}$CO line and the optically much thinner lines from its isotopologues ($^{13}$CO and C$^{18}$O) are powerful to reveal to the spatial distributions of molecular gas in a wide range of density from diffuse outskirts to dense cores forming young stars. The C$^{18}$O emission line is generally optically thin, and is useful to investigate dense regions. It is even possible to quantify the velocity field of dynamically infalling gas toward young stars \citep[e.g.,][]{Ohashi1997,Momose1998,Shimoikura2016}. However, we often face difficulties to reveal structures of such dense regions only with the C$^{18}$O emission line, because it has a rather low critical density ($\sim10^{2-3}$ cm$^{-3}$), and thus its spectra tend to be contaminated by a large amount of emission from the diffuse background and foreground along the line-of-sight (LOS). In addition, the C$^{18}$O molecules tend to be depleted onto dust grains in cold and dense regions \citep[e.g.,][]{Bergin2002}, which makes a dip in the C$^{18}$O spectra unrelated to the velocity field. The density and velocity distributions around the densest regions in molecular clouds are of our particular interest in terms of star formation, which should be better probed by other molecular lines with much higher critical densities. If we use two or more such molecular lines with different optical depths, it should also be possible to reveal the structure of dense regions along the LOS. In this paper, we introduce a useful method to probe into the LOS structures of dense regions in molecular clouds, and report results of its applications to the Taurus Molecular Cloud 1 (TMC-1). The method utilizes the combination of the CCS$(J_N=4_3-3_2)$ and HC$_3$N$(J=5-4)$ emission lines at 45 GHz whose critical densities are $\sim 10^{4-5}$ cm$^{-3}$. The basic idea of the method is to identify distinct velocity components based on the optically thin hyperfine lines of the HC$_3$N emission, and to determine the order of the detected components along the LOS by solving the radiative transfer to reproduce the observed CCS emission line as well as the other hyperfine lines of the HC$_3$N emission which are much optically thicker and thus sensitive to the absorption by the gas in the foreground. TMC-1 is a well-known filamentary condensation and a part of an extended molecular cloud in Taurus. In this paper, we will refer to the condensation as the TMC-1 filament, or more simply, TMC-1. Carbon-chain molecules such as CCS and HC$_3$N are known to be abundant in TMC-1 and in other young star forming regions \citep[e.g.,][]{Hirahara1992,Suzuki1992,Hirota2009,Shimoikura2012, Taniguchi2016a,Taniguchi2016b,Taniguchi2017a,Taniguchi2017b,TaniguchiSaito2017}, and thus the emission lines from these molecules can be a useful tool to investigate the initial conditions of star formation. In addition, the CCS and HC$_3$N emission lines at 45 GHz are relatively intense compared to those from other carbon-chain molecules, and they also have an advantage to be observable simultaneously with a millimeter-wave radio telescope like the 45m telescope at the Nobeyama Radio Observatory (NRO), which we actually used to obtain the data for this paper, because their rest frequencies differ only by $\sim10$ MHz. The CCS and HC$_3$N spectral data used in this paper were obtained using the NRO 45m telescope. The data were obtained for the purpose to detect the Zeeman splitting of the CCS line to measure the magnetic field in TMC-1. About 30 hours integration provided us with wonderful spectral data having an extremely high velocity resolution of 0.0004 km s$^{-1}$. Detection of the Zeeman splitting and the related analyses of the magnetic field will be presented in a subsequent publication (Nakamura et al. 2018, in preparation). In this paper, we will concentrate on the analyses of the LOS structures of TMC-1 inferred from the spectral data. We describe the observational procedures in Section \ref{sec:obs}. In Section \ref{sec:results}, we will explain how we analyzed the spectral data, and present results of the analyses. We have detected four distinct velocity components, and found that they are lying on the LOS from far side to near side to the observer in the order of increasing radial velocities, suggesting the inward motion of TMC-1. In Section \ref{sec:dis}, we further discuss the global inward motion of TMC-1 on the basis of the $^{13}$CO and C$^{18}$O data available on the website of NRO. Summary of this paper is given in Section \ref{sec:conclusions}. | \label{sec:conclusions} We have observed the cyanopolyyne peak in the TMC-1 filament with the CCS($J_N=4_3-3_2$) and HC$_3$N($J=5-4$) emission lines at 45 GHz utilizing the Z45 receiver and PolariS spectrometer installed in the 45m telescope at the Nobeyama Radio Observatory (NRO). Thirty-hours integration with these instruments provided us with wonderful spectral data having a very high velocity resolution of 0.0004 km s$^{-1}$ and a noise level of 0.04 K. Analyses of the spectra infer various information on the structure of TMC-1 along the line-of-sight (LOS). Main conclusions of this paper are summarized in the following: \begin{enumerate} \item Based on the analyses of the optically thin $F=4-4$ and $5-5$ hyperfine lines of the HC$_3$N emission, we identified four velocity components along the LOS of TMC-1. These components are separated by $0.1-0.2$ km s$^{-1}$ and have a line width of $0.05-0.09$ km s$^{-1}$. We regard that the four velocity components represent subfilaments in the TMC-1 filament, and we named them A, B, C, and D in the order of increasing radial velocity. \item To investigate the order of the velocity components along the LOS, we solved radiative transfer for the CCS and HC$_3$N spectra taking into account the effect of absorption by the components lying in the foreground of the other components. Results infer that the four components are located in the order of A, B, C, and D from the further to closer positions to the observer. The fact that the components with lower radial velocity are lying at further positions from the observer indicates that these velocity components are moving inward toward the center of the TMC-1 filament. \item Results of the analyses of the optically thick $F=6-5$, $5-4$, and $4-3$ hyperfine lines of the HC$_3$N emission infer the existence of an additional component which should be diffuse gas lying in the foreground of the other components and contribute to the observed HC$_3$N spectrum only as the absorber without emitting the molecular line by itself. The component that we call E follows the same velocity trend as the other components A--D. \item We investigated the $^{13}$CO and C$^{18}$O spectra downloaded from the data archive of NRO, and found an evidence of the global inward motion of the TMC-1 filament. To better understand the global structure of the filament, we made a simple model which can account for the main features of the $^{13}$CO and C$^{18}$O spectra as well as the origin of the additional component E. \item Virial analyses infer that the TMC-1 filament can collapse by the self-gravity within a time scale of $\sim2 \times 10^5$ years, unless it is supported by the magnetic field of an order of $\sim100$ $\mu$G. Because there is no YSO forming there, we suggest that TMC-1 is in the gravitational equilibrium being supported by the magnetic field, and that the observed inward motion may represent oscillation of the filament. \item Our model infers a possibility that the observed velocity components A--D might be the emission from different parts of a shrinking or oscillating single filament which does not necessarily need subfilaments. In the case of TMC-1, we regard that the four components represent real subfilaments, because they show asymmetric velocity distributions in the optically thin lines with respect to the systemic velocity. However, in oder to check how much this picture is realistic, it is crucial to resolve TMC-1 at a very high angular resolution that can be achieved by a large interferometer such as ALMA. \end{enumerate} | 18 | 8 | 1808.01802 |
1808 | 1808.01669_arXiv.txt | Assuming a non-gravitational interaction amongst the dark fluids of our universe namely, the dark matter and dark energy, we study a specific interaction model in the background of a spatially flat Friedmann-Lema\^itre-Robertson-Walker geometry. The interaction model, as we found, solves the background evolution in an analytic way when the dark energy takes a constant barotropic equation of state, $w_x$. In particular, we analyze two separate interaction scenarios, namely, when the dark energy is a fluid other than the vacuum energy (i.e., $w_x \neq -1$) and when it is vacuum energy itself (i.e., $w_x = -1$). We found that the interacting model with $w_x \neq -1$ produces stable perturbation at large scales for $w_x < -1$ with the coupling strength $\xi <0$. Both the scenarios have been constrained with the latest astronomical data having distinct origin. The analyses show that a very small interaction with coupling strength is allowed and within 68.3\% confidence-region, $\xi =0$ is recovered. The analyses further show that a large coupling strength significantly affects the large scale dynamics of the universe while according to the observational data the interaction models are very well consistent with the $\Lambda$-cosmology. Furthermore, we observe that for the vacuum interaction scenario, the tension on $H_0$ is not released while for the interacting dark energy scenario with $w_x < -1$, the tension on $H_0$ seems to be released partially because of the high error bars in $H_0$. Finally, we close the work with the Bayesian evidence which shows that the $\Lambda$CDM cosmology is favored over the two interacting scenarios. | \label{Intro} The theory of non-gravitational interactions between dark matter and dark energy is the main concern of this work. The origin of such interacting theory did not appear suddenly in the cosmological scheme. It has a well motivated history that we shall discuss here. However, before that, we need a basic introduction about the dark matter and dark energy. According to latest observational suggestions \cite{Ade:2015xua} dark matter (DM) and dark energy (DE) are the main influential sources of the total energy budget of the universe where the dark matter contributes around 26\% of its total energy density, is pressureless and unseen while the dark energy, a hypothetical fluid occupying 68\% of the total energy density of the universe is accelerating the expansion history of the universe. The best description for such observational information is the $\Lambda$CDM cosmology where $\Lambda$ acts as the dark energy fluid and the CDM is the cold dark matter that is pressureless. This is a non-interacting scenario in the sense that both $\Lambda$ and CDM are conserved separately. Despite of great success of $\Lambda$CDM cosmology, the cosmological constant problem \cite{Weinberg} still lacks a satisfactory explanation. The cosmological constant problem is basically confined with the mismatched value of the cosmological constant predicted from the high and low energy scales of the universe. In the following we shall discuss how the interacting dynamics is closely related to the cosmological constant problem. In fact this coupling mechanics was originated because of the cosmological constant problem and finally it became very useful to explain some other issues. Let us move to the next section for an elaborative discussion on the origin of interacting dark matter and dark energy. In the late eighties, there was no concept of dark energy but the discrepancy in $\Lambda$ was remaining to be a serious issue for modern cosmology. To account of such issue, one of the attempts was to consider a toy model where scalar field is coupled to gravity \cite{Wetterich:1994bg}. The energy-momentum tensor of such coupled scalar field introduces a time dependent cosmological constant and consequently, it became a possible solution to the cosmological constant problem since the objection on the time-independent cosmological constant is naturally solved due to having a variable nature of the cosmological constant. After the official introduction of dynamical dark energy models in several forms (see \cite{Copeland:2006wr, AT, Bamba:2012cp} for a detailed survey on them), it was found that they automatically induce coincidence problem \cite{Zlatev:1998tr}. We note that the cosmological constant being time-independent cannot escape from the same problem. Quite interestingly, it was reported in \cite{Amendola:1999er} that if dark energy and dark matter are allowed to interact non-gravitationally with each other, the coincidence problem can be solved. Following this, a series of works with coupled dark matter and dark energy had the same conclusions \cite{Chimento:2003iea,Cai:2004dk, Hu:2006ar, delCampo:2008sr,delCampo:2008jx}. However, some recent results fueled the investigations of coupled dark matter and dark energy with the claim that the observational data favor an interaction in the dark sector \cite{Salvatelli:2014zta,Nunes:2016dlj,Kumar:2016zpg,Yang:2016evp,vandeBruck:2016hpz,Yang:2017yme,Kumar:2017dnp,Yang:2017zjs,Kumar:2017bpv,Yang:2017ccc}. Additionally, again some very recent investigations in this direction strongly claim that the tension on the local Hubble constant can be released if the interaction in the dark sector is allowed. However, the most important question in the coupling dynamics is, what should be the energy transfer rate between the dark sectors? Before we look for an appropriate transfer rate we recall that the nature of both dark matter and dark energy is unknown. On this ground the sensible approach is to consider some well motivated phenomenological transfer rates, or interaction functions and test the expansion history with the available astronomical data. A number of different interaction rates between dark matter and dark energy have been studied in the last several years \cite{Billyard:2000bh,Gumjudpai:2005ry,Barrow:2006hia,% Zimdahl:2006yq,Amendola:2006dg,CalderaCabral:2008bx,% Chimento:2009hj,Quartin:2008px,Valiviita:2009nu,% Clemson:2011an,Thorsrud:2012mu,Pan:2013rha,% Yang:2014hea,Faraoni:2014vra,Yang:2014gza,Nunes:2014qoa,% Pan:2014afa,Chen:2011cy,G:2014mea,Pan:2012ki,Li:2013bya,Duniya:2015nva,% Valiviita:2015dfa,Sola:2016ecz,Mukherjee:2016shl,% Pan:2016ngu,Dutta:2017kch,Cai:2017yww,Odintsov:2017icc,Pan:2017ent,% Yang:2018pej, Yang:2018ubt, Yang:2018euj}. For a comprehensive review on different interaction rates, we refer to \cite{Bolotin:2013jpa, Wang:2016lxa}. We also note that the interaction between the dark sectors has also been examined in a more general framework where the geometry of the universe is inhomogeneous \cite{Izquierdo:2017igb,Izquierdo:2017pnp}. In this work we concentrate on the spatially flat Friedmann-Lema\^itre-Robertson-Walker universe where we introduce an interaction between dark energy and pressureless dark matter that exactly solves the background evolution. That means the evolution equations for dark matter and dark energy are analytically solved. The appearance of analytic structure of the background evolution makes the cosmological model quite interesting because the cosmological parameters associated with this model take analytic forms too. However, this is not new because the analytic structure for such interaction model has already been reported by some of the authors in a previous work \cite{Sharov:2017iue}. But the motivation of the present work is completely different. Here we aim to test the large-scale stability of the model which is very important because without stable perturbations there will be no such structure formation of the universe. The analysis of structure formation in models of DE and DM, from the point of view of the cosmological perturbations theory, plays an essential role when the different models are confronted with the observations-data \cite{Hwang:2009zj}. As it is well known, these dark scenarios imprint a signature on the cosmic microwave background (CMB) power spectrum \cite{Bean:2003fb,Weller:2003hw}. Thus, the study of the cosmological perturbations is important and also need to be well-analyzed. In particular, for models with interaction between DE and DM, with adiabatic initial conditions and the perturbation theory were studied in Ref. \cite{Bean:2007ny}, see also \cite{Herrera:2016uci, delCampo:2013hka, Wang:2016lxa}. Also an analysis in models with an interacting DE-DM with a constant equation of state, was analyzed in \cite{Valiviita:2008iv}. Here, the authors demonstrated that perturbations were unstable together with a rapid growth of DE fluctuations. In this sense the test the large-scale stability is fundamental, since without stable perturbations there will be no such structure formation. We organize the work in the following way. In section \ref{sec-eqns} we describe the basic equations for the interacting model both at background and perturbative levels. The analytical solutions are discussed in section \ref{sec-solutions}. The section \ref{sec-results} details the results of the analysis following the observational data used in this work. Finally, we close our work with a brief summary in section \ref{sec-conclu}. | \label{sec-conclu} An interacting scenario between a pressureless dark matter and a dark energy fluid availing constant barotropic equation of state has been considered. The underlying geometry of the universe is characterized by the spatially flat FLRW line element and the interaction rate $Q = Q (\rho_t^\prime) = Q (\rho_c, \rho_x)$ has been given explicitly in eqn. (\ref{interaction}) or eqn. (\ref{eq-int}). This interaction rate is very appealing in the sense that the evolution equations for the dark sectors (cold dark matter and dark energy) can be exactly solved, and thus, one can directly measure their deviation from the standard evolution laws of the dark fluids with no-interaction. We note that initially this kind of interaction was introduced by Chimento \cite{Chimento:2009hj} where the author proposed a very general interaction rate that recovers the interaction in eqn. (\ref{eq-int}) and discussed its theoretical implications. Later on its observational viability was tested when dark energy is the cosmological constant but at the background level \cite{G:2014mea}, consequently, in a recent article \cite{Sharov:2017iue}, the authors generalized this study for both $w_x = -1$ and $w_x \neq -1$ at the background level with the recent observational data. However, it is quite certain that the dynamics of such interaction models at the large scale of the universe, is promising for a better understanding of the entire scenario. That means, the most important question related with the interaction model is, how the structure formation of the universe depends when such interaction is included in the cosmological scenario. Thus, in the present work we discuss the perturbations and the structure formation of the universe when such interaction is present between the dark fluids. Now, in order to test the resulting cosmological scenarios with the available observational data, we use \texttt{cosmomc}, a markov chain monte carlo package that extracts the model parameters with a sufficient convergence following the Gelman-Rubin statistics \cite{Gelman-Rubin}. The observational data include cosmic microwave background radiation, baryon acoustic oscillations, redshift space distortions, local Hubble constant, supernovae type Ia from joint light curve analysis, Hubble parameter values at different redshifts from cosmic chronometers and finally the weak gravitational lensing data. For a better analysis, we have considered two distinct interacting scenarios, namely when the dark energy is other than the cosmological constant (i.e., $w_x \neq -1$) and the other one is the cosmological constant itself. For IDE scenario, the constraints on the model parameters have been summarized in Table \ref{tab:results-I} where we present the 95.4\% limits (lower) on the coupling parameter $\xi$. And in Fig. \ref{fig-contour1}, we show the contour plots for different combinations of the model parameters at 68.3\% and 95.4\% confidence levels. The right corners of Fig. \ref{fig-contour1} also shows the one-dimensional posterior distributions for some selected model parameters as well. From the observational constraints on the coupling parameter, $\xi$, summarized in the last row of the Table \ref{tab:results-I}, we find that $\xi =0$ is consistent with the observational data. Moreover, from the constraints on the dark energy equation of state, $w_x$, one can see that it is actually very very close to the cosmological constant boundary. Thus, we see that the interaction model is actually equivalent to the non-interacting $\Lambda$CDM background. However, in the large scale distribution, the interaction model may exhibit some differences even for a very small coupling strength. From the imprints on the CMB TT spectra (see the right panel of the Fig. \ref{fig:CMB-ide}) and also from the matter power spectra (see the right panel of the Fig. \ref{fig:Mpower-ide}), it is evident that for a very small coupling strength ($\xi =-0.0001$), the model presents a very minimal deviation from the non-interacting $\Lambda$CDM cosmology. Now, for the interacting cosmological constant (labeled as IVS), the results have been summarized in Table \ref{tab:results-II}. The corresponding contour plots at 68.3\% and 95.4\% confidence-levels are also shown in Fig. \ref{fig-contour-vacuum1} with the one-dimensional posterior distributions for some selected parameters of this model. From the estimation of the coupling strength shown in Table \ref{tab:results-II}, one can see that $\xi$ is concistent with the non-interaction limit (i.e., $\xi =0$), at least according to the current observational data. In fact, for this model we have realized a similar trend as in IDE. For instance, from Fig. \ref{fig-scattered-ivs}, similar to IDE model, we find that lower values of the Hubble parameter allow non-zero interaction in the dark sector. The deviation of this interaction scenario from the non-interacting $\Lambda$CDM cosmology is also found to be insensitive (see the right panels of Fig. \ref{fig:CMB-ivs} and Fig. \ref{fig:Mpower-ivs}) unlike the IDE scenario where although the deviation is small but they are detectable. We also raise one interesting point that has become a hot issue at current cosmological research $-$ the observed tension on the $H_0$ parameter from its global \cite{Ade:2015xua} and local measurements \cite{Riess:2016jrr}. We found that the allowance of the interaction increases the error bars on the Hubble parameter measurements, and consequently, the parameters space for $H_0$ is increased. This becomes effective to release the tension partially and is reflected from some combinatons for IDE only. While the interacting vacuum model is not suitable to release the tension. One may argue that the allowance of extra degrees of freedom in the parameters space of the interacting dark energy models (for IDE, the number of parameters is 8 while for IVS this number is 7) might be suitable to alleviate such tension. Similar results have been reported in some recent works \cite{Kumar:2017dnp, DiValentino:2017iww}, but however, since the theory of interaction is phenomenological and hence its conclusions too, therefore, the analysis with a different interaction model might be perhaps important to see whether the model can avail the same property or not. The relation of the extra degrees of freedom to the tension on $H_0$, in the interacting dark energy models surely needs further attention. Finally, we computed the Bayesian evidence for each interacting scenario with respect to the non-interacting $\Lambda$CDM model (see Table \ref{tab:bayesian}). Our analysis shows that the non-interacting $\Lambda$CDM is preferred over the two interacting dark energy scenarios, at least according to the current observational data sets. | 18 | 8 | 1808.01669 |
1808 | 1808.04566_arXiv.txt | { Deep mid-infrared (MIR) surveys have revealed numerous strongly star-forming galaxies at redshift $z\lesssim2$. Their MIR fluxes are produced by a combination of continuum and Polycyclic Aromatic Hydrocarbon (PAH) emission features. The PAH features can dominate the total MIR flux, but are difficult to measure without spectroscopy. } { We aim to study star-forming galaxies by using a blind spectroscopic survey at MIR wavelengths to understand evolution of their star formation rate (SFR) and specific SFR (SFR per stellar mass) up to $z\simeq0.5$, by paying particular attention to their PAH properties. } { We conducted a low-resolution ($R\simeq 50$) slitless spectroscopic survey at 5--$13\ \mu$m of $9\ \mu$m flux-selected sources ($>0.3$~mJy) around the North Ecliptic Pole with the Infrared Camera (IRC) onboard AKARI. After removing 11 AGN candidates by using the IRC photometry, we identified 48 PAH galaxies with PAH 6.2, 7.7, and $8.6\ \mu$m features at $z<0.5$. The rest-frame optical--MIR spectral energy distributions (SEDs) based on CFHT and AKARI/IRC imaging covering 0.37--$18\ \mu$m were produced, and analysed in conjunction with the PAH spectroscopy. We defined the PAH enhancement by using the luminosity ratio of the $7.7\ \mu$m PAH feature over the $3.5\ \mu$m stellar component of the SEDs. } { The rest-frame SEDs of all PAH galaxies have a universal shape with stellar and $7.7\ \mu$m bumps, except that the PAH enhancement significantly varies as a function of the PAH luminosities. We identified a PAH-enhanced population at $z\gtrsim0.35$, whose SEDs and luminosities are typical of luminous infrared galaxies. They show particularly larger PAH enhancement at high luminosity, implying that they are vigorous star-forming galaxies with elevated specific SFR. Our composite starburst model that combines a very young and optically very thick starburst with a very old population can successfully reproduce most of their SED characteristics, although we could not confirm this optically thick component from our spectral analysis. } {} | Mid-infrared (MIR) extragalactic studies have been providing new insights about galaxies in the distant universe, for three main reasons: first, about half of the total energy throughout cosmic history is emitted between the MIR and far-infrared (FIR) wavelengths (e.g., \citealt{elbaz02,lefloch05,dole06,caputi07,goto11a}). Second, the effect of dust extinction is much less prominent at MIR wavelengths when compared to optical (OPT) and near-infrared (NIR) wavelengths, and is a good spectral region for measuring activity from star formation as well as active galaxy nuclei (AGNs), even in the presence of copious amounts of dust. Third, under limited technology at the time of AKARI \citep{murakami07} and before, in particular about the large cryogenic space telescope for sharper diffraction-limited resolution and the sensitive detector system, the MIR spectral region has been more sensitive to flux from distant astronomical sources than the FIR one. The importance of deep MIR extragalactic surveys was first recognised by the discovery of strong evolution from $15\ \mu$m source counts by using ISOCAM \citep{isocam} onboard the {\it Infrared Space Observatory} ({\it ISO}; \citealt{iso}). The rapidly evolving population was found as an excess of $15\ \mu$m sources at a flux of 0.1--0.5~mJy (e.g., \citealt{elbaz99,serjeant00}; see also \citealt{lagache04,wada08,pearson10}). Later, extensive studies at MIR and other wavelengths helped to define the global spectral energy distribution (SED) shapes across the OPT--NIR--MIR--FIR for various types of luminous galaxies (such as AGNs, starburst galaxies, luminous infrared galaxies (LIRGs), and ultra-luminous infrared galaxies (ULIRGs); e.g., \citealt{spinoglio95, pearson01,lagache04,rowan04,lefloch05}). Such studies clearly indicated that most of these galaxies, excluding AGNs, show prominent emission features in their MIR spectra, which are believed to originate in Polycyclic Aromatic Hydrocarbons (PAHs; e.g., \citealt{lutz98,xu98}). The luminosity of the PAH features, as well as that of the underlying hot dust continuum, has been used as a measure of star formation rate (SFR), using conversions from MIR luminosity to the FIR one, where the bulk of the energy from star-forming regions is emitted (e.g., \citealt{genzel98,rigopoulou99,farrah07,shipley16}). The unprecedented sensitivity of {\it Spitzer} at MIR wavelengths has greatly improved our understanding of cosmic galaxy evolution, with help of various diagnostics of galaxies for their activities up to $z\sim 2$--4 provided by both MIPS imager \citep{rieke04} at $24\ \mu$m and IRS spectrometer (\citealt{hock04}; e.g., \citealt{sajina07,yan07,pope08,wu10}; see also \citealt{spoon07}). Cosmic star formation history, or star formation rate density (SFRD), in galaxies and/or AGNs has been particularly examined (e.g., \citealt{menendez07,farrah08,pope08,nordon12}. See also \citealt{goto10,goto11a,goto11b}). In many studies, the analysis has relied on SED templates and/or models for scaling from the MIR wavelengths to the bulk of the dust emission in the FIR wavelengths. The scaling is based mostly on studies of nearby galaxies and AGNs; there is no guarantee that it is appropriate at higher redshifts. In some rare cases, extremely deep {\it Spitzer} FIR photometry was used to directly find global SEDs even at higher redshifts ($z\sim 2$--3; e.g., \citealt{lefloch05,bavouzet08,murphy11}). Recent {\it Hershel} \citep{pilbratt10} FIR photometry has improved the situation (e.g., \citealt{berta11,elbaz11,gruppioni13,magnelli13}). It turned out that their NIR--MIR--FIR SEDs are systematically different from scaled-up versions of local SED templates of presumably the same activity type (e.g., \citealt{menendez07,farrah08,pope08,elbaz11,murphy11,magdis12,nordon12}). \cite{nordon12} argued that this introduces significant offsets in measuring SFRD. They claimed that the MIR spectral features are not simple tracers of SFR, but that their power is modulated by changing physical conditions in the Inter Stellar Medium (ISM) or in Photo-Dissociation Regions (PDRs) (see also \citealt{elbaz11}). This appears reasonable because such an effect has been indeed observed in spatially resolved SINGS studies of local galaxies \citep{sings}, as well as GOALS studies of luminous infrared galaxies \citep{goals}, where a range of MIR spectral features is seen within individual objects \citep{dale06}. We note, however, that their integrated properties over the galaxy scale greatly smear out these local variations \citep{bavouzet08}. Another serious problem in analysing galaxy evolution has been the uncertainty in the assumed $K$-correction used to derive rest-frame quantities (e.g., rest-frame MIR luminosity function and SFRD, \citealt{lefloch05,caputi07,bavouzet08,nordon12}). The $K$-correction to rest-frame monochromatic MIR luminosity is particularly large for redshifted ($z\gtrsim 1$) galaxies because of the contribution of the PAH features. {\it Spitzer} observed mainly in its IRAC \citep{irac} $8\ \mu$m and MIPS $24\ \mu$m bands, which miss the most prominent PAH features around $7.7\ \mu$m at $0.3<z<1.8$. Without completely reliable redshift information, interpreting the observed in-band fluxes is not straightforward at MIR. This also introduced a complicated selection function for {\it Spitzer} colour selection for higher-$z$ sources. For example, sensitive MIPS $24\ \mu$m surveys favour both $z\sim 2$ star-forming galaxies with the redshifted prominent PAH~$7.7\ \mu$m feature and AGNs with very red continuum emission. Distinguishing these possibilities has called for IRS spectroscopy (e.g., \citealt{yan04,yan07,pope08}). Although the SED templates have been empirically calibrated to reproduce various observed correlations among broad-band photometric data (e.g., \citealt{chary01,lagache04}), the uncertainties in $K$-corrections are still large. In particular, to make the MIR part of the templates realistic showing complex MIR spectral features, observed MIR spectra of small number of representative galaxies are often implanted on empirical low-resolution SED templates that are based on a simple synthetic model of dust emission \citep{dale01} or a galaxy stellar evolutionary model in combination with radiative transfer calculation through dusty circumstellar regions \citep{polletta07}. AKARI was a cryogenic space infrared telescope that observed in the NIR, MIR, and FIR spectral regions \citep{murakami07}. In addition to its primary mission to perform an all-sky survey \citep{murakami07,ishihara10}, some time was spent on the deeper pointing-mode observations of some specified targets during an intermittence of the scanning. Multi-band deep extragalactic imaging surveys at NIR and MIR (2--$24\ \mu$m) were performed toward the North Ecliptic Pole (NEP) region by using the pointing mode (the AKARI NEP surveys; \citealt{maruma06} for the summary) with the wide-field Infrared Camera (IRC; \citealt{onaka07}). The AKARI NEP surveys included as many contiguous bands as possible, with the data covering the entire wavelength range continuously with 9 filters at a spectral resolution of $R\simeq 5$. With this filter set, one can discriminate range of SED types, including AGNs with red continuum-dominated SED, normal and starburst galaxies with hot dust continuum and PAH features, luminous infrared galaxies with strong silicate absorption (peaking at $9.7\ \mu$m), up to $z\sim 2$ (see, e.g., \citealt{takagi05,wada08,takagi10,hanami12}). As \cite{takagi10} demonstrated (see also \citealt{hanami12,kim12}), the MIR colours of redshifted sources show extreme diversity, due to combination of the MIR features from a range of the SED types and redshift, and such rich multi-colour information can be used to extract various information about the nature of the sources. \cite{takagi07,takagi10} and \cite{hanami12} have utilised SED fitting techniques to extract the activity type, redshift, extinction, MIR and FIR luminosities, and so on, from the complex SEDs with much less ambiguity than before. During the course, they also showed that some observed SEDs could not be well reproduced by simple models. Even if the fit seems successful, we need to be cautious, because it relies on local SEDs or very simple physical models to fit observations of redshifted galaxies in which physical conditions may be different from the local ones. Because of the rapid evolution of galaxies peaking at $z\sim 2$ (e.g., \citealt{elbaz99,serjeant00,lagache04}), evolution of the SEDs should be examined and taken into account to interpret various observables, such as source counts and monochromatic MIR luminosities. Spectroscopy of the MIR photometric sample of galaxies in the similar wavelength range would provide a recipe to properly interpret the photometric properties. The IRC had not only multi-band imaging capability, but also a wide-band low-resolution spectroscopy capability \citep{ohyama07}. This was possible because the IRC includes transmissive direct-view dispersers (a prism and grisms) in the filter wheel, as well as the broad-band imaging filters. In the spectroscopy mode, by using the short slits at an edge of the field of view (FOV) of the IRC, spectroscopic studies of active galaxies have been extensively done. Especially, the $3.3\ \mu$m PAH feature has been utilised to trace the star formation activities and to diagnose the AGN activity \citep{imanishi08,imanishi10,woo12,castro14,ichikawa14,yano16}. In addition to the regular slit spectroscopy, the IRC could perform slitless spectroscopy: all sources that can be imaged within its $\simeq 10'\times 10'$ FOV are dispersed by either a prism or grism. This slitless mode was particularly well-suited to blind spectroscopic surveys of point-like sources: The survey can be unbiased because there is no pre-selection of sources from, e.g., colour or flux at other wavelengths. Instead, the sources are simply selected after the spectroscopy observations on the basis of their fluxes at the same wavelengths as for the spectroscopy. This is particularly important for studying the MIR evolution of galaxies, because, as noted earlier, it is quite difficult to define a fair sample that is independent of types (including AGN) and strengths of the activities, and/or redshift for statistical studies. Sensitive observations with this mode at MIR wavelengths are available only from space to avoid high atmospheric background, and the IRC was the first instrument that provided us with this unique observing opportunity. In this paper, we first describe the design and observation of our survey program in Sect.~\ref{survey_section}, and data reduction procedure in Sect.~\ref{data_reduction_section}. Our spectral PAH fit is described in Sect.~\ref{pahfit_section}, and the results are analysed in Sect.~\ref{spectral_characteristics_results}. The OPT--NIR--MIR broad-band photometry is also compiled for the spectroscopic sample in Sect.~\ref{photo_data}, and its basic characteristics are analysed in Sect.~\ref{photo_characteristics_results}. In particular, we photometrically classify activity types, normal/starburst galaxies and AGNs, in Sect.~\ref{photo_classification_results}. For galaxies with the PAH features detected in their spectra, we analyse their colour-redshift relations in Sect.~\ref{colour_redshift_diagrams}, and construct their rest-frame SEDs in Sect.~\ref{rest_frame_seds}. We then compare their spectroscopic and photometric properties in Sect.~\ref{spec_photo_comp}. In particular, we compare PAH luminosities measured in both photometric and spectroscopic ways in Sect.~\ref{pah_luminosity_comp_section}, and characterise the variation of the rest-frame SED shape by the spectroscopic properties in Sect.~\ref{sed_variation_result}. We here identify the PAH-enhanced population at $z\simeq 0.35$--0.5 as a distinctive subgroup of the PAH galaxies. Next we compared the observed rest-frame SEDs with various SED templates and models in Sect.~\ref{SED_fit_result}. We discuss implications of the SED variation of the PAH galaxies, in particular the PAH-enhanced population, for the star-formation properties in Sect.~\ref{SFR_sSFR_discussion}. We finally summarise advantages and limits of our slitless spectroscopic and photometric data analysis in Sect.~\ref{method_advantage_limit}. Conclusions are given in Sect.~\ref{conclusion_section}. We use $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\rm m}=0.3$, and $\Omega_{\rm \lambda}=0.7$ throughout this paper. | In this paper, we have reported results from our MIR spectroscopic survey toward the NEP with the IRC on the AKARI space telescope. The survey, known as ``slitless SpectroscoPIC surveY of galaxies'' ({\it SPICY}), is a slitless spectroscopic survey of a $9\ \mu$m flux-selected sample. This is the very first blind spectroscopic survey conducted at MIR wavelengths. The biggest advantage of the slitless spectroscopy is that the survey can be unbiased, because sources are selected simply based on their flux at the same wavelengths as for the spectroscopy. This enables us to study the MIR evolution of galaxies in more robust way. We analysed the {\it SPICY} MIR spectra at 5--$13\ \mu$m, together with multi (13)-band photometric data at OPT--NIR--MIR wavelengths (0.37--$18\ \mu$m) from the combined AKARI/CFHT NEP photometric catalogue, to investigate star formation properties of galaxies at $z=0.0$--0.5. We summarise the main results and conclusions below. \begin{enumerate} \item We collected 5--$13\ \mu$m low-resolution ($R\simeq 50$) spectra of $S9W\ (9.0\ \mu\mathrm{m})>0.3$~mJy sources within 14 IRC FOVs ($\simeq 14\times 10'\times 10'$ regions). We performed simple PAH fit for the 6.2, 7.7, 8.6, and $11.3\ \mu$m PAH features on the detected extragalactic sources, excluding sources with damaged spectra (due to contamination and/or truncation that can randomly happen on the slitless spectroscopy images). We then found 48 galaxies with typical PAH features for star-forming galaxies (the PAH galaxies), with redshifts ranging $z=0.0$--0.5. We also identified 11 AGN candidates by using red NIR colours at 2--$3\ \mu$m (and also red MIR$/$NIR colours at 3--$7\ \mu$m as a secondary condition), so that the colours are consistent with those of the red continuum-dominated AGN SED template and all PAH galaxies are classified as non-AGNs. We confirmed that the AGN candidates show red continuum-dominated SEDs at NIR--MIR wavelengths. The remaining non-AGN and non-PAH galaxies are mostly E--Sa with intrinsically weak PAH features, or faint Sc--starburst galaxies below the spectroscopy sensitivity, according to their broad-band colours. Basic information about the PAH galaxies, including their source positions, OPT--NIR--MIR photometry, as well as the PAH fit results (redshifts and luminosities of the PAH $6.2\ \mu$m and $7.7\ \mu$m), were reported. Similar basic information about the bright AGN candidates were also reported. \item We constructed rest-frame SEDs of the PAH galaxies at OPT--NIR--MIR wavelengths, with redshifts from the PAH fit. We found that the rest-frame SEDs look much simpler and universal within the sample, although the observed colours show extreme diversity as functions of spectral types (Sc, starburst, or LIRG) and redshift. This made our analysis easier and probably more fundamental, in contrast to ones based on the observed-frame information. The SED is composed of an OPT--NIR (0.4--$4\ \mu$m) bump (the stellar component) and a MIR (5--$18\ \mu$m) one with a peak around $7.7\ \mu$m. Both the NIR slope of the OPT--NIR bump and shape of the MIR bump are almost identical within the sample. To characterise the SED shape, we defined and measured rest-frame flux ratios of $F_{\rm rest\ 11.3\ \mu \rm m}/F_{\rm rest\ 7.7\ \mu \rm m}$, $F_{\rm rest\ 7.7\ \mu \rm m}/F_{\rm rest\ 3.5\ \mu \rm m}$, and $F_{\rm rest\ 3.5\ \mu \rm m}/F_{\rm rest\ 2.0\ \mu \rm m}$. The $F_{\rm rest\ 7.7\ \mu \rm m}/F_{\rm rest\ 3.5\ \mu \rm m}$ ratio, which represents the relative strength of the MIR bump over the OPT--NIR one, systematically changes with redshift in a range of 2--14. This ratio increases with spectroscopic PAH luminosities, $L_\mathrm{PAH}$ ($7.7\ \mu$m) and $L_\mathrm{PAH}$ ($6.2\ \mu$m), indicating that the MIR bump is essentially composed of the PAH features around $7.7\ \mu$m, and specific SFR (SFR per stellar mass; sSFR) is higher for sources with higher SFR. In contrast, the $F_{\rm rest\ 3.5\ \mu \rm m}/F_{\rm rest\ 2.0\ \mu \rm m}$ ratio shows tight distribution and is consistent with the stellar colour. The $F_{\rm rest\ 11.3\ \mu \rm m}/F_{\rm rest\ 7.7\ \mu \rm m}$ ratio decreases by a factor of $\lesssim 2$ as $F_{\rm rest\ 7.7\ \mu \rm m}/F_{\rm rest\ 3.5\ \mu \rm m}$ increases by a factor of $\sim 2$ at $z\gtrsim 0.35$. \item We photometrically measured the monochromatic luminosities at the peak of the PAH~$7.7\ \mu$m ($\nu L_{\rm \nu\ \mathrm{photo}}$ ($7.7\ \mu$m)), and compared them with those of the spectroscopy (spectroscopic PAH~$7.7\ \mu$m luminosity, $L_\mathrm{PAH}$ ($7.7\ \mu$m), and spectroscopic monochromatic luminosities at the peak of the PAH~$7.7\ \mu$m, $\nu L_{\rm \nu\ \mathrm{spec}}$ ($7.7\ \mu$m)). We found tight correlations among them, with little systematic offset as a function of redshift, and reported the scaling relations. \item We identified PAH-enhanced population showing $\log \nu L_{\nu\ \mathrm{photo}}\ (7.7\ \mu\mathrm{m})>0.50+\log \nu L_{\nu\ \mathrm{photo}}\ (3.5\ \mu\mathrm{m})$ (or, equivalently, $F_{\rm rest\ 7.7\ \mu \rm m}/F_{\rm rest\ 3.5\ \mu \rm m}>7.0$) at $z\gtrsim 0.35$. Here, $\nu L_{\nu\ \mathrm{photo}}\ (3.5\ \mu\mathrm{m})$ is a monochromatic photometric luminosities at $3.5\ \mu$m. They show elevated SFR ($\gtrsim 30\ M_{\rm \sun}\ \mathrm{yr}^{-1}$), being comparable to LIRGs, and also elevated sSFR (by about a factor of two when compared to the upper envelope of the sSFR distribution at $z<0.35$). We found no such PAH-enhanced population at lower redshifts, and we argued that this is neither due to source selection biases nor by chance due to small number statistic. \item We found that the SWIRE SED templates of normal/star-forming galaxies, Sc, starburst (M82), and LIRG (NGC~6090), can mostly reproduce both the rest-frame SEDs and the flux ratios of the PAH galaxies. We also found that the SBURT starburst models \citep{takagi03} for older ($t_{\rm burst}\gtrsim 200$~Myr) starburst with modest optical depth can mostly reproduce the observations. For the PAH-luminous and PAH-enhanced population at $z\gtrsim 0.35$ in particular, both the NGC~6090 template and the SBURT model for middle-aged ($t_{\rm burst}\simeq 200$~Myr) starburst with modest optical depth can reproduce the observed NIR--MIR SEDs. This SBURT model, however, has a problem in reproducing relatively red OPT SEDs at the same time as the enhanced MIR bump. We developed a composite SED that combines a very young ($\simeq 70$~Myr), compact ($\Theta=1.0$) and, hence, optically very thick SBURT model and a very old (13~Gyr old) SWIRE template, and found that it can reproduce the overall OPT--NIR--MIR SEDs of the PAH-enhanced population. This suggests that the PAH-enhanced population shows very young and vigorous star-formation activities on the old stellar system, although we could not confirm the presence of the optically thick component from our spectral analysis. We note that the positive correlation between SFR and sSFR for the PAH galaxies in general can also be explained by the composite models, by changing relative contribution of the very young component over the old one. Both this composite SED and the NGC~6090 SED template still have difficulties in reproducing the relatively depressed $F_{\rm rest\ 11.3\ \mu \rm m}/F_{\rm rest\ 7.7\ \mu \rm m}$ for the PAH-enhanced population. \end{enumerate} | 18 | 8 | 1808.04566 |
1808 | 1808.05947_arXiv.txt | We apply an order-of-magnitude model of gas-assisted growth, known as pebble accretion, in a turbulent medium to suggest a reason why some systems form wide orbital separation gas giants while others do not. In contrast to traditional growth by ballistic collisions with planetesimals, growth by pebble accretion is not necessarily limited by doubling times at the highest core mass. Turbulence, in particular, can cause growth to bottleneck at lower core masses. We demonstrate how a combination of growth by planetesimal and pebble accretion limits the maximum semi-major axis where gas giants can form. We find that, for fiducial disk parameters, strong turbulence ($\alpha \gtrsim 10^{-2}$) restricts gas giant cores to form interior to $a \lesssim 40 \, \text{AU}$, while for weak turbulence gas giants can form out to $a \lesssim 70 \, \text{AU}$. The correspondence between $\alpha$ and semi-major axis depends on the sizes of small bodies available for growth. This dependence on turbulence and small-body size distribution may explain the paucity of wide orbital separation gas giants. We also show that while lower levels of turbulence ($\alpha \lesssim 10^{-4}$) can produce gas giants far out in the disk, we expect these gas giants to be low-mass ($M \lesssim \, 1 M_J$). These planets are not luminous enough to have been observed with the current generation of direct-imaging surveys, which could explain why wide orbital separation gas giants are currently observed only around A stars. | In the traditional ``core accretion" model of planet formation, growth of planets proceeds in a bottom-up manner. Planets begin their growth as rocky cores, or protoplanets. If these protoplanets reach sufficient size within the lifetime of the gas disk, they will be able to trigger runaway gas accretion, resulting in a gas giant (\citealt{pollack_gas_giants}). This runaway occurs when $M_{\rm{atm}} \sim M_{\rm{core}}$, where $M_{\rm{atm}}$ is the mass of the planet's atmosphere and $M_{\rm{core}}$ is the mass of the planet in solids. The critical core mass, $M_{\rm{crit}}$, where this occurs is usually quoted as $M_{\rm{crit}} \sim 10 M_\oplus$, though the actual mass depends on the disk parameters, especially the opacity and the core's accretion rate (see, e.g. \citealt{raf06}, \citealt{pymc_2015}). A gas giant will not form if the planet cannot reach $M_{\rm{crit}}$ within the lifetime of the gas disk, $\tau_{\rm{disk}}$, which is $\sim 2.5 \, \text{Myr}$ for G stars (\citealt{lifetimes_mamajek}, \citealt{lifetimes_ribas}). Traditional models rely on gravitational focusing to increase the effective radius for collisions. These models, which we will refer to as ``canonical core accretion" or ``planetesimal accretion" models, give growth timescales that are generally fast enough to reach critical core mass for $a\lesssim10\,\text{AU}$, but become longer than the disk dispersal timescale past this distance. (See \citealt{gold}, hereafter GLS, for a review of gas-free regimes.) Observations of exoplanetary systems have challenged this canonical core accretion model in a number of ways. Here we focus on the existence of systems that feature gas giants at wide orbital separations (see, e.g. \citealt{bowler_DI_review} for a review). Of particular note is the planetary system surrounding the star HR 8799, which exhibits a nonhierarchical, multiplanet structure: HR 8799 consists of four gas giant planets ($M\sim10 \, M_J$) at extremely wide projected separations: 14, 24, 38, and 68 AU (\citealt{HR8799_orig}, \citealt{HR8799_fourth}). HR 8799 poses a serious challenge to canonical core accretion models because the last doubling timescale for growth at these distances is far too long for a core to reach the critical mass necessary to trigger runaway growth within $\tau_{\rm{disk}}$. Additional effects, such as gas drag from the planet's atmosphere (\citealt{ii_03}) or damping of the planetesimals' random motions by the nebular gas (\citealt{raf}), can increase the cross section for collisions further. Neither of these effects, however, are sufficient to allow the \textit{in situ} formation of gas giants at $70 \, \text{AU}$. A number of alternative formation scenarios have been proposed to explain the formation of HR 8799. One commonly suggested explanation is that HR 8799 is evidence of an alternative formation scenario known as ``gravitational instability," wherein the gaseous component of the protoplanetary disk becomes unstable to gravitational collapse and subsequently fragments into the observed gas giant planets (\citealt{boss_1997}; see also \citealt{kl_2016} for a more recent review). However, \cite{kratter_gas_giants} pointed out that it is difficult to form fragments of the sizes seen in HR 8799 without having these ``planets" grow to brown dwarf or even M-star masses. The lack of observed brown dwarfs at wide orbital separations provides some evidence against this hypothesis, but additional statistical work is needed (\citealt{bowler_DI_review}). Outward scattering after formation at smaller orbital separations is another possibility, but $N$-body simulations by \cite{dodson-robinson_gas_giants} find that scattering is unlikely to produce systems with the multiplanet architecture of HR 8799. In recent years, a third possibility has emerged: a modification to the theory of core accretion commonly referred to as ``pebble accretion," which we will also refer to as ``gas-assisted growth" (\citealt{OK10}, \citealt{pmc11}, \citealt{OK12}, \citealt{lj12}, \citealt{LJ14}, \citealt{lkd_2015}, \citealt{mljb15}, \citealt{vo_2016}, \citealt{igm16}, \citealt{xbmc_2017}, \citealt{rmp_2018}). In pebble accretion, the interaction between solid bodies and the gas disk is considered in detail when determining the growth rates of planets. In particular, gas drag can enhance growth rates by removing energy from small bodies. Particles that deplete their kinetic energy within the gravitational sphere of influence of a larger body can become bound to this parent body, which will eventually lead to accretion of the particle by the growing protoplanet. This process can occur at larger impact parameters than are required for the particle to collide with the core, which in turn increases the accretion cross section. This interaction often affects mm-cm-sized bodies the most strongly. Note, however, that for low-density, ``fluffy" aggregates, the radius of bodies most substantially affected by gas drag can be substantially larger. For gas-assisted growth to operate, a reservoir of pebble-sized objects must exist in the protoplanetary disk. Because the sizes of these pebbles are comparable to the $\sim$ mm wavelengths used to measure dust surface densities in the outer regions of protoplanetary disks, observations can directly probe the surface densities in the small solids that fuel gas-assisted growth. These observations find large reservoirs of small, pebble-sized solids (\citealt{andrews_09}, \citealt{andrews}). An example is shown in Figure \ref{fig:andrews09}, which presents disk surface densities measured by \cite{andrews_09}. The figure shows the surface density in particles of radius $0.1 \, \text{mm} - 1 \, \text{mm}$, which is inferred by integrating the size distribution used in the paper ($d N / d r_s \propto r_s^{-3.5}$) from 0.1 to $1 \, \text{mm}$. Performing the integration gives the fraction of the measured solid surface density contained in this size range ($\sim 70\%$). \begin{figure} [h] \centering \includegraphics[width=\linewidth]{rev_sd_comp_3} \caption{Colored lines show the dust surface density in $0.1 \, \text{mm} - 1 \, \text{mm}$ sized particles taken from $870$ $\mu \text{m}$ continuum emission observations of protoplanetary disks done by \cite{andrews_09}. See text for details. Also shown for reference is the value of the solid surface density in the minimum-mass solar nebula (MMSN), appropriate for the outer disk, $30 \, (a/\text{AU})^{-3/2} \, \text{g} \, \text{cm}^{-2}$ (\citealt{weid_mmsn}, \citealt{hay_mmsn}), as well as the fiducial surface density used in this work to match the observations. In the gray shaded region the values of the curves are extrapolations to scales smaller than the observations can resolve.} \label{fig:andrews09} \end{figure} Given this observed reservoir of small solids, pebble accretion dramatically increases the expected growth rate of large cores. Under fiducial conditions, the timescale for a core's last doubling to canonical values of $M_{\rm{crit}}$ is below the disk lifetime, even at many tens of AU separations. Though fast accretion of solids deposits enough energy to delay the onset of runaway accretion of a gas envelope, once a core has reached several Earth masses, finely tuned disk conditions are required to slow atmospheric growth enough to prevent runaway from ultimately occurring. Thus, growth via pebble accretion seems to predict that wide orbital separation gas giants should be common. However, direct imaging surveys show that planets $\gtrsim 2-5 \, M_J$ are rare at distances $> 30 \, \text{AU}$ (\citealt{brandt_di}, \citealt{chauvin_di}, \citealt{bowler_DI_review}, \citealt{gal_gas_giant_freq}). One possibility for solving this problem is the presence of turbulence in the nebular gas. In this work, by ``turbulence," we generally mean any anomalous root mean square (RMS) velocity of the nebular gas that is not due to the laminar velocity that arises from radial pressure support in the disk. The main effect of turbulence on pebble accretion is to increase the velocity dispersion of the pebbles due to their coupling with the gas; it is only in Section \ref{final_mass} that we connect our parameterization of the turbulent RMS velocity to the transport of angular momentum in the disk. Turbulence can both increase the kinetic energy of an incoming particle and decrease the core's gravitational sphere of influence. Turbulence also drives particles vertically, reducing the overall densities of small bodies and slowing accretion. Turbulence is usually only included in models of protoplanetary growth by pebble accretion by increasing the particle scale height and hence reducing the mass density of solids. Some models of the early stages of planetesimal growth discuss the effects of turbulence (e.g. \citealt{gio_2014}, \citealt{hgb_2016}), but these models are concerned with accretion at cross sections comparable to the core's geometric cross section; i.e. they neglect the effects of the core's gravity. In this paper, we use an order-of-magnitude model of pebble accretion (\citealt{rmp_2018}, hereafter R18) to propose a criterion for the formation of gas giants via gas-assisted growth. In particular, R18 investigated how turbulence affects the growth of gas giant cores as a function of core mass. High-mass cores ($\gtrsim 10^{-2}-10^{-1} M_\oplus$) can grow on timescales less than the lifetime of the gas disk, even in strong turbulence. However, for lower-mass cores and stronger turbulence, the range of pebble sizes available for growth is restricted. In this case, the pebble sizes for which growth is most efficient often cannot be accreted, and growth can ``stall" at low core masses. In effect, a core must first achieve a minimum mass before it can quickly grow to $M_{\rm{crit}}$ via gas-assisted growth. In this paper, for our fiducial calculation, we assume that growth to this minimum mass happens by canonical core accretion, which allows us to place semi-major axis limits on where gas giant growth is possible. We also calculate values for the core mass needed at a given semi-major axis for pebble accretion to be rapid, which apply regardless of how the early stages of growth proceed. The assumption that low-mass growth is fueled by planetesimal accretion requires that, in addition to the reservoir of small pebbles, a substantial population of larger planetesimals has formed. We discuss the ramifications of varying the mass in planetesimals in Section \ref{upper_lim}. Close to the central star, planetesimal accretion can dominate the early growth of planets, with pebble accretion setting the growth timescale for high-mass cores. Far from the central star, however, planetesimal accretion is less efficient, limiting its ability to grow cores to high enough masses that pebble accretion kicks in. Thus, turbulence can set the maximum distance at which gas giant formation is possible via pebble accretion. We find that for quiescent disks, gas giants can form far out in the disk ($a \lesssim 70 \, \text{AU}$), but for stronger turbulence, this maximum distance is smaller (e.g. $a \lesssim 40 \, \text{AU}$ for $\alpha \gtrsim 10^{-2}$). Furthermore, while disks with weaker turbulence can have gas giants at wider orbital separation, the weaker viscosities in these disks mean that the masses of the gas giants formed are likely lower ($\lesssim 2 \, M_J$), which would preclude them from being detected by the current generation of direct-imaging surveys. Therefore, there may exist a population of wide orbital separation gas giants that have yet to be found due to their low luminosities. In Section \ref{overview}, we review our model, which is discussed in detail in R18. In Section \ref{wide_sep}, we discuss how gas-assisted growth operates at wide orbital separation, contrasting the rapid growth at high core mass with the slower growth for low-mass cores. In Section \ref{gas_giants}, we explore how turbulence can place limits on the semi-major axes where gas giants can form. In Section \ref{final_mass}, we investigate the implications for the final masses of gas giants if turbulence plays a role in gap opening in addition to early core growth. Finally, in Section \ref{summary}, we summarize our results and give our conclusions. | \label{summary} In this paper, we have used our previously discussed model of gas-assisted growth in a turbulent disk to study the problem of growth of gas giants at wide orbital separations. At these large distances, last doubling timescales for growth by planetesimal accretion are far longer than the disk dispersal timescale of the gas, making growth of gas giants by canonical core accretion extremely difficult. Gas-assisted growth allows cores to easily complete their last doubling time to critical core mass, even in strong turbulence. The maximal growth rate provided by pebble accretion, $t_{\rm{Hill}}$, is extremely rapid, even in the outer disk. For massive cores. even strong turbulence does not substantially inhibit growth. The same is not true for smaller core masses. however. Growth of gas giants at large distances can easily stall at smaller core masses. By integrating our growth rates over small-body size we obtained the minimum mass past which pebble accretion timescales drop below the lifetime of the gas disk, $M_{\rm{peb}}$. By assuming that the early stages of growth are set by gravitational focusing of planetesimals, we were able to translate these minimum masses into limits on the semi-major axes where gas giant growth is possible. We demonstrated that as the disk becomes more turbulent, the range of semi-major axes where gas giants can grow is sharply reduced. These effects may play a large role in the paucity of gas giants at wide orbital separations found by direct-imaging surveys; if disks are not quiescent enough, then pebble accretion may simply produce smaller planets that are unable to accrete sufficient mass in small bodies to go critical. In addition, our mass limits are relevant regardless of how early growth proceeds -- for example, if a body were scattered from the inner disk and it exceeded our minimum mass, it could grow to $M_{\rm{crit}}$ and trigger runaway gas accretion. We also presented approximate analytic expressions for both $M_{\rm{peb}}$ and the upper distance limit where gas giants can form, $a_{\rm{upper}}$. In addition to the strength of turbulence, we find that the available particle sizes and abundance of planetesimals are major factors in where gas giants can form. Gas-assisted growth is sensitive to the Stokes numbers of the pebbles, as opposed to their absolute size. Thus, if particles of the ``correct" range of Stokes numbers are not available, then gas-assisted growth timescales can be quite slow. Furthermore, if planetesimals are not abundant enough, then the early stages of growth via planetesimal accretion can take too long for subsequent growth via pebble accretion to occur on rapid timescales. Finally, we examined the role that viscosity plays in determining the final mass that gas giants reach, in addition to setting where a critical mass core can form. We find that, regardless of the quantitative metric used to determine the final gas giant mass, higher-viscosity disks should feature higher-mass gas giants but at smaller orbital separations. More quantitatively, at the lower $\alpha$ values needed to produce gas giants out to $a \gtrsim 70$ AU, the gas giants formed will be too low-mass to have been observed by direct-imaging surveys. Thus, there may lurk a population of wide orbital separation gas giants that the current generation of imaging surveys has yet to detect. Thus, if growth of gas giants at wide orbital separations proceeds by gas-assisted growth, and if gap opening sets the final masses of gas giants, we can make qualitative predictions about the observed population of gas giant planets. For a given stellar mass, we would expect that higher-mass gas giants should be observed closer in to the central star, as the disks with higher levels of turbulence will produce higher-mass gas giants at smaller orbital separations. We note that this conclusion may be altered if final planet masses depend on disk surface density, which is not the case for the gap-opening criterion we use. For different stellar masses, we expect higher-mass stars to exhibit higher levels of ionizing radiation. Thus, larger stars may have disks with higher $\alpha$ values and consequently host more massive gas giants. In addition, the larger disk masses exhibited by more massive stars could push the limits for growth past the distances given here for our fiducial surface density, allowing planets to form at larger distances in high-$\alpha$ disks. We suggest that this propensity to produce massive planets in high-turbulence disks may be the reason that most currently observed directly imaged gas giants have been found orbiting A stars. \vspace{1mm} \noindent The authors wish to thank Diana Powell, Renata Frelikh, and John McCann for useful discussions, and Eugene Chiang for his thoughtful suggestions on the manuscript. We also thank the anonymous referee for their helpful comments that improved the quality of the manuscript. MMR and RMC acknowledge support from NSF CAREER grant number AST-1555385. \appendix In this Appendix, we describe in more detail how we calculate the parameters that are used to determine $t_{\rm{grow}}$. For more information, see R18. | 18 | 8 | 1808.05947 |
1808 | 1808.00279_arXiv.txt | We present the first-ever images of the circumstellar environment of the carbon-rich AGB star II~Lup in the infrared and sub-mm wavelengths, and the discovery of the envelope's non-spherical morphology with the use of high-angular resolution imaging techniques with the sparse aperture masking mode on NACO/VLT (that enables diffraction limited resolution from a single telescope) and with ALMA. We have successfully recovered images in $Ks$ (2.18\micron), $L'$ (3.80\micron) and $M'$ (4.78\micron), that revealed the non-spherical morphology of the circumstellar envelope around II~Lup. The stellar surface of the AGB star is unresolved (i.e. $\leq30$ mas in $Ks$) however the detected structure extends up to 110 mas from the star in all filters. Clumps have been found in the $Ks$ maps, while at lower emission levels a hook-like structure appears to extend counter-clockwise from the south. At larger spatial scales, the circumstellar envelope extends up to approximately 23 arcsec, while its shape suggests a spiral at four different molecules, namely CO, SiO, CS and HC$_3$N, with an average arm spacing of 1.7 arcsec which would imply an orbital period of 128 years for a distance of 590pc. | In the last decade significant steps have been made in understanding stellar mass loss during the asymptotic giant branch (AGB) stage of low- to intermediate-mass stars ($M_\text{init}=0.8-8$~M$_\text{\sun}$). Although the precise mechanisms for mass loss remain contentious, some interplay between pulsation, dust formation and radiative driving is highly likely. This has been demonstrated, for carbon stars, by dynamic models of \citet{hoefner2003}, \citet{nowotny2005,nowotny2011}, \citet{mattsson2010}, \citet{liljegren2016}, and their applications to high-angular interferometric observations in, e.g., \citet{wittkowski2017}, \citet{rau2017,rau2015}, \citet{ohnaka2007,ohnaka2015}. For oxygen-rich stars, the interplay has been demonstrated by \citet{bladt2015}, and the interferometric observations by \citet{wittkowski2018}, \citet{karovicova2013} and \citet{ohnaka2012}. For the case of carbon-rich stars in particular ($M\leq4$~M$_\text{\sun}$), it has been suggested that these stars produce high mass-loss yields, and are therefore important players in the chemical enrichment of the interstellar medium in the solar neighbourhood \citep[e.g.,][]{mattsson2010b,schroeder2001}. However, such predictions are also dependent on the metallicity and the model in use \citep[e.g.,][]{karakas2016}. ~\citet[][and references therein]{nowotny2013} have shown that increasing mass-loss rates in carbon-rich AGB stars will result in higher circumstellar reddening, and therefore these stars will appear more obscured \citep[see also][]{liljegren2016} . On the subject of what shapes the stellar ejecta, the two most-favoured mechanisms are magnetic fields and binarity, although of course these two are not mutually exclusive and hybrid models may also be favoured. If the morphology is governed by binarity, the oxygen-rich and carbon-rich stars should show the same range of morphologies in their circumstellar structures, but confirming this will require a large sample of sources. The shape of the stellar ejecta beyond the AGB phase, as witnessed in the majority of imaging surveys \citep[e.g.,][]{castro-carrizo2010,lagadec2011}, often departs from spherical symmetry (hereafter, asymmetry). This is more evident in later evolutionary stages, such as the planetary nebula phase where only about 20\% of planetary nebulae have been found to be spherically symmetric \citep{parker2006}. With the advancement of observational techniques, it is now possible to detect such morphological changes even within two stellar radii from the surface of AGB stars. As part of an on-going effort in detecting asymmetries in the ejecta of AGB stars \citep[e.g.,][]{lykou2015,lagadec2011,tuthill2000,tuthill2000b}, we present here our investigation of one such star, II~Lup (Section~\ref{thestar}). This work is based on single-dish and on interferometric observations (Section~\ref{obs}) and the results are compared with historical findings from speckle imaging (Section~\ref{mysizes}). We present the first-ever images of II~Lup in the near-infrared (Section~\ref{imaging}) and the sub-mm (Section~\ref{alma}) wavelengths. We derive the physical parameters of II~Lup in Section~\ref{physics}, discuss our imaging results in Section~\ref{discussion}, and present our final conclusions in Section~\ref{conclusions}. \begin{table} \centering \caption{II~Lup astrometry.}\label{positions} \begin{tabular}{lllc}\hline Survey & RA (J2000) & Dec (J2000) & Epoch\\ \hline USNO-B1.0 & 15 23 05.154& -51 25 58.65&1986.7 \\ USNO-B1.0 &15 23 05.214& -51 25 59.19 & 1987.4\\ PPMXL & 15 23 05.091& -51 25 58.79 &1989.91\\ PPMXL & 15 23 05.106& -51 25 58.94&1991.43\\ GSC2.2 &15 23 05.077 & -51 25 58.58&1997.288\\ GSC2.3 & 15 23 05.084 & -51 25 58.82 &1992.56\\ 2MASS & 15 23 05.075 &-51 25 58.73 &1999.441 \\ DENIS &15 23 05.072 &-51 25 58.84 &1998.458\\ SHS & 15 23 05.083& -51 25 58.89&1999.505\\ HST$^a$ &15 23 04.999 &-51 25 58.29 &2004.611\\ {\it Gaia} DR2 & 15 23 05.0744 & -51 25 58.881 & 2015.5\\ \hline \multicolumn{4}{l}{$^a$ Position derived from {\sc sextractor} photometry in the {\em Hubble}}\\ \multicolumn{4}{l}{Legacy Archive.}\\ \end{tabular} \end{table} \begin{figure} \centering \includegraphics[width=0.4\textwidth]{./figure1.pdf} \caption{A field image of II~Lup taken with the {\em Hubble} Space Telescope's ACS/HRC camera and the $F606W$ filter. II~Lup is the brightest star in the field, and the region is overlayed with the pointing of several surveys, namely {\it Gaia} DR2 (blue boxes), 2MASS (red circle), GSC2.3 (green triangles), PPMXL (magenta diamonds) and USNO-B.1 (black boxes).} \label{hst} \end{figure} \begin{figure} \centering \includegraphics[width=0.45\textwidth]{./figure2.pdf} \caption{A three-colour image of the field around II~Lup from the AAVSO Photometric All-Sky Survey (APASS) in 2011. The colours blue, green and red stand for the filters Johnson $B$, $V$ and Sloan $i'$, respectively. The entire field-of-view shown in Fig.~\ref{hst} is indicated by a white box for comparison. II~Lup was not detected in filters $B$ and $V$ (probably fainter than 16.5 mag in $V$). The oblate morphology of the source within the box is in part due to the poor angular resolution and focus of the APASS survey, and due to the nearby, eastern field stars in Fig.~\ref{hst} which are also bright in $i'$. The position of II~Lup in the {\it Gaia} system is marked with a green circle.} \label{apass} \end{figure} | The carbon star II~Lup is known to exhibit an obscuration event in addition to a primary period of 575 days. According to \citet{feast2003} this can be explained either by the ejection of dust clumps or by the presence of a companion. After re-analysing the infrared lightcurve of II~Lup, we confirm a secondary period of approximately 19 years. However, due to the incompleteness of the sample, the true secondary period may be significantly longer. We attempted to estimate the physical parameters of II~Lup (luminosity, effective temperature and mass), however we find that the results are inconclusive due to the large uncertainties of the measurements used. Nevertheless, we obtain a dust temperature of $\sim$1200 K from the $Ks$ angular size of II~Lup. This work presents the first-ever images of the dusty envelope of II~Lup in near-infrared and sub-mm wavelengths. We have used an aperture masking technique to obtain diffraction limited images from a single-dish telescope in $Ks$, $L'$ and $M'$. The interferometric data revealed that: \begin{itemize} \item[(i)] the angular size of the source, $\theta_{Ks}$, is less than 35 mas, and therefore the stellar radius is much smaller than 20 au (assuming a distance of 590pc), \item[(ii)] the morphology of the circumstellar envelope is non-spherical and that its shape is hook-like at low emission levels ($Ks$ map, Fig.~\ref{macimall}), possibly resembling the inner coil of a spiral, or associated with a counter-spiral shaped by orbital motion \citep[e.g., ][]{mohamed2012}, \item[(iii)] the morphology of the envelope is different at longer wavelengths ($L'$ and $M'$), where the envelope expands directly north from the central source. However, due to a lower resolution and signal-to-noise at those wavelengths, the fidelity of those images is lower and we can conclude that the envelope is smaller than 200 mas in these bands. \end{itemize} Our analysis of the archived ALMA channel maps from $^{12}$CO ($J=1-0$), SiO ($\nu=0, J=1-0$), CS ($J=2-1$) and HC$_3$N ($J=11-10$), reveals the true extent of the non-spherical circumstellar envelope and most importantly hints at a large-scale spiral structure in all maps. From the spacing of the spiral arms in the CO map ($\sim1.7$ arcsec), we derive an approximate orbital period of 128 years. A study of the higher CO transitions tracing the inner wind of II Lup would be beneficial in determining the history of the mass loss process. Given the complexity of this object we suggest new observations of the circumstellar environment of II~Lup in order to compare the morphology at different epochs. We strongly recommend (a) further monitoring of its variability in the visual and near-infrared, and (b) a polarimetric and spatio-kinematic study of its non-spherical dusty environment with large-aperture telescopes (e.g., GRAVITY/VLTI and SPHERE/VLT). | 18 | 8 | 1808.00279 |
1808 | 1808.06205_arXiv.txt | Interactions between the extragalactic background light (EBL) and very high energy $\gamma$ rays (VHE; $E > 10 \rm ~ GeV$) from cosmological sources alter their $\gamma$-ray spectrum. The stronger absorption of harder $\gamma$ rays causes a steepening of the observed $\gamma$-ray spectrum at the high energy tail which can be expressed by an increase of the power index. The effect provides a link between high energy astrophysics and the evolution of galaxies. In this work we develop a new hybrid EBL model by augmenting our previous analytic model with information from semi-analytic galaxy catalogues. The model allows us to study the $\gamma$-ray opacities of individual $\gamma$-ray paths through the (simulated) universe and evaluate the effect of local fluctuations in the EBL intensity along each path. We confirm an order of magnitude fluctuations in the local EBL (based on the cumulative light from model galaxy distributions within spheres of ~ $50 ~ h^{-1} \rm Mpc$). However, the effect of these fluctuations on the VHE $\gamma$-ray opacity is insignificant due to the overall very small contribution of the local EBL to the total EBL. We also investigate the effect which a galaxy crossing the line of sight of the $\gamma$-ray source may have on the VHE spectrum. We find that galaxies with stellar masses of $M >10^{11} \rm ~ M_\odot$ could have significant effect on the $\gamma$-ray absorption. It is unlikely that the observed variation in the spectral index is caused by the proximity of single galaxies to the $\gamma$-ray paths as these are extremely rare. | \label{sec:intro} The extragalactic background light (EBL) is the integrated diffuse light from stars and active galactic nuclei (AGN) emitted through the history of the Universe redshifted to longer wavelengths due to cosmic expansion. The spectrum of the EBL lies between the ultraviolet (UV) and far-infrared (FIR) with two distinct peaks: a first peak at $\sim 1 \rm ~ \mu m$ due to direct stellar emission; and a second peak at $\sim 100 \rm ~ \mu m$ caused by stellar light that has been absorbed and re-emitted by the intra-galactic dust. Since the EBL carries information of the star formation history of the Universe, its measurements can be used to provide constraints on star formation models and the baryonic matter content of the Universe \citep{Dwek2001}. Direct measurements of the EBL are difficult because of strong foreground contamination by galactic and zodiacal light. Nevertheless, efforts have been made to provide upper and lower limits on the EBL using intensity measurements of the sky and deep galaxy counts, respectively \cite[see][and references therein]{Dwek2001,Dwek2012}, and references therein). However, setting lower limits based on galaxy counts is less reliable at longer wavelength since unresolved galaxies become source of confusion. Furthermore, constraining the EBL in the $10-70\, \mu$m range is complicated by the emission from the interplanetary dust \citep{Kelsall1998}. On the other hand, indirect measurements of the EBL can be obtained by studying the spectral attenuation of distant $\gamma$-ray sources. Cosmic $\gamma$-rays can be absorbed along the way through pair production, $\gamma+\gamma'\rightarrow e^+ + e^-$, where $\gamma'$ indicates an EBL photon. This process causes spectral steepening of blazars spectra in the very high energy (VHE; $E>10 \rm ~ GeV$) regime \citep{Abramowski2013}. Since the attenuation depends on the $\gamma$-ray energy and the EBL density along the line of sight, it can be used to constrain the EBL. This method requires an understanding of the $\gamma$-ray source spectrum which is still a matter under discussion \citep[e.g.,][]{Stecker1996, Dwek2005, Aharonian2006, Mazin2007, Albert2008}. Thus, it is difficult to differentiate between the source-inherent effects and the signature of the EBL on the observed spectrum \citep{Mazin2007}. As illustration of the above we replicate a figure from \cite{Sinha2014} (Fig.~\ref{fig:slo}) which shows a positive correlation between the spectral index ($\Gamma$) of the high-frequency peaked BL Lac (HBL) sources and redshift, suggesting that more distant sources have a higher probability to interact with the EBL along the path leading to the steepening of the VHE spectra. The authors excluded extreme HBLs for homogeneity purposes which brings out more clearly the observed intrinsic systematic spectral hardening within this blazar class \citep[see also,][]{Ackermann2011}. However, even with this homogeneous sample there still is a large scatter in the $\Gamma$s for sources at the same redshift. Various models of blazars emission mechanisms indicate that blazars have a broad range of $\gamma$-ray peak frequencies, and therefore varying VHE spectral indices. In this work, we investigate a possible contribution from the EBL fluctuations to the scatter of the observed spectral indices. Several models have been developed to derive an overall spectrum of the local EBL and its evolution. These models use different strategies to determine the evolution of the comoving luminosity density as a function of redshift. Most recent models can be classified into to four types: ``backward" evolution models, ``forward" evolution models, ``cosmic chemical" evolution models, and ``semi-analytic" models \citep{Dwek2001}. Backward evolution models begin with the observed present day luminosity functions for galaxy populations in the local Universe and extrapolate them to higher redshift \cite[e.g.,][]{Malkan1998,Rowan-Robinson2001,Stecker2006}. Forward evolution models start with early structure formation scenarios to predict the galactic evolution forward in time \cite[e.g.,][]{Dwek1998,Razzaque2009,Finke2010,Kudoda2017a}. Cosmic chemical evolution models consider basic galaxy ingredients such as gas, metallicity and radiation content and follow their evolution in a large comoving volume element \cite[e.g.,][]{Pei1995,Pei1999}. Semi-analytic models apply physically motivated descriptions and recipes to simulate galaxy formation and evolution \cite[e.g.,][]{Kauffmann1993, Cole1994, Somerville1999,Gilmore2012}. Semi-analytic models are a more challenging approach, but they provide a better insight into the complex physical processes involved in the production of the EBL. The standard $\gamma$-ray absorption models assume spatially homogeneous EBL densities. However, the spatial galaxy distribution is inhomogeneous on small scales, thus one expects the EBL to fluctuate to a certain degree as well. To what extent fluctuations affect the $\gamma$-ray transmissivity is not fully understood. The effect of inhomogeneous EBL densities on $\gamma$-ray opacities has been explored by \cite{Furniss2015,Kudoda2015,Kudoda2017a} and \cite{Abdalla2017}. \cite{Kudoda2017a} (hereafter referred to as KF17) predicted maximum changes of $\pm 10 \%$ in the $\gamma$-ray transmissivity, which is in agreement with the results reported by \cite{Furniss2015} and \cite{Abdalla2017}. A change of $\pm 10 \%$ in the $\gamma$-ray transmissivity, however, translates into only marginal differences in the power law slopes of the absorbed $\gamma$-ray spectra ($\lesssim \pm 1\%$). In KF17, we developed a purely analytical framework for computing the EBL fluctuation based on the phenomenologically determined variance of the SFR. In this work we develop a new technique to model the EBL (which may be adaptable for other problems as well) by combining the forward and semi-analytic approaches in order to study the impact of the local environment on the overall EBL spectrum and its effect on the $\gamma$-ray opacity of the Universe. The forward approach, borrowed from KF17, will be used to model the isotropic contribution to the local EBL from galaxies at large distances. While, the semi-analytical galaxy catalogues from the Millennium database \citep{Lemson2006,Guo2013} will be used to create a light cylinder through the Universe to simulate the $\gamma$-ray path with a more realistic distribution of galaxies close to the line-of-sight. Combining the two approaches allows us to model $\gamma$-rays on individual paths through a locally inhomogeneous EBL. The structure of the paper is set as follows. We detail the new EBL model in Section \ref{sec:mod}. The results of the EBL model and comparison with our previous EBL model are presented in Section \ref{sec:res}. The calculations of the $\gamma$-ray opacity are discussed in Section \ref{ssc:opa}. The case of a galaxy encounter with the $\gamma$-ray path is shown Section \ref{ssc:gal}. Finally, we summarize and give a brief conclusion in Section \ref{sec:con}. In this work we used the WMAP7 cosmological parameters ($\Omega_m$, $\Omega_b$, $\Omega_\Lambda$, $h$) = ($0.272$, $0.0455$, $0.728$, $0.701 \rm ~ km/s/Mpc$). \begin{figure} \begin{center} \includegraphics[width=0.5\textwidth]{fig01.pdf} \caption{Observed spectral indices $\Gamma$ of HBL as a function of redshift (data obtained from table 1 in \protect\cite{Sinha2014}) excluding extreme HBLS and those with uncertain redshifts.} \label{fig:slo} \end{center} \end{figure} | \label{sec:con} With this work we introduce a new hybrid EBL model combining semi-analytic and forward evolution approaches to determine the EBL and its fluctuations in a realistic manner. Instances of the local EBL are calculated at the centre of a $50 ~ h^{-1} \rm Mpc$ spheres as the cumulative intensity of all galaxies within those spheres. The galaxy properties are obtained from the MR7 galaxy catalogue. The spectra of individual galaxies are generated using those galaxy properties and the stellar population synthesis spectral library BC03. The EBL originating from all sources outside the $50 ~ h^{-1} \rm Mpc$ spheres is assumed to be homogeneous and calculated analytically following KF17. We find one order of magnitude fluctuations in the local EBL. However, when we compute the total EBL by adding the homogeneous background contribution the variation is less than $\sim 1 \%$. This is expected as the local EBL contribution is only a small fraction of the total EBL intensity. Based on the hybrid model we calculate the EBL along $\sim 1000$ $\gamma$-ray paths up to $z=0.5$, i.e., enveloping the $\gamma$-ray we construct light cone cylinders of $50\,h^{-1}\rm ~ Mpc$ comoving radius through the MR7 simulation volume within which we compute the local EBL based on the local galaxy content. The EBL beyond the cylinder is determined analytically for every time step. The hybrid EBL model approach then allows us to integrate the VHE $\gamma$-ray absorption along the $\gamma$-ray paths and determine the change of the observed $\gamma$-ray spectral index due to the preferential absorption of VHE $\gamma$-rays by the EBL photons. The integrated opacities agree with the observed steepening of the hard $\gamma$-ray spectra as a function of redshift. But, we find only negligible variations of the opacities along the different $\gamma$-ray paths. Thus, our hybrid model which takes into account the contribution of individual galaxies close to the $\gamma$-ray path cannot account for the scatter in the spectral index seen in Fig. \ref{fig:slo}. To cover very close encounters of galaxies with $\gamma$-ray paths (which are not apparent in our statistical sample) we manually insert an individual galaxy of different stellar mass ($5\times10^6 \rm ~ M_\odot$ to $5\times10^{12} \rm ~ M_\odot$) very close to the line-of-sight. We then compute $\gamma$-ray absorption and the increase of the spectral index $\Gamma$ which is set to be $\Gamma=-1.5$ without absorption. We find that the presence of a single galaxy with a stellar mass $<10^9 \rm ~ M_\odot$ has no measurable impact on $\Gamma$ independent of its distance to the line-of-sight. On the other hand, high mass galaxies with $10^{10}-10^{12} \rm ~ M_\odot$ at distances less than $50 \rm ~ kpc$ from the line-of-sight cause a clearly visible steepening of the VHE $\gamma$-ray spectrum. Given the future increase of the number of observed VHE $\gamma$-ray sources (roughly a factor of $10$ with the coming CTA) those encounters may become more important. | 18 | 8 | 1808.06205 |
1808 | 1808.01613_arXiv.txt | {N/A} {N/A} {N/A} {N/A} {} | 18 | 8 | 1808.01613 |
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1808 | 1808.03284.txt | We present a targeted follow-up \textit{Hubble Space Telescope} WFC3 F160W imaging study of very massive galaxies ($\log(M_{\rm{star}}/M_{\odot})> 11.2$) selected from a combination of ground-based near-infrared galaxy surveys (UltraVISTA, NMBS-II, UKIDSS UDS) at $1.5<z<3$. We find that these galaxies are diverse in their structures, with $\sim1/3$ of the targets being composed of close pairs, and span a wide range in sizes. %, as measured by the effective half-light radii of the S\'ersic models fit to their two-dimensional light profiles. At $1.5<z<2.5$, the sizes of both star-forming and quiescent galaxies %with $\log(M_{*}/M_{\odot}$)$\geq11.2$ are consistent with the extrapolation of the stellar mass-size relations determined at lower stellar masses. At $2.5<z<3.0$, however, we find evidence that quiescent galaxies are systematically larger than expected based on the extrapolation of the relation derived using lower stellar mass galaxies. We used the observed light profiles of the blended systems to decompose their stellar masses and investigate the effect of the close pairs on the measured number densities of very massive galaxies in the early universe. We estimate correction factors to account for close-pair blends and apply them to the observed stellar mass functions measured using ground-based surveys. Given the large uncertainties associated with this extreme population of galaxies, there is currently little tension between the (blending-corrected) number density estimates and predictions from theoretical models. Although we currently lack the statistics to robustly correct for close-pair blends, we show that this is a systematic effect which can reduce the observed number density of very massive galaxies by up to a factor of $\sim1.5$, and should be accounted for in future studies of stellar mass functions. | \label{sec3-intro} In contrast to the hierarchical assembly of dark matter haloes, observations indicate that the most massive galaxies in the nearby universe were among the first to build-up their stellar mass and quench. In the nearby universe, massive galaxies are found to be older, more metal rich and to have formed their stars more rapidly and at earlier cosmic epochs compared to their lower-mass counterparts \citep{terlevich01, bernardi03, trager00, thomas05, gallazzi05, gallazzi06, yamada06, kuntschner10, mcdermid15}. Corroborating their early formation times are results from recent deep near-infrared (NIR) surveys which reveal that very massive galaxies were already in place by $z\sim4$ (merely $\sim1.5$~Gyr after the Big Bang; e.g., \citealt{marchesini10, ilbert13, muzzin13a, straatman14, duncan14, tomczak14, caputi15, grazian15, song16, davidzon17}), and spectroscopic follow-up campaigns, confirming that these massive galaxies have evolved stellar populations at $z>3$ (e.g., \citealt{marsan15, marsan17, glazebrook17, schreiber18}). Thus, the observed properties of the most massive galaxies serve as critical benchmarks to understand the detailed physical mechanisms that impact galaxy formation and evolution in the early universe. A two-phase scenario has been proposed for the evolution of massive galaxies: a rapid, compact formation at early epochs via highly dissipative processes (e.g. by experiencing gas-rich major mergers or violent disk instabilities; \citealt{hopkins06, dekel09, krumholz10, dekel14, wellons15, bournaud16}), and following the quenching of star-formation, a later phase of assembly dominated by undergoing dry minor mergers with satellite galaxies \citep{nipoti03, khochfar06,naab09b, oser10, hilz12, hilz13}. Several observables serve to corroborate this scenario: the uniform, old stellar populations of $z\sim0$ massive galaxies \citep{mcdermid15}, the build-up of stellar haloes in (central) massive galaxies (e.g. \citealt{buitrago17, huang18a, huang18b}), and the dramatic size evolution observed for the massive, quiescent galaxy population since $z\sim2$ \citep{trujillo06, buitrago08, franx08, vandokkum08b, cimatti08, bezanson09, damjanov09, kriek09b, williams10, vandokkum10a, vanderwel11, newman12, szomoru12, whitaker12a, patel13, vanderwel14a, belli14a, belli15, hill17}. The structural evolution of galaxies is sensitive to their assembly history and feedback processes, as such, the observed size and morphology of galaxies in various environment and halo mass regimes is a critical benchmark for theoretical models to reproduce (e.g., \citealt{genel17, furlong17}). A census of galaxy size has now been obtained out to $z\sim4$ across a wide range in stellar mass and star formation activity (e.g., \citealt{shen03, trujillo04, bezanson09, patel13, vandokkum14, vanderwel14a, straatman15, allen17}). However, the majority of information on the size evolution of \textit{massive} galaxies is obtained from samples with stellar masses in the range of $1-2 \times 10^{11} M_{\odot}$; as such, the size-mass relation at the extreme massive end of the galaxy population (i.e., $\log(M_{*}/M_{\odot}$)$\geq11.25$) at $z>1.5$ remains poorly constrained. Abundance matching techniques suggest that ultra-massive galaxies, those with $\log(M_{*}/M_{\odot})>11.60$ should reside in dark matter haloes of a few $\times 10^{14} M_{\odot}$ at all redshifts, implying that they are the progenitors of the Brightest Cluster Galaxies (BCGs) in the local universe. Therefore, measuring how these massive systems evolve in size compared to their (relatively) lower-mass cousins could provide valuable information on how their assembly takes place, and whether this evolution is related to their halo properties (e.g., concentration, mass, or subhalo occupation number). %\\\\\\\\\\\\ Owing to the low spatial density of these objects, %(due to the exponential decline of the stellar mass function), identifying a statistically large sample of very massive galaxies requires relatively deep and wide NIR surveys using ground-based facilities, which typically lack the spatial resolution to derive robust sizes for these compact, distant galaxies (the typical FWHM $\sim~0.8-1^{''}$ corresponds to a physical size of $\sim~6-9$~kpc at $z=1.5-3$). To this end, we have obtained follow-up \textit{HST}/WFC3 $H_{160}$ imaging for a sample of very massive ($\log(M_{*}/M_{\odot})>11.25$) galaxies at $1.5<z<3.0$ selected using relatively deep and wide-field ground based NIR surveys. The $H_{160}$ band, the reddest filter currently available for high-resolution imaging, probes the rest-frame wavelength regime just blueward of the $r$ band ($\sim 6100$ {\AA}) at $z\sim1.5$ to wavelengths just redward of the rest-frame Balmer break at $z\sim3.0$ (i.e., $\sim~3900$~{\AA}) In this study, we present the \textit{HST}/WFC3 $H_{160}$ imaging for 37 targets with stellar masses $\log(M_{*}/M_{\odot})>11.2$ at $1.5<z<3.0$ in the NMBS-II, UltraVISTA and UKIDSS UDS. In Section~\ref{sec3-data} we briefly describe the datasets used to select this sample and the targeted \textit{HST} observations. Section~\ref{sec3-analysis} presents the analysis and relevant measurements employed in this study. We present the results in Section~\ref{sec3-results} and summarize these results in Section~\ref{sec3-summary}. Throughout this paper we assume the standard $\Lambda$CDM cosmological parameters $\Omega{_M}=0.3$, $\Omega{_\Lambda}=0.7$ with $H_{0}=70$~km~s$^{-1}$~Mpc$^{-1}$ and a \citet{chabrier03} initial mass function (IMF). All magnitudes listed are in the AB system. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | \label{sec3-summary} We presented the investigation of the structural properties of very massive galaxies ($\log(M_{*}/M_{\odot})>11.2$) at $1.5<z<3.0$. Owing to their low spatial density in the distant universe, identifying and assembling a large enough sample of very massive galaxies requires large survey volumes. We selected a sample of 37 galaxies from the combined UltraVISTA, NMBS-II and UDS catalogs to perform \textit{HST} WFC3/F160W follow-up imaging in order to accurately determine their sizes and morphologies. We modeled their 2D light profiles using GALFIT and compared their size distributions with the high-$z$ sample of \citet{vanderwel14a}. Visual investigation of the $H_{160}$ imaging revealed that 13/37 targets were unresolved in the parent $K$~band catalogs. We investigated the effect of galaxy blending on the SMF at $1.5<z<3.5$ by decomposing the estimated stellar masses of the close-pair systems based on their observed $H_{160}$ fluxes. Based on this analysis the results can be summarized as follows: \begin{itemize} \item{At $1.5<z<2.5$, the sizes of both star-forming and quiescent galaxies with $\log(M_{*}/M_{\odot})>11.2$ are relatively consistent with those found in \citet{vanderwel14a}.} \item{At $2.5<z<3$, sizes for quiescent galaxies at $\log(M_{*}/M_{\odot})>11.2$ appear to be systematically larger than what is expected based on the extrapolation of the relation derived from lower stellar mass galaxies, confirming results in \citet{patel17}. } \item{We found that the effect of galaxy blending is most significant for the largest stellar mass bin ($\log(M_{*}/M_{\odot})\approx11.6$) considered at $1.5<z<2.5$, although it remains consistent with the SMF of \citet[as calculated from the maximum likelihood method]{muzzin13a}. } \item{From the comparison with theoretical predictions, we find that the Illustris simulation agrees well with the observed number density, although their simulated volume is too small to probe the most massive galaxies. Similarly good agreement at $\log(M_{*}/M_{\odot}$)$<11.5$ is found between observations and the predictions from the SAM of \citet{henriques15}. However, the observed number density of the most massive galaxies (i.e., $\log(M_{*}/M_{\odot}$)$>11.5$) is over-predicted at $1.5<z<2.5$ and under-predicted at $2.5<z<3.0$.} \end{itemize} | 18 | 8 | 1808.03284 |
1808 | 1808.06343_arXiv.txt | Early-phase Type Ia supernovae (SNe Ia), especially those with luminosity enhancement within the first few days of explosions (``early-excess SNe Ia"), play an irreplaceable role in addressing the long-standing progenitor and explosion issue of SNe Ia. In this paper, we systematically investigate 11 early-excess SNe Ia from subluminous to luminous subclasses. Eight of them are selected from 23 SNe Ia with extremely early-phase optical light curves (``golden" early-phase SNe Ia), and three of them are selected from 40 SNe Ia (including 14 golden samples) with early-phase UV/NUV light curves. We found that previously discovered early-excess SNe Ia show a clear preference for specific SN Ia subclasses. In particular, the early-excess feature shown in all six luminous (91T- and 99aa-like) SNe Ia is in conflict with the viewing angle dependence predicted by the companion-ejecta interaction scenario. Instead, such a high early-excess fraction is likely related to the explosion physics of luminous SNe Ia; i.e. a more efficient detonation happening in the progenitor of luminous SNe Ia may consequently account for the early-excess feature powered by the radiation from a $^{56}$Ni-abundant outer layer. The diversity of early-excess features shown in different SN Ia subclasses suggests multiple origins of the discovered early-excess SNe Ia, challenging their applicability as a robust progenitor indicator. Further understanding of the early-excess diversity relies not only on multiband photometry and prompt-response spectroscopy of individual early-excess SNe Ia but also on investigations of the general early-phase light-curve behavior of each SN Ia subclass, which can be realized through ongoing/forthcoming transient survey projects in the near future. | \label{sec:intro} Type Ia supernovae (SNe Ia) have been thought to originate from the thermonuclear explosion of a carbon-oxygen white dwarf (WD) in a binary system. Even though great success was achieved by using them as a cosmic distance indicator in the 1990s \citep{perlmutter97,perlmutter99,riess98}, the progenitor and the physics leading to the explosion are still under debate \citep{hillebrandt13,maoz14,Maeda16}. In the last decades, a tremendous number of SNe Ia have been found through various kinds of transient surveys, and a growing number of SNe Ia discovered within a few days of their explosions provide unique information about the progenitor system and the explosion mechanism of SNe Ia \citep{nugent11,foley11,cao15,shappee16,marion16, hosseinzadeh17,JJA2017}. \citet{kasen10} proposed that a prominent brightening in the first few days of the explosion can be observed under specific viewing directions due to the interaction between the expanding ejecta and a nondegenerate companion star, which makes SNe Ia with additional luminosity enhancement in the early time (``early-excess SNe Ia"; EExSNe Ia) a powerful indicator for the single-degenerate (SD) progenitor system \citep{kasen10,maeda14,kutsuna15}. Since then, surveys for EExSNe Ia have become particularly popular in time-domain astronomy. The first reported EExSN Ia, iPTF14atg, has been suggested as strong evidence for the companion-ejecta interaction \citep{cao15}. Soon after, the companion-interaction scenario was also proposed as the likely origin of early-excess features of SN~2012cg and the recently discovered SN~2017cbv \citep{marion16,hosseinzadeh17}. However, whether the early light-curve excesses of these SNe Ia are exclusively attributed to the companion interaction is still under debate. Theoretically, the interaction between dense circumstellar matter (CSM) and SN ejecta (\citealp{shen12,levanon15}, \citeyear{levanon17}; \citealp{tanikawa15,piro16,maeda18}) and vigorous mixing of radioactive $^{56}$Ni in the outermost region of SN ejecta \citep{piro16} may produce a similar early-excess feature to that predicted by the companion interaction. For instance, \citet{kromer16} claimed that the spectra of iPTF14atg cannot be reproduced by the SD scenario but show good consistency with the prediction of merging two sub-Chandrasekhar-mass WDs. If the progenitor system of iPTF14atg is not an SD system, then the early light-curve excess of iPTF14atg cannot be attributed to the companion-interaction scenario, and other effects should come into play. Different physical mechanisms have also been proposed to explain the early excess shown in SN~2012cg and SN~2017cbv \citep{shappee18,hosseinzadeh17,miller18}. Other than these ``companion-interaction-like" EExSNe Ia, the early light-curve excesses of several SNe Ia have been initially explained by different scenarios (e.g. iPTF14bdn, \citealt{smitka15}; iPTF16abc, \citealt{miller18}). In particular, \citet{JJA2017} reported a normal-brightness EExSN Ia, MUSSES1604D (SN~2016jhr), that shows a prominent but red early excess and strong Ti \textsc{ii} absorptions around the $B$-band maximum. The early excess of this peculiar object can be well explained by the radiation emitted from short-lived radioactive elements that were generated by a precursory detonation at a thin helium shell of its primary WD, which is the first robust evidence of the multiple origins of EExSNe Ia. As opposed to the EEx discoveries in several SN Ia subclasses, there is no clear evidence of EEx detection in normal SNe Ia so far. For example, a smooth-rising light curve of SN~2011fe from the very beginning (with brightness $\sim$1000 times fainter than its peak) provides a significant constraint on the companion type \citep{nugent11,li11}. The nondetection of EEx in normal SNe Ia can be explained by either an intrinsically faint EEx due to a small nondegenerate companion star or an unfavorable viewing angle by chance under the companion-interaction scenario. However, such an observational fact is in contrast to the high EEx fraction\footnote{For SNe Ia in each subclass, an EEx fraction is defined as the number fraction of EExSNe Ia over EExSNe Ia and well-observed early-phase SNe Ia that do not show EEx features (non-EExSNe Ia; selected from papers and public resources) in the same subclass.} discovered in 91T/99aa-like SNe Ia, which may suggest differences in their natures between some SN Ia subclasses (see \S4.2). With the largest sample of early-phase SNe Ia discovered until 2018, we present new evidence to further prove multiple origins of EExSNe Ia and discuss the explosion mechanism and progenitor of specific SN Ia subclasses from the EEx perspective. This paper is organized as follows. General information on 23 well-observed early-phase SNe Ia (``golden early-phase SNe Ia") from the subluminous to luminous subclass are summarized in \S2, and further investigations of previously reported and unnoticed EExSNe Ia are shown in \S3. Discussions about the multiple origins of EEx and their associated subclasses are given in \S4, and our conclusions are summarized in \S5. | \label{sec:Conclusions} In this paper, we present general information on published, well-observed early-phase SNe Ia so far and summarize the characteristics of 11 (six reported and five unnoticed) EExSNe Ia in different subclasses. In particular, by investigating the connections and differences between these EExSNe Ia, new evidence of multiple origins of early light-curve excess and the implication for the explosion mechanism of 91T/99aa-like SNe Ia are presented. We found that a 100\% EEx detection in six early-phase 91T/99aa-like SNe Ia is significantly in conflict with the prediction of the companion-ejecta interaction scenario but can be promisingly explained by the radioactive decay of a $^{56}$Ni-abundant outer layer or, alternatively (but less likely), interacting with spherically distributed CSM. In addition, the discovered correlation between the Si \textsc{ii} absorption feature and the strength of EEx also suggests an intrinsic connection between the explosion mechanism and the early light-curve excess of 91T/99aa-like SNe Ia. By investigating the early-excess behavior, post-EEx photometric/spectroscopic properties, and explosion models proposed for luminous SNe Ia, the surface-$^{56}$Ni-decay scenario is preferred for interpreting 91T/99aa-like EExSNe Ia. Specifically, we argue that the gravitationally confined detonation is a promising scenario for producing 91T/99aa-like SNe Ia with such a high EEx fraction. Also, spectral and early-excess differences between 91T- and 99aa-like SNe Ia and their intrinsic luminosity scatters could be qualitatively explained by taking into account the variation of the off-center distance of the initial ignition point and the viewing angle effect. In contrast to the high EEx fraction of luminous SNe Ia, early light-curve excess discovered in a normal SN Ia, SN~2017erp, and a possible candidate, SN~2015ak, so far suggests that EEx may accompany with a fraction of normal SNe Ia. Whether the EEx shown in normal SNe Ia is attributed to surface $^{56}$Ni decay through the traditional delayed-detonation mechanism requires further investigation. Even though about a dozen of EExSNe Ia were successfully discovered in various SN Ia subclasses, we have not found any crucial evidence to support the companion-interaction scenario, even for the most promising candidate, iPTF14atg. The multiple origins of early light-curve excess suggest that EExSNe Ia may not be a superior indicator of the SD progenitor system as we originally expected, and we need to be more cautious when interpreting any newly discovered EExSNe Ia. Further understanding of the early-excess scenarios relies not only on individual studies of well-observed EExSNe Ia but also on systematical investigations of the early-phase light-curve behavior of each SN Ia subclass, which can be realized with ongoing survey projects such as the Zwicky Transient Facility (ZTF; \citealp{smith14}), the MUlti-band Subaru Survey for Early-phase SNe Ia (MUSSES; \citealt{JJA2017,miyazaki18}), and forthcoming transient surveys with the Tomo-e Gozen Camera mounted on the 1.05-m Kiso Schmidt telescope \citep{sako18} and the Large Synoptic Survey Telescope (LSST; \citealp{ivezic08}) in the near future. | 18 | 8 | 1808.06343 |
1808 | 1808.04493_arXiv.txt | The Simons Observatory (SO) will make precise temperature and polarization measurements of the cosmic microwave background (CMB) using a set of telescopes which will cover angular scales between 1 arcminute and tens of degrees, contain over 60,000 detectors, and observe at frequencies between 27 and 270 GHz. SO will consist of a 6\,m aperture telescope coupled to over 30,000 transition-edge sensor bolometers along with three 42\,cm aperture refractive telescopes, coupled to an additional 30,000+ detectors, all of which will be located in the Atacama Desert at an altitude of 5190 m. The powerful combination of large and small apertures in a CMB observatory will allow us to sample a wide range of angular scales over a common survey area. SO will measure fundamental cosmological parameters of our universe, constrain primordial fluctuations, find high redshift clusters via the Sunyaev-Zel’dovich effect, constrain properties of neutrinos, and trace the density and velocity of the matter in the universe over cosmic time. The complex set of technical and science requirements for this experiment has led to innovative instrumentation solutions which we will discuss. The large aperture telescope will couple to a cryogenic receiver that is 2.4\,m in diameter and nearly 3\,m long, creating a number of technical challenges. Concurrently, we are designing the array of cryogenic receivers housing the 42\,cm aperture telescopes. We will discuss the sensor technology SO will use and we will give an overview of the drivers for and designs of the SO telescopes and receivers, with their cold optical components and detector arrays. | \label{sec:intro} % The cosmic microwave background (CMB) has emerged as one of the most powerful probes of the early universe. Measurements of temperature anisotropies on the level of ten parts per million have brought cosmology into a precision era, and have placed tight constraints on the fundamental properties of the Universe\cite{Planck2018}. Beyond temperature anisotropies, CMB polarization anisotropies not only enrich our understanding of our cosmological model, but could potentially provide clues to the very beginning of the universe via the detection (or non-detection) of primordial gravitational waves. To provide a complete picture of cosmology, measurements at multiple frequencies of both large and small angular scales are important. % This is the goal of the Simons Observatory (SO). SO will field a 6\,m diameter crossed Dragone\cite{Dragone1978} large aperture telescope (LAT) coupled to the large aperture telescope receiver (LATR) (Sec. \ref{sec:lat}). The LAT is designed to have a large field of view (FOV)\cite{Niemack2016,Parshley2018} (7.2$^\circ$ at 90 GHz) capable of supporting a cryostat with up to 19 tubes, each containing three lenses and a focal plane of detector arrays (henceforth optics tubes). To reduce the development risk of such a large cryostat, the LATR is designed with capacity for 13 optics tubes. During the initial deployment, we plan to populate seven optics tubes with three detector wafers in each for a total of over 30,000 detectors. Note that each optics tube could be upgraded to support four wafers for a $\sim$33\% increase in the number of detectors per optics tube. With this upgrade and the deployment of 19 optics tubes, the LAT focal plane could support roughly 125,000 detectors at 90/150\,GHz. SO will also deploy an array of 42 cm aperture small aperture telescopes (SATs) with seven detector wafers in each coupled to over 30,000 detectors (Sec. \ref{sec:sat}). SO will observe with three frequency pairs in order to observe the CMB peak signal and constrain polarized foreground contamination from galactic synchotron and dust emission at lower and higher frequencies. The SO frequency bands are: 27/39 GHz, low frequency(LF); 90/150 GHz, mid-frequency(MF); and 220/270 GHz, ultra-high frequency(UHF). The LATR will initially be populated with four MF optics tubes, two UHF optics tubes, and one LF optics tube, while the SAT array will initially be composed of two MF SATs and one UHF SAT with a fourth LF SAT to follow. Details of the SO observing strategy with the two classes of telescope can be found in Stevens et al. 2018\cite{Stevens2018}. A description of the sensitivity calculator used to optimize the SO instrument design appears in Hill et al. 2018\cite{Hill2018}. Details of the calibration strategy planned for SO can be found in Bryan et al. 2018\cite{Bryan2018}. Additional details of the forecasting studies which led to the SO instrument configuration as well as a more in depth description of the observatory will be in a pair of papers in preparation. In this paper we will introduce the basic design of the detector systems, telescopes, and receivers, many of which are discussed in more detail in a set of complementary papers referenced herein. Sec. \ref{sec:sat} describes the SAT in additional detail as this paper serves as the most comprehensive reference for the SAT instrument description. | 18 | 8 | 1808.04493 |
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1808 | 1808.03058_arXiv.txt | We mapped \twCO~\Jotz, \twCO~\Jtto, \thCO~\Jotz, and \thCO~\Jtto\ lines toward supernova remnant (SNR) Cassiopeia~A with the IRAM~30m telescope. The molecular clouds (MCs) along the line of sight of \snr\ do not show optically thin, shock-broadened \twCO\ lines ($\Delta V\le 7\km\ps$ toward \snr), or high-temperature features from shock heating ($\Tk\le 22~\K$ toward \snr). Therefore, we suggest that there is no physical evidence to support that the SNR is impacting the molecular gas. All the detected MCs are likely in front of \snr, as implied by the \HCOp\ absorption line detected in the same velocity ranges. These MCs contribute H$_2$ column densities of $5\E{21}~\cm^{-2}$, $5\E{21}~\cm^{-2}$, and $2\E{21}~\cm^{-2}$ in the west, south, and center of the SNR, respectively. The 20~K warm gas at $\VLSR\sim -47~\km\ps$ is distributed along a large-scale molecular ridge in the south of Cas~A. Part of the gas is projected onto Cas~A, providing a foreground H$_2$ mass of $\sim 200 (d/3~{\rm kpc})^2~\Msun$, consistent with the mass of cold dust (15--20~K; 2--4~$\Msun$) found in front of the SNR. We suggest that the 20~K warm gas is heated by background cosmic rays with an ionization rate of $\zeta({\rm H_2})\sim 2\E{-16} \ps$. The cosmic rays and X-ray emission from \snr\ are excluded as the heating sources of the clouds. | \label{sec:intro} Cassiopeia~A, a.k.a., Cas~A (G$111.7-2.1$), is the remnant of a young supernova that exploded around AD $1681\pm 19$ \citep{fesen06} at a distance of $3.4^{+0.3}_{-0.1}$~kpc \citep{reed95}. Due to its youth and brightness, \snr\ is among the best-observed supernova remnants(SNRs) and the prototype for many aspects of SNR studies. The current manifestations of an SNR are the combined results of the intrinsic supernova properties and the environment it evolves in. The study of the environment not only provides crucial information of the SNR itself, but also helps us to understand how supernova explosions affect the interstellar medium of the galaxies. In spite of a wealth of observations of \snr\ across the electromagnetic wavelengths, there is no consensus on the immediate SNR environment. It is a common view that \snr\ evolves in the stellar winds of the progenitor star. The early suggestion of the idea was based on the detections of the quasi-stationary flocculi \citep{vandenbergh73}. The circumstellar medium of \snr\ was likely produced by slow and dense winds, characterized by a $\rho_{\rm w}\propto r^{-2}$ density profile \citep[e.g.,][]{vink04a}. This picture well explains the positions of the SNR's forward and reverse shocks \citep{chevalier03}, the X-ray properties of shocked ejecta knots \citep{laming03,hwang09}, and the characteristics of the shocked ambient gas \citep{lee14}. The preshock density estimated from the X-ray analysis is around $0.9\pm0.3~\cm^{-3}$, without strong variation at different position angles \citep{lee14}. Outside the SNR boundary the environmental gas is cooler and its relation with \snr\ less clear. The difficulty is how to distinguish the immediate environment and the interstellar medium along the line of sight. The dense gas along the line of sight of \snr\ is mainly distributed in the velocity range of $-50$ to $-30\km\ps$ (associated with the Perseus arm), and at $\sim 0~\km\ps$ (associated the Orion spur), as observed in molecular emission and absorption lines \citep{dejager78,batrla83,batrla84,goss84,bieging86,wilson93,anantharamaiah94,reynoso02,kilpatrick14}, \ion{H}{1} absorption \citep{mebold75,bieging91,schwarz97}, [\ion{C}{1}] \citep{mookerjea06}, and carbon recombination lines \citep{payne89,anantharamaiah94,kantharia98, mookerjea06,oonk17, salas17,salas18}. \citet{reynoso97b} found some small absorption \ion{H}{1} features at the lower LSR velocity of $-69$ to $-62\km\ps$, and suggested that they are from neutral knots driven away by the progenitor wind of \snr. If this velocity of foreground \ion{H}{1} structures represents the upper limit of \snr's systemic velocity, it may indicate that \snr\ is not associated with any of the clouds in the velocity range of $-50$ to $-30~\km\ps$. This idea is supported by the carbon recombination line studies, which suggest that the gas at $\VLSR\lesssim -47\km\ps$ is at least 100~pc in front of \snr\ \citep{kantharia98, salas17}. However, since the molecular gas is traced by molecular lines, the relationship between the molecular gas and \snr\ should be studied with molecular observations. Warm molecular gas with a kinetic temperature of 20~K was found at $\VLSR\sim -47\km\ps$ \citep{wilson93}, using \thCO\ and \twCO\ observations with IRAM 30~m, while the typical temperature of an interstellar giant molecular cloud (MC) is $\sim 10~\K$. The association between \snr\ and molecular clouds (MCs) was proposed in a recent study using the SMT 12~m telescope \citep{kilpatrick14}, based on a subtle line broadening (line width of $\sim 6\km\ps$) of the \twCO\ emission at $\VLSR\sim -37\km\ps$ and $\sim -47 \km\ps$ in the west and south of the SNR. If true, the MCs could provide dense targets for the shock and cosmic-ray (CR) protons to interact, making the properties of the MCs important constraints for the studies of the SNR's shock and CRs. However, an association between \snr\ and MC is inconsistent with the low preshock density found in X-rays \citep{lee14} and the suggestion that the clouds at $\VLSR=-50$--$0\km\ps$ are all foreground gas. Motivated by the aforementioned problems, we performed new molecular line observations toward \snr\ and its environment, aiming to answer the following questions: Is the SNR interacting with MCs? what are the properties of the MCs? We also examine what is the heating source of the 20~K MC at $-47\km\ps$. Our molecular mapping observations have the best angular resolution of these lines to date for this SNR ($11''$ at \twCO\ \Jtto), and good sky coverage (83 arcmin$^2$ at \twCO\ \Jtto). | We have performed mapping observations of \twCO~\Jotz, \twCO~\Jtto, \thCO~\Jotz, and \thCO~\Jtto\ lines and \HCOp~\Jotz\ observation toward the SNR \snr\ with the IRAM~30 m telescope. Our main conclusions are summarized as follows. \begin{enumerate} \item \twCO, \thCO, and \HCOp\ emission is detected in the velocity range of $-50$ to $-30\km\ps$, and at $\sim -2\km\ps$. The MCs at $\VLSR=-47\km\ps$ have a kinetic temperature of 20~K, narrow line widths of 1--$3 \km\ps$, column densities up to $6\E{21}~\cm^{-2}$, and density of the order of $10^3\cm^{-3}$. Some clouds are distributed in a wide velocity interval of $-44$ to $-30\km\ps$, with a mean temperature of $13~\K$, column density $<5\E{21}~\cm^{-2}$, and density of a few $\times 10^{2}~\cm^{-3}$. \item The MCs at $\VLSR=-50$ to $-30\km\ps$ are clumpy. The identified \twCO\ \Jtto\ clumps have main-beam temperatures of 2.2 -- 16.7~K, velocity dispersions of 0.3--2.5~\kms, and sizes from subparsec to 3 pc. \item The MCs toward \snr\ do not show any of the properties that are typical for shocked MCs: (1) optically thin broad line with FWHM much larger than that of the environmental gas, (2) broadened CO lines along with \twCO\ \Jtto\ to \Jotz ($R_{21/10})>1$, and (3) MCs in post-shock regions much hotter than in the preshock regions. Therefore, we suggest that there is no physical evidence to support that the SNR is impacting molecular gas. \item We detect an absorption line of \HCOp\ \Jotz\ at $\VLSR=-50$ to $-33\km\ps$ near the radio peak west of the SNR, suggesting that all the detected molecular gas is foreground gas. \item The foreground MCs result in a high absorption of the SNR emission at the west, south, and center of the SNR. The 20~K warm gas contributes a foreground mass of $\sim 200 d_3^2~\Msun$ for \snr, which can explain the cold dust (15--20~K; 2--4~$\Msun$) found in front of \snr. The cooler MCs in the velocity range $-44$ to $-30\km\ps$ contribute a foreground H$_2$ mass of $\sim 280d_3^2 \Msun$. \item The gas at $\VLSR\sim -47\km\ps$ has a kinetic temperature of 20~K, warmer than MCs at other velocities. The warm gas is extended over 10~pc (projected distance) south of the SNR. CRs are likely the heating source, and the required CR ionization rate is $\zeta({\rm H_2})\sim 2\E{-16}\ps$. The CRs are provided by the local Galactic CR background, but not by \snr, since the high-energy CRs from \snr\ do not provide enough energy to heat the gas, and its low-energy CRs cannot diffuse far enough to the MCs. The X-ray emission of \snr\ is insufficient to heat the MCs. \end {enumerate} | 18 | 8 | 1808.03058 |
1808 | 1808.03572_arXiv.txt | { We monitored the Seyfert-1 galaxy \C between September 2014 and March 2015 at the \USB near Cerro Armazones in $BVRIJK$ and a narrowband filter covering the redshifted \Ha line. In addition we obtained a single contemporary spectrum with the spectrograph FAST at Mt. Hopkins. Compared to earlier epochs \C is about a factor of three brighter, allowing us to study the shape of the broad line region (BLR) and the dust torus in a high luminosity phase. { The analysis of the light curves yields that the dust echo is rather sharp and symmetric in contrast to the more complex broad H$\alpha$ BLR echo. We investigated how far this supports an optically thick bowl-shaped BLR and dust torus geometry.% The comparison with several parameterizations of these models supports the following geometry: The BLR clouds lie inside the bowl closely above the bowl rim up to a halfcovering angle $0^\circ < \theta < 40^\circ$ (measured against the equatorial plane). Then the BLR is spread over many isodelay surfaces, yielding a smeared and structured echo as observed. Furthermore, if the BLR clouds shield the bottom of the bowl rim against radiation from the nucleus, the hot dust emission comes essentially from the top edge of the bowl ($40^\circ < \theta < 45^\circ$). Then, for small inclinations as for 3C120, the top dust edge forms a ring that largely coincides with a narrow range of isodelay surfaces, yielding the observed sharp dust echo. The scale height of the BLR increases with radial distance from the black hole (BH). This leads to luminosity dependent foreshortening effects of the lag. We discuss the implications and possible corrections of the foreshortening for the BH mass determination and consequences for the lag (size) -- luminosity relationships and the difference from interferometric torus sizes. } } | \label{sec:introduction} The quasar paradigm comprises a supermassive BH, a central X-ray source, an accretion disk (AD) surrounded by a broad line region (BLR), and a molecular dusty torus farther out. These formal components serve as working frame: the AD, BLR, and torus may have smooth transitions rather than being separated entities with sharp boundaries. As the inner quasar regions cannot be resolved by conventional imaging techniques, reverberation mapping (RM) is the main tool of the trade (\citealt{1972ApJ...171..467B,1993PASP..105..247P,2004PASP..116..465H}). The RM technique traces the response of irradiated regions to the light fluctuations of continuum emission from the inner AD. In particular, RM studies of the BLR trace the response of line emission to continuum variations and determine the time lag, $\tau$, between their signals, with the BLR size, \rblr$\sim c \tau$ ($c$ is the speed of light). Reverberation mapping can reach a spatial resolution of better than $10^{-5}$\,pc, clearly surpassing the capabilities of current imaging instruments. Line-to-continuum RM studies find that $R_{\rm BLR}$ ranges from a few light days to several light months. At low redshift, a relatively tight size-luminosity relation for the H$\beta$-emitting region has been identified: \rblr$({\rm H}\beta)\propto L^\alpha$ with $\alpha \simeq 0.5$ and monochromatic active galactic nucleus (AGN) luminosity $L$ at 5100\AA~ (\citealt{2000ApJ...533..631K,2013ApJ...767..149B}). Conversely, $L$ may be inferred from \rblr or \rdust measurements, opening the possibility of using AGNs as standard candles for cosmological distance studies (\citealt{1999OAP....12...99O,1999MNRAS.302L..24C,2011ApJ...740L..49W, 2011A&A...535A..73H,2013A&A...556A..97C,2014ApJ...784L..11Y,2017MNRAS.464.1693H}). To characterize the line profiles, various studies using the ratio full width at half maximum (FWHM) to line dispersion $\sigma_{v}$ revealed a large diversity between sources with flat-topped lines (showing large FWHM/$\sigma_{v}$) and sources with narrow-peaked lines with extended wings (showing small FWHM/$\sigma_{v}$); the latter class shows a trend toward higher accretion rates as measured by Eddington ratios, while inclination appears to play a minor role in FWHM/$\sigma_{v}$ (\citealt{2000ApJ...536L...5S, 2004A&A...426..797C,2006A&A...456...75C}). Alternatively, BLR properties could be luminosity dependent, such that the use of $\sigma_{v}$ as a proxy for $v_{\rm rot}$ leads to an overestimate of $M_{\rm BH}$ at high$-z$ or high luminosity. From RM studies of individual Seyferts during phases of low and high luminosity \cite{2011Natur.470..366K,2013A&A...549A.100K,2013A&A...558A..26K} find an indication that both the scale height above the equatorial plane and the ratio of turbulent to rotation velocity $v_{\rm turb} / v_{\rm rot}$ of the BLR increase with luminosity. A potential explanation is that brightening of an individual AGN makes the BLR visible at larger $R_{\rm BLR}$, where $v_{\rm rot}$ should be smaller (Kepler's law); if $v_{\rm turb}$ remains constant then the ratio $v_{\rm turb} / v_{\rm rot}$ increases with $R_{\rm BLR}$. If the height-to-radius ratio $H/R $ is proportional to $ v_{\rm turb} / v_{\rm rot}$, this implies a growing vertical extent of the BLR above the disk plane. It is clear that geometry has a substantial influence on the calculation of the central BH mass. Well-sampled light curves allow us to restore information about the spatially unresolved BLR geometry (\citealt{1991ApJ...367L...5H,1991ApJ...367..493M,2012ApJ...754...49P}). In a nutshell, tightly localized gas distributions (e.g., nearly face-on flat disks) produce a sharp echo of the continuum light curve, while more isotropic gas configurations (e.g., spheres) lead to smoothed light curves for the line emission band. The line-integrated transfer function depends on the BLR geometry, i.e., the location of the BLR clouds at the time they respond; BLR kinematic information is not necessary. It is intriguing to check if the geometry of the visible BLR changes with luminosity (or accretion rate \mdot) in a nonisomorphic manner. For the archetypal Sy-1 galaxy NGC\,5548, \cite{2007arXiv0711.1025G} analyzed the AGN energy budget and derived crucial conclusions concerning the geometry of the BLR and dust torus: the BLR has a likely covering factor about 40\%, which translates to a half-covering angle $\theta \approx 40^{\circ}$. The BLR shields a substantial fraction of the dust torus from direct illumination by the AD, allowing for the observed relatively small near-infrared (NIR) contribution to the AGN energy budget. At the same time, the BLR obscuration also removes the problem that the dust torus covering factor is greater than the BLR covering factor, and is consistent with the observed fraction of obscured AGNs. The flux reduction at the torus also reduces the problem of AGN dust reverberation lags giving sizes smaller than the dust sublimation radii. Near-infrared RM studies of the dusty torus find $\tau_{\rm dust} \approx 4 \times \tau_{H\beta}$ between hot-dust continuum emission and optical continuum fluctuations (\citealt{2014ApJ...788..159K}). However, $R_{\rm dust} = c \cdot \tau_{\rm dust}$ is three times smaller than the dust sublimation radius, $R_{sub}$, inferred from the UV luminosity (\citealt{2006ApJ...639...46S,2007A&A...476..713K}). {To resolve this conflict, \cite{2010ApJ...724L.183K, 2011ApJ...737..105K} proposed a bowl-shaped dust torus that smoothly continues into the central AD. The AD emission is highest at the polar region and lowest at the equatorial region. The anisotropy of the AD emission controls the angle dependent dust sublimation radius and thus the concave rim of the bowl. For the bowl-shaped torus, assuming that the hot dust emission arises from the entire bowl surface, the dust transfer functions show about three times faster responses than for a bagel-shaped torus (see, e.g., Fig.~1 of \citealt{1995PASP..107..803U}), in agreement with observations. Notably, for small inclinations (i.e., rather face-on than edge-on view) the transfer functions of the dust emission show a sharp peak. Kawaguchi \& Mori's studies pioneered the bowl-shaped dust torus, but did not yet consider the BLR. The relation between BLR and dust torus has long been debated. For instance, \cite{2011A&A...525L...8C} proposed that the AD reaches far out until it meets the torus and that the BLR clouds are launched from the outer AD in a dusty wind at that radius, where the temperature in the AD matches the dust sublimation temperature. In this scenario, however, how far the reverberation-based BLR and dust lags should essentially be the same is an open issue contrary to the factor 3-5 larger dust lags observed so far. Consequently, in continuation of Kawaguchi \& Mori's work, \cite{2012MNRAS.426.3086G} modeled BLRs, assuming that they are confined by a bowl-shaped torus geometry, and suggested that the BLR clouds lie above the bowl surface (see their Fig.~1). Goad et al. successfully tested the bowl-confined BLR models with reverberation data of NGC\,5548. Notably, the BLR clouds may shield part of the dust from the AD radiation, bringing the hot dust covering fraction closer to that deduced by \cite{2011MNRAS.414..218L} and \cite{2011ApJ...737L..36M}. Then the hot dust emission would arise essentially from the edge, i.e., top rim, of the bowl. For the Seyfert-1 WPVS48, \cite{2014A&A...561L...8P} found an exceptionally sharp NIR echo, which led these authors to favor the theory that the hot dust is essentially located at the edge rather than the entire rim. Simultaneous BLR and dust reverberation studies might be able to shed further light on the geometry and the interplay between BLR and dust. Therefore, we have embarked on such a study of 3C\,120. } The BLR galaxy \C belongs to the most frequently monitored Seyfert-1 galaxies and results from several campaigns have been reported:\ spectroscopic RM by \cite{1998PASP..110..660P} covering the H$\beta$ region, velocity resolved RM covering H$\gamma$, H$\beta,$ and H$\alpha$ by \cite{2012ApJ...755...60G,2013ApJ...764...47G} and \cite{2014A&A...566A.106K}. Photometric reverberation mapping {(PRM)} of H$\beta$ has been reported by \cite{2012A&A...545A..84P,2014A&A...568A..36P}. During the spectroscopic campaigns before 1998 and in 2008 and 2009 (Kollatschny) the average $\lambda log(L_{\lambda 5100})$ AGN luminosity at 5100\AA\ was about ${44.01}$ and ${44.12}$. In 2009 and 2010 (Pozo Nu\~nez) this dropped to ${43.84}$\,erg\,s$^{-1}$, similar as in 2010 and 2011 (Grier, ${43.87}$\,erg\,s$^{-1}$). Between September 2014 and March 2015 we continued photometric monitoring; \cite{2015A&A...581A..93R} found a brightening by a factor three (${44.32}$\,erg\,s$^{-1}$) from the $B$ and $V$ band light curves. The brightening occurred between January and August 2014, as revealed by a spectrum in December 2013 taken at the Asiago Observatory where $\lambda log(L_{\lambda 5100}) \approx 43.82$\,erg\,s$^{-1}$ (PhD thesis, private communication, data in \citealt{2015A&A...578A..28B}). We report on the analysis of the entire data set in $BVRIJK$ and a narrow band (NB) at 680\,nm and contemporaneous spectrum taken in November 2014. We assumed a luminosity distance of 138\,Mpc for a cosmological model with $\Omega_m = 0.27$, $\Omega_v = 0.73$ and $H_0 = 73$\,km\,s$^{-1}$Mpc$^{-1}$. | \begin{enumerate} \item During our epoch, \C has passed a high brightness state with $log(\lambda L_{\lambda 5100}) = {44.32}$~erg\,s$^{-1}$, which is three times brighter than in 2009/2010. \item The CCF between $B$ and \Ha yields a mean rest-frame delay of $\tau_{\mathrm{BLR}}=68.9^{+12.4}_{-13.3}\,$ld. The CCF is broader than indicated by the rather small uncertainties; it shows a broad plateau between 55 and 95 days and a clear substructure peak at 60 days. This indicates that the BLR is spread over a large range of isodelay surfaces. \item The NIR $J$ and $K$ band light curves allow us to measure a dust emission delay ($\tau_{K}= 94.4 ^{+4.1}_{-6.9}\,$ld), which is unusually short compared to the BLR delay. Likewise the CCF between $B$ and the NIR bands is surprisingly sharper than between $B$ and \Ha. This indicates that the hot dust emission arises from a confined volume, which is covered by a narrow range of isodelay surfaces. \item Our data reach good consistency with the model by \cite{2012MNRAS.426.3086G}, where the BLR clouds shield the bowl-shaped dust rim up to a half-covering angle of $\theta \approx 40^\circ$ from the nuclear radiation. Such a BLR reproduces the observed structured broad $B$ -- \Ha CCF with a pronounced low-$\tau$ peak predicted from the models. In parallel, the hot dust emission arises from the small part of the rim at $40^\circ < \theta < 45^\circ$ facing the observer, reproducing the sharp and symmetric $B$ -- NIR CCF; because this hot dust edge lies closer to the observer than the equatorial plane this also explains the short lag by foreshortening of the light travel time. \item In a bowl-shaped geometry the calculation of the viral black hole mass M$_{\rm vir}$ may be affected by the foreshortening of the observed BLR lag, which leads to an underestimate of the BLR distance from the BH and of M$_{\rm vir}$. The foreshortening effect increases with increasing distance of the BLR from the BH. For the 3C\,120 data, a tentative correction of the foreshortening effect by spectroscopic or geometric $H/R_x$ estimates yields consistent M$_{\rm vir}$ values for different Balmer lines (\Ha, \Hb) and different AGN luminosities. \item If AGN in general exhibit a bowl-shaped BLR-dust torus geometry, the foreshortening effects naturally contribute to the scatter in the relationship between AGN luminosity and lag-inferred BLR size. The foreshortening effects may also explain why reverberation-based torus sizes are on average about a factor two smaller compared with interferometric torus sizes. Hysteresis effects may also contribute to the scatter in the relationship between AGN luminosity and dust lag. \end{enumerate} Combined with spectral and infrared data, our analysis indicates a vertically extended BLR region that can be explained as an envelope of a thick dust torus with a bowl-shaped inner border. Further light curves and emission line spectra for a sample of objects are required to corroborate this approach. The monitoring of the BLR emission lines should ideally cover distinct luminosity epochs of a single source, which will reveal how the vertical structure of a particular emission line grows with luminosity. This is especially interesting for the strong and thus easily measurable \Ha emission line that likely has the largest variation in scale height.} | 18 | 8 | 1808.03572 |
1808 | 1808.06027_arXiv.txt | We present chemical abundance determinations of two \ion{H}{II} regions in the dIrr galaxy Leo A, from GTC OSIRIS long-slit spectra. Both \ion{H}{II} regions are of low excitation and seem to be ionised by stars later than O8V spectral type. In one of the \ion{H}{II} regions we used the direct method: O$^{+2}$ ionic abundance was calculated using an electronic temperature determined from the [\ion{O}{III}] $\lambda\lambda$4363/5007 line ratio; ionic abundances of O$^+$, N$^+$, and S$^+$ were calculated using a temperature derived from a parameterised formula. O, N and S total abundances were calculated using Ionisation Correction Factors from the literature for each element. Chemical abundances using strong-line methods were also determined, with similar results. For the second \ion{H}{II} region, no electron temperature was determined thus the direct method cannot be used. We computed photoionisation structure models for both \ion{H}{II} regions in order to determine their chemical composition from the best-fitted models. It is confirmed that Leo A in a very low metallicity galaxy, with 12+log(O/H)=7.4$\pm$0.2, log(N/O)=$-$1.6, and log(S/O)=$-$1.1. Emission lines of the only PN detected in Leo A were reanalysed and a photoionisation model was computed. This PN shows 12+log(O/H) very similar to the ones of the \ion{H}{II} regions and a low N abundance, although its log(N/O) ratio is much larger than the values of the \ion{H}{ii} regions. Its central star seems to have had an initial mass lower than 2 M$_\odot$. | In the nearby Universe, the most common galaxies are the dwarf ellipticals and irregulars. The analysis of these small galaxies allow us to study the history of galaxy formation and the amount and distribution of dark matter, as they seem to be dominated by it. Observations and analysis of the chemistry in the interstellar medium (ISM) in these galaxies may help us to understand their star formation and evolution \citep[and references therein]{Hernandez-MartinezPena2009}. The study of \ion{H}{ii} regions and planetary nebulae (PNe) in irregular galaxies provides important information on the ISM metallicity. By the analysis of the chemical abundances in both types of nebulae we can infer the evolution of the interstellar metallicity and hence the chemical evolution of the galaxy. This can be done owing to the chemical abundances obtained from PNe (produced by low and intermediate mass stars with initial mass M $<$ 8 M$_\odot$) represent the abundance of the progenitor cloud through the elements not modified by stellar nucleosynthesis, whereas the chemical abundance of \ion{H}{II} regions (associated to massive stars with M $>$ 8 M$_\odot$) represent the present-day chemical content of the ISM. Leo A, with coordinates $\alpha$(J2000.0) = 09:59:24.8, $\delta$(J2000.0) = +30:44:57, is an isolated dIrr galaxy in the Local Group, located at about 800 kpc from the Milky Way and at 1200 kpc from M31, and it presents a radial velocity of 22.3 km s$^{-1}$ \citep{DolphinSaha2002,McConnachie2012}. Leo A is a low mass galaxy with M$_{dyn}$ $\leq$ 2.5$\times$10$^7$ M$_\odot$, M$_{stars}$= 6$\times$10$^6$ M$_{\odot}$ and M$_{HI}$=1.1$\times $10$^7$ M$_{\odot}$ \citep[and references therein]{McConnachie2012} and shows recent star formation. Its stellar population contains old and young population components, although it conspicuously lacks a prominent ancient population, older than 10 Gyr \citep{TolstoyGallagher1998}. \citet{ColeSkillman2007} derived a star formation history (SFH) for Leo A founding that its first 5 Gyr were of quiescence, most of the star formation happened more recently than 6 Gyr ago and the fraction of ancient stars was conclusively determined to be less than 10\%, although there is a small number of RR Lyr \citep{DolphinSaha2002,BernardMonelli2013}. Some \ion{H}{II} regions were reported by \citet{StrobelHodge1991} and one PN was discovered by \citet{SkillmanKennicutt1989} from imaging later published by \citet{StrobelHodge1991}. From spectroscopic data obtained with the 5-m Palomar telescope \citet{vanZeeSkillman2006} derived chemical abundances of four diffuse \ion{H}{II} regions and the PN. In spite of their observational effort (2 hrs of exposure time for the deeper spectra) these authors could not detect the auroral lines [\ion{O}{III}] $\lambda$4363 nor [\ion{N}{ii}] $\lambda$5755, in the \ion{H}{II} regions; these lines are needed to compute the electron temperature of the plasma and, hence, they only obtained some abundance ratios by using semi-empirical methods. They derived an average oxygen abundance 12+log(O/H)=7.38$\pm$0.10 for the \ion{H}{II} regions. For the PN, they measured the [\ion{O}{III}] electron temperature and obtained 12 + log(O/H) = 7.30$\pm$0.05, by just adding the O$^+$/H$^+$ and the O$^{+2}$/H$^+$ ionic abundances. Owing to its low metallicity, Leo A is an ideal laboratory to study the nature of the metal poor ISM. The analysis of chemical abundances in \ion{H}{II} regions and PNe in metal poor environments can give hints about stellar evolution (whether stellar 3rd dredge-up episode is present or not), if dIrr galaxies are chemically homogeneous, even it is possible to explore the primordial helium abundance. A well constrained 'one-zone' chemical evolution model (CEM) for Leo A has been performed by \citet[][submitted, hereinafter H-M2018]{Hernandez-MartinezCarigi2018}, aiming to analyse its formation and evolution in a cosmological context and exploring the possibility of outflows driven by star formation. The main observational constraints that they included in their models are the star formation rate and the metallicity evolution obtained by \citet{GallartMonelli2015} which is part of the Local Cosmology from Isolated Dwarfs project Group (LCID)\footnote{http:\\www.iac.es/proyecto/LCID}. In addition, as the ISM abundances are very important to constrain the chemical evolution of the galaxy, they included the oxygen abundances measured by \citet{vanZeeSkillman2006} for the \ion{H}{II} regions and for the PN. \defcitealias{Hernandez-MartinezCarigi2018}{H-M2018} % \citetalias{Hernandez-MartinezCarigi2018} were able to reproduce the metallicity evolution and, within uncertainties, the past component of the gas for the PNe. However the O abundances predicted by their CEM are below the observational value for the \ion{H}{II} regions (present component). They argue that the chemistry of the four \ion{H}{II} regions analised by \citet{vanZeeSkillman2006} were computed by semi-empirical and strong-line methods and their uncertainties are quite large, so it is necessary to have better determinations. Our main goal in this work is to obtain accurate chemical abundances of two \ion{H}{II} regions in Leo A. The observations, obtained with the 10.4-m Gran Telescopio Canarias (GTC), and data reduction are described in \S 2. Line intensities, physical conditions and ionic abundances are presented in \S 3, where we have included the observations for the PN as given by \citet{vanZeeSkillman2006} to calculate its physical conditions using the same methodology as for \ion{H}{II} regions. In \S 4 the total abundances derived for the \ion{H}{II} regions and the PN, from different methods (direct method and strong line methods) are shown. Photoionisation structure models computed with the code CLOUDY for the observed \ion{H}{II} regions and the PN are described in \S 5 and our results are discussed in \S 6. | In this work we have derived the chemical composition in two \ion{H}{II} regions in Leo A, with different methods: the direct method (by using a very uncertain electron temperature measured in one of the \ion{H}{II} regions), the ONS method of strong lines \citep{PilyuginVilchez2010} and through the computation of photoionisation models. For the highest excitation \ion{H}{II} region (\ion{H}{ii} West) all the methods provide similar chemical abundances for O and N with values 12+log(O/H) $\sim$ 7.4 - 7.5 and log(N/O) $\sim -$1.7. The fainter \ion{H}{ii} region East does not show [\ion{O}{iii}] lines, therefore no direct method or ONS method can be applied, and its chemical composition was derived from a photoionisation model, giving 12+log(O/H) $\sim$ 7.4 and log(N/O) $\sim$ $-$1.6. Thus, it is confirmed that Leo A is one of the most metal poor galaxies in the Local Group, where its present ISM has 12+log(O/H) $\sim$ 7.4$\pm$0.2. Conclusively this galaxy has evolved very slowly in agreement with the star formation history derived by \citet{ColeSkillman2007} which includes initial 5 Gyr of quiescence. The log(N/O) = -1.6 is extremely low compared with the values for evolved galaxies as the Milky Way which presents log(N/O) of about -0.8 \citep{EstebanGarcia-Rojas2018}, but it is consistent with values found in irregular galaxies. From a sample of more than 30 low-metallicity irregular galaxies, \citet{Garnett1990} found that the ratio log(N/O) varies from $-$1.7 to $-$1.4, with an average log $<$N/O$>$ = $-$1.47$\pm$0.10, therefore the low N/O ratio in Leo A is similar to the values for low-metallicity irregular galaxies. The photoionisation models for the two analised \ion{H}{II} regions were computed with the code CLOUDY and the best fit models show that the ionising stars in these \ion{H}{II} regions correspond to a star with T$_{eff}$ of about 35600 K, and ionising photon rate of 10$^{48.9}$ photons s$^{-1}$ for the more ionised \ion{H}{ii} region West, and to a star of T$_{eff}$=28000 K and ionising photons rate of 10$^{48.0}$ photons s$^{-1}$ for the less ionised \ion{H}{ii} region East. % It is interesting to notice that \ion{H}{ii} West is the highest excitation \ion{H}{ii} region in Leo A, considering the 4 \ion{H}{ii} regions studied by \citeauthor{vanZeeSkillman2006}, and its central star has a mass of about 25 M$_\odot$. This indicates that a low-mass galaxy like Leo A, with a present SFR of about 10$^{-4}$ M$_\odot$ yr$^{-1}$ \citep{ColeSkillman2007} is capable of forming high mass stars, similarly to what occurs in the very low-mass galaxy Leo P \citep{McQuinnSkillman2015}. However and although we have not analysed the whole sample of \ion{H}{ii} regions in the galaxy, in the present starburst in Leo A it appears that the most massive star formed has mass greater or equal 25 M$_\odot$ and this is in agreement with the chemical evolution model by \citetalias{Hernandez-MartinezCarigi2018} where the upper limit required for the IMF for Leo A is 40 M$_\odot$. \medskip The only PN known in this galaxy is also a very metal poor object. Its O/H abundance as derived from the direct method and the photoionisation model is 12+log(O/H) $\sim$ 7.34 -- 7.45. The best model required a central star with T$_{eff}$ of 125,000 K, log(g) = 7.0 and ionising photon rate of 10$^{48}$ photons s$^{-1}$. The PN O/H abundance ratio is similar to the abundances of the \ion{H}{ii} regions and the same is true for other elements like Ne, Ar, and S. However the N/O abundance ratio is larger than in \ion{H}{ii} regions, with a value log(N/O) of $-$0.45. It is common to find that planetary nebulae have enriched their N abundances (also He and C) due to dredge-up processes or Hot Bottom Burning. In this case the enrichment is similar to the N enrichment of a disc PN in our galaxy \citep[see e.g.,][]{KingsburghBarlow1994}. The N/O value is not sufficiently high to declare this PN as a Peimbert Type I PN which are PNe with N/O larger than 0.5 and also seem to be He-enriched \citep{Peimbert1978}. The chemical abundances for this PN can be compared to the values predicted by stellar evolution models for low-intermediate mass stars by, e.g., \citet{Karakas2010} and \citet{FishlockKarakas2014} and it is found that the central star should have had an initial mass lower than about 2.0 M$_{\odot}$ in order to produce the observed chemical abundances in the PN. | 18 | 8 | 1808.06027 |
1808 | 1808.04508_arXiv.txt | We have used integral field spectroscopy to study the internal kinematics of the H\,{\sc ii} galaxies CTS 1020 and UM 461. We based our analysis on the velocity and velocity dispersion maps, and intensity-velocity dispersion ($I-\sigma$) and velocity-velocity dispersion ($V_r-\sigma$) diagrams. We found that the motion in both star-forming knots of UM 461 has different patterns, suggesting a weak kinematical connection between the knots. The overall kinematics of the galaxy is probably affected by stellar feedback. CTS 1020 has an ordered motion with a gradient compatible with a disc rotating at $\sim50$ km s$^{-1}$, though the velocity field is disturbed. In both galaxies the highest and lowest $\sigma$ values are distributed in the outer parts and are associated with the diffuse gas that permeates the galaxies. UM 461 has a ring-like structure with small regions of increasing $\sigma$ in the eastern knot, which resemble what we could expect in a collect and collapse scenario of star formation. We found that UM 461 seems to be more susceptible to stellar feedback, whereas in CTS 1020 the gravitational potential dominates. | H{\,\sc ii} galaxies, also known as \it blue compact dwarf \rm depending on classification criteria \citep{Melnick1985}, are a subclass of dwarf galaxies characterized by its compactness ($\sim1$ kpc), high star formation rate and a spectrum dominated by intense emission lines superimposed to a weak stellar continuum, which resembles that observed in giant H{\,\sc ii} regions in spiral galaxies \citep{Sargent1970}. Among the observed emission lines are the hydrogen Balmer series and the forbidden lines of oxygen ([O\,{\sc iii}] $\lambda\lambda4959$, 5007, 4363, [O\,{\sc ii}] $\lambda\lambda3726$, 3729), nitrogen ([N\,{\sc ii}] $\lambda\lambda6548$, 6583), and sulfur ([S\,{\sc ii}] $\lambda\lambda6716$, 6731). The fact of being gas-rich and metal-poor objects raised the hypothesis that these galaxies were young systems forming their first stars \citep{Searle1972}, but the idea of being old systems with intermittent star formation bursts interleaved by quiescent periods has been supported by observations of an underlying old stellar population \citep{Thuan1983,Telles1997a,Westera2004,Corbin2006}. An important characteristic of these galaxies is the supersonic nature of their emission line profile, which is broader than that observed in typical H{\,\sc ii} regions \citep{Smith1970,Smith1971}. \citet{Terlevich1981} proposed that the gravitational potential is responsible for this supersonic line widths, as they found a correlation between H$\beta$ luminosity and velocity dispersion ($L\sim\sigma^{4}$) and radius and velocity dispersion ($R\sim\sigma^{2}$) in giant H{\,\sc ii} regions. This scenario was further supported by \citet{TenorioTagle1993}, who proposed a model to explain and maintain the supersonic motion, which is given by the constant passage of low-mass stars producing bow shocks. Alternatively, the supersonic motions may be maintained by the mechanical energy injected into the interstellar medium from the ongoing star formation activity, stellar winds, radiation pressure and supernovae explosions, all of which contribute to increase the turbulence \citep{Green2010,Moiseev2012,Moiseev2015}. These scenarios have difficulties to explain the $L\sim\sigma^{4}$ and $R\sim\sigma^{2}$ relations. Which of these mechanisms is dominant remains an open problem. \citet{Gallagher1983} suggested that the dominant mechanism depends on the system scale, being the gravitational potential in regions of hundreds of parsecs (supergiant H{\,\sc ii} regions) and energy from massive young stellar populations in regions of tens of parsecs (giant H{\,\sc ii} regions). Furthermore, \citet{TenorioTagle1996} demonstrated that the mechanical energy from massive stars is not sufficient to explain the observed line broadening, in agreement with \citet{Yang1996} who found that the stellar winds and supernova explosions act increasing the dispersion caused by the gravitational potential. Kinematics and dynamics of H{\,\sc ii} galaxies have been investigated, early on, by looking at their velocity dispersion. Rapidly a scale relation between emission line and velocity dispersion have been established \citep{Melnick1988}. These scale relations have been used early on as distance indicator. Taking advantage of larger survey and better data, recent studies could minimise errors in these relations \citep{Bordalo2011, Chavez2014}. These studies also show that such relations are subject to evolutionary effects, responsible, according to authors, for part of the dispersion of such relations. More recently, high redshift H{\,\sc ii} galaxies have been used as tools for precise cosmology \citep{Terlevich2015}. Kinematical studies of H\,{\sc ii} galaxies with 2D mapping instrumentation were first made by \citet{Ostlin1999,Ostlin2001} and focused on the analysis of velocity field to determine the mass distribution using rotation curves. The results showed a disturbed velocity field and supersonic velocity dispersion in a small sample of H\,\sc ii \rm galaxies, what pointed out that velocity dispersion dominates the gravitational potential. However, as the morphology of the galaxies suggested interaction or merger, the authors concluded that those galaxies were not systems in equilibrium, but in a merger process \citep{Ostlin2001}. Recently, using integral field spectroscopy, \citet{Lagos2016} and \citet{Kumari2017} came up with the same conclusion for the galaxies Tol 65 and NGC 4449, respectively, with a merger being responsible for trigger the star formation. The kinematics of both objects seems to be affected by stellar feedback \citep{Lagos2016,Kumari2017}. Further evidence of stellar feedback on H\,\sc ii \rm galaxies is presented by \citet{Cairos2017a,Cairos2017b} that found supersonic velocity dispersion in areas surrounding H\,\sc ii \rm regions and in the outskirts of the galaxies. In order to disentangle the line broadening mechanisms, we used diagnostic diagrams, such as $I - \sigma $, $I - V_r$ and $V_r - \sigma$, that have been revealed to be precious tools to find signatures of peculiar motions, as expanding shells, radial motions \citep{MunozTunon1996,Bordalo2009,Plana2017} and of turbulent ISM and \hbox{H\,{\sc ii}} region \citep{Moiseev2012}. The paper is organized as follows. In Section~\ref{sec:OR} we present observations and the reduction techniques used. The ionized gas kinematics is presented in \S~\ref{sec:IFUresults}, discussion of the diagnostic diagrams of H\,{\textsc{ii}} complexes is presented in \S~\ref{sec:diagrams} along with a statistical analysis of specific diagrams in section \S~\ref{sec:statistics}. The section \S~\ref{sec:PCA} is dedicated to a Principal Component Analysis of the data cubes. In \S~\ref{sec:discussion} our results are discussed. Finally in \S~\ref{sec:conclusion}, we give the summary and draw general conclusion. | \label{sec:conclusion} In this work, we have studied the H{\,\sc ii} galaxies UM 461 and CTS 1020 based on integral field spectroscopy (Gemini GMOS-IFU). Taking advantage of monochromatic, velocity and velocity dispersion maps, we embarque in a kinematical analysis using different diagnostic diagrams (like $I-\sigma$ and $V_r-\sigma$) to investigate the nature of the internal kinematics for both objects. % As mentioned before, velocity dispersion of ionized gas plays a major role in H{\,\sc ii} galaxies dynamics. The $L~ vs~ \sigma$ relation, based on single measurements, is interpreted as gravity being the main mechanism causing the supersonic broadening of emission profiles \citep{Chavez2014}. The main result of our study is that the kinematics of ionized gas of these two galaxies is different, but it also shows similarities. Differences come from the velocity and velocity dispersion maps themselves: in UM 461 no ordered motion is present, only velocity gradient; in CTS 1020 a disk like rotation pattern can be seen, even if a larger field of view is necessary to confirm it. Velocity dispersion maps show the same differences: in UM 461 regions of low dispersion correspond to high intensity regions, and CTS 1020 shows high dispersion areas where the intensity is the hightest and where intensity is low as well. $I-\sigma$ diagrams for both galaxies offer some differences. UM 461 diagram shows, according \citet{Moiseev2012}, signature of H\,{\sc ii} regions, constant velocity dispersion and high monochromatic emission, in both knots centers. It also shows the presence of turbulent diffuse gas. On the other hand, despite thte fact that, in CTS 1020 case, is still possible to identify the turbulent diffuse gas surrounding the galaxy, high $\sigma$ is related to high intensity and seems to decrease outwards, suggesting that $\sigma$ is driven by virial motions. Applying statistical methods, a closer analysis of the $V_r - \sigma$ diagrams, shows that several independent regions with weak and moderate correlation are consistent with systemic motions toward or away the observer. When reported on a geographic map, these regions are consistent with low and high velocities on the velocity maps. In the case of CTS 1020, it might means that the rotating disk can also be interpreted as large regions animated of opposite motions. A large field of view will be needed in order to find out if the velocity field really represents a rotating disk in that case. Finally, we also performed a PCA analyis of the data cube in order to improve the SNR. Our results show that data have been improved where the SNR was low, but also show that PCA seems to have modified the shape of the reconstructed profiles, resulting in more symetrical ones. | 18 | 8 | 1808.04508 |
1808 | 1808.06816_arXiv.txt | We consider the model of composite dark matter assuming stable particles of charge $-2$ bound with primordial helium nuclei by the Coulomb force in $O$He atoms. We study capture of such dark atoms in matter and propose the possibility of the existence of stable $O$-enriched superheavy nuclei and $O$-nuclearites, in which heavy $O$-dark matter fermions are bound by electromagnetic forces with ordinary nuclear matter. $O$He atoms accumulation in stars and its possible effect on stellar evolution is also considered, extending the set of indirect probes for composite dark matter. | \label{sec:sec1} There is overwhelming evidence for the presence of a dark matter (DM) in the Universe \cite{Ade:2015xua} and together with most popular, but still elusive weakly interacting massive particle (WIMP) \cite{Arcadi.2018}, there exist numerous theoretical models including axions, sterile neutrinos, primordial black holes \cite{Belotsky.2014,Carr.2016,Carr.PRD.2010}, strongly interacting massive particles and superweakly interacting particles (see Refs.~\cite{DMRev,Aprile:2009zzd,Feng:2010gw} for review and references). Even electromagnetically interacting massive particle (EIMP) candidates are possibly hidden in neutral atomlike states. Dark $O$He atoms, in which hypothetical $-2$ charged particles are bound with primordial helium nuclei, occupy a special place on this list. Such models involve only one free parameter of new physics --- the mass of $-2$ charged EIMPs --- so many features of this type of dark matter can be described by the known nuclear and atomic physics. In 2005, Glashow \cite{Glashow:2005jy} proposed a kind of EIMP model, according to which stable teraquarks $U$ (of mass of the order of tera-electron-volts and of electric charge $+2/3$) form a $UUU$ baryon bound with tera-electrons $E$ of charge $-1$ in the neutral $(UUUEE)$ atom. However, the primordial He formed in the big bang nucleosynthesis captures all the free $E$ in positively charged (He$E$)$^+$ ions, preventing a required suppression of the positively charged particles that can bind with electrons in atoms of anomalous hydrogen. In general, stable single charged EIMPs form anomalous hydrogen either directly binding with ordinary electrons ($+1$ charged EIMPs), or indirectly ($-1$ charged EIMPs) forming first $+1$ charge ion with primordial helium and then anomalous hydrogen with ordinary electrons \cite{BKSR1}. Therefore, anomalous hydrogen overproduction excludes any significant amount of stable single charged EIMPs. Nevertheless, there are several models that predict stable double charged particles without stable single charged particles. In particular, the hypothesis of the heavy stable quark of the fourth family may provide a solution, if an excess of $\bar{U}$ antiquarks with charge $(-2/3)$ is generated in the early Universe. Excessive $\bar{U}$ antiquarks then form $\bar{U}\bar{U}\bar{U}$ antibaryons with the electric charge $-2$, which are captured by He forming $O^{--}$He$^{++}$ ($O$He) atoms \cite{Khlopov:2005ew} right after the appearance of the He nuclei in the big bang nucleosynthesis. This hypothesis has found implementations in the model of almost commutative geometry as well as in models of walking technicolor and has been extensively discussed in the literature; see Refs.~\cite{Khlopov:2010ik,Khlopov:2011me,gKhlopov,khlopov.proc.2014,khlopov.proc.2015,khlopov.ijmpd.2015,Wallemacq:2015cjr} and references therein. The model is particularly predictive since the only parameter that one needs to know is the mass of the $O$-particle. The model can explain the observed excess of the positronium annihilation line in the galactic bulge and excessive fraction of high-energy cosmic-ray positrons, if the mass of this particle does not exceed 1.3 TeV, challenging the direct test of this explanation in searches for stable double charged particles at the LHC \cite{probes}. Charge conservation implies the existence of $+2$ charged particle $O^{++}$ together with $O^{--}$. To avoid overproduction of anomalous isotopes by $O^{++}$, $O$He-dominated dark matter should be asymmetric with strongly suppressed $+2$ charged particles. In the walking technicolor model \cite{KK,KK1}, due to sphaleron transitions, such excess is related to the baryon excess, giving the observed dark matter/baryon matter density ratio for a reasonable choice of parameters. In the early Universe when temperature fell below 1 keV, the rate of expansion started to exceed the rate of energy and momentum transfer from plasma to $O$He gas (see, e.g., Ref.~\cite{gKhlopov} for review and references). As a result, $O$He decoupled from plasma and radiation and played the role of dark matter on the matter-dominated stage. Before decoupling from plasma and radiation, $O$He density fluctuations convert in sound waves. It leads to the suppression of small-scale fluctuations. Thereby $O$He dark matter was called warmer than cold dark matter for an $O$He mass about 1 TeV, typical for cold dark matter particles \cite{khlopov.proc.2014}. The averaged baryonic density in the course of structure formation and in galaxies is sufficiently low making baryonic matter at large scales transparent for $O$He. So, for a galaxy with mass $M = 10^{10}M_\odot$ and radius $R = 10^{23}$ cm, $n\sigma R = 8 \cdot 10^{-5} \ll 1$, where $n = M/4\pi R^3$ and $\sigma = 2 \cdot 10^{-25}$ cm$^2$ is the geometrical cross section for $O$He collisions. For that reason, in the period of formation of the first objects, $O$He does not follow the condensation of baryonic matter, so the $O$He model avoids constraints from the cosmic microwave background \cite{CMB} and formation of the first stars \cite{FS}. In galaxies and galaxy clusters, $O$He behaves like collisionless gas avoiding constraints from Bullet Cluster observations \cite{BC}. Only dense matter objects like stars or planets are opaque for it. The protostellar cloud with the solar mass becomes opaque for $O$He when it contracts within $8 \cdot 10^{15}$ cm. Correspondingly, the protoplanet cloud of the mass of the Earth becomes opaque when it contracts to $10^{13}$ cm. Because of the nuclear interaction cross section of elastic collisions with terrestrial matter, $O$He is slowed down to thermal velocity in the matter of underground detectors. It leads to negligible nuclear recoil in $O$He collisions with nuclei in direct-detection experiments. Positive results of DAMA/NaI and DAMA/LIBRA and negative results of other groups are explained in the $O$He model by annual modulation of the rate of low-energy binding of $O$He with intermediate mass nuclei \cite{Khlopov:2010ik,Khlopov:2011me,gKhlopov}. Open problems of this explanation related with the existence and role of the dipole potential barrier in $O$He-nucleus interaction are discussed in Refs.~\cite{khlopov.proc.2014,khlopov.proc.2015,khlopov.ijmpd.2015,Wallemacq:2015cjr}. On the other hand, various hypotheses of the existence of superheavy nuclei with the atomic numbers essentially higher than that of ordinary atomic nuclei have been explored. In 1971, Migdal suggested the possibility of superdense nuclei glued by a pion condensate \cite{Migdal:1971cu,Migdal:1974yx,Migdal:1978az,Migdal:1990vm}. Lee and Wick conjectured $\sigma$-condensate superheavy nuclei \cite{Lee:1974ma,Lee:1974kn}. Bodmer proposed collapsed quark nuclei \cite{Bodmer:1971we}. Reference \cite{Migdal:1977rn} demonstrated that the interior of a nucleus with a charge $Z\gg 1/e^3$, $e$ is the charge of the electron, $\hbar=c=1$, is electrically neutral and Refs.~\cite{Voskresensky:1977mz,Voskresensky:1978uf,Migdal:1990vm} suggested the possibility of existence of nuclei stars of the atomic number $(10^2-10^3)\le A\le 10^{57}$, the electric charge of which is compensated by the negatively charged pion condensate and the electrons. References \cite{Migdal:1990vm,Voskresensky1977,Kolomeitsev:2002gd} argued that if there existed negatively charged light bosons of mass less than $(30-32)$ MeV there would exist exotic objects, nuclei stars, of arbitrary size (until the effects of gravity can be neglected) with density typical for normal atomic nuclei, bound by strong and electromagnetic interactions. Witten \cite{Witten:1984rs} suggested the possible existence of quark nuggets, constructed from up, down, and strange quarks, with the atomic number between $(3\cdot 10^2-10^3)\le A\le 10^{57}$, see Ref.~\cite{Alcock:1986hz}, as candidates for the DM in the Universe. De Rujula and Glashow \cite{DeRujula:1984axn} called these stable drops ``nuclearites'' and discussed conditions for their feasible detection in terrestrial conditions. They have also discussed charged massive particles (CHAMPs) \cite{DeRujula:1989fe}. They argued that negative CHAMPs may bind to protons in superheavy isotopes. Superheavy nuclei and nuclearites may exist in the Galaxy as debris from the big bang, supernovae explosions, star collisions, and other astrophysical catastrophes. Numerous subsequent works focused on the consideration of the strange stars as a new family of compact stars. Besides that, exotic matter like the pion condensates and the quark matter in various phases may exist in the interiors of some neutron stars \cite{Weber,Ivanov:2005be,Boeckel:2010hm,Dondi:2016yjl}. The other side of the problem is the possible influence of dark matter captured by stars on the stellar structure and evolution. In particular, it can lead to observable effects in neutron stars \cite{Kouvaris.PRD.2014}. Below, we assume that the DM may consist of $O$-particles bound in $O$He atoms. Colliding with the ordinary atomic nuclei, $O$He atoms may undergo fusion reactions with the formation of superheavy $O$-nuclei. However, the simplest from the viewpoint of new physics and principally being the subject of the complete quantum mechanical treatment of $O$He interaction with matter, such a description still remains an open question of the $O$He model. Putting aside this uncertainty, we suggest the idea of the possibility of the existence of $O$-nuclearites, constructed of self-bound nuclear matter at the density typical for the nuclear saturation, in which the positive electric charge of protons is compensated by negatively charged $O^{--}$. Such nuclearites might be formed in $O$He interaction with nuclei, and we study their effect in astrophysical conditions. The paper is organized as follows. In Sec.~\ref{sec:sec2}, we formulate the idea of the existence of $O$-nuclearites. In Sec.~\ref{sec:sec3} we take into account the effects of gravitation. Then, Sec.~\ref{sec:sec4} presents some estimates for the $O$-nuclearite accumulation during star evolutions. Finally, Sec.~\ref{sec:sec5} contains some concluding remarks. | \label{sec:sec5} With the lack of evidence for WIMPs in direct and indirect searches for dark matter, the fields of study of possible dark matter physics should be strongly extended. Dark atoms of $O$He are of special interest in view of the minimal involvement of new physics in their properties. The hypothesis on stable double charged particle constituents of dark atoms sheds new light on the strategy of dark matter studies, offering a nontrivial explanation for the puzzles of direct and indirect dark matter searches. In particular, in the context of this hypothesis, collider searches for dark matter are not related to the effect of missing mass, momentum, and energy, but are related to the search for stable double charged particles. Astrophysical indirect effects of $O$He dark matter are related to radiation from $O$He excitation in collisions in the center of Galaxy. It can explain the excess of the positronium annihilation line, observed by INTEGRAL in the galactic bulge, provided that the mass of the double charged $O$ particle is near 1.25 TeV, that is within the reach of search for such particle at the LHC. However simple in the description of new physics, the old-fashioned and seemingly well-known nuclear and atomic physics turn out to be nontrivial and rather complicated in the description of dark atoms and their interaction with matter. Nuclear physics of $O$He atoms is still unclear and remains an open problem of this approach. The crucial point is the existence of a potential barrier in the interaction of $O$He with nuclei. If such a barrier exists in the $O$He interaction with sodium nuclei, the capture of the Na nucleus by $O$He to a low-energy bound state beyond nuclear radii can explain the positive effect of direct dark matter searches for the annual modulation signal in DAMA/NaI and DAMA/LIBRA experiments. Annual modulation follows in this explanation from the annual modulation of the $O$He concentration in the matter of the detector, while small recoil energy explains the absence of positive effects in other experiments. The rate of capture is determined by electric dipole transition, which is strongly suppressed in cryogenic detectors, while the absence of a low-energy bond state in $O$He interaction with heavy nuclei makes it impossible to test this hypothesis in detectors with heavy element content, like liquid xenon. On the other hand, if such barrier does not exist or is not efficient, inelastic collisions dominate in $O$He-nucleus interactions and overproduction of anomalous isotopes inevitably rules out the $O$He dark matter hypothesis. The formation of an $O$H$^{-}$ ion in proton capture by $O^{--}$ may lead to another potential problem for the $O$He scenario. The abundance of such ions is severely constrained by searches for stable charged massive particles and anomalous isotopes in sea water \cite{Smith.NPB.1979,Smith.NPB.1982,Hemmick.PRD.1990,Verkerk.PRL.1992,Yamagata.PRD.1993,Kudo.PLB.2001}. Production of such ions in the early Universe is strongly suppressed, since all the free $O^{--}$ are captured by primordial helium before proton capture becomes possible. However, in the Galaxy, $O$He destructions in stars and in cosmic rays can release free $O^{--}$, which can be captured by protons, forming $O$H$^{-}$ ions and an anomalous $-2$ charged component of cosmic rays. In principle, the capture of such components by Earth can lead to a dangerous amount of anomalous isotopes in sea water, but the corresponding analysis, involving detailed study of $O$He evolution in the Galaxy, goes beyond the scope of the present work. Putting aside these problems, we turn here to the extension of studies of possible effects of $O$He in nuclear matter and astrophysical conditions. We proposed the possibility of the existence of stable $O$-nuclearites and discussed various mechanisms for their formation. Observation of $O$-nuclearites, in which dark matter is bound with the normal nuclear matter, would be an important event that could provide us additional information on the possibility of the existence of dark $O$He atoms of dark matter and their properties. | 18 | 8 | 1808.06816 |
1808 | 1808.08726_arXiv.txt | {Current models that explain giant (type II) X-ray outbursts in Be/X-ray binaries (BeXB), are based on the idea of highly distorted disks. They are believed to occur when a misaligned and warped disk becomes eccentric, allowing the neutron star to capture a large amount of material. The BeXB 4U\,0115+63 underwent two major outbursts in 2015 and 2017.} {Our aim is to investigate whether the structural changes in the disk expected during type II outbursts can be detected through optical polarimetry.} {We present the first optical polarimetric observations and new optical spectra of the BeXB 4U 0115+63 covering the period 2013--2017. We study in detail the shape of the H$\alpha$ line profile and the polarization parameters before, during, and after the occurrence of a type II X-ray outburst. } {We find significant changes in polarization degree and polarization angle and highly distorted line profiles during the 2017 X-ray outburst. The degree of polarization decreased by $\sim$ 1\%, while the polarization angle, which is supposed to be related with the disk orientation, first increased by $\sim 10^{\circ}$ in about two months and then decreased by a similar amount and on a similar timescale once the X-ray activity ceased.} {We interpret the polarimetric and spectroscopic variability as evidence for the presence of a warped disk.} | The X-ray source \src\ is a Be/X-ray binary \citep[BeXB,][]{ziolkowski02,paul11,reig11}. These systems contain a neutron star and a Be star in a relatively wide ($P_{\rm orb}\sim$ few tens of days) and eccentric orbit. The most prominent feature of a Be star is the gaseous equatorial disk around its equator. The disk is in Keplerian rotation, is geometrically thin, and is in vertical hydrostatic equilibrium \citep{rivinius13}. The disk is of paramount importance to understanding the behavior of BeXBs as it constitutes the main source of variability and is responsible for the three main observational properties of Be stars: emission lines, infrared excess, and polarization. The emission lines and infrared excess are formed by recombination in the disk. Linear polarization results from Thomson scattering, when photons from the Be star scatter with electrons in the Be disk \citep{poeckert79,wood96,yudin01,halonen13a,haubois14}. The effect of Thomson scattering is to reduce the component of the electric vector parallel to the scattering plane by $\cos^2\chi$, where $\chi$ is the scattering angle, while the intensity of the perpendicular component remains unaltered after scattering. Thus the light becomes polarized perpendicularly to the scattering plane, which roughly coincides with the plane of the disk. The polarization angle then gives information about the orientation of the disk \citep{wood96, quirrenbach97}. The polarization degree increases with the optical depth (or density) of the gas in the disk and with the inclination with respect to the observer \citep{wood96,halonen13a}. BeXBs are transient X-ray sources that spend most of the time in a dormant state, although persistent X-ray sources also exist \citep{reig99}. When active, transient BeXBs exhibit two types of X-ray outburst. Normal or type I outbursts show a moderate increase in X-ray flux ($L_X \simless 10^{37}$ \ergs), occur near periastron passage, and last for a fraction of the orbit. Giant or type II outbursts are significantly brighter ($L_X \simmore 10^{37}$ \ergs), do not occur at any preferential orbital phase and may last for several orbits. Current models that explain type II outbursts are based on the idea that a highly misaligned disk becomes warped and eccentric, allowing the neutron star to capture a large amount of material \citep{martin11,okazaki13,martin14a}. \citet{okazaki13} performed numerical simulations and showed that the warped shape of a misaligned disk leads to enhanced mass accretion if the warped part gets across the orbit of the neutron star. They also showed that the accretion rate is much higher for small tilt angles ($\beta \simless 20^{\circ}$) than for large tilt angles ($\beta \simmore 40^{\circ}$). The reason is that at low tilt angles, the tidal torque of the neutron star favors the creation of dense (denser than the disk) mass streams even before the actual collision between the disk and the neutron star. When the neutron star interacts with the stream, the mass accretion rate is enhanced. In highly tilted systems, the stream forms more slowly because the tidal torque is weaker. Therefore the disk itself has to be several times denser than in less tilted systems to supply the same amount of gas. \citet{martin14a} showed that giant outbursts can also occur in highly tilted systems if the disk is highly eccentric. By increasing the eccentricity of the disk, the neutron star can capture matter from its outer parts more easily than in a circular disk. There are several ways to make an initially circular disk eccentric but the most promising one appears to be the Kozai-Lidov mechanism \citep{kozai62,lidov62,martin14b,fu15a}. The key idea in this mechanism is that the product of the disk inclination and eccentricity remains constant. Thus, a test particle that is initially on a circular orbit in a misaligned disk undergoes a series of oscillations where the inclination and eccentricity interchange periodically. During the oscillation, the disk can attain large eccentricities after a few orbital periods. Simulations show that the more tilted the disk is, the larger the disk can grow and the more eccentric it becomes \citep{martin14a}. Since the polarimetric parameters depend so strongly on the internal (density) and external (orientation) conditions in the disk and the models that explain type II outbursts require warped, misaligned, eccentric precessing disks, then we should observe significant changes in the polarization degree and angle during a giant outburst. This work examines this idea. \begin{table*} \caption{Results of the polarization observations of the optical counterpart to \src\ in the $R$ band. The uncertainties were determined following the prescription given in \citet{king14}.} \label{pol} \begin{center} \begin{tabular}{ccccccccc} \noalign{\smallskip} \hline\noalign{\smallskip} JD (2,400,00+) & P[\%] & $\sigma_P$[\%] & PA[deg] & $\sigma_{PA}$[deg] & $q$ &$\sigma_q$ & $u$& $\sigma_u$ \\ \noalign{\smallskip}\hline\noalign{\smallskip} 56549.5671 & 2.8 & 0.7 & -72.6 & 7.4 & -0.0228 & 0.0075 & -0.0158 & 0.0070 \\ 56576.4997 & 3.4 & 0.5 & -67.3 & 3.9 & -0.0238 & 0.0047 & -0.0242 & 0.0046 \\ 56592.4652 & 3.9 & 0.5 & -67.3 & 3.8 & -0.0272 & 0.0050 & -0.0275 & 0.0052 \\ 56623.3618 & 3.0 & 0.3 & -68.8 & 3.1 & -0.0222 & 0.0034 & -0.0202 & 0.0033 \\ 56854.5923 & 3.9 & 0.4 & -69.0 & 2.6 & -0.0291 & 0.0035 & -0.0262 & 0.0036 \\ 56872.5402 & 4.4 & 0.5 & -66.6 & 3.3 & -0.0301 & 0.0051 & -0.0320 & 0.0050 \\ 56885.5121 & 4.2 & 0.5 & -65.3 & 3.5 & -0.0277 & 0.0051 & -0.0321 & 0.0051 \\ 56903.5471 & 3.9 & 0.5 & -63.7 & 3.8 & -0.0238 & 0.0053 & -0.0309 & 0.0052 \\ 57264.5125 & 3.9 & 0.4 & -70.0 & 2.9 & -0.0301 & 0.0043 & -0.0252 & 0.0037 \\ 57346.3532 & 4.2 & 0.2 & -69.0 & 1.8 & -0.0316 & 0.0026 & -0.0284 & 0.0026 \\ 57555.5471 & 4.3 & 0.2 & -70.0 & 1.4 & -0.0328 & 0.0022 & -0.0276 & 0.0019 \\ 57641.5372 & 3.8 & 0.2 & -68.5 & 1.2 & -0.0279 & 0.0017 & -0.0261 & 0.0016 \\ 57665.4385 & 4.0 & 0.1 & -68.1 & 1.0 & -0.0292 & 0.0014 & -0.0279 & 0.0014 \\ 57695.4444 & 4.0 & 0.1 & -67.0 & 0.9 & -0.0280 & 0.0013 & -0.0288 & 0.0013 \\ 57929.5633 & 4.0 & 0.1 & -70.2 & 0.9 & -0.0312 & 0.0014 & -0.0258 & 0.0012 \\ 57957.4767 & 4.0 & 0.1 & -64.9 & 1.0 & -0.0255 & 0.0014 & -0.0306 & 0.0013 \\ 57966.4957 & 4.2 & 0.2 & -63.5 & 1.1 & -0.0252 & 0.0017 & -0.0335 & 0.0016 \\ 57970.5607 & 4.1 & 0.1 & -66.0 & 1.0 & -0.0277 & 0.0015 & -0.0307 & 0.0015 \\ 57985.6152 & 4.2 & 0.2 & -61.5 & 1.4 & -0.0228 & 0.0019 & -0.0352 & 0.0021 \\ 57986.5391 & 4.2 & 0.2 & -61.2 & 1.4 & -0.0223 & 0.0020 & -0.0352 & 0.0019 \\ 57994.5950 & 3.8 & 0.2 & -62.2 & 1.3 & -0.0215 & 0.0018 & -0.0314 & 0.0018 \\ 58015.4837 & 3.9 & 0.3 & -64.3 & 1.9 & -0.0246 & 0.0027 & -0.0309 & 0.0025 \\ 58020.5472 & 3.5 & 0.3 & -64.7 & 2.4 & -0.0221 & 0.0029 & -0.0270 & 0.0030 \\ 58043.5234 & 3.1 & 0.3 & -70.8 & 3.1 & -0.0246 & 0.0034 & -0.0195 & 0.0034 \\ \noalign{\smallskip} \hline \end{tabular} \end{center} \end{table*} \begin{table*} \caption{Polarization degree in various bands.} \label{multipol} \begin{center} \begin{tabular}{cccccc} \noalign{\smallskip} \hline\noalign{\smallskip} Date & JD (2,400,000+) & $B$ & $V$ & $R$ & $I$ \\ & & (\%) &(\%) &(\%) &(\%) \\ \noalign{\smallskip} \hline\noalign{\smallskip} 16-06-2016 &57555.56 &4.14$\pm$1.45 &3.89$\pm$0.36 &4.29$\pm$0.21 &3.24$\pm$0.26 \\ 10-09-2016 &57641.55 &4.06$\pm$1.43 &4.00$\pm$0.34 &3.82$\pm$0.17 &3.40$\pm$0.15 \\ 02-11-2016 &57695.46 &4.12$\pm$0.59 &3.89$\pm$0.23 &4.02$\pm$0.13 &3.62$\pm$0.15 \\ 01-08-2017 &57966.53 &4.12$\pm$0.76 &4.13$\pm$0.25 &4.20$\pm$0.16 &3.88$\pm$0.14 \\ 17-10-2017 &58043.54 &4.85$\pm$1.61 &3.07$\pm$0.70 &3.14$\pm$0.34 &2.97$\pm$0.46 \\ \noalign{\smallskip} \hline \end{tabular} \end{center} \end{table*} \begin{table*} \caption{Calibrated photometric magnitudes. The errors correspond to the standard deviation of the difference between the measured and the cataloged value of the standard stars.} \label{phot} \begin{center} \begin{tabular}{cccccc} \noalign{\smallskip} \hline \noalign{\smallskip} Date &JD (2,400,000+) & B &V &R &I \\ \noalign{\smallskip} \hline \noalign{\smallskip} 29-07-2013 &56503.54 & 16.95$\pm$0.02 & 15.41$\pm$0.02 & 14.48$\pm$0.02 & 13.45$\pm$0.02 \\ 29-08-2013 &56534.56 & 17.07$\pm$0.02 & 15.53$\pm$0.02 & 14.62$\pm$0.02 & 13.58$\pm$0.04 \\ 20-08-2014 &56890.57 & 16.97$\pm$0.02 & 15.49$\pm$0.02 & 14.61$\pm$0.02 & 13.64$\pm$0.03 \\ 14-09-2014 &56915.49 & 16.97$\pm$0.02 & 15.48$\pm$0.01 & 14.58$\pm$0.02 & 13.56$\pm$0.03 \\ 22-07-2015 &57226.56 & 16.99$\pm$0.02 & 15.37$\pm$0.01 & 14.37$\pm$0.01 & 13.24$\pm$0.01 \\ 18-11-2015 &57345.41 & 16.88$\pm$0.02 & 15.24$\pm$0.01 & 14.22$\pm$0.02 & 13.11$\pm$0.02 \\ 08-09-2016 &57640.54 & 16.73$\pm$0.09 & 15.00$\pm$0.08 & 13.91$\pm$0.09 & 12.78$\pm$0.09 \\ 06-10-2016 &57668.51 & 16.55$\pm$0.01 & 14.84$\pm$0.01 & 13.76$\pm$0.01 & 12.60$\pm$0.01 \\ 03-11-2016 &57696.49 & 16.61$\pm$0.02 & 14.88$\pm$0.02 & 13.79$\pm$0.02 & 12.62$\pm$0.02 \\ 25-06-2017 &57930.58 & 16.35$\pm$0.03 & 14.59$\pm$0.03 & 13.47$\pm$0.06 & 12.32$\pm$0.07 \\ 11-07-2017 &57946.54 & 16.37$\pm$0.02 & 14.63$\pm$0.02 & 13.57$\pm$0.02 & 12.44$\pm$0.03 \\ 28-08-2017 &57994.47 & 16.77$\pm$0.02 & 15.15$\pm$0.02 & 14.12$\pm$0.02 & 13.02$\pm$0.03 \\ \noalign{\smallskip} \hline \end{tabular} \end{center} \end{table*} | We report for the first time changes in the optical polarization parameters during a giant X-ray outburst in the BeXB \src. Most notably, the polarization angle, and the profile of the \ha\ line, and therefore the orientation of the Be disk, changed on timescales comparable to the orbital period. We interpret this variability as evidence for a warped disk, supporting models that predict highly perturbed disks as the origin of type II outbursts in BeXB. | 18 | 8 | 1808.08726 |
1808 | 1808.09516_arXiv.txt | { Using the same lens galaxies, the ratios of tangential shears for different source galaxy redshifts is equal to the ratios of their corresponding angular-diameter distances. This is the so-called shear-ratio test (SRT) and it is valid when effects induced by the intervening large-scale structure (LSS) can be neglected. The dominant LSS effect is magnification bias which, on the one hand, induces an additional shear, and on the other hand, causes a magnification of the lens population. Our objective is to quantify the magnification bias for the SRT and show an easy-to-apply mitigation strategy that does not rely on additional observations. We use ray-tracing data through the Millennium simulation to measure the influence of magnification on the SRT and test our mitigation strategy. Using the SRT as a null-test we find deviations from zero up to \(10 \%\) for a flux-limited sample of lens galaxies, which is a strong function of lens redshift and the lens-source line-of-sight separation. Using our mitigation strategy we can improve the null-test by a factor of \(\sim \!100\). } | 18 | 8 | 1808.09516 |
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1808 | 1808.10682_arXiv.txt | We present new constraints on gas-phase C, N, and O abundances in the molecular layer of the IM Lup protoplanetary disk. Building on previous physical and chemical modeling of this disk, we use new ALMA observations of C$_2$H to constrain the C/O ratio in the molecular layer to be $\sim0.8$, i.e., higher than the solar value of $\sim0.54$. We use archival ALMA observations of HCN and H$^{13}$CN to show that no depletion of N is required (assuming an interstellar abundance of $7.5\times10^{-5}$ per H). These results suggest that an appreciable fraction of O is sequestered in water ice in large grains settled to the disk mid-plane. Similarly, a fraction of the available C is locked up in less volatile molecules. By contrast, N remains largely unprocessed, likely as N$_2$. This pattern of depletion suggests the presence of true abundance variations in this disk, and not a simple overall depletion of gas mass. If these results hold more generally, then combined CO, C$_2$H, and HCN observations of disks may provide a promising path for constraining gas-phase C/O and N/O during planet-formation. Together, these tracers offer the opportunity to link the volatile compositions of disks to the atmospheres of planets formed from them. | Gas-rich circumstellar disks around young stars provide a window to study the materials that are incorporated into forming planetary systems. Chemistry and dynamics in the disk can shift the balance of gas- versus ice-phase volatiles containing, e.g., carbon, nitrogen, and oxygen. In the core accretion paradigm \citep{pollack1996}, those materials that end up as rocks or ices are incorporated into the ``solid'' planetesimals, while the remaining gas can be accreted into natal planets' atmospheres. Correspondingly, it is essential to understand the form volatiles containing carbon, nitrogen, and oxygen take in the disk spatially and over time, to understand the initial elemental compositions of forming planets \citep[e.g.,][]{oberg2011co,piso2016,cridland2016,cridland2017}. Many processes can alter the gas versus solid abundances in the disk, including snow lines \citep[e.g.,][]{oberg2011co}, chemistry \citep{bergin2014far,furuya2014,eistrup2016,schwarz2018}, mixing and/or diffusion of gas \citep[e.g.,][]{semenov2011,kama2016}, and redistribution of ices as dust grows and evolves \citep[e.g.,][]{hogerheijde2011,piso2016,krijt2016,oberg2016}. To date, observations of various carbon and oxygen carriers have suggested a substantial ``missing'' volatile mass within the disk molecular layer as traced by CO in the submillimeter \citep{favre2013,cleeves2015tw,zhang2017} and H$_2$O vapor in the far-infrared with {\em Herschel} \citep{bergin2010,hogerheijde2011,du2017}, though see also \citet{kamp2013} regarding model dependencies. For the few disks where estimates for both exist, more water in the disk surface is ``missing'' than CO when compared to interstellar abundances, and both have abundances lower than what simple desorption (thermal and non-thermal) models predict \citep[e.g., in TW Hya's disk;][]{hogerheijde2011,favre2013,schwarz2016,kama2016,cleeves2015tw}. Absolute elemental abundances are often difficult to estimate due to uncertain hydrogen disk masses. In this context, relative elemental abundances, such as C/O and N/O are promising avenues toward robustly characterizing disk volatile compositions. Recently, \citet{bergin2016} reported bright hydrocarbon rings of C$_2$H and $c-$C$_3$H$_2$ in the TW Hya \citep[see also][]{kastner2015} and DM Tau disks. Using chemical models, \citet{bergin2016} found the abundances of these small hydrocarbons were especially sensitive to the gas-phase C/O ratio of the disk. The observations required high C/O values, $>1$, to reproduce the observed line intensities \citep[see also][]{du2015}. Such prospects for measuring C/O in disks are now especially exciting as we enter an era where the elemental compositions, including C/O, of exoplanets' atmospheres \citep[e.g.,][]{madhusudhan2011,kreidberg2014,macintosh2015,bonnefoy2016,lavie2017}. In contrast to carbon and oxygen abundance estimates, there are few constraints on total nitrogen abundances or N/O ratios in disks owing in large part to the difficulty of observing the likely primary nitrogen carrier, N$_2$, \citep[e.g.,][and references therein]{schwarz2016}. Abundant nitrogen bearing species, such as N$_2$H$^+$ and HCN, are sensitive to other disk parameters than total N abundance, such as temperature structure, carbon abundance, and ionization rate. As such, interpreting these species in the context of bulk nitrogen abundance requires detailed knowledge of the source. In this work, we constrain the carbon, nitrogen, and oxygen content of the warm molecular layer in the IM Lup protoplanetary disk using results from our previous study of CO and its isotopologues \citep{cleeves2016im} and new and archival ALMA observations of C$_2$H and HCN and H$^{13}$CN. The solar mass star IM Lup harbors a massive gas rich disk, $M_{\rm disk} \sim0.1-0.2$~M$_{\odot}$ based upon continuum (SED and resolved millimeter images) and CO multi-isotopologue multi-line data presented in \citep{cleeves2016im}. The source is relatively young at an age of $0.5-1$~Myr \citep{mawet2012}. In \citet{cleeves2016im}, we found that IM Lup's CO is under-abundant by a factor of $\sim20$ compared to an interstellar CO abundance of $1.4\times10^{-4}$ per H based upon the dust-inferred disk mass. For comparison, this younger object appears to be ``missing'' less CO than the older TW Hya system, whose CO abundance is $\sim3-5\times$ less abundant than IM Lup \citep[e.g.,][]{favre2013}. However, based on this data alone it is difficult to tell whether the observed ``missing'' gas-phase CO is a result of missing carbon or missing oxygen or both; or, alternatively, missing total gas mass compared to dust. In the present paper, we explore what C/O and N/O abundance ratios are required in the disk's warm molecular layer to reproduce the observed C$_2$H, HCN, and HC$^{13}$N line intensities. | \label{sec:conclusion} Using detailed physical and chemical models constrained by CO and dust observations, along with C$_2$H and HCN data from ALMA, we constrain the C/O and N/O ratios in the molecular layer of the IM Lup disk. Our observations trace the properties of the disk where C$_2$H and HCN emit, primarily at normalized heights above $z/r\gtrsim0.2$. In these layers, the C$_2$H observations favor a super-solar elemental C/O ratio of $\sim0.8$. This high ratio is consistent with preferential loss of water ice from the surface, e.g., due to sequestration by an evolving and growing population of grains. We do not need to sequester any nitrogen for our best fit model, such that the N/O ratio is also super-solar, $\sim10$. The gas phase values of C/O and N/O vary spatially depending on the degree of UV shielding at a given location (Section~\ref{sec:cono}). While grain sequestration of ices tends to remove volatiles from the surface, it implicitly carries them to the midplane and into the inner disk through settling and radial drift \citep[e.g.][]{oberg2016,piso2016,krijt2016}. If this interpretation is correct, this process would result in a large enhancement over interstellar values of water ice in the inner disk midplane, a mild enhancement of carbon-bearing ice, and relatively little nitrogen ice transport in the solid phase. Correspondingly, we expect the C/O and N/O ratios in midplane solids to both be lower than solar, with N/O much less than solar in the disk midplane. Where settled ices eventually end up radially is still an open question. Radial drift of solids is thought to be quite efficient \citep[see review of][]{testi2014}, but this process may be slowed by the emergence of pressure variations in the disk that can effectively trap solids \citep[e.g.,][]{weidenschilling1980}. Now with ALMA, ringed radial structures are being observed, and may even be common \citep[e.g.,][]{alma2015,andrews2016,isella2016,loomis2017,huang2018,fedele2018}. In the absence of pressure traps, these grains should travel inward, thermally desorb, and enhance primarily oxygen, followed by carbon, and relatively little nitrogen. \citet{salyk2011} indeed found very low N/O ratios, $5\times10^{-4}$, in the inner disk with {\em Spitzer}; however, the nitrogen ``correction factors'' in the inner disk gas from HCN to total N are uncertain and require additional chemical modeling to constrain. These results have interesting consequences for the debate regarding abundance measurements in disks \citep[see summary of][]{bergin2017}. At the low gas temperatures typical of disks, H$_2$ does not emit appreciably \citep{bergin2013hd,mcclure2016,bergin2017}. As a result, other uncertain mass tracers are typically used, such as the total dust mass multiplied by a conversion factor, or even optically thin CO emission itself. From our modeling, we find that we do not need to deplete nitrogen significantly, with a factor of $4-20\times$ difference between the CO depletion factor and that for nitrogen. These results would point to abundance variations between these species rather than an overall under-accounting of disk mass, which would impact all volatile abundances at similar if not equal levels. We of course cannot rule out with these data alone some missing gas mass, since for the higher CR models, a small amount of nitrogen depletion (a factor of a few) is allowed, even though these are not the global best fits. However, missing gas alone cannot explain the higher degree of depletion needed for CO and water ice. In the future, additional observations of nitrogen-bearing molecules like HCN may help to break this CO mass / gas mass degeneracy, where in larger disk surveys CO masses appear globally low relative to the dust masses scaled by the interstellar gas-to-dust ratio \citep{ansdell2016,miotello2016,long2017}. Going forward, these results show that readily observable molecular tracers like CO, C$_2$H, and HCN, and their isotopologues combined with astrochemical models can be used to constrain the gas-phase C/O and N/O ratios within planet-forming disks. Such measurements may furthermore shed light on the inferred volatile composition of the ices, once we better understand the mechanism(s) of volatile loss from the warm molecular layer \citep[e.g.,][]{furuya2014,kama2016}. Future high sensitivity, resolved observations of many sources may further help shed light on ``typical'' gas-phase C/O and N/O ratios, how they spatially vary, and how these ratios may vary with time as planets are forming in the disk. Together, these can one day be compared with C/O measurements of exoplanet atmospheres, to eventually help unravel their formation locations and histories, where observations have already shown a wide range of exoplanet C/O values, even within a single planetary system \citep{bonnefoy2016,lavie2017}. | 18 | 8 | 1808.10682 |
1808 | 1808.07484_arXiv.txt | We present an open access grid of 3930 calculations of externally evaporating protoplanetary discs. This spans a range of disc sizes (1--400\,AU), disc masses, UV field strengths (10--10$^4$\,G$_0$) and stellar masses (0.05--1.9\,M$_\odot$). The grid is publicly available for download, and offers a means of cheaply including external photoevaporation in disc evolutionary calculations. It can also be queried using an online tool for quick estimates of instantaneous mass loss rates (e.g for convenient evaluation of real observed systems). \textsc{fried} itself illustrates that for discs around stars $\leq0.3$\,M$_\odot$ external photoevaporation is effective down to small radii ($<50$\,AU) down to UV fields at least as weak as 10\,G$_0$. At the other end of the scale, in a $10^4$\,G$_0$ environment photoevaporation is effective down to 1\,AU even for stellar masses at least as high as 1.9\,M$_\odot$. We also illustrate in which regimes CO survives in the photoevaporative outflow for significant mass loss rates; marking a system a good candidate to detect external photoevaporation in weak--intermediate UV environments through sub--Keplerian rotation. Finally we make illustrative mass loss rate estimates for discs in Taurus based on the \cite{2011A&A...529A.105G} star--disc parameters, finding that around half are expected to have both significant mass loss and retain CO in the photoevaporative outflow. | Planets are now known to exist around most stars (at least in the relatively local Milky Way) and exhibit a diverse range of architectures \citep{2015ARA&A..53..409W}. One of the key drivers in modern astrophysics is to understand the reason for this diversity, as well as how our own Solar system fits in to the wider population. To do so we must understand how the circumstellar ``protoplanetary'' discs of material around young stars not only ubiquitously give rise to planet formation, but also do so in a way that leads to high diversity in the resulting planetary parameters. To this end, substantial advances in our observational and theoretical capabilities have been made in recent years. Multi-wavelength observations from instruments like SPHERE \citep[e.g.][]{2017Msngr.169...32G}, ALMA \citep[e.g.][]{2015ApJ...808L...3A, 2017ApJ...851L..23C, 2018A&A...610A..24F} and the GPI \citep[e.g.][]{2041-8205-815-2-L26, 2041-8205-814-2-L27} are giving us the best observational insight yet into the inner workings of protoplanetary discs and planet formation \citep[for reviews and discussion on future advances see][]{2011ARA&A..49...67W, 2015PASP..127..961A, 2016PASA...33...59S}. Similarly, theoretical models are making rapid advances to capture the rich physics of planet-forming discs, which includes chemistry, magnetic fields, dust and gas dynamics and radiation transport \citep[for reviews and discussion of current and future advances in modelling of protoplanetary discs see][]{2016PASA...33...53H, 2016JGRE..121.1962M}. However, the problem is complicated further in that protoplanetary discs are found around young stars \citep[the discs are typically dispersed well before 10Myr, e.g.][]{2015A&A...576A..52R} and young stars are typically still in the clusters from which they formed. There are therefore \textit{environmental} factors that need to be accounted for. These are: one-off gravitational encounters, a binary (or tertiary, etc.) companion, and irradiation of a disc by other stellar members of the young cluster. There is growing evidence that gravitational encounters are generally of secondary importance to photoevaporation. For example \cite{2001MNRAS.325..449S} demonstrated this when comparing dynamical and radiative distruption of discs in an Orion like environment. Furthermore, \cite{2018MNRAS.475.2314W} and \cite{2018MNRAS.478.2700W} demonstrate that dynamical interactions are, statistically speaking, always a secondary effect in sculpting a disc population compared to photoevaporation, even in much weaker UV environments. Interactions have to be very close in order to see a significant effect on the disc evolution \citep{2018MNRAS.475.2314W}. Protoplanetary discs can more obviously be affected by a binary companion, potentially being significantly disrupted as in the case of RW Aurigae \citep[][]{2006A&A...452..897C, 2015MNRAS.449.1996D, 2018ApJ...859..150R} truncated \citep{1980ApJ...241..425G} or warped \citep{1983MNRAS.202.1181P, 2018MNRAS.473.4459F}. There is also evidence that this affects the resulting planetary populations \citep[e.g.][]{2002ApJ...568L.113Z, 2007A&A...462..345D}. The external photoevaporation of protoplanetary discs has been a challenging effect to gauge observationally. For many years, only the proplyds -- discs within close proximity of O stars -- were obviously observed to be photoevaporating \citep[e.g.][]{1996AJ....111.1977M, 1998AJ....115..263O, 2001AJ....122.2662O, 2002ApJ...566..315H}. These are typically irradiated by a UV field of order $10^5$\,G$_0$\footnote{G$_0$ is the Habing unit of UV radiation, which is $1.6\times 10^{-3}$ erg cm$^{-2}$ s$^{-1}$ over the wavelength range ($912$\AA$<\lambda<2400$\AA) \cite{1968BAN....19..421H}}. The effects of photoevaporation in the vicinity of O stars have also been studied recently by, for example, \cite{2017AJ....153..240A}, \cite{2018ApJ...860...77E} using the spatial distribution of disc properties. Some direct measures of the mass loss rate from discs near O stars have been made, for example by \cite{1987ApJ...321..516C}, \cite{1999AJ....118.2350H} and \cite{2002ApJ...566..315H}, finding mass loss rates of order $10^{-6}$\,M$_\odot$\,yr$^{-1}$. In recent years there has been growing evidence of external photoevaporation in weaker radiation environments. \cite{2016ApJ...826L..15K} found proplyds in a $\sim3\times10^3$\,G$_0$ environment and \cite{2017MNRAS.468L.108H} used numerical models to propose external photoevaporation as the reason for the large CO halo around IM Lup, which is only in a $\sim4$\,G$_0$ environment \citep{2016ApJ...832..110C}. Nevertheless, external photoevaporation in weak--intermediate radiation environments is generally unconstrained, in part because it has not yet been actively searched for since the signatures of external photoevaporation aren't well known. As most stars are not in such a strong UV environment as the proplyds \citep{2008ApJ...675.1361F} understanding the evolution of discs in weaker environments is important. Modelling the external photoevaporation of discs is difficult because the thermodynamic properties of the flow are set by photodissociation physics for the UV fields that the majority of star--discs are exposed to \citep[in the limit of being very close to a strong EUV source, photoionisation dominates, see Figure 12 of ][]{2018MNRAS.478.2700W}. Computing the thermal structure of a photodissociation region (PDR) requires the solution of a chemical network that is also sensitive to the non-local distribution of matter. That is, the temperature at one point in the flow is sensitive to the rest of the flow structure because this sets the cooling by the escape of line photons and also the attenuation of the UV field by dust/molecular self-shielding in outward lying regions of the photoevaporative flow. For this reason, for a long time only semi-analytic models of the flow structure and hence mass loss rate could be produced \citep[e.g.][]{1994ApJ...428..654H, 1998ApJ...499..758J, 2004ApJ...611..360A, 2016MNRAS.457.3593F}. These are quick to compute but are only able to obtain solutions in certain subsets of parameter space. To date they also all only consider 1\,M$_\odot$ stars. Computing solutions for arbitrary parts of the parameter space requires full photochemical-dynamical models that iteratively solve the PDR chemistry/temperature with the dynamics. This is both difficult to implement and computationally expensive, but has now been achieved by \cite{2016MNRAS.463.3616H} using the \textsc{torus-3dpdr} code (discussed in section \ref{sec:num_meth}). Because these mass loss rates are difficult to compute, and expensive, having a large grid of publicly available pre-computed models would therefore open up consideration of external photoevaporation to the wider community. The value of pre-computed mass loss rates from models such as the above is that they can either be used to estimate the instantaneous mass loss rate for real systems, or can be applied to viscous evolutionary models of discs \citep[][]{2007MNRAS.376.1350C, 2013ApJ...774....9A, 2015ApJ...815..112K, 2017MNRAS.468L.108H, 2018MNRAS.475.5460H, 2018MNRAS.478.2700W}. In this paper we present the results of a large grid of external photoevaporation models as described above, which we refer to as the \textsc{fried} (\textbf{F}UV \textbf{R}adiation \textbf{I}nduced \textbf{E}vaporation of \textbf{D}iscs) grid. This covers a wide parameter space of stellar mass, disc mass, disc radius and UV field. It is publicly available for direct download, but we also provide an online tool for making quick mass loss rate estimates. The rest of this paper is as follows. In section \ref{sec:construction} we discuss how the \textsc{fried} grid is constructed, in section \ref{sec:theFriedGrid} we provide an overview of the resulting grid and the online resources. Finally in section \ref{sec:discussion} we apply the grid to an illustrative population of discs. | \label{sec:discussion} \begin{figure*} \centering \includegraphics[width=19cm]{./CO5Transition_diffmass_correctedGrab.pdf} \caption{The ratio of the extent out to which CO is abundant to the extent of the Keplerian disc, as a function of disc size and the ratio of disc--to--stellar mass. The dashed contours denote the point at which CO is depleted at the disc outer edge. Models with mass loss rates less than $10^{-9}$\,M$_\odot$\,yr$^{-1}$ are not included and the region below the solid black line is not included in our parameter space. The left and right hand panels are for UV fields of 10 and $10^3$\,G$_0$ respectively and the stellar mass 0.1, 0.5, 1 and 1.6\,$M_\odot$ from top to bottom. If CO extends beyond the disc outer edge, then it is possible it could be used to detect the photoevaporative wind, for example through a sub-Keplerian rotation profile or a radially increasing temperature measured with CO line ratios. {Note that since the grid extent is typically 1000\,AU, some of the ratios presented here may be lower limits if the grid size is what limits the CO extent. }} \label{fig:COextents} \end{figure*} \subsection{Constraints on detecting external photoevaporation using CO in weak--intermediate UV regimes} The PDR thermal physics in our models also contains information on the chemical structure of the flow which could provide clues as to the best scenarios in which to search for external photoevaporation in action. CO observations towards discs are common, but typically focus on rings and gaps in the main body of the disc, rather than the outer disc where external photoevaporation is expected to be strongest. Two possible signatures of external photoevaporation in CO are sub-Keplerian rotation in the outer disc \citep{2016MNRAS.457.3593F, 2016MNRAS.463.3616H} and a radially increasing temperature profile. Sub-Keplerian rotation in the candidate evaporating disc IM Lup \citep{2016ApJ...832..110C, 2017MNRAS.468L.108H} has been detected by \cite{2018A&A...609A..47P}. A radially increasing temperature profile could be probed using CO line ratios. Given the above and current lack of other diagnostics, of specific interest are regions of the parameter space where a kinematic and thermal tracer such as CO is observable and the mass loss rate is non-negligible (and, ideally, the deviation from Keplerian rotation is spectrally resolvable). {We therefore compute the ratio of the radius at which CO is reduced in abundance ($R_{\textrm{CO}}$) to the disc outer radius ($R_{\textrm{disc}}$). This transition happens quickly, so we mark the point at which the CO abundance drops to less than $10^{-5}$. Note that the role of different reactions as well as cosmic rays and photodissociation on setting the CO abundance are well known and the reader is directed to \cite{1986ApJS...62..109V} and \cite{1988ApJ...334..771V} for more information.} Values of $R_{\textrm{CO}}/R_{\textrm{disc}}>1$ therefore have CO abundant in the photoevaporative wind. Of these models, we then only retain those in which the mass loss rate is greater than $10^{-9}$\,M$_{\odot}$\,yr$^{-1}$. This leaves us with a grid of models that summarise the best regimes in which it might be possible to probe a photoevaporative wind using CO. Figure \ref{fig:COextents} summarises this dataset for a 0.1, 0.5, 1 and 1.6\,M$_\odot$ stars in 10 and $10^3$\,G$_0$ environments. In the low UV case if the stellar mass is low then CO is abundant in the flow and there is still significant mass loss down to small disc outer radii ($<50$\,AU). However at higher stellar masses there is a lower limit on the radius of the disc at which the mass loss rate is actually significant. At higher UV field strengths in the low stellar mass scenario CO is depleted in the photoevaporative wind except for very compact discs. Conversely for higher stellar masses CO is abundant \textit{and} has significant mass loss down to smaller disc outer radii than in the 10\,G$_0$ case. Discs around $1-2$\,M$_\odot$ stars in $\sim10^3$\,G$_0$ environments therefore seem to provide a good opportunity to identify external photoevaporation in action. Another interesting point to note is that the typical factor 2 or more extent of CO relative to the Keplerian disc outer edge is large enough to account for the observed extent of the gas relative to dust in Lupus discs by \cite{2018ApJ...859...21A}. Since only small grains are entrained in a photoevaporative wind \citep{2016MNRAS.457.3593F} and hence will not be detected in the millimetre continuum this raises the possibility that external photoevaporation might be a contributing factor in setting the observed relative gas--dust disc sizes. However this has the caveat that we cannot robustly constrain the surface brightness in CO from 1D models so cannot currently predict the observable CO extent, just that out to which it is abundant. \subsection{Illustrative instantaneous mass loss rate estimates for real systems: Taurus discs} \label{sec:illustrativeReal} The \textsc{fried} grid makes it trivial to estimate mass loss rates for real systems based on estimates of the stellar mass, disc mass, disc radius and incident UV field strength. In cases where not all of these are known, the grid can be used to assess ranges of plausible values. \begin{table*} \caption{Illustrative mass loss rates for real systems for a range of UV field strengths. All star-disc parameters are from Guilloteau et al. (2011). } \label{tab:realApp} \begin{tabular}{c c c c c c c c c } \hline System & $M_* (M_{\odot})$ & $M_d (10^{-3}\,M_\odot)$ & $R_d$ (AU) & $\dot{M} (M_\odot\,\textrm{yr}^{-1})$ & CO in flow?& High/Medium/Low \\ & & & & $10G_0$ & (Figure 1) & Mass loss rate \\ \hline BP Tau & 0.78 & 5.4 & 57. & $10^{-10}$ & -- & L \\ CI Tau & 0.76 & 37 & 201. & $4.6\times10^{-7}$ & Y &H \\ CQ Tau & 1.7 & 6.3 & 188.0 & $2.6\times10^{-10}$ & -- & L \\ CY Tau & 0.48 & 16.5 & 92. & $9.9\times10^{-8}$ & Y & H \\ DG Tau & 0.7 & 36.0 & 198. & $5.2\times10^{-7}$ & Y &H \\ DL Tau & 0.7 & 49.0 & 179. & $1.8\times10^{-7}$ & Y &H \\ DM Tau & 0.47 & 31.1 & 274. & $2.8\times10^{-6}$ & Y &H \\ DQ Tau & 0.55 & 12.1 & 439. & $3.6\times10^{-7}$ & -- & H \\ GM Aur & 1.37 & 27.0 & 578. & $8.9\times10^{-7}$ & -- &H \\ Lk Ca 15 & 1.12 & 28.4 & 178. & $2.3\times10^{-8}$ & Y &H \\ MWC480 & 1.8 & 182.3 & 155. & ? & -- &? \\ MWC758 & 1.8 & 10.6 & 187. & $2.3\times10^{-9}$ & -- &M \\ HL Tau & 0.7 & 90.6 & 280.16 & $3.7\times10^{-6}$ & Y &H \\ HH 30 & 0.25 & 8.1 & 123. & $7.5\times10^{-7}$ & Y & H \\ DG Tau B & 3.0 & 67.9 & 303. & ?& -- & ? \\ T Tau N & 1.9 & 0.1 & 67. &$10^{-10}$ & -- & L \\ Haro6-13 & 0.55 & 0.6 & 90. & $4.9\times10^{-10}$ & -- & L \\ Haro6-33 & 0.55 & 0.5 & 439. & $2.6\times10^{-8}$ & -- & H \\ \hline \end{tabular} \end{table*} \label{lastpage} In Table \ref{tab:realApp} we calculate the mass loss rate for the dust disc extents and masses in Taurus inferred by \cite{2011A&A...529A.105G}. PDR modelling of Taurus requires an average UV field of around 10\,G$_0$ according to \cite{2013MNRAS.429.3584H} to give consistent H\,\textsc{i} and CO observations, so we adopt this UV field strength for our illustrative assessment. For discs larger than 400\,AU (the upper limit on our grid) we set the disc radius to 400\,AU, which is, if anything, an underestimate of the mass loss rate (see Figure 1). We do not similarly limit the disc mass. These estimates are based on dust disc radii, so given the gas is generally found to be more extended \citep[e.g.][]{2013A&A...557A.133D, 2017A&A...605A..16F} and mass loss rates are higher for larger discs it is quite a conservative estimate. In the sixth column of Table \ref{tab:realApp} we note whether each disc has CO in the flow from Figure \ref{fig:COextents}. In the final column of Table \ref{tab:realApp} we denote any system in which the mass loss rate is above $10^{-8}$\,M$_\odot$\,yr$^{-1}$ a high mass loss rate (H), otherwise if the mass loss rate is above $10^{-9}$\,M$_\odot$\,yr$^{-1}$ we denote it ``medium'' (M), else the mass loss rate is low (L). In two cases, MWC\,480 and DG Tau B, the disc mass is extremely high, so the \textsc{fried} grid is unable to compute a reliable mass loss rate estimate. At least half of the discs in our illustrative application have significant mass loss rates ($>10^{-8}$\,M$_\odot$\,yr$^{-1}$) in a 10\,G$_0$ environment. Of these discs with high mass loss rates there are 8 objects in which CO is predicted to exist in the flow over a significant radial range and therefore make possible targets for probing subtle sub-Keplerian rotation and/or radially increasing temperature profile in the outer disc (there may be other signatures of external photoevaporation that are yet to be identified). | 18 | 8 | 1808.07484 |
1808 | 1808.09450_arXiv.txt | One of the most important problems in the context of cataclysmic variables (CVs) is the lack of observations of systems with periods between 2 and 3.12 hours, known as the period gap. The orbital evolution of CVs with periods shorter than those in the gap is dominated by gravitational radiation while for periods exceeding those of the gap it is dominated by magnetic braking of the secondary star. Spruit \& Ritter (1983) showed that as periods approach 3 hours and secondary stars become fully convective a sharp decline in magnetic dynamo and braking efficiency would result in such a gap. Recent X-ray observations finding coronal magnetic energy dissipation is similar in fully convective and partly radiative M dwarfs cast this theory into doubt. In this work, we use Zeeman-Doppler imaging observations culled from the literature to show that the complexity of the surface magnetic fields of rapidly rotating M dwarfs increases with decreasing rotation period. Garraffo et al.\ (2018) have shown that the efficiency of angular momentum loss of cool stars declines strongly with increasing complexity of their surface magnetic field. We explore the idea of \citet{Taam.Spruit:89} that magnetic complexity might then explain the period gap. By generating synthetic CV populations using a schematic CV evolutionary approach, we show that the CV period gap can naturally arise as a consequence of a rise in secondary star magnetic complexity near the long period edge of the gap that renders a sharp decline in their angular momentum loss rate. | \label{s:intro} One of the most challenging puzzles in stellar evolution that emerged in the 1970s and early 1980s was the {\it cataclysmic variable period gap}. Cataclysmic variables (CVs) are interacting binary stars comprising a white dwarf accreting from either a main-sequence or slightly evolved star, or from a brown dwarf. As the population of known CVs grew, it became clear that systems with periods between 2 and 3.12 hours were rarely observed compared with objects with shorter and longer periods \citep{Robinson:83,Ritter:84, Knigge.etal:11}. While some of the underlying physics giving rise to this period gap are still debated, current models invoke changing rates of angular momentum loss as secondary stars are whittled down to lower and lower masses by attritional accretion onto their compact companion. The orbital evolution of CVs with periods shorter than those in the gap is dominated by gravitational radiation \citep{Faulkner:71, Paczynski.Sienkiewicz:81}, while for periods exceeding those of the gap it is dominated by magnetic braking \citep{Eggleton:76, Verbunt.Zwaan:81}. In the latter regime, systems lose angular momentum via the magnetized winds of the non-degenerate companion and, as a consequence, their orbital separation is reduced and they spin up. With mass from the secondary star being lost to the primary, the secondary star drifts to later and later spectral type. By the time the system reaches the upper boundary of the period gap ($\sim 3.12$ hours), the secondary star has been reduced to the mass of a fully-convective M dwarf (\citealt{Robinson.etal:81,Spruit.Ritter:83}, and references therein). \cite{Spruit.Ritter:83} and \citet{Rappaport.etal:83} showed that a fast decrease in angular momentum loss of $\sim 90\%$ as the secondary approaches the fully-convective limit would result in strong suppression of the mass accretion from the secondary star onto its companion that would explain the appearance of the period gap. A detailed review of this evolutionary scenario has been provided by \citet{Knigge.etal:11}. Typically, CVs are discovered through UV or X-ray emission from heated accreting material: if the accretion stops the systems cannot be easily discerned. The underlying motivation for a fairly abrupt angular momentum loss reduction at the fully convective limit stems from arguments that magnetic dynamo action in Sun-like stars occurs at the interface between the convection zone and the radiative interior---the ``tachocline''; \citep[see, e.g.,][]{Spruit.Ritter:83}. The fact that the secondary star becomes fully convective near the edge of the period gap led to the idea that there is a fundamental change in the effectiveness of the dynamo at this limit. This supposed demise of the dynamo thus results in a magnetic braking ``disruption''. However, evidence for such a dynamo demise has historically been weak or lacking. Recently, \cite{Wright.Drake:16} have found that the X-ray emission level as a function of stellar rotation period is essentially the same for both fully convective, and partly convective and more Sun-like stars for both rapid and slow rotators. X-ray activity has been shown to be a good proxy for surface magnetic flux \citep[e.g.][]{Pevtsov.etal:03}. \citet{Wright.Drake:16} then argued that the invariance of X-ray behavior regardless of the presence or absence of a radiative zone supports the arguments advanced by \citet{Spruit:11} that magnetic dynamo action is instead distributed in the convection zone. The collateral effect of the \citet{Wright.Drake:16} results is that, at face value, there is no change in the surface magnetic activity that drives magnetic braking: the {\it disrupted magnetic braking} theory for the CV period gap is broken. \citet{Taam.Spruit:89} had anticipated this and suggested that the magnetic field of the secondary star might grow in complexity at the edge of the period gap, reducing the number of open field lines and the subsequent angular momentum loss. Recently, it has been shown using different types of magnetohydrodynamic (MHD) wind models that the efficiency of angular momentum loss of cool stars strongly depends on the complexity of their stellar surface magnetic fields \citep[][from hereon CG16]{Reville.etal:15a, Garraffo.etal:15, Garraffo.etal:16a}, echoing the earlier analytical conclusions of \citet{Taam.Spruit:89}. \citet[][hereafter CG18]{Garraffo.etal:18} provided a new predictive spin-down model based on sophistcated MHD wind modeling that accounts for this magnetic modulation, and assumes that the complexity of the magnetic field is a function of Rossby number {\it Ro} ({\it Ro} $ = P_{rot}/\tau$, where $\tau$ is convective turnover time). This assumption is well supported by Zeeman-Doppler imaging (ZDI) observations of the magnetic fields on the surfaces of Sun-like stars showing that faster rotating stars store a larger fraction of their magnetic flux in higher order multipole components of the field \citep[e.g.,][]{Donati:03, Donati.Landstreet:09, Marsden.etal:11a, Waite.etal:11, Waite.etal:15}. As a consequence, they lose angular momentum much less efficiently. Magnetic braking of CVs is usually modeled considering the secondary star as a single star and assuming it has a simple, fixed magnetic configuration such as a dipole. In this work, we study the available ZDI observations of late M-dwarfs by \cite{Morin.etal:10} to infer the underlying geometry of their magnetic fields as a function of rotation period. We follow the idea of \citet{Taam.Spruit:89} and explore the possibility that there is an disruption of angular momentum loss near the upper boundary of the period gap ($\sim 3.12$) resulting from an increasing magnetic complexity of the secondary star. % Using the observed magnetic geometry of M Dwarfs as a function of rotation period derived here and the spin-down model provided in CG18, we use the schematic method of \citet{Spruit.Ritter:83} to reconstruct the orbital evolution of CVs near the gap period. Using this evolutionary recipe and the expected formation rate of CVs, we generate synthetic populations and predict the fraction of systems and their visibility as a function of orbital period. The paper is organized as follows. In Section~\ref{mdwarfs} we compute the complexity of the available ZDI observations for late M Dwarfs and compare our data to the complexity function assumed in CG18. In Section~\ref{aml} we use that function and the prescription for angular momentum loss from CG18 to describe a single system's orbital evolution. In Section~\ref{ps} we generate and evolve synthetic populations and compare them with observations. In Section~\ref{results} we discuss our findings and summarize our main conclusions. | \label{results} We have found that existing ZDI observations of late M dwarfs support a picture of increasing surface magnetic field complexity with decreasing rotation period for values of the period of several hours. This is consistent with the stellar spin-down model presented by CG18 that explains the bimodal rotation period distributions of stars in young open clusters in terms of evolving surface magnetic complexity. Greater magnetic complexity leads to suppression of mass loss and a shortening of the magnetic ``lever" that acts as the rotational brake in late-type stars. Consequently, as CVs evolve toward shorter periods they experience a reduction in angular momentum loss rate and a reduction in the accretion rate driven by magnetic braking, as conjectured by \citet{Taam.Spruit:89}. We have modeled a synthetic population of CVs using the standard CV evolution equations \citep{Spruit.Ritter:83,Knigge.etal:11} together with the magnetic braking prescription provided in CG18. As periods approach 3.2 hours, the reduction in angular momentum loss is so rapid and efficient ($\sim 90 \%$) that the accretion rate in most systems drops sufficiently to allow the puffed-up donor star to shrink back into thermal equilibrium. These are just the conditions that \citet{Spruit.Ritter:83} pointed out would produce the CV period gap and that \citet{Taam.Spruit:89} found could arise from an increase in surface magnetic complexity. The secondary star no longer fills its Roche lobe and mass accretion stops, rendering it observationally inconspicuous. However, the system continues to lose angular momentum through magnetic braking at this slower rate and eventually (at $P \approx 2$ hours) the orbital separation decreases enough for accretion to resume and consequently for the system to become conspicuous again. Our model predicts the presence of a few systems accreting within the gap, consistent with observations. The explanation of the period gap in terms of an increase in magnetic complexity of the secondary as systems approach the gap as first suggested by \citet{Taam.Spruit:89}, rather than the dynamo itself shutting down, is fully consistent with X-ray observations indicating that magnetic field generation is equally efficient above and below the M dwarf fully-convective limit. The disrupted magnetic braking idea of \citet{Spruit.Ritter:83} and \citet{Rappaport.etal:83} is not broken. | 18 | 8 | 1808.09450 |
1808 | 1808.08276_arXiv.txt | We present recalibrations of the {\tt GALFORM} semi-analytical model of galaxy formation in a new N-body simulation with the Planck cosmology. The Planck Millennium simulation uses more than 128 billion particles to resolve the matter distribution in a cube of $800$ Mpc on a side, which contains more than 77 million dark matter haloes with mass greater than $2.12 \times 10^{9} h^{-1} {\rm M_{\odot}}$ at the present day. Only minor changes to a very small number of model parameters are required in the recalibration. We present predictions for the atomic hydrogen content (HI) of dark matter halos, which is a key input into the calculation of the HI intensity mapping signal expected from the large-scale structure of the Universe. We find that the HI mass $-$ halo mass relation displays a clear break at the halo mass above which AGN heating suppresses gas cooling, $\approx 3 \times 10^{11} h^{-1} M_{\rm \odot}$. Below this halo mass, the HI content of haloes is dominated by the central galaxy; above this mass it is the combined HI content of satellites that prevails. We find that the HI mass - halo mass relation changes little with redshift up to $z=3$. The bias of HI sources shows a scale dependence that gets more pronounced with increasing redshift. | Measuring fluctuations in the intensity of $21\,$cm line emission offers a novel way to map the large-scale structure of the Universe that is competitive with the largest planned optical galaxy redshift surveys \citep{Bull:2015}. The $21\,$cm line is a forbidden transition between hyperfine structure in the ground state of atomic hydrogen. As a consequence, redshift surveys of galaxies detected through their weak HI emission at best currently contain thousands rather than the hundreds of thousands or even millions of galaxies reached by optically selected surveys (for a review of extragalactic HI astronomy see \citealt{Giovanelli:2015}). The next generation of surveys, such as the Widefield ASKAP L-band Legacy All-sky Blind surveY, WALLABY, the Australian Square Kilometer Array Pathfinder, will measure HI for over half a million galaxies \citep{Johnston:2008,Duffy:2012}. However, such HI surveys will still be limited to the local Universe. A solution to this problem is to exploit the finite angular and frequency resolution of radio telescopes, to effectively stack the emission from many galaxies in a single pointing and hence boost the HI signal to a measurable level. Measuring the intensity of HI emission from all the sources within some volume bypasses the challenge of detecting the emission from individual sources, allowing a view of the large-scale structure of the Universe to be obtained that is smoothed on small scales \citep{Battye:2004,Zaldarriaga:2004, Peterson:2006,Pritchard:2012, KOvetz:2017}. The power spectrum of HI intensity fluctuations has already been measured in spite of the cosmological signal being much smaller than the galactic foreground \citep{Switzer:2013}. The fact that this measurement was made at a much higher redshift ($z \sim 0.8$) than that for which estimates of the HI mass function are available by the measurement of emission from single galaxies ($z \sim 0.05$) illustrates the potential of intensity mapping. Encouraged by this, a number of HI intensity mapping experiments are either under construction or proposed (e.g. BAO from Integrated Neutral Gas Observations (BINGO) \citealt{Battye:2016}; the Canadian Hydrogen Intensity Mapping Experiment (CHIME) Pathfinder \citealt{Bandura:2014}, the Five-hundred-meter Aperture Spherical radio Telescope (FAST) \citealt{Bigot-Sazy:2016}, MeerKAT \citealt{Pourtsidou:2017} and the Square Kilometre Array (SKA) \citealt{Santos:2015}). A key input into the prediction for HI intensity fluctuations is the HI content of dark matter halos, combining the contribution from the central galaxy with that of all of the satellite galaxies in the halo. Many empirical approaches have been proposed to describe the HI content of dark matter haloes including: 1) simple scalings with halo mass \citep{Santos:2015}, 2) arguments based on the effective circular velocity of halos that contain HI, limited at low circular velocities by photo-ionization heating of the intergalactic medium and at high velocities as a result of the central galaxy tending, with increasing halo mass, to become bulge rather than disk dominated, and hence gas poor (\citealt{Barnes:2009, Bagla:2010}), 3) more sophisticated scalings with halo mass, with a broken power law and variable amplitude, constrained to fit various observations of the abundance and clustering of HI galaxies \citep{Padmanabhan:2015,Pad:2016,Padmanabhan:2017a,Padmanabhan:2017b,Obuljen:2018,Paul:2018}. Several studies using hydrodynamical simulations have yielded predictions for the HI galaxy mass $-$ halo mass relation by post-processing the simulation results to divide the cold gas mass of a galaxy into atomic and molecular components\footnote{By `atomic' and `molecular' gas we mean phases of the interstellar medium in which the hydrogen is preferentially in atomic (HI) or molecular (H$_{2}$) form, respectively.} \citep{Dave:2013,Villa:2014,Crain:2017,Villaescusa:2018,Ando:2018}. Here we use a different approach to predict the form of the HI mass $-$ halo mass relation and its evolution: a physically motivated model of galaxy formation in which the atomic and molecular gas contents of model galaxies are tracked at all times. Semi-analytical models calculate the transfer of baryons between different reservoirs within dark matter halos that are growing hierarchically \citep{Baugh:2006, Benson:2010,SomervilleDave2015}. Comparisons between the predictions made by models developed by different groups, which are publicly available, reveal that the models have reached a level of maturity such that they can give robust predictions for the baryonic content of dark matter haloes and the clustering of samples defined by different properties, such as stellar mass, cold gas mass and star formation rate\footnote{This is true for samples defined by galaxy number density, after ranking the galaxies by the value of a property such as stellar mass or cold gas mass. The models differ in the observations used to set their parameters, and so some distribution functions will agree between models (such as the stellar mass function) better than others (such as the cold gas mass function).} \citep{Contreras:2013,Guo:2016,Pujol:2017,Lagos:2018}. \cite{Contreras:2015} examined the dependence of galaxy properties on the mass of the host halo (or the mass of the host subhalo at infall for the case of satellites). This study revealed that some properties, such as stellar mass, display a simple dependence on host halo mass, albeit with considerable scatter (see also the review by \citealt{Wechsler:2018}). Other properties, however, such as cold gas mass and star formation rate, are predicted to have a complicated dependence on halo mass. Hence we cannot simply translate trends uncovered by the analysis of stellar mass selected samples and assume that these hold for HI-selected samples; a physical model is needed to connect the HI mass of a galaxy to its host halo mass. Here, we use the {\tt GALFORM} semi-analytical galaxy formation model \citep{Cole:2000, Baugh:2005, Bower:2006,GonzalezPerez:2014,Lacey:2016}. The treatment of star formation in {\tt GALFORM} was extended by \cite{Lagos:2011} to model the atomic and molecular hydrogen contents of galaxies, rather than just the total cold gas mass that was considered in earlier versions of the model (note that some other semi-analytical codes now also have this capability: \citealt{Fu:2010,Somerville:2015,Stevens:2017,Xie:2017,Lagos:2018}). A comprehensive overview of {\tt GALFORM} and the way in which the model predictions respond to varying the galaxy formation parameters can be found in \cite{Lacey:2016}. \cite{Kim:2015} used the {\tt GALFORM} model of \cite{Lagos:2012} to investigate the physics behind the low mass end of the HI mass function, arguing that the photoionization heating of the intergalactic medium was the key process shaping this prediction. The same model was used by \cite{Kim:2017} to predict the HI content of halos and the HI intensity fluctuation power spectrum. Here we update the background cosmological model used in {\tt GALFORM} with a new N-body simulation, the Planck Millennium. We recalibrate the versions of {\tt GALFORM} calibrated using halo merger trees extracted from an N-body simulation run with the WMAP-7 cosmological parameters \citep{Guo:2013}. In particular, we consider the models introduced by \cite{Lacey:2016} and \cite{GonzalezPerez:2014} (see also \citealt{Guo:2016,GP:2018}). We recalibrate the galaxy formation parameters which define these models to reflect the change in the cosmological parameters, the improved mass and time resolution of the Planck Millennium simulation outputs and the incorporation of an improved treatment of galaxy mergers (\citealt{Simha:2017}, see also \citealt{Campbell:2015}). We then use the recalibrated models to predict the HI content of dark matter halos and its evolution. The recalibrated models described here have been used by \cite{Cowley:2018} to make predictions for the counts and redshift distributions of galaxies that will be seen by the James Webb Space Telescope and by \cite{Nuala:2017} to test the assumptions behind halo occupation distribution modelling of galaxy clustering. The Planck Millennium simulation has also been used by \cite{Veena:2018} to study the spin and shape alignments of haloes in the cosmic web. The layout of this paper is as follows. In Section 2, we introduce the galaxy formation model, describing the Planck Millennium N-body simulation (\S~2.1), giving an outline of {\tt GALFORM} (\S~2.2), explaining the differences and similarities of the two variants of {\tt GALFORM} considered (\S~2.3) and closing with the recalibration of these models (\S~2.4). The predictions for the HI content of dark matter halos are given in \S~3, along with comparisons to previous results, and our conclusions are given in \S~4. Some other predictions relating to intensity mapping using tracers other than HI are presented in Appendix~\ref{sec:IM}, and a variant of the recalibrated models with gradual ram pressure stripping of the hot gas in satellite galaxy halos is discussed in Appendix~\ref{sec:ram}. Note that throughout we quote masses in units of $h^{-1} \, {\rm M_{\odot}}$ and lengths in units of $h^{-1} {\rm Mpc}$, retaining the reduced Hubble parameter, $h$, where $H_{0} = 100 \, h \, {\rm km} \, {\rm s}^{-1} \,{\rm Mpc}^{-1}$. | We have presented implementations of the {\tt GALFORM} models introduced by \cite{GonzalezPerez:2014} and \cite{Lacey:2016} in a new, high resolution N-body simulation, the P-Millennium or PMILL run. The models required a minor recalibration due to the improved mass resolution of the halo merger trees extracted from the PMILL, compared to those available from the simulation originally used to calibrate the models. The change in cosmology from a WMAP-7 model to a Planck cosmology does not require a significant change in the model parameters. We also took this opportunity to update the treatment of galaxy mergers in {\tt GALFORM}, using the model introduced by \cite{Simha:2017}. In the end, only minor changes were required to two model parameters in each case to obtain a similar level of agreement with the observational data used to set the model parameters. One clear application of the improved halo mass resolution in the P-Millennium is to make predictions for the atomic hydrogen content of dark matter haloes. Observational determinations of the HI mass function are in their infancy and show significant disagreement at high masses. Nevertheless, current estimates do agree with one another around the break in the mass function and suggest that the global HI density at the present day is dominated by galaxies with HI masses $\sim 10^{9.5} h^{-2} {\rm M_{\odot}}$. Our model predictions show that these are central galaxies in halos with mass $\approx 10^{11.5} h^{-1} {\rm M_{\odot}}$, which is approximately the halo mass above which the suppression of gas cooling by AGN heating becomes important in the models. In the PMILL, such haloes are resolved by $\sim 3000$ particles and have reliable merger histories. A calculation made using the same galaxy formation parameters but with dark matter halo merger trees restricted to the same mass resolution as the WM7 simulation resolves around half of the global HI mass recovered at the PMILL resolution. There have been a large number of recent studies that have proposed empirical parametric forms for the HI mass $-$ halo mass relation, which is a key input for HI intensity mapping predictions (\citealt{Santos:2015, Padmanabhan:2015,Padmanabhan:2016,Padmanabhan:2017a,Padmanabhan:2017b}). Our ab initio predictions show clear differences from the results of those studies. We find a sharp break in the HI mass $-$ halo mass relation above the halo mass for which AGN heating stops gas cooling onto central galaxies. Also, the form of the predicted relation shows remarkably little dependence on redshift over the interval $z=0-3$. The break in the HI mass $-$ halo mass relation marks a shift from central galaxies dominating the HI content of low mass halos, to the combined satellite population becoming more important in high mass halos. The depth of the dip at the break is reduced somewhat if the gradual ram pressure stripping of the hot gas halos of satellite galaxies is allowed. It is not straightforward to compare the predictions of our model to others in the literature, as most calculations follow the total cold gas mass rather than considering the atomic and molecular hydrogen contents of galaxies separately. For example, \cite{Martindale:2017} calculated the HI mass of galaxies in the {\tt L-GALAXIES} model in post-processing and used the HI mass function as a constraint on the model parameters. This resulted in an improved fit to the low mass end of the HI mass function, compared to the prediction of the \cite{Henriques:2015} model. As demonstrated by \cite{Lagos:2011b}, however, post-processing model predictions to compute the HI mass of galaxies can give very different predictions to self-consistently changing the star formation law and computing the evolution of the atomic and molecular hydrogen contents of galaxies. A small number of models do track the atomic and molecular hydrogen contents of galaxies self consistently (\citealt{Fu:2010, Lagos:2011, Lagos:2012, Popping:2014, Xie:2017,Stevens:2017,Lagos:2018}). The strength of the break in the HI mass $-$ halo mass relation could be sensitive to the way in which different processes are modelled in {\tt GALFORM}, such as AGN feedback or the cooling of gas in satellite galaxies. The treatment of gas cooling in satellites and its impact on the model predictions is discussed in Appendix~\ref{sec:ram}. Regarding the modelling of AGN feedback in {\tt GALFORM}, once a halo satisfies the conditions for AGN heating to affect gas cooling (see Eqns.~\ref{eq:alpha_cool} and \ref{eq:f_Edd}), the cooling flow is turned off completely. In the {\tt L-GALAXIES} semi-analytical model, for example, the suppression of cooling sets in more gradually (\citealt{Croton:2006,Henriques:2017}). \cite{Zoldan:2017} compared the HI mass $-$ halo mass relations predicted by a range of semi-analytical models, including an earlier version of the {\tt GALFORM} model by \cite{Bower:2006}. Zoldan et~al. made a similar plot to our Fig.~\ref{fig:himhalo1}, but did not go on to examine the total gas content of dark matter haloes. The comparison of Zoldan et~al. shows that, out of the models considered, AGN feedback is most efficient at stopping gas cooling in the Bower et~al. model. Nevertheless, the typical satellite masses are very similar between models, suggesting a total HI mass $-$ halo mass relation that would be similar to the one presented here. Recently, hydrodynamic simulations of cosmologically representative volumes have been able to reproduce the observed stellar mass function and other observables \citep{Vogelsberger:2014, Schaye:2015}. \cite{Crain:2017} present predictions for the HI content of galaxies in the EAGLE simulation of \cite{Schaye:2015}. The HI masses are calculated in post-processing. Crain et~al. state that the predictions for the HI mass function are poorly converged, with a substantial change on improving the mass resolution. Curiously, the fiducial EAGLE run does reproduce the cosmic abundance of HI with redshift reasonably well, despite not matching the present day HI mass function \citep{Rahmati:2015}. This is an observable that semi-analytical models tend to struggle to match at $z>0$ \citep{Popping:2014, Crighton:2015}. \cite{Villaescusa:2018} present predictions for the HI content of halos by post-processing the Illustris TNG simulation described by \cite{Nelson:2018}. Villaescusa et~al. find that there is no break in the HI $-$ halo mass relation, although the slope does get shallower for halos in which AGN feedback is important. These authors also find that satellites dominate the HI content of massive halos. However, the mass at which this transition occurs is somewhat higher than in our predictions. The predictions presented here extend those of \cite{Hank:2017}, who considered a version of the \cite{Lagos:2012} model with a new treatment of the suppression of gas cooling in low mass haloes due to photoionisation heating of the intergalactic medium. The model presented here uses a recently determined cosmology and a higher resolution N-body simulation, and has been recalibrated to reproduce selected observations of the galaxy population. The predictions that we have presented for the HI mass $-$ halo mass relation and its evolution will help to guide forecasts for the performance of HI intensity mapping experiments to probe the nature of dark energy. | 18 | 8 | 1808.08276 |
1808 | 1808.01901.txt | {Simultaneous observations of \psr\ with ESA's \xmm, the Giant Metrewave Radio Telescope and international stations of the Low Frequency Array revealed synchronous X-ray/radio switching between a radio-bright (B) mode and a radio-quiet (Q) mode. During the B mode we detected \psr\ in 0.2$-$2 keV X-rays and discovered pulsed emission with a broad sinusoidal pulse, lagging the radio main pulse by 0.208$\pm$0.012 in phase, with high pulsed fraction of 70$-$80\%. During the Q mode \psr\ was not detected in X-rays (2$\sigma$ upper limit a factor $\sim$9 below the B-mode flux). The total X-ray spectrum, pulse profile and pulsed fraction can globally be reproduced with a magnetized partially ionized hydrogen atmosphere model with three emission components: a primary small hot spot ($T$$\sim$3.6$\times10^6$ K, $R$$\sim$17 m), a larger cooler concentric ring ($T$$\sim$1.1$\times10^6$ K, $R$$\sim$280 m) and an antipodal hot spot ($T$$\sim$1.1$\times10^6 $ K, $R$$\sim$100 m), for the angle between the rotation axis and line of sight direction $\sim66^\circ$. The latter is in conflict with the radio derived value of $(84\pm0.7)^\circ$. The average X-ray flux within hours-long B-mode intervals varied by a factor $\pm$20\%, possibly correlated with variations in the frequency and lengths of short radio nulls or short durations of weak emission. The correlated X-ray/radio moding of \psr\ is compared with the anti-correlated moding of \psrb, and the lack of X-ray moding of \psrc. We speculate that the X-ray/radio switches of \psr\ are due to variations in the rate of accretion of material from the interstellar medium through which it is passing. } | So far, synchronous X-ray and radio mode switching has only been observed for the old and nearly aligned \psrb\ \citep{hermsen2013}. When this paragon of pulsar-mode switching \citep{suleymanova1984, rankin2006} switches from a radio ``bright'' (B) mode to a radio ``quiet'' (Q) mode, the X-ray flux increases by a factor of $\sim$2.4 and the X-ray pulsed fraction changes in magnitude and as a function of energy \citep{mereghetti2016}. The results are broadly consistent with the predictions of the partially screened gap model for rotation-powered pulsar emission (e.g. \citet{szary2015}), but they might also imply global magnetospheric rearrangements to explain the mode switching, as proposed in studies of intermittent pulsars \citep{kramer2006, lyne2010, camilo2012, lorimer2012}. It is imperative to find other moding pulsars in which this phenomenon can be precisely characterised and the results compared with \psrb. In a first attempt to find another example, we have conducted a similar campaign on the radio mode-switching pulsar \psrc\ \citep{fowler1981, fowler1982}. This pulsar exhibits mode switching on timescales of minutes, much shorter than the characteristic timescale of hours seen for \psrb. We found no evidence for simultaneous radio/X-ray mode switching of \psrc\ \citep{hermsen2017}. Though, we did, surprisingly, find a correlation between the mode durations and regular modulations in the pulse intensity. For this pulsar, it is possible that the very different geometry between the magnetic and spin axes (nearly orthogonal) may be hiding any visible X-ray mode changes, or that the physics of the short moding is quite different. In this work we report on a new observational campaign targeting the moding pulsar \psr. This pulsar is one of the brightest radio pulsars in the Northern sky, located at a distance of only $\sim$320 pc. It has a period $P$ of $\sim530$ ms, a spin-down age of $\sim 4.9 \times 10^6$ yr and an inferred magnetic field of $\sim 9.8 \times 10^{11}$ G. It is by no means an ordinary pulsar, showing main-pulse, inter-pulse and post-cursor emission components \citep{backer1973}. The pulsar was found to exhibit nulling (i.e. abrupt cessation and re-activation of its radio emission; see e.g. \citet{rathnasree1995}. More recently \citet{young2012}) found the nulls to be both short-term ($\sim$ min) as well as unusually long-term ($\sim$ hours or more). During bursts in its B mode \psr\ was found to exhibit single-pulse modulations with a repetition period $P_3$$\sim5P$ \citep{backer1973a, weltevrede2006a, weltevrede2007}. \citet{sobey2015} reported the discovery with LOFAR of a very weak quiet (Q) emission mode during these apparent nulling intervals. The transition between the Q mode and B mode occurs simultaneously for the main pulse, inter-pulse and post-cursor within about a single rotation of the neutron star. From the latter work we estimate the pulsar to be in the Q/null mode for $\sim$30\% of the time and the remaining $\sim$70\% in the B mode. \psr\ was also detected with {\it XMM-Newton} in X-rays in April 2002 \citep{becker2004}, with an effective EPIC Pn exposure of 33.7 ks. The detected source count rate in the Pn CCD was $3.5 \times 10^{-3}$ counts s$^{-1}$ for energies 0.3$-$10 keV, or an X-ray luminosity (0.5$-$10 keV) of $\sim 1.0 \times 10^{29}$ erg s$^{-1}$. Furthermore, there was a 2.2$\sigma$ indication for the detection of a pulsed signal. Considerably more exposure is apparently required to confirm the indication for pulsed emission, as well as for significantly constraining the spectral parameters of the X-ray emission. We carried out an X-ray/radio campaign on \psr\ with 6 $\times$ 25 ks of {\it XMM-Newton} observations, and simultaneous radio observations primarily with the GMRT at 325 MHz, supported by LOFAR UK, FR, SE and DE International stations in standalone mode at 150 MHz. From Lovell Telescope radio monitoring of \psr\ we derived an ephemeris to facilitate phase folding of the arrival times of the X-ray events. In Section 2, we present the radio observations and how we defined the radio modes. In the subsequent sections, we present the {\it XMM-Newton} X-ray observations (Section 3), the spatial analysis of the sky maps for the six X-ray observations, leading to the discovery of synchronous X-ray/radio mode switching (Section 4), the first detection of the X-ray pulsed signal from \psr\ (Section 5), followed by the spectral characterisation of the total and pulsed emissions (Section 6). In Section 7, we summarise our findings, followed in Section 8 by a discussion of them in comparison with the results from the previous campaigns on \psrb\ and \psrc. Finally, our overall conclusions are presented in Section 10. | We observed the radio-mode switching \psr\ for $\sim$39 hours simultaneously in X-rays and the radio band and report the discovery of synchronous {\it correlated} X-ray \& radio mode switching: when \psr\ is radio bright (B), we detect the pulsar in X-rays in the energy band 0.2$-$2 keV with a high pulsed fraction of 70$-$80\% due to a sinusoidal pulse that is lagging the radio main pulse by $\sim$ 0.2 in phase. During the radio-quiet or null mode (Q) we do not detect \psr\ in X-rays with an upper limit almost an order of magnitude lower than the reported flux in the B mode. This result is surprising, because the only pulsar for which synchronous X-ray \& radio moding was found so far, \psrb, showed {\it anti-correlated} mode switches \citep{hermsen2013}. The latter pulsar exhibits thermal pulsed as well as non-thermal unpulsed X-ray emission, both varying by a factor $\sim$ 2 in flux between similar radio B and Q modes \citep{mereghetti2016}. \psrb\ is classified as a nearly aligned rotator (near alignment of the magnetic and rotation axis) and \psr\ as a near-perpendicular rotator, but these pulsars have in common their broad thermal X-ray pulses with high pulsed fractions. Such high pulsed fractions can not be generated modelling black-body emission isotropically emitted from a hot polar cap. For \psr, we showed that its X-ray spectrum, pulse shape and pulsed fraction can de reproduced with a magnetized partially ionized hydrogen atmosphere model, with the two polar caps having the same temperatures but somewhat different sizes. However, this atmosphere-model fit to our X-ray data requires the angle between the rotation axis and line of sight direction to be $\sim66^\circ$, in conflict with the accurately radio-derived value of $\sim84^\circ$. \citet{storch2014} applied a similar atmosphere model to the X-ray data of \psrb\ and approximately reproduced also for this nearly aligned pulsar the pulse profile and high pulsed fraction, but in this case the radio-derived angles could be used in the model. In conclusion, we do not have an explanation yet for the high X-ray pulsed fractions that is consistently in agreement with radio-derived parameters. In this work we discovered, in addition to the synchronous X-ray \& radio mode switching, a new type of X-ray variability: the average X-ray flux within B-mode intervals of duration $\sim$ 7 hours varied by a factor $\pm$20\%. We speculate that the X-ray generation follows very closely the variations in frequency and lengths of short radio nulls or short durations of weak emission seen during the B mode, possibly on time scales of even few pulses. This possibility offers interesting constraints on the interpretation of what is causing mode switching: are we dealing with a local or global magnetospheric phenomenon? A follow-up investigation of the available single-pulse radio data is required to investigate such a possible tight X-ray/radio correlation. In our discussion on the possible nature of mode changing in \psr, we discussed in detail the radio and X-ray characteristics of \psr\ compared to those of the also radio-mode switching \psrc. For the latter pulsar an earlier X-ray/radio campaign did not reveal synchronous moding \citep{hermsen2017}. Both pulsars are near-perpendicular rotators. We are speculating that in \psr\ we are not seeing `true' mode-changing but the sudden appearance of strong bursts whose intensities follow a self-similar (\ie fractal) distribution over a wide range of timescales. Such a system could be identified as exhibiting self-organized criticality. In this context, we speculate that \psr\ is accreting material from a debris disk or the interstellar medium through which it is passing, to explain some of its X-ray characteristics. Further study of rotation-powered radio pulsars using simultaneous X-ray/radio data is needed to test and develop these various hypotheses. | 18 | 8 | 1808.01901 |
1808 | 1808.05929_arXiv.txt | Previous findings show that the existence of dense cores or bulges is the prerequisite for quenching a galaxy, leading to a proposed two-step quenching scenario: compaction and quenching. In this scenario, galaxies first grow their cores to a stellar mass surface density threshold and subsequently quenching occurs, suggesting that galaxies evolve from extended star-forming galaxies (eSFGs), through compact star-forming galaxies (cSFGs), to quenched population. In this work, we aim at examining the possible evolutionary link between eSFGs and cSFGs by identifying the trends in star formation rate (SFR), gas-phase metallicity and \HI\ content, since one would naturally expect that galaxies evolve along the track of cold gas consumption and metal enhancement. We select a volume-limited sample of 15,933 galaxies with stellar mass above $10^{9.5}$\msolar\ and redshift of $0.02<z<0.05$ from the NASA-Sloan-Atlas catalog within the ALFALFA footprint. The cSFGs on average exhibit similar or slightly higher SFRs of $\sim$0.06 dex and significantly higher gas-phase metallicity (up to 0.2 dex at low mass) with respect to the eSFGs, while the cSFGs dominate the galaxy population of the most intense star formation activities. More importantly, overall the median \HI\ content and gas depletion time of cSFGs are about half of eSFGs. Our result supports the compaction and quenching scenario that galaxies evolve and grow their cores along the track of cold gas consumption and metal enhancement. The environments of eSFGs and cSFGs are indistinguishable, suggesting that the compaction process is independent of any environmental effects at least for low-redshift universe. | \label{sec:introduction} The bimodality of color or star formation rate for galaxies has been found in decades by large imaging and spectroscopic surveys \citep[e.g.][]{Strateva-01, Baldry-04, Bell-04, Blanton-05, Faber-07, Wetzel-Tinker-Conroy-12}. Galaxies thus are naturally separated into two populations: the star-forming (SF) galaxies and quenched galaxies (QGs). The typical SF galaxies are actively forming stars with prominent disk-like morphology, whereas the quenched galaxies usually have no no-going star formation with pronounced spheroid-like morphology \citep[e.g.][]{Strateva-01, Kauffmann-03, Baldry-04, Brinchmann-04, Li-06, Muzzin-13, Barro-17, Wang-18a}. Deep narrow-field surveys show that this bimodality persists up to redshift of 2.5 \citep[e.g.][]{Bundy-06, Brammer-09, Huang-13}, and the prevalence of quenched population has dramatically increased since redshift of 1 \citep[e.g.][]{Muzzin-13, Tomczak-14}, indicating that quenching is one of the major themes in galaxy evolution over the past 8 Gyr. However, how SF galaxies become to be quenched is still not well understood. Many theoretical processes have been proposed to explain the star formation quenching, which can mainly be categorized into two classes: the internal processes and external processes. The feedback from starbursts and active galactic nuclei \citep[AGN;][]{McNamara-00, Nulsen-05, McNamara-Nulsen-07, Dunn-10, Fabian-12, Cicone-14}, belongs to the former class, act to heat the surrounding gas and/or strip the cold gas away from host galaxies. Another example of internal process is so-called ``morphological quenching" \citep{Martig-09}, i.e. the presence of a dominant bulge stabilizes the gas disk against gravitational instabilities needed for star formation. The external processes include a series of the environmental effects, such as major/minor mergers \citep[e.g.][]{Conselice-03, Cox-06, Smethurst-15}, and tidal/ram-pressure stripping \citep{Gunn-Gott-72,Moore-96, Abadi-Moore-Bower-99, Poggianti-17}, shock-excited heating \citep[e.g.][]{Rees-Ostriker-77, Birnboim-Dekel-03, Keres-05, Cattaneo-06} and strangulation \citep[e.g.][]{Weinmann-09, Peng-Maiolino-Cochrane-15, vandeVoort-17}, which act to either rapid consume the cold gas and/or expel the cold gas, or prevent the cold gas accretion. However, some of these internal processes may be closely related to environmental effects, which makes the star formation cessation picture to be rather complicated. For instance, the AGN activities can be triggered by galaxy-galaxy interactions or mergers \citep{Urrutia-Lacy-Becker-08, Ellison-11, Kocevski-12, Satyapal-14}. Since all the quenching mechanisms are working on the cold gas of galaxies, understanding the cold gas content in galaxies is essential to uncover the stage of galaxy evolution along the track of star formation quenching. Observationally, the link between quenching and structural properties has caught attention. Many works proposed that massive bulge is the key factor for quenching star formation in local galaxies \citep[e.g.][]{Bell-12, Cheung-12, Wake-vanDokkum-Franx-12, Fang-13, Bluck-14, Bluck-16}. For instance, \cite{Fang-13} found that the existence of a dense bulge is necessary but not sufficient to quench a galaxy by analyzing the stellar surface density profiles for SF and quenched central galaxies. More recently, by applying an artificial neural network approach for pattern recognition of quiescent systems on $\sim$400,000 central galaxies taken from SDSS, \cite{Teimoorinia-Bluck-Ellison-16} argued that the central velocity dispersion, bulge mass and the bulge-to-total stellar mass ratio are excellent quenching predictors, indicating that properties related to the central mass of the galaxy are most closely linked to the star formation cessation. The link between quenching and structural properties has also been established for high-redshift galaxies \citep[e.g.][]{Tacchella-15, Barro-17}. The fundamental structural differences for SF and quiescent galaxies are presented in the mass-size relation of different redshifts \citep[e.g.][]{Toft-07, Williams-10, Newman-12, vanderWel-14, Shibuya-Ouchi-Harikane-15}, where quiescent galaxies always exhibit a much higher stellar surface density than SF galaxies. These suggest that SF galaxies must grow dense cores before quenching. Recently, a two-step star formation quenching scenario has been proposed: compaction and quenching \citep[e.g.][]{Fang-13, Dekel-Burkert-14, Tacchella-15, Tacchella-16a, Tacchella-16b, Barro-17}. The term compaction, as used in the present work, can be caused by the shrinkage of galaxies with radial migration of stars and/or the growth of the cores or bulges without a significant change in overall radius (see Section \ref{subsec:compaction} for details). Under this scenario, in high redshift universe, the extended SF galaxies firstly contract via a dissipative process to loss energy and angular momentum, such as major merger, accretion of counter-rotating streams or recycled gas, usually associated with violent disc instability \citep{Dekel-Burkert-14, Zolotov-15}, subsequently become compact SF galaxies, then consume and/or lose their cold gas and turn to quiescent galaxies as a whole. In low redshift universe, the compaction of extended SF galaxies may be via minor merger, interaction with neighboring galaxies and/or bar-driven secular evolution to form the compact SF galaxies \citep[e.g.][]{Moore-96, Moore-Lake-Katz-98, DiMatteo-07, Bournaud-Jog-Combes-07, Wang-12, Barro-17, Lin-17}. Observationally, there are indicative identifications of the progenitors of the compact quiescent galaxies in the form of compact SF galaxies, whose masses, kinematics and morphologies resemble those of the compact quiescent galaxies \citep{Barro-13a, Barro-14b, Bruce-14, Nelson-14, Williams-14}, but appear to be different from other SF galaxies that have irregular or clumpy features. However, whether the extended SF galaxies are necessarily to be quenched through compaction, how the compaction process occurs and what happens in this process are still poorly understood. In this paper, we aim at testing the two-step quenching scenario, by comparing the extended SF galaxies and compact SF galaxies to find if there are any clues for the supposed evolution from eSFGs and cSFGs under the two-step quenching scenario. Following the work of \cite{Fang-13}, we use the stellar surface mass density within a radius of 1 kpc, \sgm, to quantify the growth of bulge as galaxies evolve. They have presented an existence of \sgm\ threshold, above which galaxies begin to shut down their star formation. This result persists up to at least redshift of 3 \citep{Barro-17, Whitaker-17, Mosleh-17}. While there is also a significant fraction of SF galaxies with \sgm\ greater than the threshold, studying their properties and connecting them with extended SF galaxies would probably shed light on the process of compaction and quenching. Thus, in this work, we select the sample from NASA Sloan Atlas and separate them into extended SF galaxies and compact SF galaxies\footnote{This definition of compact SF galaxies differs from previous works, where compact is an absolute term to identify the smallest galaxies \citep[e.g.][]{vanDokkum-15}.} according to \sgm-\mstar\ diagram. Although the sample galaxies we selected are limited to the low-redshift universe, this enables us to present detailed investigation on the chemical abundance, \HI\ content, and the environment for eSFGs and cSFGs, and further to examine whether there is a trend for cold gas consumption and metal enhancement from eSFGs to cSFGs. This paper is structured as follows. In Section \ref{sec:data}, we present the details on sample selection and definition of extended SF and compact SF galaxies, as well as the description of relevant physical parameters. In Section \ref{sec:results}, we investigate the star formation rates, gas-phase metallicities, \HI\ detection rate and environments as a function of stellar mass for extended SF and compact SF galaxies. In Section \ref{sec:discussion}, we discuss the gas depletion time along the star formation main sequence, the prevalence of AGN for the two populations, and the implications for our result. We summarize the result in Section \ref{sec:summary}. Throughout this paper, we assume a flat cold dark matter cosmology model with $\Omega_m=0.3$, $\Omega_\Lambda=0.7$ and $h=0.7$ when computing distance-dependent parameters. | \label{sec:summary} In this work, we select a volume-limited galaxy sample to have $0.02<z<0.05$, \lgmstar$>$9.5 and minor-to-major axis ratio greater than 0.5 from NSA galaxies located in the ALFALFA footprint. We then match the selected galaxies with MPA-JHU catalog, which results in 15,933 galaxies as our final sample. Further, we classify the galaxy sample into SF and quenched population according to the SFR-\mstar\ relation, and classify the SF galaxies into cSFGs and eSFGs based on the \sgm-\mstar\ diagram. We investigate and compare the properties of cSFGs and eSFGs, including the SFR, gas-phase metallicity, \HI\ content and environment. This is helpful to examine the possible evolution from eSFGs to cSFGs with naturally assuming that galaxies evolve along the track of cold gas consumption and metal enhancement. The main results are listed below. \begin{itemize} \item The cSFGs on average exhibit similar or slightly higher SFR with respect to eSFGs at fixed stellar mass, while the compact SF galaxies dominate the galaxy population with most intense star formation activities (0.5 dex above the normal SFMS). \item The cSFGs appear to have higher gas-phase metallicity at low stellar mass end than eSFGs, while the metallicity difference between the two populations vanishes at stellar mass higher than $\sim10^{10.5}$\msolar. \item The \HI\ detection rate of cSFGs is significantly lower than that of eSFGs at given stellar mass, implying that cSFGs are more \HI-deficient than eSFGs. Indeed, this result is confirmed by adopting the estimated \HI\ mass for \HI\ undetected galaxies in ALFALFA. \item The cSFGs and eSFGs of similar stellar mass appear to reside in similar environment, quantified by the host halo mass and local overdensity. \item The cSFGs are more frequently to host an AGN than eSFGs at given stellar mass, suggesting that the feeding of central massive black hole is more efficient in cSFGs than in eSFGs. \end{itemize} We examine the \HI\ depletion time of cSFGs and eSFGs, and find that cSFGs are more close to be quenched than eSFGs with assuming the current SFR and neglecting the replenishment of cold gas. All these findings support the compaction and quenching scenario, when assuming that galaxies evolve following the track of cold gas consumption and metal enhancement. Under this scenario, the compaction is essential to the build-up of stellar cores (or bulges) along with strong cold gas consumption and metal enhancement. In addition, the environmental effects appear to do not play a key role in the compaction process. We also note that there is no necessity of the evolution from eSFGs and cSFGs, since the two populations can be formed to be different due to different formation and gas accretion histories. | 18 | 8 | 1808.05929 |
1808 | 1808.10294_arXiv.txt | The secular evolution of a collisional star cluster of $N$-'point' stars have been conventionally discussed based on cumulative two-body relaxation process. The relaxation process requires a cut-off on the range of two-body encounter between stars in physical space and the relaxation time is characterized by Coulomb logarithm $\ln[N]$; the conventional cut-off on the encounter distance in the literature gives "dominant" effect. In addition, incorrect cut-offs exposed a mathematical "infinite-density" problem in the late stage of core-collapse. The present paper shows these are merely the results due to incorrect cut-off process. If one correctly constrains the cut-off on interaction range between stars based on truncated BBGKY hierarchy, one must introduce a self-consistent 'truncated' Newtonian mean-field (m.f.) acceleration of star at position $\bmath{r}$ and time $t$ due to a phase-space distribution function $f\left(\bmath{r}', \bmath{p}',t\right)$ for stars \begin{align} &\bmath{A}^{\triangle}(\bmath{r},t)=-Gm\left(1-\frac{1}{N}\right)\int_{\mid\bmath{r}-\bmath{r}' \mid > \triangle}\frac{\bmath{r}-\bmath{r}'}{\mid \bmath{r}-\bmath{r}' \mid^{3}} f\left(\bmath{r}',\bmath{p}',t\right)\text{d}^{3}{\bmath{r}'}\text{d}^{3}{\bmath{p}'},\nonumber \end{align} where $G$ is the gravitational constant and $m$ the mass of stars. The lower limit $\triangle$ of the distance between two stars is order of the Landau distance. The present paper shows the effect of total number on the structure of finite star cluster (in which stars undergo only two-body encounters) for the first time after establishing a mathematical formulation of a generalized Landau kinetic equation that includes the cut-off effect on both of collision term and m.f. potential by employing a BBGKY hierarchy for truncated DF stars. The cut-off effect increases the typical relaxation time by a few of percentage, which means the effect of cut-off itself is "non-dominant" on the relaxation time. On the other hand, the cut-off on m.f. acceleration is necessary to avoid "infinite-density" problem at the center of the system; the effect of total number $N$ on density profile and m.f. acceleration are shown by applying the truncated-DF BBGKY hierarchy to a toy model (a quasi-static modified Hubble density profile) for a core-halo structure of a star cluster at the late stage of evolution. | \label{sec:intro} A general point of view to understand statistical dynamics of dense star clusters is to introduce the effect of 'discreteness' of the clusters. The discreteness means the finiteness of total number $N$ of stars in a dense star cluster, say $N\approx 10^{5}\sim 10^{7}$. In the present paper, the system of concern is collisional star clusters, e.g. globular clusters and collisional nuclear star clusters without super massive black holes. As a first approximation $(N\to\infty)$, the system can be assumed smooth and its evolution is dominated by a self-consistent mean field (m.f.) potential. The effect of m.f. potential is of significance on a few of dynamical-time scales and may freeze the system into a quasi-stationary state due to rapid fluctuations in m.f. field potential (i.e. violent relaxation). The evolutions of long-lived star clusters might have been driven by less probable relaxation process, two-body close encounters \citep[e.g.][]{Goodman_1983,Goodman_1984} ), and 'slow' many-body relaxations, statistical acceleration of stars and gravitational polarization \citep{Gilbert_1968}, in addition to the effect of m.f. potential\footnote{The present work focuses on systems modeled by kinetics of one-body distribution function of stars ('point particles' interacting via pair-wise Newtonian forces), neglecting the effect of triple encounters and some realistic effects (gas/dust/dark-mater dynamics, stellar evolution, inelastic direct collisions, formation of stars and binaries, stellar mass distribution, ...)}. The most fundamental relaxation process in the evolution of collisional star clusters is arguably the statistical acceleration that stands for a non-collective relaxation and mathematically modeled by generalized Landau kinetic equation \citep{Kandrup_1981,Chavanis_2013a}. The statistical acceleration originates from the deviation of the actual force on 'test' star due to ($N-1$)-'field' stars from the smooth force due to the m.f. potential \citep{Kandrup_1988}. The statistical acceleration may be considered in association with the effect of stochastic many-body encounters \citep{Kandrup_1981_a}. Conventionally, the effect of many-body encounters approximately gives place to that of cumulative \emph{two-body} encounters between stars \citep{Chandra_1942}. While the cumulative two-body encounters become more probable on larger-space scales due to the long-range nature of Newtonian pair-wise potential, the statistical acceleration becomes greater in magnitude on smaller scales. This implies the relaxation effects on intermediate-space scales are of significance in evolution of the system. As a matter of fact, the basic assumption made in use of stochastic kinetic equations\footnote{The stochastic kinetic equations here mean collision-Boltzmann \citep{Ipser_1983}, forward Komologouv- Feller \citep{Kandrup_1980}, master\citep{Heggie_2003,Binney_2011,Merritt_2013}, Fokker-Planck\citep{Henon_1961} kinetic equations whose collision terms describe local two-body encounter in physical spaces with typically homogeneous background approximation.} is that 'test' star does not approach a 'field' star closer than the Landau distance and not go away far from the system size. The cut-off on the range of effectve encounter-distance gives the follwoing estimation for the order of relaxation time\citep[e.g.][]{Ambartsumian_1938,Cohen_1950,Spitzer_1988} \begin{align} \frac{t_\text{r}}{t_\text{d}}\approx \frac{N}{\ln\left[p_\text{max}/p_\text{min}\right]}\sim \frac{N}{[\ln{ N}]} .\label{Eq.lnN_loc} \end{align} where $p_\text{max}$ is the maximum parameter, a typical size of clusters (tidal radius, King radius, Jean length ...), and $p_\text{min}$ is the minimum impact parameter, the Landau distance (typically independent of relative velocity between two stars). The factor $[\ln{ N}]$ stands for a parameter of relaxation time scale and corresponds with `Coulomb logarithm'. The logarithm has been employed as a measure of `finite-N' effect and many-body encounter \citep{Aarseth_1998} and the corresponding Fokker-Planck models have been a sucess in sense that it is a simple numerical method for stellar dynamics \citep{Heggie_2003,Binney_2011} \subsection{A cut-off problem in use of kinetic theories} Since the stochastic kinetic theoreis do not self-consistently include the effects of inhomgeneity (even m.f. potential) and collective effects of star clusters, for correct treatment of the matter, one must resort to the first principles; BBGKY hierarchy \citep{Gilbert_1968,Gilbert_1971,Chavanis_2013a} and Klimontovich-Dupree equation \citep{Chavanis_2012}. The formulation based on the first principles shows a Coulomb Logarithm in relaxation time for local encounters but in term of wavenumber \citep[e.g.][]{Severne_1976, Kandrup_1981,Chavanis_2013a} \begin{align} \ln\left[\frac{k_\text{min}}{k_\text{max}}\right]\approx \ln[N].\label{Eq.lnN_non_loc} \end{align} where conventionally the following relations are assumed \begin{subequations} \begin{align} &k_\text{min}\approx p_\text{max}\\ &k_\text{max}\approx p_\text{min} \end{align} \end{subequations} Yet, the fundamental assumption, equation \eqref{Eq.lnN_loc}, made for stochastic kinetic theory has not been `converted' into a self-consistent kinetic equation. The motivation for this work originates not only from the generalization work of the previous works but from some doubt for equations derived from stochastic kinetic theories and first principles. Use of stochastic kinetic equation allows one to employ typical m.f. potential that is smooth limitless in physical space; this is obviously inconsistent with the cut-off, equation \eqref{Eq.lnN_loc}. There are two kinds of test star exists in the system; a `uncorrelated' test star can approach a field star limitlessly forming a m.f. potential while `correlated' test star can not approach closer than the Landau distance to avoid close encounters. the outcome is the mathematical production of infinite density due to core collapse at the late stage of two-body relaxation evolution. On one hand, the equations derived based on the first principles also have a inconsistency. To find the relation \eqref{Eq.lnN_non_loc} the previous works assumes that test star can approach field star limitlessly while this is against assumptions of weak-coupling approximation and to be cut-off on scales of the Landau distance. The purpose of the present work is to derive based on a first principle a kinetic equation that correctly `cut-off' the encounter distance in physical space and the resolve the inconsistence of the existing kinetic theories. To do so the basic `target' of kinetic equation is the g-Landau kinetic equation. This is since the equation is known to correctly take into account the inhomogeneity effect and one does not have to assign cut-off on the maximum encounter-distance \citep{Kandrup_1988,Chavanis_2013a}. Yet, one needs to assign a lower cut-off for the encounter distance. One can resort to the use of BBGKY hierarchy truncated DF invented by \citep{Grad_1958} that can isolate a physical space where the short-range interaction between particles dominates from where weak-interaction occurs. The present work employes the truncated DF to cut-off the encounter-distance at Landau distance. This corresponds to a direct extension work of \citep{Takase_1950} where the Holtsmark distribution of force fields are employed and the strong-two body encounters and the formation of binaries are neglected by truncating the DF at order of Landau distance, termed as 'rough approximation'. The present paper is organized as follows. In section \ref{sec:truncated_BBGKY} the truncated DF and the BBGKY hierarchy are explained. In sections \ref{sec:complete_WC} the g-Landau kinetic equation for the truncated DF of stars for a weakly-coupled star cluster is derived. In section \ref{sec:grainess_effect} the effect of cut-off on the m.f. potential and collision term, Coulomb logarithm, is discussed. Section \ref{sec:conclusion} is Conclusion. | 18 | 8 | 1808.10294 |
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1808 | 1808.09046_arXiv.txt | We present the results of \psrpi, a large astrometric project targeting radio pulsars using the Very Long Baseline Array (VLBA). From our astrometric database of 60 pulsars, we have obtained parallax-based distance measurements for all but 3, with a parallax precision that is typically $\sim$45 \uas\ and approaches 10 \uas\ in the best cases. Our full sample doubles the number of radio pulsars with a reliable ($\gtrsim$5$\sigma$) model-independent distance constraint. Importantly, many of the newly measured pulsars are well outside the solar neighborhood, and so \psrpi\ brings a near-tenfold increase in the number of pulsars with a reliable model-independent distance at $d>2$ kpc. Our results show that both widely-used Galactic electron density distribution models contain significant shortcomings, particularly at high Galactic latitudes. When comparing our results to pulsar timing, two of the four millisecond pulsars in our sample exhibit significant discrepancies in their proper motion estimates. With additional VLBI observations that extend our sample and improve the absolute positional accuracy of our reference sources, we will be able to additionally compare pulsar absolute reference positions between VLBI and timing, which will provide a much more sensitive test of the correctness of the solar system ephemerides used for pulsar timing. Finally, we use our large sample to estimate the typical accuracy attainable for differential VLBA astrometry of pulsars, showing that for sufficiently bright targets observed 8 times over 18 months, a parallax uncertainty of 4 \uas\ per arcminute of separation between the pulsar and calibrator can be expected. | \label{sec:introduction} With magnetic field strengths exceeding $10^{14}$~\hbox{G}, rotation rates approaching 1000~Hz, central densities exceeding $10^{14}$~g~cm${}^{3}$, and surface gravitational field potentials of order 40\% of that of a comparable mass black hole, neutron stars have proven to be powerful physical laboratories. With their large moments of inertia, when detected as radio pulsars, their pulses provide a highly regular clock. Studies of pulsars have placed strong constraints on the equation of state of neutron stars \citep{demorest10a}, provided the first detection of extrasolar planets \citep{wolszczan92a}, and provided the first observational evidence for the existence of gravitational waves \citep{taylor89a}. In many cases, these results have been obtained despite considerable uncertainty in the distance of the pulsar (or pulsars). It has not proven possible to relate a pulsar's radio luminosity to any other intrinsic physical quantity that would provide an independent distance estimate \citep[][]{szary14a}---however, it is instead possible to make use of the pulsar's dispersion measure (DM) and a model of the Galactic electron density distribution to provide this distance estimate. However, the difficulty of modeling all the small-scale structure of the ionized component of the Milky Way means that the fidelity of Galactic electron density distribution models is generally rather low. Accordingly, the reliability of DM-based distance estimates for individual pulsars is generally quite low, and errors of a factor of several are not rare \citep{deller09b,chatterjee09a}. While some pulsar science use cases are relatively unaffected by such errors, there are others for which knowing the distance is vital and the distance uncertainty becomes the limiting factor in the measurement. For instance, studies of the pulsar velocity distribution and hence supernova kicks can be biased by distance errors \citep[e.g.,][and references therein]{verbunt17a}, while studies of pulsar gamma-ray emission cannot build an accurate energy budget without a correct calibration of high-energy flux into luminosity \citep[e.g.,][]{abdo13a}. Various methods exist to obtain non-DM--based estimates of pulsar distances. These include measurements of annual orbital parallax via pulsar timing \citep[e.g.,][]{matthews16a}, visible wavelength observations \citep[e.g.,][]{2001ApJ...561..930C}, or Very Long Baseline Interferometry \citep{chatterjee09a}, or via model-dependent approaches such as \ion{H}{1} absorption limits \citep[e.g.,][]{1969Natur.221..249G,2008ApJ...676.1189M}. Of these, VLBI astrometry is the most robust. In addition to being dependent upon a model for Galactic rotation, \ion{H}{1} absorption formally provides only a lower limit; the spectra of pulsars are such that few pulsars are detected at wavelengths shorter than radio and angular resolutions are typically poorer than can be achieved with \hbox{VLBI}; and pulsar timing parallaxes are generally only achieved with millisecond pulsars. The \psrpi\ campaign was conceived as a successor to previous intensive VLBI campaigns \citep{brisken02a,chatterjee09a,deller09b} that would treble the number of radio pulsars with a distance measurement having a precision of better than 10\% and use the result to constrain the characteristics of the radio pulsar population (e.g., velocity, luminosity) as well as improving models of the Galactic electron density distribution. A subset of \psrpi\ results for two binary millisecond pulsars has been previously presented \citep{deller16a}, and in this paper we present the results for the full sample of~60 pulsars. Section~\ref{sec:obsdataproc} describes the observations, data reduction, and position extraction, while Section~\ref{sec:results} describes the astrometric results and error analysis. Section~\ref{sec:discussion} contains an analysis of both individual pulsars and parameters of the pulsar population, an evaluation of different Galactic electron density distribution models, a comparison of the VLBI results to pulsar timing, and a forward look to future observations for reference frame ties with radio pulsars. Section~\ref{sec:conclusions} contains our conclusions. | 18 | 8 | 1808.09046 |
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1808 | 1808.02448_arXiv.txt | The origin of Mercury's high iron-to-rock ratio is still unknown. In this work we investigate Mercury's formation via giant impacts and consider the possibilities of a single giant impact, a hit-and-run, and multiple collisions in one theoretical framework. We study the standard collision parameters (impact velocity, mass ratio, impact parameter), along with the impactor's composition and the cooling of the target. It is found that the impactor's composition affects the iron distribution within the planet and the final mass of the target by up to 15\%, although the resulting mean iron fraction is similar. We suggest that an efficient giant impact requires to be head-on with high velocities, while in the hit-and-run case the impact can occur closer to the most probable collision angle (45$^{\circ}$). It is also shown that Mercury's current iron-to-rock ratio can be a result of multiple-collisions, with their exact number depending on the collision parameters. Mass loss is found to be more significant when the collisions are tight in time. | Mercury has a unique composition. Its mean density is similar to the Earth's but Mercury is 20 times lighter and cannot be subject to the same self-compression. This suggests the existence of a large metallic core consisting $\sim$ 70\% of the planet's mass, i.e., a large iron-to-rock ratio (hereafter $Z_{Fe}$) which is about twice the proto-solar abundance (e.g., \citealt{Spohn}, \citealt{Hauck2013}). Several scenarios have been proposed to explain Mercury's high $Z_{Fe}$ and they cover different stages of the planet formation process and rely on different chemical and physical mechanisms. One class of mechanisms is linked to the separation of metals and silicates in the solar nebula. This can be a result of different condensation temperatures for metals and silicates \citep{Lewis}, their different conductive properties and reaction to photophoretic forces \citep{Wurm13}, or a result of different gravitation and drag force balance \citep{Weidenschilling}. These mechanisms however, typically require specific disk architecture and conditions. A second class of scenarios suggests that Mercury lost a large fraction of its rocky mantle after its formation by evaporation followed by solar wind \citep{Cameron} or mantle stripping by a giant impact \citep{Benz1988}. Gravitational collisions between bodies of similar sizes are very common in the final stages of planet formation, (e.g., \citealt{Chambers},\citealt{Quintana}). Collisions can include violent giant impacts that are energetic enough to strip away part of the mantle, and one giant impact is sufficient to reach Mercury's current $Z_{Fe}$ \citep{Benz2007}. However, the exact conditions that lead to Mercury's formation via a giant impact are still unknown. Recent scenarios proposed that Mercury might have collided with another body as large as Earth or Venus \citep{Asphaug}. Most simulations of the late formation stage start with embryos of Mars size as it is presently still computationally challenging to resolve the inner disk in such simulations. To our knowledge only \citet{Lykawka} have investigated the formation of Mercury in an inner disk simulation. In the best terrestrial planet systems, the analogs tend to be slightly heavier than Mercury's current mass leaving open the possibility of a giant impact formation scenario for Mercury. The giant impact hypothesis has been under debate since the MESSENGER mission results (\citealt{Peplowski}, \citealt{Ebel2017}) where relatively high abundance of potassium (K), thorium (Th) and uranium (U) were measured. The initial interpretation of these measurements suggested that any scenario involving high temperatures, or high energies, which would remove the volatiles from Mercury is excluded. This argument is based on our knowledge of the Moon's composition. The Moon, which is also thought to form via a giant impact, is volatile depleted with a K/Th ratio $\sim$5 times lower than the Earth's. However, the volatile depletion of the Moon is linked to recondensation and not vaporization (\citealt{Stewart2016}, \citealt{Lock2018}). Unlike Mercury, the Moon probably formed from a debris disk that accumulated to form the satellite. By analogy to the Earth-Moon system, Mercury is the Earth remnant, which retains a significant amount of volatiles. In addition, the disrupted silicate material in Mercury could be well-mixed and preserve its original composition \citep{Nittler2017}. Fractionation of volatiles within the condensed silicates is not expected, but some volatiles could be lost from atmospheres/oceans \citep{Schlichting2015, Genda2005}. There may also be transfer between impactors in hit-and-run events \citep{Burger2018}. Generally, giant impacts might not lead to volatile depletion for terrestrial planets, as argued by \cite{Ebel2017}, despite their very different impact histories, they seem to have very similar K/Th and K/U ratios. In this paper we investigate the giant impact hypothesis. We consider (1) a single giant impact (2) a hit-and-run and (3) multiple collisions in one numerical framework. We investigate a large parameter space for individual collisions to understand the outcome possibilities, as well as the sensitivity to the impactor's composition and proto-Mercury's initial state. We find that all three options can lead to the formation of a Mercury-like planet, although each scenario requires different impact conditions. This paper is organized as follows. In Section \ref{methods} we describe the methods used to model the planetary bodies and their respective collisions. We also explain the tools for the analysis of the simulation outcomes such as the clump finder. In Section \ref{Paramterstudy} we present the results of the collisions of the parameter space we have explored. In Section \ref{resultsmulti} we discuss our approach to the multiple collision scenario. In section \ref{discussion} we briefly discuss the importance of following the evolution of the ejecta and impactor, and compare our results with previous studies. In Section \ref{conclusion} we summarize and discuss the results. | } We investigated formation paths of Mercury including giant impact, hit-and-run, and multiple collisions. We presented a large parameter study for these three cases, considering different collision parameters, the sensitivity of the results to impactor's composition, and different initial states (inflated, rotating) of the target. We find that the two end-members of the range of successful giant impacts are with $b=0.5-0.7$ and $v_{imp}\sim30$~km/s (5-6~$v_{esc}$) and with $b=0.2-0.3$ and $v_{imp}\sim15$~km/s (3-4~$v_{esc}$). In the first case, the constraints on both the velocity and the angle are tight, and are not very likely. The latter requires a very small impact angle, but with a more moderate velocity ($v_{imp}\sim$15~km/s). In Case-2, the hit-and-run scenario, the impact velocity is closer to the escape velocity. % On the other hand, a massive object on a highly eccentric orbit is also somewhat rare, and its origin as well as its fate post-impact still need to be investigated and justified. In Case-2, we also find that there is a larger parameter space of possibilities to form Mercury-like planets, but the proto-Mercury has to be 4-5 times more massive than its present mass. For the same impact parameters, disruptive collisions are also more likely than in Case-1 since they are more energetic. It is difficult to assert which of the cases is more probable. Future investigations of N-body simulations could put limits on the collision rates and statistics. % Finally, we also show that Mercury can form via several collisions with less extreme conditions each time. The closer to the most likely statistical values for the impact velocity and angle, the more impacts that are required. If the next collision occurs shortly after the first one, more mantle and cloud-like material mostly composed of silicate but with a small iron fraction from the target) can be stripped. A few collisions happening tightly in time is a more effective scenario for reaching Mercury's high $Z_{Fe}$ and is furthermore also consistent with its surface volatile-rich composition. Our results can be summarized as follows: \begin{itemize} \item A single giant impact or hit-and-run impact require highly tuned impact parameters and velocities to reproduce Mercury's mass and $Z_{Fe}$. There is a somewhat larger parameter space of possibilities in the hit-and-run scenario. \item The impactor's composition affects the resulting final mass and post-impact iron distribution. \item The pre-impact state of the target affects the resulting final mass. \item A multiple collision scenario escapes the fine-tuning of the geometrical parameters but is constrained by the timing and the volatile-rich composition of Mercury's surface. \item Forming Mercury by giant impacts is feasible but difficult. \end{itemize} Mercury's origin remains poorly understood as it combines physical, chemical and dynamical processes that must be coupled. % The low frequency of metal-rich exoplanets (Mercury-analogs) suggests that forming metal-rich planets requires unique circumstances. Therefore, understanding the formation of Mercury can reflect on our understanding of exoplanetary systems. | 18 | 8 | 1808.02448 |
1808 | 1808.10384_arXiv.txt | {The dwarf planet Ceres and the asteroid Vesta have been studied by the Dawn space mission. They are the two heaviest bodies of the main asteroid belt and have different characteristics. Notably, Vesta appears to be dry and inactive with two large basins at its south pole. Ceres is an ice-rich body with signs of cryovolcanic activity.} {The aim of this paper is to determine the obliquity variations of Ceres and Vesta and to study their rotational stability.} {The orbital and rotational motions have been integrated by symplectic integration. The rotational stability has been studied by integrating secular equations and by computing the diffusion of the precession frequency.} {The obliquity variations of Ceres over $[-20:0]\MYR$ are between $2$ and $20\degree$ and the obliquity variations of Vesta are between $21$ and $45\degree$. The two giant impacts suffered by Vesta modified the precession constant and could have put Vesta closer to the resonance with the orbital frequency $2s_6-s_V$. Given the uncertainty on the polar moment of inertia, the present Vesta could be in this resonance, where the obliquity variations can vary between $17$ and $48\degree$.} {Although Ceres and Vesta have precession frequencies close to the secular orbital frequencies of the inner planets, their long-term rotations are relatively stable. The perturbations of Jupiter and Saturn dominate the secular orbital dynamics of Ceres and Vesta and the perturbations of the inner planets are much weaker. The secular resonances with the inner planets also have smaller widths and do not overlap contrary to the case of the inner planets.} | Ceres and Vesta are the two heaviest bodies of the main asteroid belt. They have been studied by the Dawn space mission, which has allowed to determine notably their shape, gravity field, surface composition, spin rate and orientation \citep{russell2012,russell2016}. However the precession frequency of their spin axes has not be determined and there are still uncertainties about their internal structures \citep[e.g.][]{park2014,ermakov2014,park2016,ermakov2017b,konopliv2018}. No satellites have been detected from observations with the Hubble Space Telescope and the Dawn space mission around these bodies \citep{mcfadden2012,mcfadden2015,demario2016}. The long-term rotation of the bodies of the solar system can be studied with secular equations \citep{kinoshita1977,laskar1986,laskarrobutel1993} or with a symplectic integration of the orbital and rotational motions \citep{toumawisdom1994}. Secular equations are averaged over the mean longitude and over the proper rotation, which is generally fast for the bodies of the solar system, and their integration is much faster. They were used by \cite{laskarjoutelrobutel1993} and \cite{laskarrobutel1993} to study the stability of the planets of the solar system. The method of \cite{laskarrobutel1993} has been applied by \cite{skoglov1996} to study the stability of the rotation and the variations of the obliquity for Ceres and nine asteroids including Vesta. However at this time, the initial conditions for the spin axes were not determined precisely and knowledge on the internal structure lacked to constrain the precession frequencies. \cite{skoglov1996} assumed that the bodies are homogeneous and concluded that their long-term rotations are relatively stable. By using secular equations and a secular model for the orbital motion, \cite{bills2017} determined the obliquity variations of Ceres. \cite{ermakov2017a} obtained the obliquity variations of Ceres for different polar moments of inertia by realizing the symplectic integration of the rotational and orbital motions. Asteroid impacts and close encounters can influence the long-term rotation of bodies in the main asteroid belt. Vesta has suffered two giants impacts \citep{marchi2012,schenk2012}, which have significantly modified its shape and its spin rate \citep{fu2014,ermakov2014}. \cite{laskar2011a} obtained an orbital solution of Ceres and Vesta, called La2010, which takes into account mutual interactions between bodies of the main asteroid belt and \cite{laskar2011b} showed that close encounters in the solution La2010 are the cause of the chaotic nature of the orbits of Ceres and Vesta. These close encounters can affect their long-term rotation. For Ceres, the obliquity drives the ice distribution on and under the surface. Ceres possesses cold trap regions, which do not receive sunlight during a full orbit. This prevents the sublimation of the ice, which can accumulate \citep{platz2016}. The surface area of these cold traps depends on the value of the obliquity. \cite{ermakov2017a} determined that the obliquity of Ceres varies between $2$ and $20\degree$ and that the cold trap areas for an obliquity of $20\degree$ correspond to bright crater floor deposits that are likely water ice deposits. \cite{platz2016} determined that one bright deposit near a shadowed crater is water ice. In addition, the Dawn mission gave evidence of the presence of ice under the surface of Ceres from the nuclear spectroscopy instrument \citep{prettyman2017} and from the morphology of the terrains \citep{schmidt2017}. The ice distribution and the burial depth with respect to the latitude depends on the history of the obliquity \citep{schorghofer2008,schorghofer2016}. For Vesta, studies of the long-term evolution of the obliquity were not performed with the initial conditions of the spin axis and the physical characteristics determined by the Dawn space mission. The main purpose of this article is to investigate the long-term evolution of the rotational motions of Ceres and Vesta. First, we explore the obliquity variations of these bodies for a range of possible precession constants obtained from the data of the Dawn mission. Then, the stability of their spin axes is studied. In this paper, we consider for the orbital motion the solutions \La{} and La2010 \citep{laskar2011a}, which do not include the rotation of Ceres and Vesta. To compute the obliquity variations, we follow the symplectic method of \cite{farago2009} by averaging the fast proper rotation. This method avoids to integrate the fast rotation and allows to use a large step to reduce the computation time. We call \Ce{} the long-term rotational solution obtained. The orbital and rotational equations are integrated simultaneously in a symplectic way and the effects of the rotation on the orbital motions are considered. We consider the close encounters of Ceres and Vesta with the bodies of the main asteroid belt used in \cite{laskar2011b} and estimate with a statistical approach their effects on the long-term rotation of Ceres and Vesta. In order to determine the secular frequencies and identify the possible secular resonances on the orbital and rotational motions, the solutions are studied by the method of the frequency map analysis \citep{laskar1988,laskar1990,laskar1992,laskar1993,laskar2003}. Moreover to study the effects of the close secular orbital resonances, we compute a secular Hamiltonian from the method of \cite{laskarrobutel1995}. We obtain a secular model, which reproduces the secular evolution of the solution \La{} and allows to investigate the effects of the secular resonances. The stability of the spin axes is studied by using secular equations with a secular orbital solution obtained from the frequency analysis of the solution \La{}. We verify beforehand that they allow to reproduce the obliquity variations computed by the symplectic method and have the same stability properties. We study the stability of the rotation in the vicinity of the range of possible precession constants to identify the secular resonances between the orbital and rotational motions. Vesta has suffered two giant impacts which have changed its shape and its spin rate \citep{fu2014,ermakov2014} and also its precession constant. We investigate if this possible evolution of precession constant changed the stability properties. Following the method of \cite{laskarrobutel1993}, we finally realize a stability map of the spin axes of Ceres and Vesta. In section \ref{SEC:methods}, we present the methods used in this paper to obtain the long-term rotation. In section \ref{SEC:prec}, we estimate the precession constants deduced from Dawn space mission and their possible variations during the history of Ceres and Vesta. In section \ref{SEC:orbobliquite}, we analyse the long-term solutions obtained for the orbital and rotational motions. In section \ref{SEC:secularmodels}, we study the effects of the orbital secular resonances with a secular Hamiltonian model. In section \ref{SEC:stab}, we study the stability of the rotation axes from the secular equations of the rotation. | 18 | 8 | 1808.10384 |
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1808 | 1808.09967_arXiv.txt | The infrared dust emission from the white dwarf GD\,56 is found to rise and fall by 20\% peak-to-peak over 11.2\,yr, and is consistent with ongoing dust production and depletion. It is hypothesized that the dust is produced via collisions associated with an evolving dust disk, temporarily increasing the emitting surface of warm debris, and is subsequently destroyed or assimilated within a few years. The variations are consistent with debris that does not change temperature, indicating that dust is produced and depleted within a fixed range of orbital radii. Gas produced in collisions may rapidly re-condense onto grains, or may accrete onto the white dwarf surface on viscous timescales that are considerably longer than Poynting-Robertson drag for micron-sized dust. This potential delay in mass accretion rate change is consistent with multi-epoch spectra of the unchanging Ca\,{\sc ii} and Mg\,{\sc ii} absorption features in GD\,56 over 15\,yr, although the sampling is sparse. Overall these results indicate that collisions are likely to be the source of dust and gas, either inferred or observed, orbiting most or all polluted white dwarfs. | Observable and real-time changes in exoplanetary systems hold important clues for dynamical processes during their birth and long-term evolution, where the number of variable systems is likely to increase owing to large ground- and space-based surveys. Sensitive monitoring of giant planetary and substellar atmospheres can reveal periodic features such as rotation, atmospheric wind speeds, and global weather patterns \citep{sne14,lou15,apa17}. Smaller major and minor exoplanetary bodies are typically out of reach for real-time monitoring, but those that actively produce debris can generate sufficient area to be detected via absorption or emission and studied over time. Such systems have the potential to constrain the bigger picture of planet formation and evolution, especially if they exhibit secular changes, non-periodic events, or signposts of important evolutionary phases such as the Late Heavy Bombardment. Active exo-cometary populations are known in a handful of systems via transient absorption in optical and ultraviolet spectra \citep{kie14a,kie14b,wel18}, while {\em Kepler} data now includes convincing transits of individual exo-comets in at least one system \citep{rap18}. The dramatic and irregular flux changes measured towards KIC\,8462852 are also broadly consistent with an exo-cometary origin \citep{boy18,wya18}. The inner and terrestrial planet-analog regions are more challenging to detect in general due to their host star proximity, but sensitivity to compact orbits has enabled {\em Kepler} to detect transits from a trio of rocky planets via tailing debris clouds associated with their short periods and thus irradiation-driven mass loss \citep{van17}. Analogous processes may contribute to the transit light curves of the polluted white dwarf WD\,1145+017 (hereafter WD\,1145; \citealt{van15,gan16,rap16}). Exoplanetary system variability has also been observed in light curves via emission, including a few spectacular examples of infrared flux changes associated with the terrestrial-planet forming regions \citep{mel12,men14}. And within the former terrestrial zones of A-type and similar stars, there is myriad evidence for rocky planetesimal activity via the evolved planetary systems orbiting and polluting white dwarf stars (see \citealt{far16} and references therein). The first clear evidence of any variability in these systems emerged via changes in disk line emission \citep{gan08}, and preceded transit detections by several years. Similar changes have now been documented for five disks via gas emission or absorption \citep{wil14, man16a,man16b,den18}, including WD\,1145 \citep{red17,cau18}, but to date only a single white dwarf system has been shown to vary in the infrared. SDSS\,J095904.69-020047.6 (hereafter SDSS\,0959) was reported to show a drop in $3-5\,\upmu$m infrared flux by around 35\% in less than 300\,d \citep{xu14} based on the comparison of its warm {\em Spitzer} IRAC discovery fluxes \citep{far12a} to measurements made with {\em WISE}. \citet{xu14} surmised the drop in flux was likely due to a change in the inner disk radius from an impact or instability, causing a few percent of the entire disk mass to be accreted within the duration of the observed change, and suggest that the resulting accretion rate increase could lead to an observable difference in the photospheric metal abundance on similarly-short timescales. This paper reports long-term, $3-5\,\upmu$m infrared flux variations in the polluted white dwarf GD\,56 (= WD\,0408--041), a hydrogen atmosphere (DA-type) star with $T_{\rm eff}\approx15\,000$\,K and a cooling age around 200\,Myr \citep{gia11}. The variability in the infrared is presented together with optical spectroscopy that reveals constant, photospheric metal absorption over a similar timescale. Infrared data are comprised of nine {\em Spitzer} observational epochs spanning 11.2\,yr from 2006 to 2017, and are supplemented by multi-epoch {\em WISE} data from 2010 and 2014--2017. There has been one substantial increase of at least 20\%, and what appear to be two decaying trends of similar magnitude, all taking place over several years each. These changes are interpreted as the production (increase) of dust clouds, and their subsequent depletion (decrease), where possible scenarios may also account for the flux decrease at SDSS\,0959. The rest of the paper is organized as follows. Section two describes the infrared data, where the analysis includes both absolute flux measurements and differential photometry using field stars, and the spectroscopic data that indicate no appreciable changes to the inferred metal accretion rate. Section three describes the (weak) constraints on disk models based on the data, as well as theoretical considerations, and discusses possible model families for GD\,56 and dusty white dwarfs in general. Section four gives the summary and outlook. | The changes in absolute infrared flux of GD\,56 are clear and compelling, but at the same time the flux ratios between the two shortest wavelength channels in both {\em Spitzer} and {\em WISE} are consistent with remaining constant (Figures \ref{fig1}--\ref{fig3}). Concurrently, there is no indication that the photospheric line strengths or metal abundances are changing, although the sampling is relatively sparse. Based on diffusion theory, the expected sinking timescales for calcium and magnesium in a hydrogen-atmosphere white dwarf with $T_{\rm eff}=15\,000$\,K, $\log\,g =8.0$ are 1.0 and 0.9\,d respectively \citep{koe09}. Thus, while the stellar photosphere should track any accretion rate changes on daily timescales or longer, none are apparent in the observations. These results are comparable to those obtained via multi-epoch spectral observations of the dusty white dwarf prototype G29-38 \citep{von07,deb08}. These observations provide empirical constraints on the possible underlying processes in the circumstellar disk orbiting GD\,56. In this section, the infrared variability and apparently constant accretion rate are discussed in the context of relevant timescales, geometry, and possible physical models. These ideas are considered in the context of other dusty systems for consistency, and with respect to disk evolution models to assess their wider applicability. The inferred, ongoing accretion rate is relevant for the following discussion, where diffusion theory predicts rates of $1.0\times10^8$\,g\,s$^{-1}$ and $0.6\times10^8$\,g\,s$^{-1}$ based on the magnesium and calcium abundances, respectively, assuming they are deposited at their bulk Earth mass fractions \citep{all01,koe09}. \begin{figure} \includegraphics[width=84mm]{camgs.eps} \caption{Multi-epoch optical spectroscopy of GD\,56 in the region of the Ca\,{\sc ii} K and Mg\,{\sc ii} 4481\,\AA \ lines, with details provided in Table \ref{tbl2}. All spectra are normalized and vertically offset for clarity, with chronological ordering from top down. The 2011 Mar 23 spectrum exhibits weaker and narrower lines, but is the result of a single exposure taken during twilight, with bright sky emission throughout the wavelength regions of interest. Hence, this may have affected the fidelity of the data during the spectral extraction process. The velocities of all absorption features are consistent within the wavelength calibration errors. \label{fig4}} \end{figure} \begin{figure} \includegraphics[width=84mm]{ewidths.eps} \caption{The measured equivalent widths of the two metal absorption features at each epoch, color-coded identically to the spectra in Figure \ref{fig4}. These measurements are shown with $2\upsigma$ error bars and indicate that no appreciable changes are observed in metal abundance, and are thus consistent with a constant accretion rate. The purple data points corresponding to the 2011 Mar 23 data may have suffered from high background contamination and problematic sky removal. \label{fig5}} \end{figure} \subsection{Previous Infrared Modeling of GD\,56} It is worthwhile to first review previous data and modeling of the infrared emission in the GD\,56 system, as it remains unusual and to date has the largest, observed, fractional infrared luminosity of any dusty white dwarf \citep{roc15}. Over a decade ago when GD\,56 was first observed with {\em Spitzer}, it was noted that the infrared emission could not be reproduced by a flat disk alone \citep{jur07}. In Figure \ref{fig6} these initial infrared observations are reproduced with better short wavelength data and atmospheric modeling, together with a flat disk model that fails to reproduce the IRAC flux measurements. Because it is now clear that the $3-5\,\upmu$m fluxes vary over time, any disk model must be viewed with caution when applied to infrared data that are not taken simultaneously, such as those in the figure. Regardless, no flat disk model can reach the plotted fluxes, and thus this model is insufficient. A modestly warped segment was later added to the flat disk model for GD\,56, and this was able to account for the overall infrared emission including its strong silicate feature \citep{jur09}. The source of any warping in an otherwise flat disk is unclear, but possibilities include gravitational perturbations by orbiting bodies, radiative instability \citep{jur07}, and dust that follows the scale height of gas. In the lower panel of Figure \ref{fig6} the stellar photosphere has been subtracted and only the face-on (flat) disk model remains with the initial IRAC data, the sole MIPS 24\,$\upmu$m observation, and the $HK$ photometry. The shaded regions give the extent of the observed variability to date in all bands with multiple measurements, where the largest variation is the 0.26\,mJy peak-to-peak flux change at 4.5\,$\upmu$m (between 2013 Oct and 2017 Dec). While the temporal coverage in the infrared has (multi) year-long gaps, the overall behavior of the fluxes in Figure \ref{fig1} suggests the variable dust emission is either stochastic, or possibly a periodic brightening followed by dimming. While radiative warping may be responsible for additional emitting surface that exceeds the intrinsic infrared brightness of a flat disk configuration, it does not necessarily follow that such a structure can cause the observed variability. \begin{figure} \includegraphics[width=84mm]{seds.eps} \caption{The upper panel plots the photometric spectral energy distribution (SED) of GD\,56 from the ultraviolet through the infrared, including the first epoch of (cryogenic) IRAC data, and MIPS 24\,$\upmu$m flux. The short wavelength data are fitted with a stellar atmosphere model of 15\,000\,K (dotted line) and added to this is a flat disk with inner temperature 1400\,K, outer temperature of 600\,K, and zero inclination. This flat disk model is representative of the best possible fit to the data with this restricted geometry, yet cannot account for the IRAC fluxes. In the lower panel the stellar photosphere has been subtracted from the photometry, and this excess flux is plotted linearly with the same flat disk model, as well as a blackbody with $T=1000$\,K. The shaded regions represent the range of observed infrared fluxes over 11.2\,yr, where the minima at 3.6 and 4.5\,$\upmu$m are conceivably (just) consistent with a face-on, flat disk. \label{fig6}} \end{figure} \subsection{Timescale and Accretion Rate Considerations} The simplest model for the observed long-term infrared variability towards GD\,56 is the production and depletion of dust clouds that provide emitting surface areas in excess of -- {\em or in lieu of} -- a flat disk configuration. Such clouds must not change orbital radius drastically, as this would change the dust temperature and hence the flux ratio between 3.5 and 4.5\,$\upmu$m. A priori, any transient or periodic dust structure can be optically thick or optically thin, but it is noteworthy that blackbody dust with $T=1000$\,K \citep{jur07} will orbit near 1.46\,$R_{\odot}$ and thus somewhat outside the stellar Roche limit (for GD\,56, asteroid-like bodies with density 3\,g\,cm$^{-3}$ should disrupt only within 0.9\,$R_{\odot}$). Before exploring possible mechanisms for dust production and depletion around GD\,56 and other dusty white dwarfs, it is best to establish some theoretical context. A good starting point is to consider the relevant physical timescales for particulate and gaseous debris in the vicinity of the Roche limit for white dwarfs, where the orbital radii are $r\approx 1.0-1.5\,R_{\odot}$. The timescales that govern five processes are important: Keplerian orbits, condensation, collisions, Poynting-Robertson (PR) drag, and viscous spreading. The orbital period at these radii for a $0.6\,M_{\odot}$ star is between 3 and 7\,h, and hence solids will orbit at (circular) speeds exceeding 250\,km\,s$^{-1}$. Condensation of metallic gas onto grains is likely to be efficient in the case of identical chemical compositions, and take place on orbital timescales or less \citep{met12}. The collisional timescale for dust in a disk-like configuration is directly comparable to the orbital period (i.e.\ $t_{\rm coll} \sim P/4\uppi\uptau$), with the optical depth $\uptau \gtrsim1$ in the case of a vertically opaque disk, and $\uptau \simeq L_{\rm IR}/L_*$ if the disk is optically thin in the radial direction. A disk may be optically thick in the radial direction (due to grazing starlight), yet remain vertically optically thin and experience less frequent collisions than implied by $f = L_{\rm IR}/L_*$. For white dwarfs with well-determined fractional dust luminosities \citep{roc15}, if their vertical optical depths are comparable to $f$, then collision timescales will be on the order of $1-20$\,d. This is at least two orders of magnitude shorter than the timescale for angular momentum loss by solids due to PR drag, where for 1\,$\upmu$m size grains with density 3\,g\,cm$^{-3}$, the timescales for GD\,56 are between 6 and 14\,yr. For the population of known dusty white dwarfs, all with $L_{\rm IR}/L_* \gtrsim 10^{-3}$, it can be shown that for $\uptau > 10^{-5}$, the collisional timescale is {\em always shorter than PR drag} \citep{far08}. Before discussing the timescale for the viscous spreading of gas, the implications of collisions are considered. In the restricted case of disks that are completely optically thin, the relevant timescale for solids is thus the collisional timescale. For GD\,56 this is a few orbital periods and less than 8\,h within 1.5\,$R_{\odot}$, requiring that the entire disk be replenished every few days or so, and with a dust mass production rate sufficient to maintain $L_{\rm IR} /L_* \approx0.03$ \citep{roc15}. This entails around 10$^{18}$\,g \citep{jur09} per collision timescale, which would eventually result in a mass accretion rate onto the star that is higher than 10$^{12}$\,g\,s$^{-1}$, and consistent with numerical simulations of collisional cascades around white dwarfs \citep{ken17a}. Such high rates of (sustained or transient) disk accretion have been suggested by theory and hinted at via observations \citep{raf11b,gir12,ken17b}, but so far have not been confirmed by (limited) X-ray data \citep{far18}. Furthermore, the maximum accretion rate from PR drag alone acting on optically thin debris is set only by the stellar luminosity and the fraction of starlight intercepted by the disk \citep{boc11}; if all radiation from GD\,56 is intercepted by dust, the maximum mass accretion rate is $3\times10^{10}$\,g\,s$^{-1}$. Thus if disks are primarily optically thin, and accretion rates do exceed this value, then collisions dominate the production of gas that eventually accretes and pollutes the star, but material must be continually supplied to the disk \citep{ken17a,ken17b}. In the context of GD\,56, these potential disk implications can be compared with the stellar atmospheric data. Based on the photospheric calcium and magnesium abundances and diffusion theory \citep{koe09}, the current total stellar accretion rate is $(0.8\pm0.2)\times10^8$\,g\,s$^{-1}$ using either of these two elements at its bulk Earth mass fraction \citep{all01}. These estimates are orders of magnitude smaller than any of the above predictions for optically thin dust shell depletion -- especially where collisions may dominate the production of gas -- but agree well with model predictions for flat and opaque dust disks depleting via PR drag (i.e. \ no collisions \citep{raf11a,boc11}). Notably, this accretion rate is typical for polluted stars with metal diffusion timescales of less than a few years, where an ongoing steady state can be confidently inferred -- based on diffusion theory, no system to date has been determined to substantially exceed $10^9$\,g\,s$^{-1}$ \citep{far16b}. If correct, then any long-term ($10^4-10^7$\,yr; \citealt{gir12}) infrared emission in dusty white dwarf systems is likely to be from solids in a flat disk, and consistent with dynamical relaxation on orbital timescales, where a sufficient component is optically thick. This represents the basic utility of the canonical flat and opaque disk model; immunity to collisional annihilation, longevity, and accretion rates that broadly match those of diffusion theory \citep{jur03,raf11a,met12}. [It is noteworthy that recent simulations by \citet{ken17a,ken17b} cannot produce a vertically thin disk of material at relevant orbital distances for dusty white dwarfs, whereas analytical work \citep{boc11,met12} appears to favor this configuration.] While collisions can be avoided by efficient damping in a flat and opaque disk \citep{far08,met12}, once solids have been sublimated or collisionally vaporized, regardless of geometry, the material will no longer evolve via radiation forces. Prior to the accretion of material onto the stellar surface, the final relevant timescale may be set by gas viscosity. And while the properties of pure gas disks orbiting white dwarfs have been discussed at length in the literature \citep{jur08, raf11a,far12b,met12,ken17b}, the implications for infrared and optical variability have not been fully explored. Several estimates exist for gas viscosity in a disk dominated by metals, where the two uncertain factors are in the size of the $\upalpha$ parameter \citep{sha73} and whether the disk is ionized or neutral. For bulk Earth material, the mean atomic weight of neutral gas would be 32.8\,$m_{\rm p}$, and if this material is fully ionized then 16.4\,$m_{\rm p}$. Using these two extrema yields viscous timescales in the range $(10/\upalpha$)\,yr $\la t_{\upnu}\la(20/\upalpha$)\,yr for a typical white dwarf disk at $r=0.2\,R_{\odot}$. Only in the limit $\upalpha\approx1$ can the viscous timescale approach that of PR drag, and only then for grains larger than 1\,$\upmu$m \citep{raf11a}. Therefore, a significant time lag may result between a change in the rate of solids delivered into an inner gas disk, and a subsequent change in the mass accretion rate from this reservoir. If $\upalpha\approx0.1-0.4$ as inferred for a variety of fully ionized accretion disks \citep{kin07}, then the shortest realistic timescale would be $t_{\upnu}\approx 25-100$\,yr. Therefore any changes in dust and gas production within a disk, including the inward movement of solids that eventually sublimate, can be mediated prior to accretion onto the stellar surface if an $\upalpha$-type gas disk is present. Any such inner disk will smooth over changes on these relatively long timescales, and no changes in accretion rate would be expected on PR drag timescales. While the presence of such inner gas disks orbiting white dwarfs is empirically unconstrained, if they are ubiquitous then changes in accretion rate might not be expected on human timescales. There are a handful of circumstellar gas detections via absorption towards polluted stars; most dramatically around WD\,1145 in the optical \citep{xu16,red17}, and possibly via a weak circumstellar component in the Ca\,{\sc ii}\,H and K lines of EC\,11246--2923 \citep{deb12a}, but also in the far-ultraviolet spectra of both SDSS\,1228+1040 and PG\,0853+516 \citep{gan12}. In none of these cases do the gas velocities approach that expected for free-fall, and hence the behavior of disk material prior to reaching the stellar surface is unknown. An $\upalpha$ disk reservoir of gas is consistent with the unchanging metal lines in GD\,56 and SDSS\,0959, despite the changes observed in infrared emission \citep{xu14}, and regardless of the exact dust production or depletion mechanisms. Because this study considers non-canonical disk geometry and evolution for GD\,56, it is worthwhile to mention the one case where a non-canonical configuration has been established. The {\em circumbinary disk} orbiting the polluted white dwarf SDSS\,J155720.77+091624.6 (hereafter, SDSS\,1557) and its companion must be optically thin and vertically extended to account for the infrared data \citep{far17}. The system exhibits a strong infrared excess that cannot be reproduced by emission from a cool companion, and prior to the discovery of its duplicity, the infrared emission was well-matched by a face-on, flat disk model \citep{far12a}. However, the orbiting companion dynamically precludes material in the region where such a flat disk would reside, and thus the $T\approx1100$\,K dust emission must originate from optically thin material at $r=3.3\,R_{\odot}$ \citep{far17}. SDSS\,1557 demonstrates the plausibility of a non-flat disk of material that continuously feeds atmospheric metals onto a white dwarf for at least several years. \subsection{Infrared Brightening and Dimming at GD\,56} A plausible scenario for infrared variability in white dwarf dust disks is the following. There are clouds of dust that are transiently or quasi-periodically appearing and disappearing in a circumstellar environment with an underlying disk that is substantially more massive. Such a picture can account for the rapid decline of infrared flux from SDSS\,0959, where the discovery observations were taken during the presence of additional clouds of material, and these were quickly re-assimilated into an existing disk. Any debris cloud generated by impact or collisions will have a range of orbits, {\em but all will intersect that of any underlying disk}, unless it occurs at the extremities. For a stray body, an impact onto the innermost edge of a particulate disk within the Roche limit is the least probable location, and requires a body or bodies smaller than a km that were immune to tidal disruption on any prior periastra \citep{bro17}. Unless optically thick, any transient dust cloud that is vertically offset from an extant disk plane would quickly attain a higher temperature than debris that is partly shielded from starlight. While transient dust clouds quickly being absorbed into an existing disk may be able to account for the sparse infrared data on SDSS\,0959, it may not work for GD\,56 where the infrared dimming takes place over many years. It is possible, however, that a simple (re-)incorporation of material above or below the plane of an underlying disk may be a process that acts like a damped oscillator. If so, then dust clouds with relatively small masses may quickly become assimilated into a larger disk if damping is near critical, but the system may oscillate if damping is inefficient or the dust cloud mass is larger. In the canonical opaque disk, damping is sufficiently high that collisions are negligible, and (re-)condensation of gas onto a particulate disk should occur on orbital timescales \citep{met12}. This scenario might account for the disappearance of gas emission lines around SDSS\,J161717.04$+$162022.4 \citep{wil14}. In the case of GD\,56, an important observational clue to the dust removal process is the fact that the flux ratios between 4.5\,$\upmu$m and 3.5\,$\upmu$m do not change substantially. This implies the dust is created, and later destroyed or subsumed, within a range of orbital radii that does not change substantially with time. Another interesting empirical indication is that the gradual decay in $3-5\,\upmu$m flux around GD\,56 appears to mimic that seen towards several bright debris disks orbiting young main-sequence stars, and where this decay is thought to be due to collisions of mm-size condensates \citep{men14,men15}. In these studies, the generation of extreme infrared excesses are attributed to planetary impacts. These impacts produce vaporized debris, which re-condenses and undergoes a collisional cascade that eventually depletes grains that emit efficiently in the infrared, aided by the removal of the smallest grains by radiation pressure. While the blowout of dust is not relevant for white dwarf systems, the fact that the collisions will dominate unless damped suggests the analogy here may be compelling. Removal of small grains around white dwarfs might occur via continued collisions and sputtering, but alternatively if they are heated to temperatures sufficient for rapid sublimation as may be the case for WD\,1145 \citep{xu18}, then there could be sufficient means to destroy dust in the vicinity of where it is produced. It is noteworthy that the channel 2:1 flux ratios do not appear to change substantially between the combined IRAC and {\em WISE} epochs reported for SDSS\,0959 \citep{xu14}. In further analogy with WD\,1145, it is plausible that minor bodies are generating clouds of $T\approx1000$\,K dust around 1.46\,$R_{\odot}$ either via collisions or sublimation -- but not tidal disintegration at this distance. It is now well established that the disk of dust and gas orbiting WD\,1145 is eccentric \citep{cau18}, and the same appears to be the case for at least two other disks \citep{man16a,den18,mir18}. This implies that tidal disintegration of a large body or bodies is not likely to be ongoing in a near-circular orbit at or near the Roche limit \citep{gur17,ver17}, but instead orbiting bodies are crossing in and out of this radius. Collisions associated with circularization of a disk may produce dust periodically, some of which may be located outside of the disk plane, and later this material becomes assimilated into the evolving disk. There are a few considerations that make this scenario plausible, as long as the dust production sites are relatively localized and not varying in orbital distance on the timescales probed by the observations. First, the circularization process is not well constrained, and material may end up over a range of orbits spanning the Roche limit. This could imply that intact bodies survive longer outside the Roche limit, and are thus available as sites of dust production. Second, the Roche limit separates disrupted and intact bodies of the same composition, but there may be a transition zone where some bodies remain intact while others are fragmented due to compositional gradients. Third, small bodies can be formed by recycled ring material at the Roche limit via disk spreading \citep{van18} -- as has been suggested for the rings of Saturn \citep{cha10} -- and this process might also play a role in the rise and fall of emitting dust. This latter scenario requires a massive disk in excess of a lunar mass, but at the same time GD\,56 has the largest known infrared excess of any dusty white dwarf. While somewhat speculative, the infrared data for GD\,56 may hint at an overall trend that includes a flux increase followed by a gradual decay. If there is a periodic signal here, it would be crudely around 7\,yr and then correspond roughly to an orbital scale of 3\,AU. In this case an infrared brightening might be expected before or near 2020. It is not unreasonable that some tidally-disrupted material remains in the process of circularization, with small bodies still strung out along a range of wide orbits \citep{deb12b,ver14}, and these periodically interact with the disk around GD\,56, generating dust that is eventually destroyed or re-incorporated. Another remote possibility is if a planetary body on a wider orbit has sufficient gravitational influence to induce a disk warp during periastron, it could also be that a warp may decay gradually as a damped oscillator (but in the case of GD\,56 it would need to be highly inefficient with a decay timescale longer than $10^3$\,d). White dwarf disk precession via general relativity can be on the order of years, with a steep dependence on radial distance \citep{mir18}, but it is unclear if this would drive a periodic warp or change in dust mass. More infrared data with better temporal coverage is required before any realistic model constraints can be made. Lastly and importantly, the dimming trend at GD\,56 after MJD 56\,000 appears to be consistent with the broad behavior predicted by numerical simulations of collisional cascades around white dwarfs \citep{ken17a}, and more generally the predictions of steady state evolution of debris disks due to collisions \citep{wya07}. In the former simulations, the basic behavior of the infrared dust luminosity is stochastic, yet switching between high states of visibility, and non-detectable low states. Steady states are possible, or even likely with continued mass input into an existing disk, but without such input an infrared excess will decay to non-detectable levels. Nevertheless, the decay trend seen in the data for GD\,56 is well within the family of predictions based on these models, and while solutions are degenerate, roughly speaking the model sets that can reproduce the infrared data have dust and gas production rates in excess of $10^{12}$\,g\,s$^{-1}$ and total belt masses in excess of $10^{23}$\,g \citep{ken17a}. While such rates are consistent with the estimate made above based on destroying all the infrared emitting dust around GD\,56 within a typical collision timescale, they are far in excess of that inferred to be ongoing via diffusion theory. Such high accretion rates, if they occur, have fundamental implications for the nature of the remnant planetary systems orbiting white dwarfs. For example, if 30\% of white dwarfs with cooling ages in the range $20-200$\,Myr accrete from disks \citep{koe14} that produce accretion rates as high as $10^ {13}$\,g\,s$^{-1}$, then a given system with a duty cycle of $0.3 \times 100$\,Myr can have accreted up to $10^{28}$\,g (2\,$M_{\oplus}$) of material. This is in stark contrast to the consumption of a mass closer to $10^{25}$\,g (Pluto) based on the more sedate accretion rates inferences made from diffusion theory \citep{koe09}. This potential tension between theory and observation cannot be resolved within the context of the current work, but the infrared behavior of GD\,56 is a strong indication that collisions are of fundamental importance in the debris disks that orbit polluted white dwarfs. | 18 | 8 | 1808.09967 |
1808 | 1808.10451_arXiv.txt | We modify the chemo-dynamical code GEAR to simulate the impact of self-interacting dark matter on the observable quantities of 19 low mass dwarf galaxies with a variety star forming properties. We employ a relatively high, velocity independent cross-section of $\sigma/m=10\,\rm{cm^2/g}$ and extract, in addition to integrated quantities, the total mass density profile, the luminosity profile, the line-of-sight velocities, the chemical abundance and the star formation history. We find that despite the creation of large cores at the centre of the dark matter haloes, the impact of SIDM on the {\it observable} quantities of quenched galaxies is indiscernible, dominated mostly by the stochastic build up of the stellar matter. As such we conclude that it is impossible to make global statements on the density profile of dwarf galaxies from single or small samples. Although based mostly on quenched galaxies, this finding supports other recent work putting into question the reliability of inferred cored density profiles that are derived from observed line-of-sight velocities. | 18 | 8 | 1808.10451 |
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1808 | 1808.02302_arXiv.txt | {Exoplanet research has shown an incessant growth since the first claim of a hot giant planet around a solar-like star in the mid-1990s. Today, the new facilities {are working to spot the first habitable rocky planets} around low-mass stars as a forerunner for the detection of the long-awaited Sun-Earth analog system. All~the achievements in this field would not have been possible without the constant development of the technology and of new methods to detect more and more challenging planets. After the consolidation of a top-level instrumentation for high-resolution spectroscopy in the visible wavelength range, a huge effort is now dedicated to reaching the same precision and accuracy in the near-infrared. Actually, observations in this range present several advantages in the search for exoplanets around M dwarfs, known to be the most favorable targets to detect possible habitable planets. They are also characterized by intense stellar activity, which hampers planet detection, but~its impact on the radial velocity modulation is mitigated in the infrared. Simultaneous observations in the visible and near-infrared ranges appear to be an even more powerful technique since they provide combined and complementary information, also useful for many other exoplanetary science~cases. } \keyword{exoplanets; multi-wavelength spectroscopy; radial velocity; instrumentation % } \begin{document} | The radial velocity (or Doppler) method allowed detecting the first extrasolar planet around a solar-like star, 51 Peg b \cite{1995Natur.378..355M}, opening the era of the quest for planets. The radial velocity (RV) technique is an ``indirect'' method to detect exoplanets since it reveals the reflex motion of the host star due to a hidden companion when both revolve around the common center of mass. The periodical variation of the stellar RV, obtained through the measurement of the Doppler shift of its spectral lines, is dependent on the characteristic of its small-mass companion. The modulation can be represented by a Keplerian function with a period equal to the orbital period of the planet, \textit{P}, and a semi-amplitude, \textit{K}, defined as~follows: \begin{equation} \label{eqn:K} K {\rm [m s^{−1}]}=\frac{28.4329{\rm [m s^{−1}]}}{\sqrt{1-e^2}}\frac{m_p \sin i}{M_J}\left( \frac{M_{\star}+m_p}{M_{\odot}}\right)^{-2/3}\left(\frac{P}{{\rm 1[yr]}}\right)^{-1/3}, \end{equation} where $M_{\star}$ is the stellar mass, $e$ the orbital eccentricity of the planet and $m_p \sin i$ is the lower limit of the planetary mass since the real mass depends on the orbital inclination, \textit{i}, unknown with the RV information alone. {The actual value of the planet mass can be obtained if the planet transits in front of its host star and a proper modeling of the observed transit light curve is performed. Just like the shape of the RV variation allows measuring a set of orbital parameters with Equation (\ref{eqn:K}), the same principle is valid for the transit light modulation, which can be modeled by specific functions presented in~\cite{2002ApJ...580L.171M}. The shape of the light curve is dependent, among other things, on the inclination of the planetary orbital plane through the so-called ``impact parameter'', \textit{b} (i.e., the projected separation between the center of the stellar disk and the center of the planet disk), according to the relation $b=\frac{a \cos i}{R_{\star}}$, where \textit{a} is the semi-major axis of the orbit and $R_{\star}$ is the stellar radius (see, e.g., \cite{2010arXiv1001.2010W}). Once the inclination is measured, the real mass of the planet can be obtained. This is just an example of the synergy among the different planet detection techniques (e.g., \cite{2013pss3.book..489W}).} The RV method is more sensitive to a large planetary mass and a short orbital period: actually, a~Jupiter-like planet orbiting at 1 AU from its host star produces an RV semi-amplitude of about 30~m~s$^{−1}$, while an Earth-like planet at the same distance induces an RV modulation of only 9 cm s$^{−1}$ (see also Table 1 in \cite{2010exop.book...27L} for an overview of the typical semi-amplitude values induced by different types of planets). An RV survey also requires an intense observational effort: the observing baseline should be long enough to sample at least one, up to many periods of the planetary companion. Current state-of-the-art spectrographs in the visible (VIS) range provide an RV accuracy of about 1~m~s$^{-1}$ (see Section \ref{sec:vis}), allowing one in principle to detect small-mass exoplanets around solar-type stars with relatively short orbital periods. Anyway, the search for low-mass planets is hampered by the presence of the stellar activity that induces an intrinsic RV variation with similar amplitude with respect to the amplitude of the Keplerian signal, or even larger. {Recently, M dwarfs have been identified as the most interesting objects to investigate, because of the higher planet occurrence rate in their habitable zones (Section \ref{nir}) and the more favorable ratio between planet and stellar masses. On the other hand, they are red stars, fainter in the VIS wavelength range.} To overcome the limitations imposed by the stellar activity {and to allow a proper measurement of their spectra}, observations in the near-infrared (NIR) range have been proposed, since in this band, the RV modulations caused by intrinsic stellar phenomena are expected to be mitigated with respect to the VIS. For sure, observations of M dwarfs (or active stars, in general) with dedicated NIR spectrographs provide major advantages, but obtaining simultaneous VIS-NIR observations would allow reaching the next level, in terms of the completeness and ``speed'' of the information (Section \ref{sec:goal}). Nowadays, astronomical research greatly benefits from multi-wavelength observations of the same phenomenon. In the field of exoplanets, particular attention to the NIR band has been payed in the last decade for the characterization of exoplanet atmospheres, both with the transit and direct imaging techniques (even in combination with the VIS range), while high-resolution high-precision spectroscopy in the NIR is still in its infancy, because of technical and technological issues that have been solved only recently. Today, thanks to to the efforts of coordinated teams in the creation of instruments like CARMENES or GIARPS (Section \ref{new}), it is possible to exploit combined VIS-NIR spectroscopic observations and start to investigate a parameter space of the exoplanet research previously not accessible. In this review, I will draw an overview of the historical, technical and scientific scenarios that led to the current achievements in the framework of multi-band high-resolution spectroscopy in the field of exoplanet search. Those results represent the foundations of the future new-generation~facilities. | The new scientific questions in the framework of the exoplanet search and characterization can be summarized as follow: ({i}) the search for Earth-mass rocky planets in the HZ of M dwarfs; ({ii}) the identification of the origin of planetary system diversity through the detection of planets around young stars; ({iii}) the characterization of the hot gas giant planets atmospheres as a laboratory for the future characterization of rocky habitable planets with ELTs or space-based telescopes. The simultaneous multi-band observations in the VIS and NIR ranges are expected to be the forthcoming parameter space for these new issues. The full monitoring of stellar activity represents the key point in the detection of exoplanets around targets with enhanced and peculiar activity such as the ones considered in ({i}) and ({ii}). On the other hand, the opportunity to explore molecular features of the planetary atmospheres, present over a wide interval of wavelengths, is now possible thanks to instruments covering the range between the bluer part of the VIS up to the edge of the NIR ({iii}). Instruments like CARMENES, GIARPS and the forthcoming NIRPS + HARPS configuration are expected to yield a significant contribution in these and in many other fields of astrophysics. They also represent the starting point for future generation spectrographs that are going to equip new astronomical facilities. {If we observe a planetary (period-mass) diagram} (or, alternatively, a (semi-major axis-mass) diagram, e.g., Figure 1 in \cite{2013Sci...340..577S}) comparing the properties of the known exoplanets with the planets of our Solar System, it is clear that the current technology is still not sufficiently adequate to find planets like our Venus, Mars, Neptune, etc., or the Earth itself. The need to reach the 1-m s$^{-1}$ (and less) RV accuracy is driven by the goal to find an Earth-like planet around a Sun-like star. As mentioned in this review, there are several impact factors on the achievable RV precision with high-resolution spectroscopy (both VIS and NIR): changes in the environmental conditions require a series of technical devices ensuring a high-level of instrumental stability; the wavelength calibration must rely on a very stable and reliable wavelength reference to reach the required precision; RV extraction methods must be more and more sophisticated. Other important technical elements can be the availability of state-of-the-art detectors and devices that allow a uniform illumination of the spectrograph slit. Last, but not least, from the observational point of view, the capability to conduct a proper data sampling of the RV signal helps to avoid the aliasing phenomenon that hampers the frequency (period) analysis. We are quite confident to say that all of these issues have been successfully addressed, at least in the visible range, at least to obtain the 1-m s$^{-1}$ regime. As previously stated, the technology transfer from VIS to NIR is far from immediate, but many efforts have been already made to allow precise RV also in the NIR domain, starting from the lessons learned in the last few decades. At this moment, we are approaching the new milestone of the Doppler method, since we are waiting for the precise NIR RV measurement from CARMENES (1 m s$^{-1}$ expected), on the one hand, and the outstanding RV precision of 10 cm s$^{-1}$ from ESPRESSO, on the other. On the scientific side, the challenge is to be able to manage all the other detectable effects, first of all, the stellar activity. Even in this case, valuable examples of data treatment are presented in literature, and it will be very interesting to see how they will work in the new ultra-precise RV domain. Reaching the cm s$^{-1}$ regime is a necessary step of the journey, but it is also crucial to know how to exploit this opportunity. In the near future, these new facilities will necessarily be more and more specialized for RV monitoring for the search for exoplanets (some of them are already employed in such a way); just think about the follow-up of the forthcoming space satellite dedicated to transit detection. Actually,~the~cooperation among different detection methods (and thus, instruments and projects) will probably be the best way to characterize a star-planet system. In this context, an important role will be played by all the tools that allow obtaining a more precise estimate of the stellar parameters like mass and radius (e.g.,~stellar~modeling and asteroseismology), since they are crucial to constrain the global planet properties. \vspace{6pt} \funding{This research received no external funding.}% | 18 | 8 | 1808.02302 |
1808 | 1808.02905_arXiv.txt | We show that gravitational floating orbits may exist for black holes with rotating hairs. These black hole hairs could originate from the superradiant growth of a light axion field around the rotating black holes. If a test particle rotates around the black hole, its tidal field may resonantly trigger the dynamical transition between a co-rotating state and a dissipative state of the axion cloud. A tidal bulge is generated by the beating of modes, which feeds angular momentum back to the test particle. Following this mechanism, an extreme-mass-ratio-inspiral (EMRI) system, as a source for LISA, may face delayed merger as the EMRI orbit stalls by the tidal response of the cloud, until the cloud being almost fully dissipated. If the cloud depletes slower than the average time separation between EMRI mergers, it may lead to interesting interaction between multiple EMRI objects at comparable radii. Inclined EMRIs are also expected to migrate towards the black hole equatorial plane due to the tidal coupling and gravitational-wave dissipation. Floating stellar-mass back holes or stars around the nearby intermediate-mass black holes may generate strong gravitational-wave emission detectable by LISA. | Black Hole (BH) No Hair Theorem states that any stationary black hole in Einstein-Maxwell theory can be characterized by its mass, spin and electric charge, which is possible to be tested with BH spectroscopy in the Advanced LIGO (Laser Interferometric Gravitational-Wave Observatory) era \cite{berti2016spectroscopy,yang2017black,brito2018black,Thrane_2017}. If additional { bosonic} fields are allowed in the setup, they may grow exponentially according to the BH superradiance \cite{Detweiler:1980uk,zouros1979instabilities} and saturate onto quasi-stationary configurations \cite{east2017superradiant,east2017superradiant2}. In particular, these hair fields (such as the QCD Axion \cite{weinberg1978new}, dark photons \cite{holdom1986two,cicoli2011testing} and string axiverse \cite{arvanitaki2010string}) around BHs may serve as Dark Matter candidates, and depending on their mass range, they could be dynamically important to the spin evolution of isolated BHs. The rotation of these fields may also generate continuous gravitational waves (GWs) that lie in the detection band of LIGO or LISA (Laser Interferometric Space Antenna) \cite{arvanitaki2015discovering,baryakhtar2017black,Brito:2017zvb,brito2017stochastic}. The rotating cloud can carry a significant fraction of energy/angular momentum (AM) of the host BH. Since the BH area generally increases following the superradiant growth of the cloud \cite{east2017superradiant2}, while interacting with an external agent, the cloud AM would not be entirely re-absorbed by the host BH (e.g., through the tidally-induced cloud depletion discussed in \cite{Baumann:2018vus}), or its horizon area would decrease. As a result, the external agent must acquire part of the cloud energy/AM during the interaction process. This AM transfer may give rise to {\it gravitational floating orbits} of a test particle, in which case the GW damping of the orbital energy and AM is balanced by the gravitational interaction with the cloud. Such orbits are first conjectured in \cite{press1972floating}, based on the observation that the horizon AM flux generated by a test particle orbiting around a rotating BH could be negative due to the superradiance effect. However, for Kerr BHs the AM gain from horizon is universally weaker than the loss due the GW radiation at infinity, which means that there is no gravitational floating orbit in Kerr spacetime. On the other hand, if the particle also couples to a massive scalar field besides the gravitational interaction, it has been shown \cite{Cardoso:2011xi,Ferreira:2017pth} \footnote{The argument of \cite{Ferreira:2017pth} is drawn in analogy to planetary systems, and a complete analysis including the backreaction on the cloud is required to prove the existence of positive AM transfer during resonances. The resonance studied here operates at lower frequency, and is still valid for complex scalar field.} that the scalar wave radiation can balance the GW radiation, and lead to floating orbits given suitable scalar field mass and coupling strength. \begin{figure}[tbp] \centering \includegraphics[width=0.4\textwidth]{cartoon2.pdf} \caption{A host BH of mass $M$ and dimensionless spin $a$ dressed with an axion cloud and companied by a star of mass $M_*$ at $R_*$. The axion cloud develops by superradiance, and is quasi-stationary without the companion star. Through the tidal interaction with the star, a bulge of the cloud develops, which {\it leads} the motion of the star. In return, the star extracts AM from the BH and the cloud to compensate the AM loss due to GW radiation. The star would float at this orbit until the whole cloud is depleted.} \label{fig:cartoon} \end{figure} In this paper we show that indeed the tidal interaction between a rotating cloud and a test particle could support gravitational floating orbits, without assuming additional axion field-matter interactions. Physically the test particle tidally deforms the cloud. Due to the cloud dissipation, there is a phase difference between the particle's orbit and the tidal bulge. Unlike the tidal interactions commonly seen in binary stars, the tidal bulge in the cloud actually leads the test particle's motion, and consequently AM transfers from the cloud to the particle. We examine this cloud energy/AM transfer mechanism in the context of EMRIs, which are important sources for LISA. We find that for a range of EMRI mass ratio and axion mass, the EMRI orbit stalls at finite radius until the axion cloud is depleted. Notice that this process could take longer than the inspiralling time of the EMRI, which implies interesting astrophysical effects. Unless { specified}, we set $G = c =\hbar =1$. | 18 | 8 | 1808.02905 |
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1808 | 1808.10517_arXiv.txt | Early energy injection to the Cosmic Microwave Background~(CMB) from dissipation of acoustic waves generates deviations from the blackbody spectrum not only at second-order but also at third-order in cosmological perturbations. We compute this new spectral distortion $\mathcal \kappa$ based on third-order cosmological perturbation theory. We show that $\kappa$ arises from heat conduction and shear viscosity of spectral distortions and temperature perturbations. The ensemble average of $\kappa$ can be directly sourced by integral of primordial non-Gaussianity, and thus depending on its shape. For local non-Gaussianity we roughly estimate the signal and find $\kappa=f^{\rm loc.}_{\rm NL}\times \mathcal O(10^{-18})$. The signal is incredibly tiny; however, we argue that it carries a specific frequency dependence different from other types of CMB spectral distortions. Moreover, we comment on other possible applications of our results. | 18 | 8 | 1808.10517 |
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1808 | 1808.05008_arXiv.txt | {The SPHERE ``planet finder'' is an extreme adaptive optics (AO) instrument for high resolution and high contrast observations at the Very Large Telescope (VLT). We describe the Zurich Imaging Polarimeter (ZIMPOL), the visual focal plane subsystem of SPHERE, which pushes the limits of current AO systems to shorter wavelengths, higher spatial resolution, and much improved polarimetric performance.} {We present a detailed characterization of SPHERE/ZIMPOL which should be useful for an optimal planning of observations and for improving the data reduction and calibration. We aim to provide new benchmarks for the performance of high contrast instruments, in particular for polarimetric differential imaging.} {We have analyzed SPHERE/ZIMPOL point spread functions (PSFs) and measure the normalized peak surface brightness, the encircled energy, and the full width half maximum (FWHM) for different wavelengths, atmospheric conditions, star brightness, and instrument modes. Coronagraphic images are described and the peak flux attenuation and the off-axis flux transmission are determined. Simultaneous images of the coronagraphic focal plane and the pupil plane are analyzed and the suppression of the diffraction rings by the pupil stop is investigated. We compared the performance at small separation for different coronagraphs with tests for the binary \object{$\alpha$ Hyi} with a separation of 92~mas and a contrast of $\Delta m \approx 6^m$. For the polarimetric mode we made the instrument calibrations using zero polarization and high polarization standard stars and here we give a recipe for the absolute calibration of polarimetric data. The data show a small ($<1$~mas) but disturbing differential polarimetric beam shifts, which can be explained as Goos-H\"ahnchen shifts from the inclined mirrors, and we discuss how to correct this effect. The polarimetric sensitivity is investigated with non-coronagraphic and deep, coronagraphic observations of the dust scattering around the symbiotic Mira variable \object{R Aqr}.} {SPHERE/ZIMPOL reaches routinely an angular resolution (FWHM) of $22-28$~mas, and a normalized peak surface brightness of ${\rm SB}_0-m_{\rm star}\approx -6.5^m$~arcsec$^{-2}$ for the V-, R- and I-band. The AO performance is worse for mediocre $\gapprox 1.0''$ seeing conditions, faint stars $m_R\gapprox 9^m$, or in the presence of the ''low wind'' effect (telescope seeing). The coronagraphs are effective in attenuating the PSF peak by factors of $>100$, and the suppression of the diffracted light improves the contrast performance by a factor of approximately two in the separation range $0.06''-0.20''$. The polarimetric sensitivity is $\Delta p<0.01$~\% and the polarization zero point can be calibrated to better than $\Delta p\approx 0.1$~\%. The contrast limits for differential polarimetric imaging for the 400~s I-band data of R Aqr at a separation of $\rho=0.86''$ are for the surface brightness contrast ${\rm SB}_{\rm pol}({\rho})-m_{\rm star}\approx 8^m\,{\rm arcsec}^{-2}$ and for the point source contrast $m_{\rm pol}({\rho})-m_{\rm star}\approx 15^m$ and much lower limits are achievable with deeper observations.} {SPHERE/ZIMPOL achieves imaging performances in the visual range with unprecedented characteristics, in particular very high spatial resolution and very high polarimetric contrast. This instrument opens up many new research opportunities for the detailed investigation of circumstellar dust, in scattered and therefore polarized light, for the investigation of faint companions, and for the mapping of circumstellar H$\alpha$ emission.} | \label{Introduction} The SPHERE ``Planet Finder'' instrument has been successfully installed and commissioned in 2014 at the VLT. The main task of this instrument is the search and investigation of extra-solar planets around bright stars $m_{\rm R}\lapprox 10^m$. Therefore SPHERE is optimized for high contrast and diffraction limited resolution observation in the near-IR and the visual spectral region using an extreme adaptive optics (AO) system, stellar coronagraphs, and three focal plane instruments for differential imaging. General technical descriptions of the instrument are given in \citet{Beuzit08,Kasper12}, and the SPHERE user manual and related technical websites\footnote{www.eso.org/sci/facilities/paranal/instruments/sphere} of the European Southern Observatory (ESO). SPHERE is a very powerful facility instrument which provides a broad suite of sophisticated instrument modes for the very demanding investigation of extra-solar planetary systems. Essentially all of these modes also provide unique observing opportunities for the study of the immediate circumstellar environment of bright stars. Technical results about the on-sky performance of the SPHERE instrument are given in \citet{Dohlen16}, and on-sky results for the AO-system are described in \citet{Fusco16} and \citet{Milli17}. A series of first SPHERE science papers demonstrates the performance of various observing modes of this instrument \citep[e.g.,][]{Vigan16, Maire16a, Zurlo16, Bonnefoy16}. However, the SPHERE instrument is complex and therefore it is appropriate to give more specific descriptions on individual subsystems and this is the first of a few technical papers for the visual focal plane instrument ZIMPOL. ZIMPOL, the Zurich Imaging Polarimeter, works in the spectral range from 500 nm to 900 nm and provides, thanks to the SPHERE AO system and visual coronagraph, high resolution ($\approx$ 20-30 mas) and high contrast imaging and imaging polarimetry for the immediate surroundings ($\rho<4$ arcsec) of bright stars. SPHERE/ZIMPOL includes a very innovative concept for high performance imaging polarimetry using a fast modulation - demodulation technique and it is tuned for very high contrast polarimetry of reflected light from planetary system. Beside this it can also be used as a high contrast imager offering angular differential imaging and simultaneous spectral differential imaging. Previous publications on ZIMPOL describe the science goal \citep{Schmid06a}, the expected performance \citep{Thalmann08}, and give reports about the concept of ZIMPOL \citep{Gisler04, Joos07,deJuanOvelar12}, the instrument design and component tests \citep{Roelfsema10,Roelfsema11,Pragt12,Bazzon12, Schmid12} and system testing \citep{Roelfsema14,Roelfsema16}. Some early science results based on SPHERE/ZIMPOL observations are given in \citet{Thalmann15,Garufi16,Kervella16,Stolker16,Khouri16,Avenhaus17, Ohnaka17a,Engler17}. \citet{Schmid17} also gives technical information about H$\alpha$ imaging and the flux calibration of ZIMPOL data. Many technical aspects must be considered for carrying out well optimized observations and calibrations with an instrument like ZIMPOL, which combines diffraction-limited imaging using extreme adaptive optics, coronagraphy, and differential techniques like polarimetry, or angular and spectral differential imaging. It is not possible to cover all these topics in detail in one paper and therefore we focus on a basic technical description and on aspects which are special to SPHERE/ZIMPOL when compared to other high contrast instruments. This should serve as a starting point for potential SPHERE/ZIMPOL users to carry out well optimized observations and data analyzes for exploiting the full potential of this instrument. We plan that subsequent papers will address other aspects of SPHERE/ZIMPOL, such as astrometry, precision photometry, a detailed technical assessment of the high performance polarimetry mode and more. This paper is organized as follows. The next section gives a brief overview on the SPHERE common path and a detailed description of the ZIMPOL subsystem, imaging properties, the ZIMPOL polarimetry, the detectors and detector calibrations, and the filters. Section 3 characterizes the ``typical'' point spread functions (PSFs) and describes special cases, like faint stars, poor atmospheric conditions, or particular instrumental effects. The topic of Sect.~4 is the SPHERE visual coronagraph and the comparison of coronagraphic test measurements taken with different focal plane masks. SPHERE/ZIMPOL polarimetry is described in detail in Sect.~5 including the concept for the control of the polarimetric signal and the correction of the measurements based on calibrations of the telescope, the instrument, and the detectors. Further we discuss the polarimetric differential beamshift, a disturbing effect which is new for astronomical optics and which was not anticipated in the design of this instrument. Then, we illustrate the very good polarimetric performance of ZIMPOL with test observations of the system R Aqr. We conclude in Sect.~6 with a summary of the most outstanding technical properties of SPHERE/ZIMPOL and an outline of the new research opportunities offered by this instrument. | This paper presents a detailed description of the SPHERE/ZIMPOL instrument to support and promote scientific investigations based on new or archival observations. The provided information should also be helpful to optimize observing and data reduction strategies, for an evaluation of the performance of this instrument for possible upgrades \citep[e.g.,][]{Lovis17}, or the design of more advanced systems for future telescopes like the ELT \citep{Kasper13,Keller10}. High contrast and high resolution observations from the ground require the combination of an adaptive optics system with coronagraphy, and a differential imager for speckle noise suppression. All these methods are tricky, because the resulting performance depends on the target brightness, the atmosphere, the selection of the coronagraph, the mode of the differential imager, and on the appropriate observing strategy. It is not possible to describe such a complex system and its performance in a single paper, or a compact user manual. This paper provides a comprehensive hardware description, and highlight three important aspects, which are particularly special for the SPHERE/ZIMPOL instrument when compared to other AO-imaging systems: (i) the characterization of the PSF-properties in the visual provided by the SPHERE AO system and the VLT telescope, (ii) the performance of the visual coronagraph, and (iii) the polarimetric measuring strategy and data characteristics of the ZIMPOL system. More technical information about ZIMPOL will become available in the future: a paper on the astrometric calibration is in preparation, and a description of photometric parameters is planned. \citet{Schmid17} describe many technical aspects on the H$\alpha$ imaging and absolute photometric measurements in several filters. Further, we expect that useful performance characteristics of SPHERE/ZIMPOL can be extracted from science papers, which aim for accurate measurements or push the performance of this instrument to the limits. \subsection{Key performance properties of ZIMPOL} SPHERE/VLT is one of the new generation extreme AO-systems available at large telescopes. The quality of the obtained PSFs depend strongly on the AO-system, the atmospheric conditions and the guide star magnitude $m_R$ and Table~\ref{TabPSFAO} characterizes the visible performance of the AO system for several typical cases using simple measuring parameters, e.g. the normalized PSF peak flux ct$_{\rm n6}(0)$. Good corrections with ct$_{\rm n6}(0)/10^6>0.3~\%$ (Strehl ratio $\gapprox 15$~\% in the R-band) are achieved for stars down to $m_{\rm R}\approx 8^m$ with an atmospheric seeing of $\approx 1''$. For bright stars $m_{\rm R}<7^m$ and good seeing conditions $<0.8''$ the relative peak flux is twice as good ct$_{\rm n6}(0)/10^6>0.6~\%$ and this corresponds to Strehl ratios of $\approx 30-50~\%$ in the R-band \citep[see also][]{Fusco16}. This provides for SPHERE/ZIMPOL a point source contrast performance of $\Delta m \approx 12.5^m$ outside $0.15''$ for the polarized flux (Table~\ref{TabRAqrpol}). The existing ADI-data of the PZ Tel binary \citet{Maire16a} indicate, that similar contrasts performance are also possible for the intensity signal of a point source. For extended polarized emission, the system reaches outside of $0.15''$ surface brightness contrast limits better than $C_{\rm SBpol} > 7.5^m{\rm arcsec}^{-2}$. The most outstanding property of SPHERE/ZIMPOL is the spatial resolution of up to 20~mas surpassing all other imaging instruments available at the VLT \citep{deZeeuw16}. This limit can even be improved with sparse aperture masking, which is one of the most recent upgrades of the SPHERE/ZIMPOL system \citep{Cheetham16}. Observations with even higher resolution requires currently interferometric observations. Visual extreme AO-systems are also available at other observatories, e.g. the pioneering MagAO-system at the Magellan telescope \citep{Close13}, or SCExAO at the Subaru telescope \citep{Jovanovic15}. Scientifically important features shared by all these visible AO systems is the access to the strong H$\alpha$ emission line, and the extension of high resolution observations from the traditional near/mid-IR to the visible wavelength range \citep[see][]{Close16}. There will be some healthy competition between these systems, but much more important is the mutual benefit in establishing common calibration targets, improving and checking measuring strategies and data reduction, and enhancing the science return thanks to complementary performance characteristics. Many properties of the ZIMPOL system are quite typical for high resolution imagers using AO and they are not repeated here. We list here the special features of ZIMPOL: \begin{itemize} \item{} a small detector pixel scale of $3.6~{\rm mas}\times 3.6~{\rm mas}$, giving 77160~pix arcsec$^{-2}$ for each detector and a high full well photo-electron capacity of $7\cdot 10^5$ e$^-$/pix, allowing for efficient high contrast observations of the circumstellar regions of bright stars $m_R<8^m$ in broad band filters, \item{} dual beam imaging and polarimetry for simultaneous observation with two cameras with many different options for the combination of filter bands, for example the combination of 1~nm or 5~nm H$\alpha$ filters with a H$\alpha$ continuum filter for spectral differential imaging. Alternatively, also narrow band filters can be combined with broad band filters for simultaneous, and therefore accurate flux measurements of the bright component in a high contrast system. For example, a bright star can be observed with the narrow Cnt820 or Cnt\_Ha filters in one arm, and the faint companion or circumstellar features in the second arm with the corresponding broad band filter I\_PRIM or R\_PRIM, respectively, \item{} a high resolution imaging polarimeter based on a fast modulation-demodulation technique for speckle noise suppression allowing high precision, $\Delta p \lapprox 10^{-4}$, broad-band polarimetry, \item{} an innovative concept for the compensation and control of the instrument polarization effects providing absolutely calibrated linear polarization measurements with an accuracy of about $\pm 0.1~\%$ for the relative Stokes parameters $Q/I$ and $U/I$, despite the fact that SPHERE/ZIMPOL is a polarimetrically complex instrument at the Nasmyth focus. \item{} the status of SPHERE/ZIMPOL as an ESO/VLT facility instrument, which ensures a well monitored and characterized system, steady improvements of the instrument operation, high standards for the execution of the observations and data calibration, a well established and user-friendly data archiving, and hopefully a long life-time. Another huge benefit is the large ESO user community using this system for a wide range of science targets with innovative observing strategies. \end{itemize} \subsection{Possible instrument upgrades} Some evolution of the ZIMPOL system capabilities can be expected in the near future, because they can be realized without much effort. Very desirable would be the availability of the ZIMPOL off-axis fields shown in Fig.~\ref{Figfields}, which are currently not offered as user mode. This would provide the possibility to extend the field of view for more extended objects to a diameter of $8''$, matching better the $11''\times 12.5''$ field offered by the infrared channel SPHERE/IRDIS. A very useful upgrade would be the availability of a low read-out noise mode for imaging, equivalent to the slow modulation polarimetric mode with low detector gain. Especially, low flux measurements taken in narrow band and line filters, as well as observations in the off-axis fields outside the light halo of the bright central star would profit because they are often read-out noise limited. Just taking long integrations $> 5$~min with a running AO system is critical because a few seconds of strongly reduced AO-performance or even an AO open loop, caused for example by a particularly bad atmospheric turbulence event, can degrade an entire long integration with an enhanced background of light from the central star. Another problem of long integrations are small pointing drifts of $\gapprox 10$~mas which cause image smearing and this would be avoidable with shorter integrations and a realignment of individual frames in the data reduction. Another type of quite easy instrument upgrades are new filters or coronagraphic masks in the exchange wheels of ZIMPOL or the visual coronagraph, respectively. New science cases may emerge, which call for filter changes, or better coronagraphic concepts could be implemented \citep[e.g.,][]{Patapis18}. Of course, most relevant would be any upgrades to the AO system or any other effort to improve the AO performance, like the suppression of the low wind effect with changes to the telescope \citep{Sauvage16b}. Better wavefront corrections for SPHERE based on new software or hardware would be particularly beneficial for the short-wavelength ZIMPOL subsystem where typical Strehl ratios are $\lapprox 50$~\% with quite some room for improvements. \subsection{New research opportunities offered by SPHERE/ZIMPOL} The special technical properties of SPHERE/ZIMPOL offer many new research opportunities, some are unique to this instrument and some are shared with the other visible light AO systems mentioned above. At least, the visual ZIMPOL observations are complementary to the data from the SPHERE near-IR focal plane instruments IRDIS and IFS, or other near-IR AO systems like for example Gemini/GPI \citep{Macintosh14} or Subaru SCExAO \citep{Jovanovic15}, by extending the wavelength domain of high resolution and high contrast imaging towards shorter wavelengths. In the following, we provide an incomplete list of science topics where SPHERE/ZIMPOL is already providing or will provide interesting or even very important contributions in high contrast imaging. In this discussion on science opportunities one should not forget the technical requirement that SPHERE/ZIMPOL observations need a bright central star $m_R\lapprox 9^m$ for the AO wavefront sensing and that the system provides only a quite limited field of view $\rho<4''$. \subsubsection{Search for extra-solar planets in reflected light.} ZIMPOL was selected by the European Southern Observatory for the SPHERE VLT ``planet finder'' instrument with the mandate to explore the detection limits of high contrast polarimetric imaging for the search of reflected light from extrasolar planets \citep{Schmid06a}. The aim of this unique instrument is to reach a contrast limit between the polarized flux of the planet $p_{\rm pl}\times I_{\rm pl}$ and the total flux of a star $I_{\rm star}$ of \begin{displaymath} C_{\rm pol}={{p_{\rm pl}\times I_{\rm pl}}\over I_{\rm star}}=10^{-8} \end{displaymath} within an angular separation smaller than $\rho<1''$. This limit would allow a detection of a Jupiter-sized giant planet or even terrestrial planets with a physical separation of $0.5-1$~AU around one of the nearest bright stars \citep{Thalmann08,Milli13}. This aim defined many of the ZIMPOL design decisions and therefore this imager is tuned for the detection of very faint polarization sources near very bright stars. The SPHERE-team carries out as part of their guaranteed time observations obtained for building the instrument, an investigation of the achievable detection limits of for the search of extra-solar planets. Very high contrast observations of a small number of targets are currently taken. The data confirm that ZIMPOL can reach at least for separations $\rho>0.5''$ the above mentioned detection limit. The short test observations of R Aqr described in the previous section give already an impression of the high contrast performance in polarimetric imaging. SPHERE/ZIMPOL is pioneering this technique and the achievable contrast limits are certainly of interest for other high contrast search programs targeting the reflected light from extra-solar planets and for the planning of future instruments. Whether a successful detection will be possible with SPHERE/ZIMPOL depends on the presence of favorable planets within about 5~pc and further progress in the observations and data analysis. \subsubsection{Differential polarimetric imaging of circumstellar disk.} Disks around young stars are a primary science case for the SPHERE instrument and differential polarimetric imaging is a powerful technique for high contrast disk observations \citep[e.g.,][]{Kuhn01,Perrin09,Quanz11,Muto12}. The high spatial resolution and the high polarimetric sensitivity of ZIMPOL are ideal for the mapping of faint circumstellar disks in polarized light. The visual ZIMPOL data can be combined with near-IR observations, for example from SPHERE/IRDIS, to study the color dependence of the reflected light from protoplanetary disks around young stars like for example HD 135344B or TW Hya \citep{Stolker16,vanBoekel17}. Let us compare for this important science case the pro and cons for the ZIMPOL and IRDIS polarimetry modes for the differential polarimetric imaging of circumstellar disk, focussing on detecting a disk and mapping structural features. For fainter central stars $m_R>8^m$ IRDIS polarimetry provides as important advantage a significantly better AO performance, because all light in the 500-900~nm range can be used for the wave front sensing. The ZIMPOL science channel shares this light with the WFS and therefore only 21~\% of the flux is available for the WFS, if the gray beam splitter is used, or about 80~\% for the dichroic beam splitter. Another advantage of IRDIS is the larger field of view of $11''\times 11''$, about ten times larger than the $3.6''\times 3.6''$ detector field of view of ZIMPOL, or much more efficient than using multiple field observations inside the $8''$ diameter instrument field of view. On the other side, the main advantages of ZIMPOL polarimetry when compared to IRDIS or other IR-polarimeters are the higher spatial resolution and the very good speckle noise suppression by the fast-modulation polarimetry techniques as described in this paper for R Aqr. Therefore, a higher contrast can be obtained for circumstellar disk around bright stars in the separation range $\rho\approx 0.02''-0.2''$ and a good example is the detection of the inner disk of HD 142527 by \citet{Avenhaus17}. For disks around bright stars, ZIMPOL is at least competitive for the separation range $\rho\approx 0.2-2.0$ as demonstrated for the faint debris disks HIP 79977 \citep{Engler17}, but also for the innermost regions of bright protoplanetary disks, like HD 100457 \citep{Benisty17}. In addition, ZIMPOL provides a more advanced polarimetric concept than IRDIS or other AO-assisted polarimetric imagers, allowing an easier calibration of polarimetric data and a quantitative analysis of the polarized reflectivity of the scattering dust \subsubsection{Mass loss from red giants.} The combination of high spatial resolution and sensitive polarimetry is ideal for the mapping of the light scattering from the circumstellar dust with SPHERE/ZIMPOL \citep[e.g.,][]{Kervella15,Khouri16,Ohnaka17a}. As shown in Sect.~\ref{SectRAqr} for R Aqr, the polarized light from dust scattering can be measured over a very wide separation range, and asymmetries, clumps and their evolution can be investigated in much detail. This important information from high resolution polarimetry, which was pioneered with interferometric observations by \citet{Ireland05}, or sparse aperture masking by \citep{Norris12}, is now also available with ``simple'' imaging. Such observations could be particularly useful for investigations and the modeling of the complex dust formation process in pulsating AGB stars \citep[e.g.,][]{Aronson17,Hoefner16,Hoefner08}. Light scattering observations are highly complementary to observations of the thermal emission of the circumstellar dust in the mid-IR with e.g. the VLTI/MIDI interferometer \citep{Paladini17}. AO observations in the visual achieves a comparable or even better resolution than mid-IR interferometry and is therefore well suited for the mapping of the complex distribution of circumstellar dust near the mass losing star. The visible range is ideal for the small dust particles formed around mass-losing red giants because they scatter much more efficiently short wavelength light. A technical challenge for very bright red giants are detector saturation issues. Red giants are much fainter in the visual and the ZIMPOL system is designed especially for high contrast observations of very bright targets and therefore also for the investigation of the brightest, most extended, nearby red giants. \subsubsection{Emission lines in stellar jets and outflows.} Stellar jets and outflows produce often H$\alpha$ and other emission lines from shocks or photoionized regions. ZIMPOL line filters are available for the H$\alpha$ 656~nm, [O\,I] 630~nm, and HeI/NaI 588/589~nm lines which may serve as tracer of different types of ionized or partially ionized gas. Achieving high contrast and high resolution observations in line filters is important for stellar jets from young stars \citep{Frank14} to access the innermost 10~AU where the outflow is not yet significantly perturbed by the interaction with the ambient medium. This requires line observations at separations below 70~mas for sources in nearby star-forming regions. Observing the innermost flow morphology is important to pinpoint the initial ejection site of the matter, e.g. a stellar wind, an ``X-wind'' from the inner edge of the disk, or a disk wind, which is then further accelerated and collimated by the combined action of magnetic fields and rotation \citep{Ferreira06}. A first demonstration of the ZIMPOL potential on this topic is provided in \citet{Antoniucci16} with H$\alpha$ and [O\,I] observations of the young binary Z CMa. The authors could trace the collimated jet from one of the components down to $\sim 70$~mas from the driving source, revealing a jet wiggling on time-scales of a few years, which may be induced by a non-detected close-in companion. For many astronomical objects countless imaging data allready exists of the circumstellar line emission in the visual wavelength region taken with ground-based or space telescopes. The line observations of R Aqr provide a good example for the complementarity of emission line imaging with the SPHERE/ZIMPOL AO system \citep{Schmid17}, with HST imaging \citep{Melnikov18}, and with seeing limited imaging \citep{Liimets18}. \subsubsection{Resolving the atmosphere of red giants.} The most extended red giant stars can be resolved with ``simple ZIMPOL imaging'' with a resolution of up to 20~mas as demonstrated for example for R Dor \citep{Khouri16} or $\alpha$ Ori \citep{Kervella16}. ``Simple imaging'' because one can take many images in several filters within a few minutes, select the best data and measure wavelength dependencies. Geometric features, such as large spots or polar and equatorial zones, can be investigated and flux ratio maps can be obtained with simultaneous differential imaging, in e.g. the TiO\_717 and Cnt748 band filters which sample cold and hot surfaces regions, or the N\_Ha and CntHa filter pair for possible signs of shock heating. The ``simple imaging'' is also ideal for a monitoring program of the temporal evolution of surface features in these stars. Even better spatial resolution $\rho< 20$~mas should be achievable with sparse aperture masking or with advanced data analysis techniques. For these reasons the ``simple'', $\approx 20$~mas resolution, SPHERE/ZIMPOL imaging of extended red giant atmospheres provides useful complementary information with respect to the higher resolution, but much harder to obtain interferometric data \citep[e.g.,][]{Haniff95,vanBelle96,Ohnaka17b}. \subsubsection{Close binary stars.} For very close binary stars, the $\approx 20$~mas spatial resolution of SPHERE/ZIMPOL is of course very useful for orbit determinations and the photometry of the individual components \citep[e.g.,][]{Janson18}. A particular niche for the visual ZIMPOL instrument, when compared to near-IR AO instruments, are faint, hot companions to red stars, like white dwarfs companions to Ba-star, or H$\alpha$ emitting active components to M-giants like the symbiotic system R Aqr shown in Fig.~\ref{rawcounts} \citep[see also][]{Schmid17}. The relative position between roughly equal flux ($|\log (f_1/f_2)|\lapprox 1$) hot and cold binary components can certainly be determined for separations of $\approx 10$~mas or even smaller with simultaneous measurements of the combined binary PSF in a visual and a red filter. Photocenter differences between two bands might even be measurable at the milli-arcsec level, if other stars in the field or the features of a coronagraphic mask can be used as relative astrometric reference. \subsubsection{Solar system objects.} The SPHERE AO system is capable to lock on moving solar system objects if they are bright enough $m_{\rm R}\lapprox 10^m$ and not too extended $\lapprox 2''$. This was demonstrated during the SPHERE commissioning for Titan\footnote{ESO press release www.eso.org/public/news/eso1417} which has a diameter of $\rho=0.8''$ and even Neptune with $\rho=2.4''$ \citep{Fusco16}. Thus, many bright asteroids, the Galilean moons, and Saturn's moon Titan can be imaged in the visual. Sizes, shapes and surface structures can be investigated in much detail \citep{Vernazza17}, and with enhanced resolution when compared to the near-IR range \citep[e.g.,][]{Marchis06}, including the polarimetric properties of the reflecting terrains. \subsection{Conclusions} SPHERE/ZIMPOL is a very versatile adaptive optics instrument and we therefore expect many exciting new scientific results from this instrument. The above listed technical performances, upgrade options, and science topics give only a few examples of possible observational projects with this instrument. Observational results based on adaptive optics profit a lot from the much enhanced spatial resolution and this provides since many years a continuous string of new detections \citep[see][]{Davies12}. The description of SPHERE/ZIMPOL given in this paper should help to define the best observing strategy for reaching deeper detection limits for new discoveries with AO observations at very high spatial resolution in the visual, using polarimetric imaging, angular or differential imaging with broad band, narrow band, or line filters. On the other side, accurate quantitative measurements with AO systems are often difficult, because of the strongly variable atmospheric conditions and the resulting system performance. Particularly problematic is the photometry for very faint companions or extended circumstellar features for which simultaneous or quasi-simultaneous differential measurements are impossible. This makes the accurate characterization of high contrast objects in different wavelength bands, taken often with different instruments and usually under different atmospheric conditions very difficult and often uncertain. A lot of effort is required to describe AO observations accuratly and in a reproducible way but this is required for a detailed characterization of high contrast objects. This paper provides therefore a lot of technical information for the accurate characterization and calibration of SPHERE/ZIMPOL measurements. | 18 | 8 | 1808.05008 |
1808 | 1808.03304_arXiv.txt | We summarize the red channel (2-5 micron) of the Planetary Systems Imager (PSI), a proposed second-generation instrument for the TMT. Cold exoplanets emit the majority of their light in the thermal infrared, which means these exoplanets can be detected at a more modest contrast than at other wavelengths. PSI-Red will be able to detect and characterize a wide variety of exoplanets, including radial-velocity planets on wide orbits, accreting protoplanets in nearby star-forming regions, and reflected-light planets around the nearest stars. PSI-Red will feature an imager, a low-resolution lenslet integral field spectrograph, a medium-resolution lenslet+slicer integral field spectrograph, and a fiber-fed high-resolution spectrograph. | } Extremely Large Telescopes (ELTs) will have the angular resolution to image a variety of exoplanets that are closer to their host stars than the exoplanets that can be imaged with current facilities. The Planetary System Imager (PSI) is proposed instrument that is designed to image exoplanets at a variety of wavelengths and spectral resolutions\cite{Fitzgerald_SPIE2018}. Figure 1 shows a possible architecture for PSI. Light from the telescope is corrected by an adaptive optics (AO) system, featuring a large-format woofer deformable mirror\cite{Guyon_SPIE2018}. The AO system comprises a single relay with gold-coated surfaces to limit thermal emissivity. Light from the adaptive optics system is directed to the cryogenic 2-5 micron channel (PSI-Red), and a tilted dichroic entrance window transmits the $>$2 micron light while reflecting $<$2 micron light. The shorter wavelength light is used in a downstream wavefront sensor that controls the woofer AO system. Additionally some of the shorter wavelength light is further corrected by a tweeter AO system\cite{Guyon_SPIE2018} and directed to a short wavelength science camera (PSI-Blue\cite{Mawet_SPIE2018}). PSI-Red will directly detect and characterize thermal emission from exoplanets while PSI-Blue will detect and characterize reflected light from exoplanets. \begin{figure}[htbp] \begin{center} \hbox{ \hspace{0.1in} \includegraphics[angle=0,width=1.0\linewidth]{f1.pdf}} \end{center} \vspace{0.0in} \caption{--- Schematic of the Planetary Systems Imager\cite{Fitzgerald_SPIE2018}. Light from a high-order woofer AO system\cite{Guyon_SPIE2018} passes an insertable pickoff dichroic to a 10 $\mu$m camera\cite{Marois_SPIE2018}. Light enters the PSI-Red science camera via a dichroic entrance window that transmits 2-5 $\mu$m light and reflects shorter wavelength light to a near-infrared wavefront sensor and PSI-blue\cite{Mawet_SPIE2018}. } \label{fig:PSI schematic} \end{figure} PSI-Red comprises an imager, a low-resolution integral-field spectrograph, a medium-resolution integral-field spectrograph, and a fiber-injection unit for high-resolution spectroscopy. Each of these channels, as well as a shared set of coronagraphic foreoptics, is housed in a cryogenic dewar to minimize the thermal infrared background. PSI-Red is a diffraction-limited instrument, which has an optical design and instrument volume that are independent of telescope aperture. As a result, we are planning a precursor instrument for Keck, SCALES (Santa Cruz Array of Lenslets for Exoplanet Spectroscopy), which is identical to the current PSI-Red design. SCALES could, in principle, be used on Keck while the TMT is under construction, and then integrated into the rest of PSI as a finished product. | 18 | 8 | 1808.03304 |
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1808 | 1808.08961_arXiv.txt | We discuss the origin of the anti-helium-3 and -4 events possibly detected by {\it AMS-02}. Using up-to-date semi-analytical tools, we show that spallation from primary hydrogen and helium nuclei onto the ISM predicts a $\bartHe$ flux typically one to two orders of magnitude below the sensitivity of {\it AMS-02} after 5 years, and a $\barqHe$ flux roughly 5 orders of magnitude below the {\it AMS-02} sensitivity. We argue that dark matter annihilations face similar difficulties in explaining this event. We then entertain the possibility that these events originate from anti-matter-dominated regions in the form of anti-clouds or anti-stars. In the case of anti-clouds, we show how the isotopic ratio of anti-helium nuclei might suggest that BBN has happened in an inhomogeneous manner, resulting in anti-regions with a anti-baryon-to-photon ratio $\bar{\eta}\simeq10^{-3}\eta$. We discuss properties of these regions, as well as relevant constraints on the presence of anti-clouds in our Galaxy. We present constraints from the survival of anti-clouds in the Milky-Way and in the early Universe, as well as from CMB, gamma-ray and cosmic-ray observations. In particular, these require the anti-clouds to be almost free of normal matter. We also discuss an alternative where anti-domains are dominated by surviving anti-stars. We suggest that part of the unindentified sources in the 3FGL catalog can originate from anti-clouds or anti-stars. {\it AMS-02} and {\it GAPS} data could further probe this scenario. | The origin of cosmic ray (CR) anti-matter is one of the many conundrums that {\it AMS-02} is trying to solve thanks to precise measurements of CR fluxes at the Earth. In over six years, {\it AMS-02} has accumulated several billion events, whose composition is mostly dominated by protons and helium nuclei. Moreover, positrons and antiprotons have been frequently observed and are the object of intense theoretical investigations in order to explain their spectral features. Indeed, anti-matter particles are believed to be mainly of secondary origin, i.e., they are created by primary CR nuclei (accelerated by supernova-driven shock waves) impinging onto the interstellar medium (ISM). However, deviations from these standard predictions have been observed, hinting at a possible primary component. In the case of positrons, a very significant high-energy excess has already been seen in {\it PAMELA} data \cite{Adriani:2010rc}. The main sources under investigation to explain this excess are DM and pulsars (see e.g. \cite{Bergstrom:2008gr, Cirelli:2008jk, Cirelli:2008pk, Nelson:2008hj, ArkaniHamed:2008qn, Harnik:2008uu, Fox:2008kb, Pospelov:2008jd, MarchRussell:2008tu, Dienes:2013xff, Kopp:2013eka, Yuksel:2008rf, Profumo:2008ms, Kawanaka:2009dk, Yuan:2013eja, Yin:2013vaa,Hooper:2008kg,Cholis:2008qq,Cholis:2008hb,Cholis:2013psa,Boudaud:2014dta,Boudaud:2016jvj,Hooper:2017gtd,Cholis:2017ccs,Cholis:2018izy}). In the case of antiprotons, a putative excess at the GeV-energy \cite{Cuoco:2016eej} is under discussion \cite{Reinert:2017aga}. Still, antiprotons represent one of the most promising probes to look for the presence of DM in our Galaxy through its annihilation. But the searches for anti-matter CR do not limit themselves to antiprotons and positrons. Hence, many theoretical and experimental efforts are devoted to detecting anti-deuterons, which are believed to be a very clean probe of DM annihilations especially at the lowest energies (below tens of GeV) \cite{Chardonnet:1997dv,Donato:1999gy,Duperray:2005si,vonDoetinchem:2015yva}. Similarly, measurement of the anti-helium nuclei CR flux is a very promising probe of new physics, that has been suggested to look for DM annihilations \cite{Duperray:2005si,Carlson:2014ssa,Cirelli:2014qia,Coogan:2017pwt} or other sources of primary CR, such as anti-matter stars or clouds \cite{Steigman:1976ev,Belotsky:1998kz,Bambi:2007cc}. Strikingly, {\it AMS-02} has recently reported the possible discovery of eight anti-helium events in the mass region from 0 to 10 GeV/c$^2$ with $Z= 2$ and rigidity $<50$ GV \cite{Choutko}. Six of the events are compatible with being anti-helium-3 and two events with anti-helium-4. The total event rate is roughly one anti-helium in a hundred million heliums. This preliminary sample includes one event with a momentum of $32.6\pm2.5$ GeV/c and a mass of $3.81\pm 0.29$ GeV/c$^2$ compatible with that of anti-helium-4. Earlier already, another event with a momentum of 40.3 $\pm$ 2.9 GeV and a mass compatible with anti-helium-3 had been reported \cite{TingTalk}. In this paper, we discuss various possibilities for the origin of {\it AMS-02} anti-helium events. Should these events be confirmed, their detection would be a breakthrough discovery, with immediate and considerable implications onto our current understanding of cosmology. The discovery of a single anti-helium-4 nucleus is challenging to explain in terms of known physics. In this article, we start stressing why such a discovery is unexpected. For this, we re-evaluate the secondary flux of anti-helium nuclei. In particular, we provide the first estimate of the $\barqHe$ flux at the Earth coming from the spallation of primary CR onto the ISM. We show that it is impossible to explain AMS results in terms of a pure secondary component, even though large uncertainties still affect the prediction. Moreover, we argue that the DM explanations of these events face similar difficulties, although given the virtually infinite freedom in the building of DM models, it is conceivable that a tuned scenario might succeed in explaining these events. We then discuss the implications of the anti-helium observation. We essentially suggest that the putative detection of $\bartHe$ and $\barqHe$ by {\it AMS-02} indicates the existence of an anti-world, i.e., a world made of anti-matter, in the form of anti-stars or anti-clouds. We discuss properties of these regions, as well as relevant constraints on the presence of anti-clouds in our Galaxy. We present constraints from the survival of anti-clouds in the Milky-Way and in the early Universe, as well as from CMB, gamma-ray and cosmic-ray observations. We show in particular that these require the anti-clouds to be almost free of normal matter. Moreover, we show how the isotopic ratio of anti-helium nuclei might suggest that BBN happened inhomogeneously, resulting in anti-regions with a anti-baryon-to-photon ratio $\bar{\eta}\simeq10^{-3}\,\eta$. Given the very strong constraints applying to the existence and survival of anti-clouds, we also discuss an alternative scenario in which anti-domains are dominated by anti-stars. We suggest that part of the unidentified sources in the 3FGL catalog can be anti-clouds or anti-stars. Future {\it AMS-02} and {\it GAPS} data could further probe this scenario. The paper is structured as follows. Section \ref{sec:secondary} is devoted to a thorough re-evaluation of the secondary astrophysical component from spallation within the coalescence scheme. A discussion on the possible limitations of our estimates and on the DM scenario is also provided. In section \ref{sec:antidomain}, we discuss the possibility of anti-domains in our Galaxy being responsible for {\it AMS-02} events. Properties of anti-clouds and their constraints are presented in sec~\ref{sec:anticloud}, while the alternative anti-star scenario is developed in section~\ref{sec:antistar}. Finally, we draw our conclusions in sec.~\ref{sec:discussion}. | \label{sec:discussion} In this work, we have studied the implications of the potential discovery of anti-helium-3 and -4 nuclei by the {\it AMS-02} experiment. Using up-to-date semi-analytical tools, we have shown that it is impossible to explain these events as secondaries, i.e., from the spallation of CR protons and helium nuclei onto the ISM. The $\bartHe$ is typically one to two orders of magnitude below the sensitivity of {\it AMS-02} after 5 years, and the $\barqHe$ is roughly 5 orders of magnitude below {\it AMS-02} reach. It is conceivable that $\bartHe$ has been misidentified for $\barqHe$. Still, we have argued that the pure secondary explanation would require a large increase of the coalescence momentum at low energies, a behavior that goes against theoretical considerations and experimental results. The DM scenario suffers the same difficulties. Hence, we have discussed how this detection, if confirmed, would indicate the existence of an anti-world, in the form of anti-stars or anti-clouds. We summarize what we have learned about the properties of anti-matter regions: \begin{itemize} \item[\textbullet] Taken at face value the isotopic ratio of anti-helium nuclei is puzzling. We have shown that it can be explained by anisotropic BBN in regions where $\bar{\eta}\sim1.3-6\times10^{-13}$. \item[\textbullet] The density, size and number of anti-matter domains is constrained by {\it AMS-02} observations and our knowledge of Galactic properties to verify Eq.~(\ref{eq:nbar}). The only theoretical assumption behind is that the isotopic ratio measured by {\it AMS-02} comes from BBN. Interestingly, a few highly dense clouds are sufficient to explain {\it AMS-02} measurements. \item[\textbullet] The annihilation rate of anti-matter in our Galaxy requires anti-domains to be poor in normal matter (typically a tenth or less of the normal matter density). Considering the annihilation rate in the early Universe leads to even stronger requirements, which would imply the existence of some exotic mechanism allowing segregation of matter and anti-matter domains all along cosmic evolution that makes the existence of such anti-clouds quite improbable. \item[\textbullet] Additionally, gamma rays can provide strong constraints on this scenario. Non-observation of spectral features in the form of lines with energies close to the proton mass strongly constrains the proton density in anti-matter domain, as given by Eq.~(\ref{eq:AnnConstraints}). However, this constraints apply only if anti-matter domains are numerous and homogeneously distributed within the Galactic disk. We anticipate that very competitive constraints can be obtained from non-observation of positron annihilations and/or pion decays. \item[\textbullet] Anti-clouds could produce a measurable flux of $\bar{p}$ and $\bar{d}$. Most of the parameter space evades current $\bar{p}$ constraints but could be probed by GAPS. \item[\textbullet] Alternatively (and more likely), these anti-helium events could originate from anti-star(s) whose main material is anti-helium-4, converted into anti-helium-3 via spallation in the dense environment surrounding the anti-star(s). \item[\textbullet] Part of the 3FGL unassociated point sources can be anti-clouds experiencing annihilations due to CRs propagating through them. They can also be anti-stars which experience annihilations as they propagate in the ISM. \item[\textbullet] Depending on the (unknown) acceleration mechanism, it is conceivable that a single near-by anti-star (whose distance to the Earth must be larger than $\sim$1 pc) contributes to the {\it AMS-02} observation. \end{itemize} All these hints can be used to build a scenario for their formation in the early Universe. Needless to say, the successful creation and survival of such objects within a coherent cosmological model is far from obvious. Here we just mention that there are many scenarios discussed in the literature \cite{Bambi:2007cc,Blinnikov:2014nea,Khlopov:1998uy}, including the Affleck-Dine mechanism \cite{Affleck:1984fy}, which would lead to the formation of ``bubbles'' of matter and anti-matter with arbitrarily large values of the baryon-asymmetry locally. Depending on the relation between their mass and the corresponding Jeans mass, these bubbles can then lead to the formation of anti-star-like objects, either through specific inflation scenarios with large density contrast \cite{Carr:1974nx,Carr:1975qj} on scales re-entering the horizon around the QCD phase-transition, i.e., $T\sim{\cal O}(100~{\rm MeV})$, or from peculiar dynamics of the plasma within the bubble, as described for instance in Ref.~\cite{Dolgov:1992pu}. In the latter scenario, the {\em negative} pressure perturbation inside the bubble leads to the collapse of baryons within this region. If the value of the baryon-asymmetry in the bubble is very large, it is even possible that different expansion rate (due to more non-relativistic matter inside the bubble) naturally leads to the growth of density perturbations much earlier than outside of these regions. Given the strong implications of the discovery of a single anti-helium-4 nucleus for cosmology, important theoretical and experimental efforts must be undertaken in order to assess whether the reported events could be explained by a more mundane source, such as interactions within the detector, or another source of yet unkown systematic error. Still, this potential discovery would represent an important probe of conditions prevailing in the very early Universe and should be investigated further in future work. \appendix | 18 | 8 | 1808.08961 |
1808 | 1808.01840_arXiv.txt | Current models of (exo)planet formation often rely on a large influx of so-called `pebbles' from the outer disk into the planet formation region. In this paper, we investigate how the formation of pebbles in the cold outer regions of protoplanetary disks and their subsequent migration to the inner disk can alter the gas-phase CO distribution both interior and exterior to the midplane CO snowline. By simulating the resulting CO abundances in the midplane as well as the warm surface layer, we identify observable signatures of large-scale pebble formation and migration that can be used as `smoking guns' for these important processes. Specifically, we find that after $1\mathrm{~Myr}$, the formation and settling of icy pebbles results in the removal of up to $80\%$ of the CO vapor in the warm ($T>22\mathrm{~K}$) disk layers outside the CO snowline, while the radial migration of pebbles results in the generation of a plume of CO vapor interior the snowline, increasing the CO abundance by a factor ${\sim}2{-}6$ depending on the strength of the turbulence and the sizes of the individual pebbles. The absence of this plume of CO vapor in young nearby disks could indicate efficient conversion of CO into a more refractory species, or a reduction in the radial mass flux of pebbles by, for example, disk inhomogeneities or early planetesimal formation. | Snowlines are believed to play an important role in protoplanetary disk evolution and planet formation in general. Marking the locations where major volatiles (e.g., $\water,\co,\coo$) transition from being predominantly in the gas-phase to solid as ices on grain surfaces, snowlines separate regions of the protoplanetary disk with possibly very different gas-phase and grain-surface chemistry, changes that are often assumed to be reflected in the composition of (giant) planets forming in different locations \citep[e.g.,][]{oberg2011}. The formation of planetesimals and planetary embryos is often associated with the water snowline \citep[e.g.,][]{drazkowska2017,schoonenberg2017,ormel2017}, but other snowlines could also be preferred sites \citep{ali-dib2017}. In the popular `pebble accretion' paradigm, planetesimals/embryos then grow rapidly by accreting mm/cm-size pebbles that drift in from further out in the disk \citep{ormelklahr2010,lambrechts2012,johansen2017}. While growth through pebble accretion can be very fast, only a small fraction of pebbles is usually accreted \citep{ormel2018}, and therefore the process relies on a large and long-lived radial flux of pebbles coming in from the outer regions of the protoplanetary nebula \citep{lambrechts2014}. Such a large-scale radial migration of ice-covered solids originating from the outer disk is expected to redistribute volatiles on a disk-wide scale \citep{oberg2016}, qualitatively changing the static picture presented in \citet{oberg2011}. The interaction between midplane snowlines and radial transport of solids and vapor has been studied in the past \citep{stevenson1988,cuzzi2004,ciesla2006} and has received a lot of attention in recent years \citep{stammler2017,schoonenberg2017,booth2017,drazkowska2017,bosman2017}. With radial drift being faster than turbulent mixing, these studies generally find an enhancement of volatiles interior to their snowline, the magnitude of which depends on the underlying pebble flux and ice content. Even before pebbles start drifting however, the formation of these large, settled dust particles can change the vertical distribution of gas-phase volatiles via the sequestration of ices in the midplane \citep{meijerink2009,du2015,kama2016,du2017}. Models studying vertical mixing find that this effect can decrease the gas-phase $\water$ and CO abundances in the warm molecular layer by anywhere between a factor of a few to almost 2 orders of magnitude, depending on the timescales involved and the details of the pebble formation process \citep{xu2016,krijt2016b}. For CO, this story of depletion above the surface snowline and potential enhancement in the inner disk is of particular importance because CO emission is commonly used as a tracer for bulk disk mass \citep[e.g.,][]{williams2014,ansdell2016,miotello2016,miotello2017,molyarova2017}. Hence, if the CO abundance is significantly depleted in the region of the disk that dominates the emission, this approach could be underestimating the true disk mass. For the handful of disks for which independent mass estimates can be made using HD, it appears as though CO is indeed depleted by a factor of a few to up to two orders of magnitude \citep{favre2013,mcclure2016,schwarz2016}. In addition, CO is the only molecule for which the snowline has been (directly) observed \citep{qi2013} and for which we can vertically and radially resolve abundances using a variety of isotopologues \citep{schwarz2016,zhang2017,dutrey2017,pinte2017,huang2018}. The aim of this paper is to construct a self-consistent model that describes how the formation and subsequent vertical settling and radial drift of pebbles alters CO abundances in different regions of the disk; both interior and exterior to the midplane snowline, as well as in the warmer surface layers of the outer disk. To that end, we focus on a single, invariant disk profile (Sect.~\ref{sec:model}) and model the vertical and radial transport of dust, pebbles, ices, and gas-phase CO while pebbles are continuously forming over Myr timescales (Sect.~\ref{sec:numerical}). By comparing models of increasing complexity (Sect.~\ref{sec:building}) and exploring the dependence on several parameters related to pebble formation/evolution (Sect.~\ref{sec:results}), we attempt to build a coherent story of how pebble migration affects CO abundances on a disk-wide scale. The results are discussed in Sect. \ref{sec:discussion} and conclusions presented in Sect. \ref{sec:concl}. | \label{sec:discussion} \subsection{Comparison to (resolved) CO observations}\label{sec:observations} Apparent depletions of gas-phase CO in the outer disk have been reported by several authors for a variety of disks \citep[][]{favre2013, du2015, kama2016, mcclure2016, schwarz2016}, with depletion factors ranging from a factor of a few to 2 orders of magnitude. In addition, assuming a Solar (i.e., non-depleted) value for the $\co/\hyy$ mixing ratio results in (very) low gas disk masses and unusually high dust-to-gas ratios \citep{ansdell2016,eisner2016,miotello2016,miotello2017}. The depletions we observe in the warmer parts of the outer disk (e.g., Fig.~\ref{fig:2D_gas}) are typically around 90\%, but we discuss possibilities for creating more extreme depletion factors below. Another result of our models that include drift is the formation of a plume of gas-phase CO interior to and around the midplane snowline. At least for TW Hya, such an obvious resurgence of CO is not seen by \citep{schwarz2016}. Even though there is a hint of an increase inside the snowline \citep[][Fig.~3(d)]{schwarz2016}, and the contrast between the CO abundance interior and exterior to the snowline is similar to what we predict (e.g., Fig.~\ref{fig:snapshots2}), \citeauthor{schwarz2016} find CO to be depleted on \emph{both} sides of the midplane snowline. While the tracer used by \citet{schwarz2016}, C$^{18}$O, is possibly optically thick in the inner disk, this picture of a lack of CO returning to the gas phase was confirmed by \citet{zhang2017} using the optically thin $^{13}$C$^{18}$O. \subsection{Increasing the amount of CO depletion}\label{sec:moredepletion} The models shown in this paper do not show CO depletions of more than an order of magnitude in the disk's surface layers. Here, we discuss \new{effects} that could potentially increase the depletion to reach the 2 orders of magnitude that have been reported for some disks. \emph{Evolution over longer timescales.} The snapshots shown in Figs.~\ref{fig:2D_gas} and \ref{fig:2D_mix} do not represent a steady state: pebbles are continuously forming and the degree of CO depletion in the outer disk is increasing with time. Because pebbles are continuing to form and the vertical mixing timescale in the outer disk is ${\sim}10^{5}\mathrm{~yr}$ or longer (Sect.~\ref{sec:alpha}), running the models for a longer period of time is expected to increase the depletion. \new{\emph{A vertical turbulence profile.} In this study, we have assumed a single, constant $\alpha$-value when describing the turbulent viscosity and diffusion coefficients (Sect.~\ref{sec:transport}). In reality, the strength and nature of the turbulence is expected to vary significantly between different regions in the disk \citep[e.g.,][]{turner2014}. Recent theoretical models studying the outer regions of protoplanetary disks tend to find a relatively weak turbulence in the midplane (corresponding to $\alpha \lesssim 10^{-3} $) and a stronger turbulence ($\alpha\sim10^{-2}$) in the upper layers \citep[e.g.,][and references therein]{simon2013a,simon2013b,bai2016}. The presence of such a vertical profile can significantly influence vertical transport of dust grains \citep{ciesla2010,ormel2018}, promoting the sequestration of icy bodies in the midplane and increasing the efficiency with which CO is removed from the gas-phase in the disk's upper regions \citep{xu2016}. Recent observational work, however, appears to show the turbulence in the upper layers of the disks around TW Hya and HD163296 is relatively weak \citep{teague2016,flaherty2017,flaherty2018}, implying $\alpha \sim 10^{-2}$ is not common in the surface layers of protoplanetary disks.} \emph{Dust-pebble interactions.} Our dust evolution model does not include pebble mass gain/loss through collisions with much smaller particles. If the accretion of small grains is efficient however, this sweep-up could contribute to the depletion of dust and volatiles from the warm molecular layer: In the models shown in this paper, the only way for a CO molecule to end up on a pebble in the midplane is to freeze out on a small grain which then grows into a (previously non-existing) pebble. If sweep-up is efficient, a second route becomes available, in which a molecule freezes out onto a small grain which is subsequently accreted by an already-existing pebble. In regions of the disk where this second route is more efficient than the first (i.e., regions with a low dust density and/or high pebble surface density), the volatile depletion could then be much more dramatic. However, collisions between pebbles or aggregates and small dust grains do not necessarily result in sticking but can also lead to mass loss in the form of erosion or cratering \citep{schrapler2011,seizinger2013c,krijt2015}. If erosion is efficient, it might not only limit further growth of pebbles, but also be the dominant source of small grains at later times \citep{schrapler2018}, potentially alleviating the problems dust coagulation models often have in producing enough small grains to match multi-wavelength observations \citep[e.g.,][]{dullemonddominik2005,pohl2017}. \emph{Chemistry.} Finally, we discuss the possibility of removing CO from the gas-phase by locally reprocessing CO through chemical reactions that lock the carbon in other molecules/species \citep[e.g.,][]{bergin2014,reboussin2015,yu2016,eistrup2017}. A recent comprehensive modeling study by \citet{schwarz2018} found that -- unless the cosmic ray rate is high -- it is difficult to deplete CO by an order of magnitude or more on a timescale of a million years, concluding that chemistry alone is not responsible for the majority of the observed depletions. Nonetheless, \new{several} models conducted at 100 au converted a significant fraction of CO to $\coo$-ice and $\mathrm{CH_3OH}$-ice on timescales shorter than a million years \citep[][Fig.~5]{schwarz2018}. With both mechanisms (chemical processing of CO and pebble-formation-mediated sequestration in the midplane) leading to an order of magnitude of CO depletion when acting on their own, it is tempting to imagine they can reach the observed two orders of magnitude when working together. In addition, while we focused exclusively on how dust growth impacted material transport, the coagulation of small grains into larger solids is also expected to alter the temperature profile and radiation field in the disk \citep[e.g.,][]{cleeves2016,facchini2017}. Developing models to understand how these physical and chemical processes interact will be the focus of future work. \subsection{Decreasing the pile-up of CO interior to the snowline}\label{sec:lessreturn} Observations do not appear to show a return of CO interior to the snowline \citep{schwarz2016}, the presence of which is a common outcome in our models that include both pebble formation and radial drift (Figs.~\ref{fig:2D_mix}). We briefly discuss possibilities that could prevent the CO from returning to the gas as pebbles grow and evolve. \emph{Reduced drift efficiency.} Unsurprisingly, the models that show the smallest CO enhancement in the inner disk are those for which the radial flux of solids is smallest (cf. Figs.~\ref{fig:snapshots2}(a) and \ref{fig:pebhistplot}(a)). One way to reduce the pebble flux is to have the pebbles keep relatively small Stokes numbers, which, in the context of our dust evolution model, happens when pebbles maintain a high porosity (model M2a). Pebble sizes and Stokes numbers could also be kept small if catastrophic fragmentation is a common outcome of pebble-pebble collisions in the outer disk \citep{brauer2007,birnstiel2012,pinilla2017}, \new{as would be the case for $v_f\sim\mathrm{1~m/s}$ (Fig.~\ref{fig:pbf}(b))} . Alternatively, the efficiency of radial drift can be reduced by structures in the gaseous disk such as pressure bumps or traps \citep{pinilla2012}, which cause the pressure gradient $\eta$ (see Eq.~\ref{eq:driftsettle}) to vary on relatively small radial scales. \new{\emph{Increased turbulence in the midplane.} The shape of the CO enhancement depends on the strength of the turbulence (compare models M1 and M2c in Fig.~\ref{fig:snapshots2}(a). The peak is less prominent for a higher value of $\alpha$ because $i)$ diffusion is more efficient at smearing out the deposited CO vapor and $ii)$ the individual sizes and the total radial flux of pebbles tend to decrease for higher $\alpha$ (see Fig.~\ref{fig:pebhistplot}). \citet{stammler2017} find that for turbulence strengths $\alpha \sim 10^{-2}$, the enhancement relative to the initial conditions becomes insignificant, although it is not clear if such high levels of turbulence are present in the disk midplane at radii outside ${\sim}30\mathrm{~au}$ \citep{simon2013a,simon2013b}. Alternatively, a lower Schmidt number would also increase the diffusivity and lead to a smaller peak in the CO abundance just interior to the snowline \citep[][Fig.~8]{stammler2017}.} \new{\emph{High mass accretion rate.} With the gas accreting radially, the plume of CO vapor forming just inside the snowline will advect inward at a velocity $v_r \sim 3 \nu_\mathrm{T}/2r$ and result in the enhancement of the entire inner disk on a timescale comparable to the local viscous time. For the disk model outlined in Sect.~\ref{sec:diskmodel} and \ref{sec:transport}, $\dot{M}\sim10^{-9}~M_\odot/\mathrm{yr}$ and $v_r\sim \mathrm{cm/s}$ around the CO snowline and this effect can be ignored on the timescales simulated in Sects.~\ref{sec:building} and \ref{sec:results}. In disks with a higher accretion rate however, $v_r$ can become significant, decreasing the degree of vapor enhancement in the inner disk and the efficiency of CO vapor retro-diffusing back across the snowline \citep{cuzzi2004}.} \emph{Planetesimal formation.} The only model in which we observe a depletion of CO vapor \new{inside the CO snowline} is one without any pebble migration (model M0b in Fig.~\ref{fig:snapshots1}(a)), in which case the pebbles outside the midplane snowline effectively become a sink for CO ice. While such a model does not appear to be realistic, a similar picture could arise if a large fraction of the pebbles can be converted into (stationary) planetesimals on timescales comparable to the drift timescale \citep[i.e., regime 3 of][]{cuzzi2004}. \emph{Chemistry.} The explanations offered above all rely on decreasing the radial flux of pebbles, thus decreasing the flux of CO ice. A steady influx of solids could still be \new{allowed}, however, if CO can be destroyed chemically. \citet{schwarz2016} studied the chemical destruction of CO in the inner disk (at 19 au), finding that removing CO on a Myr timescale is only feasible with high cosmic ray rate. Alternatively, CO could be reprocessed already in the outer disk, before freezing out on the grains in the form of hydrocarbons or $\coo$ for example (see last paragraph of Sect.~\ref{sec:moredepletion}). However, while this might alleviate the apparent problem of not seeing the return of CO around $r\approx30\mathrm{~au}$, putting the carbon in $\coo$ will only make a similar issue at the $\coo$ snowline more severe \citep[see][]{bosman2017}. Developing models that include pebble formation and drift (this paper), chemical reactions involving the dominant carbon carriers \citep{schwarz2016} as well as planetesimal formation, and comparing those models to spatially resolved observations of nearby young disks will be key to understanding how carbon is delivered to the (terrestrial) planet formation zone \citep{bergin2014}. | 18 | 8 | 1808.01840 |
1808 | 1808.01592_arXiv.txt | We propose a \hqe\ to measure cross correlations between gravitational lensing of the cosmic microwave background (CMB) and differential screening effects arising from fluctuations in the electron column density, such as could arise from patchy reionization. The \hqes\ are validated by simulated data sets with both Planck and CMB-Stage 4 (CMB-S4) instrumental properties and found to be able to recover the cross-power spectra with almost no biases. We apply this technique to Planck 2015 temperature data and obtain cross-power spectra between gravitational lensing and differential screening effects. Planck data alone cannot detect the patchy-reionization-induced cross-power spectrum but future experiment like CMB-S4 will be able to robustly measure the expected signal and deliver new insights on reionization. | Recombination of hydrogen atoms 380,000 years after the Big Bang left the early Universe with neutral hydrogen gas that was almost uniform but had small density fluctuations~\cite{2001ARA&A..39...19L}. These fluctuations seeded the population of the first stars that emitted ultraviolet (UV) radiation, by which electrons were stripped from neutral hydrogen atoms and scattered with cosmic microwave background (CMB) photons. Large inhomogeneities in the ionization fraction of the gas were created during the so-called epoch of reionization (EoR), leading to substantial variations in the CMB scattering optical depth. At later times there are additional modulations in the scattering optical depth that arise from fluctuations in the baryon density. These variations in the scattering optical depth cause secondary fluctuations in the CMB. The secondary CMB anisotropies generated at the EoR are extremely weak but can create excess power in CMB temperature and polarization power spectra~\cite{2009PhRvD..79j7302D}. However, it is very difficult to detect the small amount of excess power from these secondary anisotropies in the presence of the substantial fluctuation power coming from the (Gaussian) primary fluctuations and instrument noise. To get sufficient sensitivities, higher order estimators for patchy reionization have been developed with three-point~\cite{2018arXiv180105396F} and four-point~\cite{cora1,2011arXiv1106.4313S} correlation functions. However, they either rely on a high-redshift large scale structure tracer which is hard to obtain or contain significant higher order biases that introduce extra uncertainties for the measurements. Although the auto-power spectrum of patchy reionization can be recovered after subtracting model-dependent biases, a Gaussian noise that is almost six orders of magnitude higher than the signal makes it hard to detect. Building on previous work~\cite{2018arXiv180105396F} which investigated the utility of cross-correlating a relatively noisy optical depth reconstruction with higher signal-to-noise tracers of large scale structure, we consider the CMB gravitational lensing as a high-redshift tracer and construct a cross correlation between gravitational lensing ($\phi$) and \ds\ ($\tau$) effects. This is essentially a four-point correlation function, using two-point estimators for $\phi$ and $\tau$, respectively. In this \pp, we study theoretical predictions of this cross correlation, and construct a hybrid quadratic estimator to extract such a signal from CMB data. This paper is structured as follows: in Sec. II, we describe details of constructing the ``\hqe'' as well as various debiasing steps; in Sec. III, we validate the \hqe\ with simulations at noise levels appropriate for Planck and a CMB-S4-like experiment; we then apply this estimator to Planck 2015 temperature data in Sec. IV and conclude in Sec. V.\\ | \label{con} In this \pp, we propose a \hqe\ formalism for the study of fluctuations in the \atau, such as would be induced by \pr. We make mock data sets for Planck and CMB-S4-like experiments and numerically validate that the cross-power spectrum $\langle\phi\tau\rangle$ can be correctly recovered from the CMB data alone. Moreover, we measure a $\langle\phi\tau\rangle$ cross-power spectrum from Planck 2015 temperature data and obtain a new upper bound for \pr. Various systematic and foreground tests are performed and their effects are found to be negligible. For the next-generation CMB experiments, both temperature and polarization data can be used to measure the cross-power spectra $\langle\phi\tau\rangle$ with much higher signal-to-noise ratios, and the signal of patchy reionization in future experiments will be detected from CMB alone, extending CMB science to a new regime when the first stars and galaxies formed. | 18 | 8 | 1808.01592 |
1808 | 1808.00461_arXiv.txt | We investigate the long-term evolution of black hole accretion disks formed in neutron star mergers. These disks expel matter that contributes to an $r$-process kilonova, and can produce relativistic jets powering short gamma-ray bursts. Here we report the results of a three-dimensional, general-relativistic magnetohydrodynamic (GRMHD) simulation of such a disk which is evolved for long enough ($\sim 9$\,s, or $\sim 6\times 10^5 r_{\rm g}/c$) to achieve completion of mass ejection far from the disk. Our model starts with a poloidal field, and fully resolves the most unstable mode of the magnetorotational instability. We parameterize the dominant microphysics and neutrino cooling effects, and compare with axisymmetric hydrodynamic models with shear viscosity. The GRMHD model ejects mass in two ways: a prompt MHD-mediated outflow and a late-time, thermally-driven wind once the disk becomes advective. The total amount of unbound mass ejected ($0.013M_\odot$, or $\simeq 40\%$ of the initial torus mass) is twice as much as in hydrodynamic models, with higher average velocity ($0.1c$) and a broad electron fraction distribution with a lower average value ($0.16$). Scaling the ejected fractions to a disk mass of $\sim 0.1M_\odot$ can account for the red kilonova from GW170817 but underpredicts the blue component. About $\sim 10^{-3}M_\odot$ of material should undergo neutron freezout and could produce a bright kilonova precursor in the first few hours after the merger. With our idealized initial magnetic field configuration, we obtain a robust jet and sufficient ejecta with Lorentz factor $\sim 1-10$ to (over)produce the non-thermal emission from GW1708107. | The recent detection of the neutron star (NS) merger GW170817\footnote{Also known as GRB170817A, SSS17a, AT 2017gfo, and DLT17ck} in gravitational- and electromagnetic waves (\citealt{ligo_gw170817_gw,ligo_gw170817_multi-messenger}, and references therein) has advanced several outstanding issues in astrophysics. It has established neutron star mergers as an important (if not dominant) site of $r$-process element production (e.g., \citealt{kasen_2017,cote_2018, hotokezaka_2018}), provided unambiguous association between a neutron star merger and a short gamma-ray burst \citep{ligo_gw170817_grb}, and set constraints on the dense-matter equation of state (e.g., \citealt{bauswein_2017,margalit_2017,rezzolla_2018,chatziiouannou_2018,raithel_2018,de_2018,ligo_gw170817_eos}). Evidence for the $r$-process comes from the photometric and spectroscopic properties of the observed kilonova (e.g., \citealt{cowperthwaite_2017,chornock_2017, drout_2017,tanaka_2017,tanvir_2017}). This type of transient had been predicted to arise out of sub-relativistic, neutron-rich ejecta from the merger that is radioactively heated by freshly produced $r$-process elements \citep{Li&Pacyznski98,Metzger+10b,Roberts+11,tanaka2016,metzger_2017}. The optical opacity of lanthanides and actinides ($A>130$, produced by the $r$-process) is such that the transient was expected to evolve from blue optical to near infrared within a few days \citep{Kasen+13,tanaka2013,Barnes&Kasen13,fontes2015}, as observed. Also, the temporal evolution of the bolometric luminosity is consistent with the time-dependence of the radioactive heating rate from the $r$-process (e.g., \citealt{rosswog_2017}). Two main mass ejection channels operate in neutron star mergers: dynamical ejecta and outflows from the remnant accretion disk. The former is launched on the dynamical time of the merger ($\sim$\,ms) by tidal forces and hydrodynamic interactions (e.g., \citealt{bauswein2013,Hotokezaka+13}). Numerical relativity simulations predict this material to be sufficiently neutron-rich to produce mainly $A > 130$ elements, with varying amounts of lighter material depending on the equation of state (EOS) of dense matter and the treatment of neutrino physics (e.g., \citealt{wanajo2014,roberts2017,radice2016,foucart2016a,foucart2016b}). While magnetic fields are not expected to significantly alter the dynamics of the merger (e.g., \citealt{endrizzi_2016}) they can lead to some mass ejection on the dynamical time (e.g., \citealt{kiuchi2014,shibata_2017}). For the particular case of GW710817, the amount of dynamical ejecta expected is smaller than the total $r$-process mass inferred from the kilonova (e.g., \citealt{ligo_gw170817_ejecta,shibata_2017b}; however see \citealt{kawaguchi_2018} for a different kilonova mass estimate). The remnant accretion disk evolves on longer timescales ($\sim 100$~ms - $10$~s) and ejects mass through a combination of physical processes (see, e.g., \citealt{FM16} for an overview). Immediately after the merger, the disk is sufficiently hot and dense for neutrinos to be the primary cooling channel, with most of the nuclei fully dissociated into nucleons \citep{popham1999,ruffert1999,Narayan+01,Chen&Beloborodov07}. A key property of these disks is that they transition to being fully advective once the density drops and weak interactions freeze-out on a timescale of $\sim 300$\,ms to $1$\,s, making them prone to launching outflows \citep{Metzger+09a}. Our current understanding of the long-term disk evolution is based primarily on axisymmetric hydrodynamic simulations that include the required microphysics and neutrino treatment at various levels of sophistication, but which model angular momentum transport through an imposed shear stress with parameterized viscosity (e.g., \citealt{FM13,Just+15,fujibayashi_2018}). With this physics included, the outflow is driven primarily by viscous heating and nuclear recombination, with neutrino heating being sub-dominant when the central object is a black hole (BH). The amount of mass ejected in these simulations lies in the range $\sim 5-20\%$ of the initial disk mass after an evolution time of $\sim 10$\,s, with quantitative details depending primarily on the properties of the central object (a much larger fraction of the disk can be ejected if a hypermassive neutron star forms; \citealt{MF14}). The composition of the outflow involves mainly light $r$-process elements, with varying amounts of material with $A > 130$ depending on parameters such as the strength of angular momentum transport or the lifetime of a hypermassive neutron star (HMNS) \citep{Just+15,martin2015,wu2016,lippuner_2017}. It is generally accepted, however, that angular momentum transport in astrophysical accretion disks operates via magnetohydrodynamic (MHD) turbulence driven by the magnetorotational instability (MRI; \citealt{balbus1991}). Early GRMHD models of NS-NS/BH-NS merger remnant disks employed axisymmetric (2D) simulations that start with an initially poloidal field geometry, and which include the relevant neutrino processes, but (1) had too high an ambient density to allow for a significant outflow and/or (2) did not evolve the system for long enough to achieve the radiatively-inefficient state and {completion of} mass ejection \citep{shibata2007,shibata2012,janiuk2013,janiuk_2017}. In addition, it is well-known that in axisymmetry, as a consequence of the anti-dynamo theorem~\citep{cowling33}, MRI turbulence dissipates within $\sim 10$ disk orbits (e.g. \citealt{haw00}) and hence angular momentum transport cannot be sustained for the required timescales. Recently, \citet{siegel_2017a,siegel_2018} have reported the first three-dimensional (3D) GRMHD simulation of an accretion disk around a black hole remnant. The simulation uses a physical equation of state that includes recombination of nucleons into alpha particles, and accounts for neutrino cooling via a leakage scheme. They start their simulation with an equilibrium torus and an initial poloidal field, evolving the disk for $\sim 400$\,ms. Strong outflows are obtained, and by the end of their simulation 20\% of the initial disk mass is ejected as unbound matter at a radius of $10^8$\,cm, with 60\% of the disk mass accreted. Since by that time accretion onto the BH is mostly complete, they surmise that the remaining 20\% would continue to be ejected as an unbound outflow if the simulation was continued to longer times. \citet{nouri_2017} also studied the GRMHD evolution of an accretion disk mapped from a 3D non-magnetized numerical relativity simulation of a BH-NS merger, endowing the disk with a poloidal field and following its evolution for $60$\,ms, reaching a fully developed MRI and comparing with the non-magnetized case. While there is an existing body of work on the long-term evolution of black hole accretion disks, including a number of 3D GRMHD studies (e.g., \citealt{mckinney2012}), work has focused primarily on systems arising in X-ray binaries and active galactic nuclei, for which thermodynamic conditions, disk size, and the effect of photons and neutrinos are very different than for NS-NS/NS-BH mergers. Because of these differences, the results of previous work that focused on the sub-relativistic disk outflow (e.g., \citealt{narayan2012,sadowski2013}) are not directly applicable to the merger problem . In this paper we close the gap in disk evolution time by performing long-term GRMHD simulations of NS merger accretion disks that for the first time achieve {completion of} mass ejection (i.e., most of the initial disk material either accreted or ejected). In order to minimize the computational cost and evolve our simulations for as long as possible, we employ a number of approximations to the microphysics and neutrino treatment. We focus on understanding the basic properties of the sub-relativistic outflow when MHD turbulence transports angular momentum. To carry out our simulations, we extend the GRMHD code {\tt HARMPI}\footnote{Available at https://github.com/atchekho/harmpi \label{fn:harmpi}} to include the dominant microphysics and neutrino source terms. Simulations are evolved for long enough (several times $10^5 r_g/c$, with $r_g = GM_{\rm bh}/c^2$ the gravitational radius of a black hole of mass $M_{\rm bh}$) to achieve the advective state in the disk evolution and to reach {completion of} mass ejection. While our models resolve the MRI, we consider our work to be exploratory in nature, because not only more spatial resolution but also more physics and realistic initial conditions are required to make quantitative predictions on the wind contribution to the kilonova and $r$-process nucleosynthesis, particularly due to the sensitivity of the latter to the exact outflow composition. {Given that until now only hydrodynamic disk models have been evolved into the advective state, and that a key mass ejection mechanism (thermal energy deposition by angular momentum transport and nuclear recombination) is present in both approaches, we compare the results of our GRMHD simulation with those from hydrodynamic models that employ an alpha viscosity to transport angular momentum. The goal is to identify similarities and differences in mass ejection, thus providing a solid foundation to understand the behavior of the GRMHD model.} The structure of the paper is the following. Section~\ref{s:methods} describes our computational setup, \S\ref{s:results} presents our results, \S\ref{s:obs_implications} discusses the observational implications, and \S\ref{s:summary} closes with a summary and future prospects. The Appendix presents a detailed description of the microphysics included, and a comparison with more complete hydrodynamic models. | \label{s:summary} We have performed long-term, 3D GRMHD simulations of BH accretion disks formed during neutron star mergers. Our models start with an equilibrium torus, a strong poloidal field, and make use of a suitably calibrated gamma-law equation of state, with approximations for the temperature, neutrino cooling, and nuclear recombination that account for the dominant effects of realistic microphysics on the dynamics and composition of the flow. These approximations enable us to maximize the amount of physical time simulated and therefore achieve {completion of} mass {ejection} to large radii. To connect with previous work {and to better diagnose the GRMHD results}, we have also carried out 2D hydrodynamic simulations with shear viscosity and a pseudo-Newtonian potential, and identical treatment of other physics. Our main results are the following: \newline \noindent 1. -- When including MHD effects in general relativity, the total mass ejected from the disk is $40\%$ of the initial torus mass ($0.013M_\odot$ for an initial torus mass of $0.033M_\odot$; Table~\ref{t:models} and Figure~\ref{f:mout_hydro-mhd}). This is larger by a factor of two relative to hydrodynamic models. \newline \noindent 2. -- The ejected mass in the GRMHD model displays a broad distribution of electron fraction, entropy, and radial velocity (Figure~\ref{f:mass_histograms}). The majority of the outflow has $Y_e < 0.25$ and will thus result in a lanthanide-rich composition. When scaling our ejected fractions to a disk mass of $0.1M_\odot$, the outflow from the GRMHD model can easily account for the (nominal) mass and velocity of the red kilonova from GW170817. The disk does not eject sufficient material with $Y_e>0.25$ to account for the blue kilonova, however, despite achieving the right velocity ($0.22c$, Table~\ref{t:models}). We caution, however, that our treatment of weak interactions is approximate and does not include neutrino absorption. Also, the composition of the early outflow will depend more sensitively on the initial distribution of $Y_e$ produced during the merger, which may be on average higher than the uniform and low value we assume ($0.1$). Finally, the appearance of the kilonova can depend on the details of the spatial distribution of the lanthanide mass fraction (e.g., \citealt{kasen_2017}), which in turn depends on the degree of mixing and/or stratification of the electron fraction, not just on the bulk amounts above or below $Y_e \sim 0.25$. \newline \noindent 3. -- Mass ejection in MHD can be divided into two phases: an early ($t\leq 1$\,s as measured from $r=10^9$\,cm), magnetically-mediated phase, absent in hydrodynamic models, and a late phase following freezout of weak interactions (Figures~\ref{f:early_evolution_mhd} and \ref{f:lnu_hydro-mhd}), which operates on the angular-momentum transport timescale. The slow component ($v^r/c<0.1$) of the late-time outflow shows similar properties in both MHD and in hydronamic models (Figure~\ref{f:histogram_late}). This similarity points to a shared mechanism for mass ejection: neutrino cooling freezes out and thermal energy is deposited by turbulent dissipation or viscous heating, and by recombination of nucleons into alpha particles. The GRMHD model has an additional fast component at late times which must be mediated by magnetic processes given its absence in hydrodynamic models. The accretion history, late-time mass ejection, and neutrino luminosity of the GRMHD model is bracketed by the hydrodynamic models that use $\alpha=0.03$ and $\alpha=0.1$ (Figures~\ref{f:lnu_hydro-mhd}, \ref{f:mdot_hydro-mhd}, and \ref{f:mout_hydro-mhd}). \newline \noindent 4. -- Given our initial field geometry, which is optimized for efficient extraction of energy from the BH, we obtain a robust jet carrying $3\times 10^{50}$\,erg of electromagnetic energy (in all directions; Figure~\ref{f:edot_mhd}). A small fraction of the ejecta achieves relativistic velocities at latitudes close to the rotation axis (Figure~\ref{f:histogram_gamma-beta} and \ref{f:histogram_angle}). Comparing with models that fit the non-thermal emission from GW170817 shows that our model contains too much kinetic energy in matter with $\gamma\beta\gtrsim 1$ (\S\ref{s:relativistic_ejecta}). The angular dependence of our jet is compatible with off-axis fits to the non-thermal emission from GW170817 (Figure~\ref{f:histogram_angle}){.} We caution that this component of the disk outflow will almost certainly interact with the dynamical ejecta, and thus the final kinetic energy distribution in a realistic setting will likely differ from that in our models. \newline \noindent 5. -- A few percent of the ejecta in the GRMHD model ($7\times 10^{-4}M_\odot$) also has sufficient velocity and low $Y_e$ to generate free neutrons not captured onto nuclei during the $r$-process (\S\ref{s:obs_implications}). This component can generate early ($\sim $ hr timescale) thermal emission preceding the kilonova. The existence of this component is likely to also be sensitive to the initial magnetic field configuration, and on its location relative to the dynamical ejecta. \newline Our GRMHD model can be improved in many ways to achieve a more realistic result. The simplest modification is changing the initial field geometry, which is bound to have the largest impact on the fastest portion of the ejecta. We have adopted a strong poloidal field with a topology that maximizes energy extraction from the BH. A more realistic disk is expected to contain a significant toroidal component (e.g., \citealt{etienne2012,kiuchi2014}) that can be highly turbulent from the time of disk formation (and thus generate a different thermal evolution than our equilibrium initial conditions). A more challenging improvement involves a realistic treatment of neutrinos (emission and absorption) and microphysics (EOS with full nuclear recombination). For the case of a promptly-formed BH, the effect of neutrino absorption on the bulk of the outflow dynamics is likely to be secondary, given that (1) in hydrodynamic models, for which the outflow is slow and similar to the slow MHD component at late times, absorption is unimportant, and (2) inclusion of magnetic fields adds a fast component which neutrinos have a smaller chance to modify given its rapid expansion. Nevertheless, neutrino absorption is crucial for determining a reliable composition and therefore detailed $r$-process yields for comparison with observations (as demonstrated by \citealt{siegel_2018} who assess the effect of neutrino absorption in post-processing). Also, in the case of weak magnetic fields, neutrino energy deposition can combine with magnetic stresses in launching a successful jet (e.g., \citealt{just_2016,perego_2017}), and the magnitude of this energy deposition is very sensitive to the neutrino scheme adopted \citep{foucart_2018}. If a HMNS survives for longer than a dynamical time, the effect of neutrinos become dynamically important and proper neutrino transport is essential for a complete description. Finally, more realistic initial conditions for density, temperature, and electron fraction can be obtained by mapping from a dynamical merger simulation (as in \citealt{nouri_2017}). While the details of the initial thermodynamics and composition have a relatively minor ($\sim 10\%$) effect on the late-time slow disk outflow composition \citep{fernandez_2017}, the prompt MHD outflow preserves the initial composition of the disk. Also, the mapping from binary parameters (total mass, mass ratio, spins, etc.) to initial disk mass and field geometry is non-trivial and essential for detailed comparison with observations. Finally, evaluating the feasibility of models that rely on the interaction of a jet with the dynamical ejecta (such as models involving cocoon emission; e.g. \citealt{nagakura_2014,lazzati_2017,gottlieb_2018}) requires proper initial conditions for all merger remnant components. | 18 | 8 | 1808.00461 |
1808 | 1808.02187_arXiv.txt | The Transiting Exoplanet Survey Satellite (TESS, launched early 2018) is expected to find a multitude of new transiting planet candidates around the nearest and brightest stars. Timely high-precision follow-up observations from the ground are essential in confirming and further characterizing the planet candidates that TESS will find. However, achieving extreme photometric precisions from the ground is challenging, as ground-based telescopes are subject to numerous deleterious atmospheric effects. Beam-shaping diffusers are emerging as a low-cost technology to achieve hitherto unachievable differential photometric precisions from the ground. These diffusers mold the focal plane image of a star into a broad and stable top-hat shape, minimizing photometric errors due to non-uniform pixel response, atmospheric seeing effects, imperfect guiding, and telescope-induced variable aberrations seen in defocusing. In this paper, we expand on our previous work (Stefansson et al. 2017; Stefansson et al. 2018 [submitted]), providing a further detailed discussion of key guidelines when sizing a diffuser for use on a telescope. Furthermore, we present our open source Python package \texttt{iDiffuse} which can calculate the expected PSF size of a diffuser in a telescope system, along with its expected on-sky diffuser-assisted photometric precision for a host star of a given magnitude. We use \texttt{iDiffuse} to show that most ($\sim$80\%) of the planet hosts that TESS will find will be scintillation limited in transit observations from the ground. Although \texttt{iDiffuse} has primarily been developed to plan challenging transit observations using the diffuser on the ARCTIC imager on the ARC 3.5m Telescope at Apache Point observatory, \texttt{iDiffuse} is modular and can be easily extended to calculate the expected diffuser-assisted photometric precisions on other telescopes. | The Transiting Exoplanet Survey Satellite (TESS) \cite{ricker2015}, launched in early 2018, will survey the whole sky for transiting planets around the nearest and brightest stars. The predicted planet yield of TESS has been studied by numerous groups (e.g., \cite{sullivan2015,ballard2018,barclay2018}) that predict that TESS will detect a multitude of planets---including hundreds of Neptune-sized planets and dozens of terrestrial Earth-sized planets---orbiting stars that are sufficiently bright for detailed follow-up characterization studies from the ground and with future space-based observatories such as JWST. While TESS will transform our understanding of the diversity of exoplanet systems around the nearest and brightest stars, ground-based follow-up observations will be critical in confirming the planetary nature of the transiting signals TESS finds, and will play a key role in maximizing the scientific yield of the mission. These ground-based follow-up observations include adaptive-optics and seeing-limited imaging, high precision radial velocity observations, and further ground-based photometry. In this paper, we restrict our focus on achieving precision photometry from the ground. High precision timely ground-based photometry will play a key role in recovering the ephemerides of a number of planets, which will be essential for efficient scheduling of the most important candidates for further study with JWST\cite{benneke2017,stefansson2018}. Furthermore, high precision ground based photometry can be used as a planet confirmation tool by studying transit depths of planet candidates in different bands \cite{tingley2014}. Such observations can now be performed efficiently by simultaneously observing in a number of different bands with high-precision high-cadence specialized instruments such as MUSCAT\cite{narita2015}, and HiPERCAM\cite{dhillon2016}. Furthermore, high precision photometry can be used to detect transit time variations (TTVs), which can directly give us a handle on the mass of planets in certain planetary system architectures\cite{mazeh2013}. Although high precision photometric observations have been done successfully from space by many groups\cite{beichman2016,benneke2017}, the large number of expected planet candidates from TESS places a great need on the availability of precision photometric instruments from the ground. However, high precision photometry is difficult to achieve from the ground due to deleterious effects from the atmosphere including differential atmospheric extinction, scintillation, transparency fluctuations and seeing effects. Additionally, inter-pixel sensitivity, telescope guiding effects, and the day-night cycle all affect the photometric precision from the ground. Beam-shaping diffusers are emerging as an inexpensive technology to achieve high precision photometry from the ground. Fabricated using specialized nano-fabrication techniques, diffusers are capable of deterministically molding starlight into a broad and stabilized Point Spread Function (PSF) on a telescope imaging array. By spreading out the light over many pixels, allows the observer to increase their exposure times, allowing them to gather more photons and increasing the overall observing efficiency. This increase in observing efficiency ensures that scintillation errors are averaged down, but scintillation errors are the dominating source of error when observing bright stars on large telescopes from the ground \cite{stefansson2017}. These diffusers can be installed on telescopes large and small, both of which play an important role in follow-up efforts from the ground. Smaller telescopes have larger fields of views, allowing them to observe a larger number of reference stars---which becomes especially important when observing bright planet hosts, as nearby bright reference stars are often not available. With their larger diameters, the large telescope will be able to achieve better ultimate photometric precisions due to their enhanced light-collecting area, to both gather more photons, but also better average over scintillation errors. The use of diffusers for high precision ground-based photometric applications has been discussed in detail in \cite{stefansson2017,stefansson2018}. In these proceedings, we expand the discussion on the adaptation of diffusers in different telescopes, discussiong how to size them to yield a required PSF size. To better plan for diffuser-assisted observations from the ground, we present our open-source code \texttt{iDiffuse}, which can be used to calculate the expected diffuser-assisted photometric precision of a telescope system for given host-star. Although currently mostly designed to calculate the expected photometric precision for the diffuser on the ARCTIC imager \cite{huehnerhoff2016} on the ARC 3.5m \cite{stefansson2017}, \texttt{iDiffuse} can be easily extended to calculate the diffused photometric preicsions of other systems. We use \texttt{iDiffuse} to inform what role diffusers can play in the TESS era. This paper is structured as follows. Section \ref{sec:diffusers} gives a short description of diffusers, and how they can be used to achieve high-precision ground-based photometry, giving key best-practices in how to adapt diffusers for use on telescopes for high-precision photometry applications from the ground. Section \ref{sec:planning} describes \texttt{iDiffuse} to better help plan diffuser-assisted follow-up observations of transits. Section \ref{sec:tess} further discusses the role that diffusers can play in the TESS era and beyond. We conclude the paper in Section \ref{sec:summary} with a short summary of the paper. | \label{sec:summary} In this paper we have described diffusers and their use for precision ground-based photometry in the TESS era. We further summarize our key points in the list below. \begin{itemize} \item Diffusers are being incorporated into a number of different observatories. \item We presented \texttt{iDiffuse}, an open source tool to calculate expected on-sky diffuser-assisted photometric precisions. Although \texttt{iDiffuse} has been written with the diffuser on ARCTIC, \texttt{iDiffuse} is modular and can be easily extended to calculate the diffuser-assisted photometric precision on othere telescopes. \item Using \texttt{iDiffuse}, we predict that an ARCTIC+diffuser-like system on a 3.5m telescope should be able to recover at high confidence 4 out of every 5 planet transits that TESS will find. \item Most TESS planet hosts will be scintillation limited in transit observations from the ground. Diffusers help minimize scintillation in conventional telescope imagers by allowing increased duty cycles. \item We expect that the highest ground-based diffuser-assisted photometric precision will be achieved by adapting diffusers on large telescopes with rapid-readout imagers and scintillation correctors. \end{itemize} | 18 | 8 | 1808.02187 |
1808 | 1808.04132_arXiv.txt | New VLA detections of the variable radio continuum source VLA J181335.1$-$174957, associated with the energetic X-ray pulsar PSR~J1813$-$1749 and the TeV source HESS J1813-178, are presented. The radio source has a right circular polarization of $\sim50\%$, and a negative spectral index of $-1.3\pm0.1$, which show that it is non-thermal. The radio pulses of the pulsar are not detected from additional Effelsberg observations at 1.4~GHz made within one week of a VLA detection. This result would appear to support the idea that the continuum radio emission detected with the VLA does not trace the time-averaged emission pulses, as had previously been suggested. We discuss other possible origins for the radio source, such as a pulsar wind, magnetospheric emission, and a low-mass star companion. However, observations made at higher frequencies by Camilo et al. (in preparation) show that the VLA source is in fact the time-averaged pulsed emission and that the detection of the pulses had not been achieved because this is the most scattered pulsar known. | \label{sec:intro} PSR J1813-1749 (=CXOU J181335.1$-$174957) is the second most energetic pulsar in the Milky Way \citep{halpern2012}. It was discovered and identified as a young pulsar { with a spin period of 44.7 ms} by \cite{GH2009} based on Chandra X-Ray observations. Later \cite{halpern2012} used a few Chandra and XMM-Newton observations to determine the spin-down rate of PSR J1813$-$1749 to be $\dot{P}=1.265\times10^{-13}$, corresponding to a spin-down luminosity of $\dot{E}= 5.6\times10^{37}$~erg~s$^{-1}$, thereby establishing it as an energetic young pulsar. These values are only below those measured for the Crab pulsar \citep[e.g.,][]{halpern2012}\footnote{\cite{halpern2012} mentioned PSR J2022+3842 as the second most energetic pulsar, however a recent revision of its spin period by \cite{arumu2014} show that it is 48 ms instead of 24 ms, thus the energy of PSR J2022+3842 is below the energy of PSR J1813$-$1749.}. PSR J1813-1749 is associated with one of the brightest and most compact objects located by the HESS Galactic Plane Survey \citep{aharonian2005}, the TeV source HESS J1813-178. This HESS source has been associated with continuum high-energy emission from X-rays to gamma rays \citep{reimer2008,ubertini2005,abdo2009,albert2006}. Within the TeV extent of this source lies G12.82-0.02, a young relatively compact shell-type radio supernova remnant (SNR) with a diameter of $\sim2'$ \citep{brogan2005}. Deep X-Ray observations revealed that a pulsar wind nebula (PWN) is embedded in the SNR, and is powered by PSR J1813$-$1749 \citep[][]{helfand2007,funk2007,GH2009}. While from observations alone, it is not clear if the TeV emission comes from the SNR or the PWN, \cite{fang2010} predict that the high energy emission is produced mostly in the PWN. \cite{halpern2012} firmly established that PSR J1813-1749 is a young pulsar (characteristic age $\sim 5600$ yr) from their multi-year X-ray observations. On the other hand, several radio observations carried out in the past have not revealed any radio pulsations. The first made at 1.4 GHz on 2005 September 8 with the Australia Telescope National Facility (ATNF) Parkes telescope set an upper limit on the flux density of the radio pulsar as $S_{1.4}<0.07$~mJy \citep{helfand2007}. \cite{halpern2012} did a new search for pulsed emission using the Green Bank Telescope (GBT) in 2009 May 25 and 2009 August 17, and placed upper limits to the period-averaged flux density at 2.0 GHz of ${S_{2.0}<0.013}$~mJy and ${S_{2.0}<0.006}$~mJy, respectively. The common explanation for the lack of radio pulses from high energy pulsars, is that the radio emission beam is narrower than those at higher energies, and thus the radio pulses miss the sightline toward Earth \citep{bra1999,watters2011}. However, it is also known that pulsars may be intermittent emitters, and so observations could miss the pulses if they are scheduled during periods of inactivity. However intermittency is normally associated with pulsars of characteristic age, $\tau_c \approx$ 1~Myr or older \citep{kramer2006,llm+2012,crc+2012}. In an unexpected result, \cite{dzib2010} reported, by the first time, the detection of a compact ($\leq0\rlap{.}''2$) radio continuum source (named VLA J181335.1$-$174957) with a flux density of $180\pm20$~$\mu$Jy at 4.86 GHz, and located within $0\rlap{.}''2$ of the Chandra position of PSR J1813-1749. The observation was made on 2006 February 25. They estimated that the probability of the compact radio source being a background source is $5\times10^{-5}$. Thus there is a high chance that the compact radio source is related to the pulsar and might correspond to the integrated emission of the radio pulses. More recently, however, \cite{dzib2010} did not detect the source during a multi-wavelength radio observation on 2009 March 24, and they were the first to suggest that VLA~J181335.1$-$174957 could be an intermittent radio pulsar. The alternative is that this low signal to noise detection is spurious. We present new Karl G. Jansky Very Large Array (VLA) observations to show that the compact source reported by \cite{dzib2010} is active again. We use these new and archival observations to constrain the nature of the radio continuum source. In addition, we report on the results from pulsar periodicity searches carried out on the observations using the Effelsberg 100m Radio Telescope. These observations were separated only by a week from the VLA observations, which increase the chance of detecting radio pulsations, if any. % \begin{figure*}[!th] \centering \includegraphics[height=0.54\textwidth, trim=43 400 80 50, clip]{PSR_SNR.pdf} \caption{{\it Background:} C-band (5.5 GHz) radio image from 2012 May 6, with the compact source VLA J181335.1-174957 located at the center of the image. Diffuse emission from supernova remnant G12.82$-0.02$ is detected. {\it Contours:} 1.4~GHz radio emission of this SNR \citep[i.e.,][]{helfand2007}. This last image was obtained from the \href{https://third.ucllnl.org/gps/index.html}{MAGPIS webpage} \citep{helfand2006}. } \label{fig:SNR} \end{figure*} | We presented new deep interferometric VLA observations and single dish Effelsberg observations of the continuum radio source VLA J181335.1$-$174957, which has been related to the high energy pulsar PSR J1813-1749 and to the TeV source HESS J1813-178. The radio continuum source was detected in 10 different epochs with the VLA. The observations showed that VLA J181335.1$-$174957 is variable, with a circular polarization of $\sim50\%$ and a spectral index of $\langle\alpha\rangle=-1.3\pm0.1$. These parameters indicate that the radio emission has a non-thermal nature. The observations with the Effelsberg telescope, however, find no pulse emission and support the idea suggested by \cite{halpern2012} that the steady VLA radio source is not the integrated radio pulse emission. We have discussed three different possibilities on the nature of VLA J181335.1$-$174957. The most promising explanations of the radio continuum emission are that it corresponds either to the pulsar wind, or to magnetospheric emission from the pulsar (also known as off-pulse emission). The former explanation is very exciting since it would indicate the first detection of a pulsar wind. However, the information of pulsar winds and the off-pulse emission is still scarce, so we cannot currently favor or discard either of these possibilities. A third explanation is that PSR J1813-1749 has a low-mass stellar companion which produces the non-thermal radio emission. In this case, VLA J181335.1$-$174957 is not directly related to the pulsar. However, there is no independent evidence (e.g.\ based on timing) that PSR J1813--1749 has a companion. The enigma may have a surprisingly simple solution: Camilo et al. (in preparation) have found that PSR J1813-1749 is the most scattered pulsar known and that observations made at higher frequencies are consistent with the radio emission being the time-averaged pulsed emission. | 18 | 8 | 1808.04132 |
1808 | 1808.06071_arXiv.txt | Magnetic reconnection is a fundamental physical process in various astrophysical, space, and laboratory environments. Many pieces of evidence for magnetic reconnection have been uncovered. However, its specific processes that could be fragmented and turbulent have been short of direct observational evidence. Here, we present observations of a super-hot current sheet during SOL2017-09-10T X8.2-class solar flare that display the fragmented and turbulent nature of magnetic reconnection. As bilateral plasmas converge toward the current sheet, significant plasma heating and non-thermal motions are detected therein. Two oppositely directed outflow jets are intermittently expelled out of the fragmenting current sheet, whose intensity shows a power-law distribution in spatial frequency domain. The intensity and velocity of the sunward outflow jets also display a power-law distribution in temporal frequency domain. The length-to-width ratio of the current sheet is estimated to be larger than theoretical threshold of and thus ensures occurrence of tearing mode instability. The observations therefore suggest fragmented and turbulent magnetic reconnection occurring in the long stretching current sheet. | Magnetic reconnection, referring to dissipation and connectivity change of magnetic field, is capable of powering plasma heating, plasma motions, and particle acceleration in relativistic jets \citep{bloom11}, accretion disks \citep{balbus98}, solar and stellar flares \citep{sturrock66}, and magnetospheres \citep{phan06}. In the past decades, abundant evidence for magnetic reconnection has been disclosed including in situ measurements near the Earth and remote sensing observations such as cusp-shaped flare loops \citep{masuda94}, inflows and downflows near the reconnection region \citep{yokoyama01,savage11,takasao12,liuw13,liurui13,xuezhike16}, double hard X-ray coronal sources \citep{sui03}, and changes of connectivity of coronal loops \citep{suyang13,yangshuhong15,lileping16_np}. Unfortunately, the specific processes involved in magnetic reconnection, in particular what occur in the reconnection region, remain mysterious. Theoretically, magnetic reconnection is believed to take place in a localised region, i.e., the so-called current sheet, that has enhanced resistivity \citep{priest14,yamada10}. In the Sweet-Parker model, the current sheet is limited to a long and thin region, in which the reconnection proceeds steadily but too slowly to interpret the real energy release rate. Through invoking slow-mode shocks extending from a shortened Sweet-Parker current sheet, the Petschek model is able to significantly boost the reconnection rate \citep{petschek64}. Nevertheless, the current sheet width in the Petschek model is of ion inertial scale, which can hardly match the detectable width in observations. Therefore, it was proposed that the current sheet can be fragmented into many magnetic islands by tearing mode instability \citep{furth63,shibata01} and develops turbulence to achieve the fast reconnection \citep{lazarian99}. However, such a picture has been short of direct observational evidence although documented by numerical simulations \citep{kowal09,shencc11,barta11} and indicated by various indirect observations such as simultaneous intermittent plasmoid ejections and hard X-ray/radio bursts \citep{asai04,nishizuka09,takasao16}, vortex above flare arcades \citep{mckenzie13,scott16}, and complex transition region line profiles with bright cores and broad wings \citep{innes15}. \begin{figure*} % \vspace{-0.0\textwidth} % \centerline{\hspace*{0.00\textwidth} \includegraphics[width=0.9\textwidth,clip=]{f1.pdf} } \caption{\textbf{Super-hot current sheet in the wake of an erupting bubble on 2017 September 10.} (a) Top: A composition of the AIA 193 {\AA} (red; temperature response peaks at $\sim$1.6 and 18 MK), 131 {\AA} (green; $\sim$0.4 and 11 MK), and 171 {\AA} (blue; $\sim$0.6 MK) images showing an erupting bubble and induced current sheet at 15:53 UT. Middle and bottom: DEM-weighted average temperature and total EM maps showing that the erupting bubble has a high temperature ($\sim$8 MK) but low density at its center. (b) Images at the Hinode-XRT Al-poly, SDO-AIA 193 {\AA} and 131 {\AA} passbands, DEM maps at the temperatures of 2 and 20 MK, and total EM map showing that the current sheet appears as a long and thin feature at 16:15 UT. The vertical dashed line in the AIA 193 {\AA} image indicates the location of the slit used to construct the AIA time-distance plot in Figure \ref{f2}a. Note that, we do not calculate the average temperature and EM of the flare loops and cross-shaped structure as shown in panels a and b because the flux is saturated there (white region). (c) The DEM of the current sheet (left) at five specific regions (small boxes in panel b) and the total EM (right) along five dashed lines as shown in the EM map of panel b.} \label{f1} \end{figure*} In this study, we present a detailed analysis of a limb solar eruption on 2017 September 10 that produced an X8.2-class flare (SOL2017-09-10T16:06UT\footnote{http://sprg.ssl.berkeley.edu/~tohban/wiki/index.php}) and a fast coronal mass ejection (CME). In particular, the presence of a thin and long hot plasma sheet underneath an erupting CME fits perfectly into the current sheet structure, as predicted in the theoretical model \citep{linjun00}, and the dynamic behaviours of the plasma within and around the current sheet provide direct and solid evidence of a turbulent and intermittent nature of magnetic reconnection. | In the models of flux-rope-induced CME/flare eruptions \citep{shibata95,chen11_review}, a pre-existing flux rope escapes away from the solar surface due to loss of equilibrium \citep{linjun00}, leading to the formation of a CME and a flare at almost the same time \citep{zhang01,cheng11_fluxrope}. Magnetic reconnection acts as strong coupling between the CME and the flare as indicated by the simultaneity between the evolution of the CME velocity and the variation of the flare emission \citep{zhang01}. The linear bright feature in the wake of the erupting flux rope has been argued to be the induced current sheet, where electric current is enhanced and magnetic field is dissipated \citep{linj15}. Previous observations of the current sheet are mostly limited by the wavelength window that only responds to relatively narrow and low temperatures and/or the field of view that is not large enough \citep{linjun05,linj09,ciaravella03,ciaravella08,savage10,ling14,seaton17}. Therefore, studies based on these observations are mostly speculative in particular on the origin of the current sheet and its relation to the CME and flare. Moreover, the observations in those works could not provide further information on the detailed physical processes occurring in the current sheet, therefore it has seldom been addressed what kind of reconnection it is. In this study, we present a solar limb eruption event, which displays a distinct picture of the CME/flare eruption with unprecedented clarity. Observations with a continuous field of view from 1 to 30 $R_\odot$ and multi-wavelengths including the white-light, EUV, and X-rays enable us to reveal the origin of and specific processes involved in magnetic reconnection. We successfully detect almost all ingredients predicted by models during a single eruption including the erupting hot flux rope, super-hot current sheet, cusp-shaped flare loops, inflows, and high-speed sunward and anti-sunward outflow jets, some of which have been detected in previous observations \citep{savage10,ling14,seaton17,yanxl18_cs,liuwei18}. The high temperature of the flux rope envelope and the cusp-shaped flare loops probably originates from the collision of the outflow jets with the local dense plasma and/or the direct heating by slow-mode shocks at both ends of the current sheet \citep{liuw13}. The high temperature of the current sheet, however, requires a local heating by magnetic energy dissipation inside the current sheet itself. The turbulent behaviour of energy release in the current sheet is also revealed. A high Lundquist number, suggested by a large length-to-width ratio ($>$16) of the current sheet, leads to the generation of magnetic islands due to tearing mode instability \citep{furth63}, which subsequently appear as intermittent sunward outflow jets and anti-sunward moving blobs when shot out of the current sheet. Simultaneously, the turbulence develops in the current sheet \citep{strauss88,lazarian99}. On the one hand, its effect helps achieve anomalous resistivity to boost magnetic dissipation rate. On the other hand, it may mediate the formation of magnetic islands with their size and energy presenting a power law distribution. This process finally makes the intensity and velocity of the sunward outflow jets exhibit a power law distribution. In particular, the spectral index of the former is found to vary from --1.2 to --1.8, which suggests that the turbulence mediate the reconnection process in the current sheet, resulting the formation of different scaled magnetic islands, consistent with previous numerical results \citep{kowal09,shencc11,barta11}. The deviation from the fully developed isotropic turbulence (with a Kolmogorov turbulence spectral index of --5/3) may be due to the role of magnetic field. The significant non-thermal motions shown in the Fe XXIV line also evidence the existence of turbulence. In summary, these observations show that the magnetic reconnection, at least in solar eruptions, does not proceed uniformly in space and time. Instead, the current sheet should be composed of fragmented structures, in which magnetic reconnection dissipates magnetic energy in a turbulent way \citep{kontar17} to heat the plasma and drive the outflow jets. | 18 | 8 | 1808.06071 |
1808 | 1808.03654_arXiv.txt | A standard prediction of galaxy formation theory is that the ionizing background suppresses galaxy formation in haloes with peak circular velocities smaller than $\Vpeak \simeq 20 \,\kms$, rendering the majority of haloes below this scale completely dark. We use a suite of cosmological zoom simulations of Milky Way-like haloes that include central Milky Way disk galaxy potentials to investigate the relationship between subhaloes and ultrafaint galaxies. We find that there are far too few subhaloes within 50 kpc of the Milky Way that had $\Vpeak \gtrsim 20\,\kms$ to account for the number of ultrafaint galaxies already known within that volume today. In order to match the observed count, we must populate subhaloes down to $\Vpeak \simeq 6\,\kms$ with ultrafaint dwarfs. The required haloes have peak virial temperatures as low as $1,500$ K, well below the atomic hydrogen cooling limit of $10^4$~K. Allowing for the possibility that the Large Magellanic Cloud contributes several of the satellites within 50 kpc could potentially raise this threshold to $10\,\kms$ ($4,000$ K), still below the atomic cooling limit and far below the nominal reionization threshold. | \label{s:intro} One of the foundational developments in near-field cosmology was the discovery of ultrafaint dwarf galaxies in the Sloan Digital Sky Survey (SDSS) (see \citealt{Willman2010} for a review). More recent efforts from DES, PanSTARRS, and MagLiteS (among other surveys) have led to many additional discoveries of ultrafaint Milky Way satellites \citep{Koposov15,ADW15,Laevens15a,Laevens15b,ADW16}; the current census of ultrafaint satellites in the Milky Way is approximately 45. These galaxies are incredibly faint, with luminosities as low as $\rm \sim 350\,L_\odot$, and heavily dark matter dominated \citep{Simon07}. As such, they may represent the long-discussed `Missing Satelltes' of the Milky Way \citep{Klypin1999,Moore99}. We expect many more such objects to exist within the virial radius of the Milky Way; only about half the sky has been surveyed and is only complete to within $\sim 30$ kpc for the faintest dwarfs \citep[][]{Walsh07}. The stars in all ultrafaint galaxies are universally old ($\ga 11 \,{\rm Gyr}$) and this lends credence to the idea that their star formation was shut down in response to reionization at high redshift \citep{Bovill09,Brown14,Weisz14,Wheeler15}. While most of these ancient ultrafaint dwarfs are satellites of larger systems, it is statistically unlikely that environmental quenching could have quenched star formation early enough in these objects to explain the absence of young stars in all of them \citep{KRW18}. Reionization suppression is an attractive mechanism for explaining the uniformly ancient stellar populations of ultrafaint dwarfs, as there should be a dark matter halo mass scale below which galaxy formation is severely limited by the ionizing background \citep{Efstathiou92}. The majority of models that have explored the reionization suppression scale have found that most dark matter haloes with peak maximum circular velocities ($\Vpeak$) smaller than $\Vpeak \simeq 20-30 \,\kms$ are unable to accrete gas after reionzation \citep{Thoul96,Gnedin00,Hoeft06,Okamoto08}. This is not unexpected, as haloes of this size have virial temperatures of $T_{\rm vir} \sim 20,000$ K, which is similar to the IGM temperature after reionization \citep[e.g.][]{McQuinn16,Onorbe17}. Suppression at this scale also naturally solves the missing (classical) satellites problem \citep{Bullock00,Benson02,Somerville02,Kim17,Read18}. More recently, \citet{Ocvirk16} have used full radiative transfer simulations of the Local Group to show that reionization suppresses galaxy formation in haloes with $M_{\rm vir}$ $\simeq$ $5\times 10^{8}$ $\rm M_{\odot}$ measured at $\it{z}$ = 5.5, which is equivalent to $V_{\rm max}$ $\simeq$ 20-25 $\rm km \,s^{-1}$ at this redshift~\footnote{Note that $V_{\rm vir} \propto (1+z)^{1/2}$ at fixed halo mass if we ignore the potential (mild) evolution in the virial overdensity definition.}. A similar quenching threshold is seen in high-resolution hydrodynamic simulations that track dwarf galaxy formation down to redshift zero \citep{Sawala16b,Munshi17,Benitez17,Fitts17,Maccio17}. For example, \citet{Fitts17} have used FIRE zoom simulations to study dwarf galaxy formation and find that the majority of haloes with peak subhalo masses below $10^9\,\Msun$ form no stars. This is equivalent to a threshold at $V_{\rm peak} \simeq 20\,\kms$. A second scale of relevance for low-mass galaxy formation is the atomic hydrogen cooling limit at $10^4$ K, which corresponds to a $V_{\rm peak} \simeq 16\,\kms$ halo. Systems smaller than this would require molecular cooling to form stars. Taken together, one might expect that most ultrafaint satellite galaxies of the Milky Way should reside within subhaloes that fell in with peak circular velocities in the range $16-30\,\kms$, though these systems would have lower maximum circular velocities ($V_{\rm max}$) {\em {today}} as a result of dark matter mass loss after infall onto the Milky Way's halo ($V_{\rm max} \le V_{\rm peak}$). In addition to tidal stripping, the destruction of dark matter subhaloes due to interactions with the potential of the central galaxy itself is a crucial physical process that must be included in any comparison to satellite galaxy counts \citep{D10,Brooks14,Zhu2016,Wetzel2016,SGK2017}. The effect of the central galaxy decreases subhalo abundances by about a factor of two within the virial radius compared to dark matter only simulations; a similar effect is seen in massive elliptical galaxy haloes \citep{Despali17,Graus18}. The enhanced destruction is particularly important for subhaloes close to the central galaxy. At the Milky Way scale, almost all haloes above the resolution limit are destroyed within 20 kpc \citep{SGK2017}. This level of depletion, combined with the pace of new discoveries at small Galacto-centric radii, leads us to ask whether there may be \textit{too many} Milky Way satellites rather than not enough \citep[see, e.g.][]{Jethwa18,Li18}. In this work, we use a new suite of cosmological zoom simulations of Milky Way-like haloes simulated with an evolving Milky Way disk (plus bulge) potential to explore the relationship between dark matter haloes and ultrafaint galaxies. We show that a conventional reionization suppression scale at $\Vpeak \simeq 20\,\kms$ drastically under-produces the count of Milky Way satellites within $50$ kpc of the Galactic center. Assuming the Milky Way is typical of our simulation suite, it appears that a significant fraction of ultrafaint galaxies must form in subhaloes with peak circular velocities less than $10\,\kms$. These haloes have virial temperatures of $T_{\rm vir}$ $< 4,000$ K, which is well below the nominal atomic hydrogen cooling limit. In section, \ref{s:sims and methods} we describe the simulations. Section \ref{s:results} presents our results. Section \ref{s:caveats} discusses some possibilities that could alter our conclusions and Section \ref{s:conclusion} provides a summary discussion. | \label{s:conclusion} In this work, we utilized a new suite of 12 cosmological simulations of Milky Way-like haloes that include a central disk potential (Kelley et al., in preparation). The haloes were chosen to match the mass of the Milky Way ($ M_{\rm vir}=0.8-2\times10^{12}$ $\rm M_{\odot}$). The galaxy is modeled as an evolving disk + bulge potential that grows to match the Milky Way at $z=0$. The inclusion of the galaxy drastically affects the subhalo abundance at small radii, as shown in Figure \ref{fig:subhalo_props}, which is consistent with past work based on hydrodynamic simulations \citep{Brooks14,Zhu2016,Wetzel2016,SGK2017,Graus18}. We compared the subhalo distributions at small radius to the current census of galaxies that exist within 50 kpc of the Milky Way disk, most of which are ultrafaint galaxies discovered in digital sky surveys \citep[e.g.][]{Willman05,Zucker06,ADW15,Koposov18}. Even though our census of small galaxies is likely incomplete, we require very low-mass haloes to host galaxies in order account for all of the presently-known galaxies (see Figure~\ref{fig:large_small_radii}). In particular, we need haloes as small as $ V_{\rm peak} = 6 \,\kms$ to host ultrafaint galaxies. These systems have virial temperatures as low as $\sim 1,500$ K, which is not only well below the typical scale where reionization is expected to suppress galaxy formation ($\Vpeak \simeq 20 ~\kms$ and $T_{\rm vir} \simeq 15,000$ K) but also much smaller than the atomic cooling limit ($\Vpeak \simeq 16 ~\kms$ and $T_{\rm vir} = 10,000$ K). We explored these results by mock-observing our Milky Way haloes over regions that mimic DES and SDSS fields using toy models that allow the fraction of haloes that host galaxies to vary from 0 to 1 at a characteristic $\Vpeak$ scale. As shown in Figure \ref{fig:V_peak_sims}, the models that work populate haloes with $\Vpeak$ values between $6$ and $16 ~\kms$ ($T_{\rm vir} = 1,500 - 10,000$ K), much smaller than would be conventionally expected. It is important to note that any new discoveries of ultrafaint dwarf galaxies close to the Milky Way would only increase the need to populate very small haloes with galaxies. There are at least three possibilities that could change these conclusions. First, several of the ultra-faint `galaxies' we include in our analysis (Table 1) could be misclassified star clusters. Ultra-faint galaxies are differentiated from star clusters by inhabiting dark matter haloes \citep{WS12}. If some fraction of the satellites in our comparison do not have dark haloes then the implied threshold $\Vpeak$ for galaxy formation would increase accordingly. A second possibility is that many, if not most dwarfs that were discovered by DES were brought in with the LMC. If this is the case, it would bias the Milky Way dwarf population to be over-abundant compared to what is typical for haloes of the Milky Way's mass. Of the 18 total dwarf galaxies within 50 kpc of the Milky Way, five were discovered in DES. Figure \ref{fig:large_small_radii} shows that if all five of these were removed from the comparison the need to populate very low-mass subhaloes would be lessened, but that there would still be tension unless we allow for most $\Vpeak \simeq 10 ~\kms$ haloes to host galaxies. Even in this fairly conservative scenario (which ignores sky coverage incompleteness) we require galaxy formation below the atomic cooling limit. A third possibility is that much of the destruction we see is a numerical artifact \citep{vdb18a}. By conventional convergence-test standards we appear to be well-resolved down to $\Vpeak = 6 ~\kms$ and certainly to $\Vpeak = 10 ~\kms$; however, by the criteria described in \citet{vdb18b}, we could be affected by numerical issues at the critical scale $\Vpeak = 20 ~\kms$. If so, then ultrafaint galaxies may reside within halos with $V_{\rm peak}$ above the canonical reionization suppression scale, and the numerical challenge facing Milky Way satellite modelers will become quite significant. Aside from the caveats discussed above, our results suggest that haloes well below the atomic cooling limit host ultrafaint galaxies. One implication of such a scenario is that we would expect $\sim 1000$ such systems within 300 kpc of the Milky Way (see the left panel of Figure ~\ref{fig:subhalo_props} at $\Vpeak \simeq 7 ~\kms$ and the last row of Table 1 in \citealt{Kim17}). This count will be testable with LSST and is significantly higher than previous completeness-correction estimates, which had relied on dark-matter only simulations. Whatever the answer, the results presented here motivate renewed efforts to understand galaxy formation in the the smallest haloes and the effect of the near-field environment on their evolution. | 18 | 8 | 1808.03654 |
1808 | 1808.03476_arXiv.txt | Recent developments in the observation and modelling of kink oscillations of coronal loops have led to heightened interest over the last few years. The modification of the Transverse Density Profile (TDP) of oscillating coronal loops by non-linear effects, in particular the Kelvin-Helmholtz Instability (KHI), is investigated. How this evolution may be detected is established, in particular, when the KHI vortices may not be observed directly. A model for the loop's TDP is used which includes a finite inhomogeneous layer and homogeneous core, with a linear transition between them. The evolution of the loop's transverse intensity profile from numerical simulations of kink oscillations is analysed. Bayesian inference and forward modelling techniques are applied to infer the evolution of the TDP from the intensity profiles, in a manner which may be applied to observations. The strongest observational evidence for the development of the KHI is found to be a widening of the loop's inhomogeneous layer, which may be inferred for sufficiently well resolved loops, i.e $>$ 15 data points across the loop. The main signatures when observing the core of the loop (for this specific loop model) during the oscillation are: a widening inhomogeneous layer, decreasing intensity, an unchanged radius, and visible fine transverse structuring when the resolution is sufficient. The appearance of these signatures are delayed for loops with wider inhomogeneous layers, and quicker for loops oscillating at higher amplitudes. These cases should also result in stronger observational signatures, with visible transverse structuring appearing for wide loops observed at SDO/AIA resolution. | \label{sec:intro} Kink (or transverse) oscillations of coronal loops have been intensively studied over the last two decades since their detection with the Transition Region And Coronal Explorer (TRACE) \citep{1999SoPh..187..229H} in 1999 \citep{1999ApJ...520..880A, 1999Sci...285..862N}. Many examples of standing kink modes have been clearly observed \citep[e.g][]{2012A&A...537A..49W, 2016A&A...585A.137G,2016SoPh..291.3269S, 2017ApJ...842...99L} with the enhanced spatial and temporal resolution of the Atmospheric Imaging Assembly (AIA) on board the Solar Dynamics Observatory (SDO) \citep{2012SoPh..275...17L}. The accepted mechanism for the rapid damping of these oscillations is resonant absorption \citep[e.g.~recent review by][]{0741-3335-58-1-014001}. Kink oscillations can be used to perform seismology and obtain estimates for the local plasma parameters \cite[e.g][]{2001A&A...372L..53N,2012RSPTA.370.3193D,2013A&A...552A.138V}, which can aid studies of other processes in the Suns atmosphere. Recently in \cite{2016A&A...589A.136P,2017A&A...600A..78P} this approach was updated to include the proposed Gaussian and exponential damping regimes \citep{2013A&A...551A..39H,2013A&A...551A..40P}, which makes the inversion problem well posed when the switch between the two regimes can be observed. This is based on a simplified model of the Transverse Density Profile (TDP) of the loop, described by a uniform core with an inhomogeneous layer where the density varies linearly between the background and internal density. This analytic description is also subject to the thin boundary layer approximation, which has been shown to vary the damping rates due to resonant absorption by a factor of 2 for thick non-uniform layers \citep{2004ApJ...606.1223V}. An agreement between the seismologically determined TDP and the TDP inferred from the transverse intensity profile was found in \cite{2017A&A...600L...7P}. This was extended in \cite{2017A&A...605A..65G} to perform a statistical study of the TDPs of 233 coronal loops. In many cases there was evidence for a Gaussian transverse density profile, or thick linear inhomogeneous layer. The Bayesian inference approach allows the different density profiles to be quantitatively compared, and also allows robust estimation of the uncertainties of the model parameters. This study indicated that loops may have thicker boundary layers than is typically assumed, or even constantly varying TDPs, subject to the limitations and simplifications discussed. Additionally in \cite{2018ApJ...860...31P} this technique was applied to an oscillating coronal loop to infer the time evolution of the density profile model parameters. Seismological studies assume that the TDP of the loop remains constant during the oscillations, making it important to understand any changes which do occur. There are many effects and non-linear mechanisms which can cause the TDP to vary, which will in turn modify the observed damping behaviour. Large amplitude kink waves have been shown to produce plasma flows along the field, and the ponderomotive force can cause accumulation of density at the loop top \cite[e.g][]{2004ApJ...610..523T, 2009PhPl...16g2115C, 2012A&A...544A.127V}. The effect of a time-varying cross-section was recently investigated analytically in \cite{2017A&A...602A..50R}. In \cite{2016A&A...590L...5G} it was observed that the quality factor of kink oscillations decreases as the oscillation amplitude increases, indicating that finite amplitude effects are playing some role in modifying the damping time and/or period. This could be due to effects which modify the structure of the loop at high amplitudes. A qualitatively similar dependence was found in \cite{2016A&A...595A..81M}, where non-linear effects such as the growth of the KHI instability were found to modify the damping profile of the kink mode at high amplitudes. The Kelvin-Helmholtz instability (KHI) \citep{1984A&A...131..283B, 1994GeoRL..21.2259O} has been shown to occur in numerical simulations within the inhomogeneous layer of oscillating loops or prominences due to the shear flows, redistributing both the density and temperature of the plasma \citep[e.g][]{2008ApJ...687L.115T, 2010ApJ...712..875S, 2014ApJ...787L..22A}. In particular \cite{2016ApJ...830L..22A, 2017ApJ...836..219A} included forward modelled EUV emission from a loop subject to the KHI instability, noting how the oscillation can appear decayless under certain circumstances, and how seismology can be performed based on the phase mixing which takes place. The loop's structure can also evolve during oscillations if they contain unresolved sub-structure or multi--threadedness. Recently \cite{2016ApJ...823...82M} showed that transverse oscillations in loops with substructure cause the strands to merge and produce a more homogeneous density structure. The KHI vortices generated in these simulations are often referred to as Transverse Wave Induced KHI (TWIKH) rolls. In this paper the technique used in \cite{2017A&A...600L...7P,2018ApJ...860...31P} and \cite{2017A&A...605A..65G} is applied to forward modelled EUV emission from numerical simulations of kink oscillations \citep{2017ApJ...836..219A}. In contrast to \cite{2017ApJ...836..219A}, the effect of the development of KHI on the parameters of a TDP model inferred from the EUV emission is investigated. The variation of the parameters may be detected even when the resolution is not sufficient to resolve the complex substructure generated by the TWIKH rolls. In addition, simulations with a larger amplitude of oscillation, and a loop with a larger inhomogeneous layer are analysed and compared. In Sect.~\ref{sec:numdat} the numerical models and data are described. In Sect.~\ref{sec:intfit} the analysis is described and the effects of resolution and noise are explored. The obtained temporal evolution is presented in Sects.~\ref{sec:m1} and \ref{sec:m2m3}, and further discussion and summary are given in Sect.~\ref{sec:disc}. \begin{figure} \centering \includegraphics[width=9.5cm]{eps_est.pdf} \caption{The initial density and temperature profile of the loop in models M1 and M3 (dotted and dashed respectively) and an approximation of the density profile using Model $L$ (solid red). The $x$ axis is given as the radial coordinate $r$ divided by the loop minor radius $R$, defined to occur halfway though the loop's inhomogeneous layer. The fit shown by the red line gives a value for the width of the inhomogeneous layer of $\epsilon$ = 0.32.} \label{dprof} \end{figure} | \label{sec:disc} The main purpose of this study is to describe the inferred evolution of the loop's Transverse Density Profile (TDP) caused by non-linear effects which occur during kink oscillations (largely due to the KHI instability). This is motivated by recent advances in kink oscillation observations \citep[e.g][]{2016A&A...585A.137G}, modelling \citep[e.g][]{2017A&A...601A.107P, 2017A&A...602A..74H} and seismology \citep[e.g][]{2017A&A...603A.101L, 2017A&A...600A..78P}. For observational analysis and theoretical works which assume the transverse structure is stationary it is important to determine if the numerically and analytically modelled processes which cause evolution of the loop's transverse structure can be detected in observations. As discussed in previous studies there are several shortcomings of the method used to infer the TDP of a coronal structure from the observed intensity profile. Since an isothermal approximation is made, i.e the temperature inside and outside the loop is assumed to be equal, temperature variation is detected as variation of the density, due to a variation of the instrumental response function. Any of the plasma which emits at temperatures not covered by the chosen AIA wavelength is not detected. The numerical data used corresponds to a loop which is far from isothermal, however in Section 3 reasonable estimates for the density profile are obtained from the intensity profile in the AIA filter which corresponds to the core of the loop. Due to isothermal approximation in our method the radius ($R_L$) and inhomogeneous layer width ($\epsilon$) are underestimated, by $\approx$ 10 and 20\% respectively. In Sect. \ref{sec:m1} the effect that downsampling the resolution and adding noise has on the final intensity profile of the loop was highlighted. A resolution of approximately 20 points across the loop, corresponding to a radius of $R_{L}$ = 3.5 Mm at AIA resolution, is seen to mask the appearance of the intensity peaks from the TWIKH rolls. This corresponds to the wider loops observed with AIA, and is therefore a best case scenario for current observations. Evolution of the loop in the TD maps can clearly be seen, as well as evolution of the density profile parameters inferred from the intensity profile. The main observational signatures when using Model $L$ for the density profile are; decreasing density enhancement ($A$), a widening inhomogeneous layer ($\epsilon$), a constant minor radius ($R_L$) and almost no visible transverse structuring. The visible decrease in intensity (and in the inferred value of $A$) was also detected in \cite{2017ApJ...836..219A}, and is due to the mixing of the internal and external plasma. The widening of the inhomogeneous layer was detected in \cite{2017ApJ...836..219A} as an increase in the loop's minor radius. Strong oscillatory behaviour is also seen in the time series of the TDP parameters, due to the effect of the oscillation itself on the TDP. In Sect. \ref{sec:m2m3} the effect of varying the width of the inhomogeneous layer and increasing the oscillation amplitude was investigated. The main difference in the former case was the delayed onset of the KHI despite the increased efficiency of the resonant absorption, as it takes longer for the sharp gradients in density and velocity to be generated. The variation of the TDP parameters after this onset is stronger however. The larger amplitude quickened the onset of the KHI and caused strong variation in the TDP of the loop, causing the appearance of multiple strands, visible even at the lower resolution $R2$. Both TD maps and time series for M2 and M3 are limited in length due to the numerical stability, the observational signatures are expected to be even stronger in reality. In \cite{2018ApJ...860...31P} the evolution of the inferred TDP over time for the analysed loop is presented, finding that the parameters showed some oscillatory behaviour, but no strong overall trend. However, this lack of KHI signatures could be due to the low oscillation amplitude. Further examples should be chosen and analysed in the same manner. The study and technique should also be extended to incorporate other EUV wavelengths or data from other instruments. This method for inferring the TDP is limited by the density profile used. The strong peaks in intensity generated in M3 meant that the uncertainties on the inferred parameters became large, as they can not be modelled by Model $L$. For observational searches of KHI in oscillating loops it will be difficult to observe the TWIKW rolls directly. This is in part due to the unknown level of substructure within coronal loops. It is often difficult to determine if there are many threads within a given coronal loop, or if they are spatially separated along the line of sight \citep[e.g][]{2013A&A...556A.104P, 2016ApJ...826L..18B, 2017ApJ...840....4A}. The results presented here can be compared to observations of coronal loops which appear homogeneous in a given EUV channel, and it should be noted that a simple cool and dense loop model has been used. More detailed analysis of the observable signatures of the KHI in different observations was given in \cite{2017ApJ...836..219A}. Further work should also be done in the context of the numerical simulations performed. The timeseries for M2 and M3 could be extended with additional numerical treatment. Additionally, the effects of reconnection could be explored, due to the turbulence induced by the TWIKH rolls, and the effect of this on the observational signatures. The changing width of the inhomogeneous layer highlighted has implications for the damping of the kink mode, as well as the seismology which is based on the damping behaviour. It also has implications for the spatial distribution of the energy disposition. Detection of significant evolution of coronal loop parameters during oscillations would increase the need for the inclusion of non-linear effects in observational analysis and theoretical modelling. | 18 | 8 | 1808.03476 |
1808 | 1808.09909_arXiv.txt | A brief review of the recent astronomical data, indicating that the universe is abundantly populated by heavy black holes (BH), is presented. Conventional astrophysics and cosmology cannot explain such a high population of BHs. A mechanism of the paper of 1963 is described, which at least qualitatively explained the observational data. In particular, the prediction that massive primordial BHs can be cosmological dark matter "particles" is discussed. | {In this talk two different but related subjects are considered:}\\ {\it I. A mechanism of massive PBH formation, which is much different from the caninocal ones.} {According to this mechanism the fundament of PBH creation is build at inflation by making large isocurvature fluctuations at relatively small scales, with practically vanishing density perturbations.} {It is achieved by a simple modification of a popular scenario of baryogenesis. Density perturbations are generated rather late after the QCD phase transition.} {The emerging universe looks like a piece of Swiss cheese, where (the black) holes are high baryonic density objects occupying a minor fraction of the universe volume.}\\ {{\it II. A brief review of new (and not only new) astronomical data which are in strong tension with the accepted standard cosmological model.}} {The data nicely fit the suggested scenario of PBH formation. More and more observational evidence in favor of massive and supermassive PBH {and other, not yet explained in the standard way phenomena} in the universe, appears practically every week. {The usual or astrophysical black holes are the results of stellar evolution after a star exhausted its nuclear fuel and collapsed into a compact object, either into a neutron star or into a black hole. This collapse into black hole happens if the mass of the star is, roughly speaking, larger than three solar masses. Such black holes can be created in sufficiently old or even contemporary universe.} {There exist also the astrophysical supermassive black holes (SMBH) with huge masses, ${M \sim(10^6 - 10^{9}) M_\odot}$, where $M_\odot \approx 2 \times 10^{33} $ g is the solar mass. They are supposed to be the products of matter accretion to smaller BHs or matter accretion to matter excess in galactic centers.} Primordial black holes (PBH) formed in the very early universe if the density excess at the cosmological horizon was large, ${\delta \rho /\rho \gtrsim 1}$, as it was first suggested by Zeldovich and Novikov~\cite{ZN}. {Normally the masses of PBH created through this mechanism were supposed to be rather low and the mass spectrum was quite sharp, close to the delta-function.} An alternative mechanism of formation of very massive PBH with wide spread mass spectrum was proposed in ref{. \cite{ADJS}, see also the subsequent paper~\cite{DKK} , where the model was further elaborated. {Heretic declaration of 1993} is turning now into the acknowledged faith with more and more astronomical data proving its truth. The conclusion that the observed mysterious phenomena are induced by {\it massive primordial} black holes allows to cure an avalanche of inconsistencies with the standard ${\Lambda}$CDM cosmology and astrophysics. In the most extreme form the suggested in 1993 scenario gives rise to:\\ {$\bullet$ {Cosmological Dark Matter fully (or predominantly) made by PBHs.\\ {$\bullet$ {Primordial creation of majority of black holes in the universe.}}\\ {$\bullet$ {Very early QSO formation}.}\\ {$\bullet$ {Early creation of metals and dust}.}\\ {$\bullet$ {Seeding of large galaxies by supermassive PBH.}}\\ $\bullet$ {Seeding of globular clusters by ${10^3 - 10^4}$ PBHs} and dwarf galaxies by ${10^4 - 10^5}$ PBHs~\cite{AD-KP}.\\ {$\bullet$ {Clouds of matter with high baryon-to-photon ratio.}}\\ $\bullet$ {A possible by-product: plenty of (compact) anti-stars, even in the Galaxy.} Due to lack of space many references are omitted here. They can be found in the reviews~\cite{monsters,AD-UFN}. together with more detailed discussion % | } The described mechanism~\cite{ADJS,DKK} nicely fits the basic trend of the surprising astronomical data reviewed in Sec.~\ref{s-rev-obs}. In a sense this picture of the universe was predicted a quarter of century ago. There are still quite many problems demanding deeper research:\\ $\bullet $ Quantitative study of the galaxy formation seeded by massive PBH. In particular it would be interesting to find how the mass of the seed influences the emerging galaxy type. Such an analysis worth doing for large galaxies, dwarfs, and globular clusters. \\ $\bullet $ Formation of galaxies in low density regions of the universe. Such a study could explain existence of small galaxies with superheavy black holes inside, or even existence of superheavy BH in almost empty space.\\ $\bullet $ The problem of high velocity stars in galaxies. What is the source of their acceleration? Are they primordial stars or are accelerated by the abundant population of the PBH with intermediate masses, $(10^3-10^5) M_\odot$?\\ $\bullet $ A possible outcome of the described here mechanism of massive PBH creation is an accompanying creation of antimatter objects, which may be abundant in the Galaxy. Their discovery would prove the validity of the mechanism. | 18 | 8 | 1808.09909 |
1808 | 1808.07313_arXiv.txt | After an eleven year observing campaign, we present the combined visual--spectroscopic orbit of the formerly unremarkable bright star HR~7345 (HD~181655, HIP~94981, GJ~754.2). Using the Separated Fringe Packet (SFP) method with the CHARA Array, we were able to determine a difficult to complete orbital period of $331.609 \pm 0.004$ days. The 11 month period causes the system to be hidden from interferometric view behind the Sun for 3 years at a time. Due to the high eccentricity orbit of about 90\% of a year, after 2018 January the periastron phase will not be observable again until late 2021. Hindered by its extremely high eccentricity of 0.9322 $\pm$ 0.0001, the double-lined spectroscopic phase of HR~7345 is observable for 15 days. Such a high eccentricity for HR~7345 places it among the most eccentric systems in catalogs of both visual and spectroscopic orbits. For this system we determine nearly identical component masses of 0.941 $\pm$ 0.076 $M_{\odot}$ and 0.926 $\pm$ 0.075 $M_{\odot}$ as well as an orbital parallax of 41.08 $\pm$ 0.77 mas. | \subsection{Observational History of HR~7345} HR~7345 (HD~181655, HIP~94981, GJ~754.2) is a 6th magnitude star in the constellation Lyra near the boundary line with Cygnus and seemed for decades to be an unremarkable system aside from it's proximity to the Sun. We discovered it to be a highly eccentric binary, even though it had not been identified in previous multiplicity surveys. It was included as part of a David Dunlap Observatory (DDO) radial velocity survey of 681 relatively bright stars for which velocities were lacking \citep{young45}. A decade later \citet{HALI55} used the same DDO spectra to calculate the luminosity and spectroscopic parallax ($\pi_{sp}$ = 79.7 mas) for the system and classified its spectrum as G8~V. Two years later, \citet{CRIS57} determined its trigonometric parallax ($\pi_{trig}$ = 39 mas) from photographic plates and found a value very close to the modern Hipparcos \citep{Gaia16} measurement ($\pi_{Hip}$ = 39.34 mas). During the next 20 years, it was spectroscopically and photometrically measured and classified, with no hint of variability, even being listed as a radial velocity standard star by \citet{BE79} from 20 measurements over three years at the Fick Observatory. The system was measured six more times between 1978 and 1983 at the McDonald Observatory 2.1-m telescope where it just barely fell outside of the 2.5~$\sigma$ error limit for their definition of a ''radial-velocity standard star" \citep{BA86}. HR~7345 was even observed in the early eighties by the primary author's thesis advisor with speckle interferometry \citep{McA1987} on the Kitt Peak 4~m telescope, which gave a null result for their single observation in 1985. The timing of that measurement was particularly unlucky, as the companion would have been just coming out of periastron but not yet separated enough for easy resolution on the 4~m telescope. Radial-velocity data were again acquired during the CORAVEL survey of \citet{DM1}, who searched for companions of solar-type stars in the solar neighborhood. Twelve measurements of HR~7345, taken over the course of 1200 days between 1983 and 1989, indicated very little velocity variation. In retrospect, their timing was most unfortunate, as they were tantalizingly close to the very short observing window when the system would have exhibited double lines between 1983 to 1985. Unfortunately, again because of the very limited window, nearly all subsequent observations of HR~7345 failed to show evidence of binarity \citep{Duf95, Feh97, Nid02, Gray03, Halb03, Nord, Val05, Holm09, Crif10, Soub13, Gaia16}. In fact, out of all the radial velocity observations collected during the past 50 years, only one, an ELODIE spectrum acquired in 2000 November, was close enough to periastron to exhibit partial separation of its double lines \citep{Prug07}. Finally, the Palomar-Testbed Interferometer observed the system over 40 times between 1998 and 2005 and saw no evidence of a companion with their 86-110m baselines and deemed it to be a suitable calibrator star \citep{GVB07}. Further inquiry into the reasons why it was not detected are ongoing, but due to orbital elements projected backwards, there were several years when the companion should have been detectable. \subsection{High Eccentricity Binaries} As is often the case, many of the most interesting systems are discovered by accident. Although not originally considered until many observations were obtained, the importance of surveying high-eccentricity binary systems cannot be overstated. Characterizing high-eccentricity systems can provide insight into the statistics of stellar formation mechanisms, multiplicity fractions and star formation rates which all lead to the Inital Mass Function (IMF) \citep{AMB1937, BATE2009, TOKO2016b}. It is well known that both visual and spectroscopic observations can easily miss a significant fraction of high-eccentricity systems \citep{DM2, RAG2010, GRI2012} by something as simple as timing where either the relative motion of the components in the visual case is very slow for most of their orbital period, or the relatively short time when the spectra would exhibit double-lines. This combined with the wide variation of inclinations and distances means even discovering high-eccentricity systems is often left up to a chance observation. Once found, these systems can help define the limits of the eccentricity distributions \citep{RAG2010}, relations between period and eccentricity \citep{FIN1936}, and the mechanism that creates such extreme systems (See Section \ref{HEBC}). | Prior to the observations taken in 2018 January, the calculated orbit left many open issues because of the unobserved part of the orbit around periastron. The value and large error in inclination due to the missing part of the orbit produced masses that were significantly lower than expected, but a large enough error bar to include the canonical masses for stars of that spectral classification. While the data and orbit were deemed solid enough to publish as they were, we decided to make one more attempt if the weather permitted to improve the combined orbit. The weather over Fairborn Observatory proved to be significantly more cooperative in obtaining spectra of the system during the 2018 periastron and filled in the missing single day of phase coverage. Observations with the CHARA Array were setup for the four days on either side of the predicted periastron passage, but due to wind, clouds, and humidity, we were only able to obtain a single data set directly on periastron. Fortunately, the results from this single night proved to be very valuable and allowed a final recalculation of the orbit, which provided significant reductions in the inclination value and associated errors. While the 2018 observations significantly improved the orbit and resulting masses, those masses, although consistent with the spectral type, have relatively large uncertainties. With the use of our orbital period and component masses, we obtained from Kepler's third law a semi-major axis of 1.156 AU. The system's large eccentricity produces a periastron separation of 0.079 AU or 16.9 $R_{\sun}$. \citet{h81} has shown that for stars in an eccentric orbit the rotational angular velocity of an individual star will tend to synchronize with that of the orbital motion at periastron. He called this situation pseudosynchronous rotation. To see whether the components of HR~7345 have achieved that state, we first determined the projected rotational velocities of the components from our rotational broadening fits and found $v$~sin~$i$ values of 2.6 $\pm$ 1.0 km~s$^{-1}$ for each component. If the rotational and orbital axes are parallel, as is generally assumed for stars in binary systems, then our orbital inclination of 29.5\arcdeg produces rotational velocities of 5.3 $\pm$ 1.0 km~s$^{-1}$. From Kepler's third law, the pseudosynchronous period is 5.88 days. If we adopt radii of 0.95~$R_{\sun}$ for the G5 dwarfs, then the predicted pseudosynchronous velocities of the components are 8.2 km~s$^{-1}$. Thus, although the rotational velocities that we have determined for the two components of HR~7345 are larger than typical values of about 2--3~km~s$^{-1}$ for most late-type stars \citep{Val05}, the components of HR~7345 have not yet reached pseudosynchronous rotation. In a review of precise masses and radii for normal stars, \citet{Torres10} cataloged 23 systems with both visual and spectroscopic orbits that produced component stellar masses determined to better than 3 percent. As listed in Table \ref{orbparam}, the masses of HR~7345 have a 9 percent uncertainty. As most of the error in the masses is tied to the inclination, more observations around periastron should refine the orbit further and decrease the mass uncertainty. With the current orbital elements, the ideal time to observe the system is the during the week before and after periastron, when the system is separated by less than 15~mas. Such observations will be attempted around 2021 September 11, during the next easily observable periastron passage. | 18 | 8 | 1808.07313 |
1808 | 1808.00705_arXiv.txt | As one of the series of papers reporting on a large reverberation mapping campaign, we apply the maximum entropy method (MEM) to 9 narrow-line Seyfert 1 galaxies with super-Eddington accretion rates observed during 2012-2013 for the velocity-delay maps of their H$\beta$ and H$\gamma$ emission lines. The maps of 6 objects are reliably reconstructed using MEM. The maps of H$\beta$ and H$\gamma$ emission lines of Mrk 335 indicate that the gas of its broad-line region (BLR) is infalling. For Mrk 142, its H$\beta$ and H$\gamma$ lines show signatures of outflow. The H$\beta$ and H$\gamma$ maps of Mrk 1044 demonstrate complex kinematics -- a virialized motion accompanied by an outflow signature, and the H$\beta$ map of IRAS F12397+3333 is consistent with a disk or a spherical shell. The \hb\ maps of Mrk~486 and MCG~+06-26-012 suggest the presence of an inflow and outflow, respectively. These super-Eddington accretors show diverse geometry and kinematics. Brief discussions of their BLRs are provided for each individual object. | Broad emission lines are prominent features in ultraviolet and optical spectra of active galactic nuclei (AGNs) and are believed to stem from the so-called broad-line regions (BLRs), which are photoionized by the ionizing radiation from accretion disks surrounding the central supermassive black holes \citep{osterbrock1989}. The broad emission lines reverberate in response to the varying ionizing continuum with a light-traveling time delay, therefore, appropriate analysis of reverberation properties of broad emission lines delivers information on the kinematics and geometry of BLRs (e.g., \citealt{bahcall1972,blandford1982,peterson1993}). Velocity-resolved time-lag analysis measures the time lags as a function of line-of-sight velocity. In practice, it divides the emission line into several velocity bins and carries out cross-correlation analysis of the light curves at different velocities. This is a preliminary step offering a glimpse into the geometry and kinematics of the BLRs, and has been applied to a number of objects (e.g., \citealt{bentz2008,bentz2009,bentz2010,denney2009a, denney2009b,denney2010,grier2013}). It has been shown that BLRs have diverse geometry and kinematics (such as outflows, inflows or virialized motion, see \citealt{gaskell1988,grier2013}; \citealt{du2016b}, hereafter \citetalias{du2016b}). The signature of inflow/outflow is identified as the mean lag being smaller on the red/blue wing of the line profile, while a disk-like BLR has a symmetric pattern with smaller lags on both the red and blue wings. However, the velocity-resolved time-lag analysis measures the mean time lags at different velocities rather than revealing the detailed response features of broad emission lines. A velocity-delay map resolves the response of broad emission line not only at different velocities but also at different time-delay, therefore embodying all the information on the BLR response to the varying continuum \citep{bentz2010,pancoast2012,grier2013,pancoast2014b,grier2017}. The maximum entropy method (MEM, \citealt{horne1991,horne1994}, see details in Section \ref{sec3}) and dynamical modeling method (\citealt{pancoast2011,pancoast2014a,li2013}) were developed for this purpose. \begin{figure*} \centering \includegraphics[angle=0,width=0.7\textwidth]{AW.pdf} \caption{\footnotesize H$\beta$ velocity-delay maps of Mrk~142 for different values of $\calA$~and $\calW$ ($\alpha$ is fixed to 2000). From the left to right panels, the value of $\calW$ increases by factors of 5 while the delay maps show less flexibility. From the top to bottom panels, the increasing $\calA$ by factors of 10 progressively strengthens the weight of entropy in the velocity direction, and smears the sub-structure in this direction. The values of $\calA$~and $\calW$ in the central panel are our choice. The color bar is shown at the lower-right corner. } \label{AW} \end{figure*} Since the fall of 2012, we have monitored a sample of high accretion rate AGNs, aiming at better understanding the role of accretion rates on BLRs (\citealt{du2014}, hereafter \citetalias{du2014}; \citealt{du2015}, hereafter \citetalias{du2015}; \citealt{du2016a}, hereafter \citetalias{du2016a}) and the physics of accretion onto black holes (\citealt{wang2013}; \citealt{wang2014b}). The observations in the first three years show that super-Eddington accreting massive black holes (SEAMBHs), with $L_{\rm Bol}/L_{\rm Edd}\sim$ a few, have the following characteristics: 1) their H$\beta$ lags are significantly shorter than anticipated from the well-known BLR radius-luminosity relationship (e.g., \citealt{kaspi2000,bentz2013}), and the amount of shortening depends on the accretion rates (\citetalias{du2015,du2016a}), 2) their luminosities turn out to be saturated, in accordance with predictions of the slim disk model (\citetalias{du2015,du2016a}), 3) the \feii\ emission shows clear reverberation in response to the varying continuum with similar lags to those of H$\beta$ lines \citep[hereafter \citetalias{hu2015}]{barth2013,hu2015}. Obviously, it is necessary to investigate the details of the BLR kinematics in those extreme objects to see how they may differ from lower accretion rate AGNs. The velocity-delay maps of the BLRs for about 11 objects have been reconstructed in the past twenty years \citep{ulrich1996,bentz2010,pancoast2012,grier2013,pancoast2014b,grier2017}. However, they are mainly AGNs with ``normal'' accretion rates, $L_{\rm Bol}/L_{\rm Edd}\sim0.1$. Using the velocity-resolved time-lag analysis, we have probed the geometry and kinematics among the BLRs in 9 SEAMBH candidates monitored between 2012-2013 (hereafter SEAMBH2012) in \citetalias{du2016b}, identifying both disk-like and inflow/outflow signatures. In this paper, we use MEM to reconstruct velocity-delay maps for these SEAMBHs. With the high cadence and homogeneous sampling of the observations, we successfully recover velocity-delay maps of the H$\beta$ line for 6 objects and of the H$\gamma$ line for 3 objects. The paper is organized as follows. Section 2 briefly describes the reverberation mapping (RM) observations and data reduction. Section 3 presents the methodology of MEM and the details of application to RM data. Section 4 summarizes the results of the obtained velocity-delay maps and presents discussions on individual objects. The conclusion is given in Section 5. Unless stated elsewise, time lags are given in the rest frame. | We have reconstructed the velocity-delay maps for 6 objects in our SEAMBH2012 campaign. The BLR of Mrk 335 is consistent with an infalling dynamics, maintaining a similar structure as in the campaign of \cite{grier2012} two years ago. Mrk 486 also shows inflow signatures in its velocity-delay map. Mrk 1044 shows a complex pattern that looks like a combination of an outflow and a symmetric component. These two components may be generated from the self-shadowing effect of slim accretion disk in SEAMBH \citep{wang2014b}. The maps for Mrk 142 and MCG +06-26-012 also show features suggestive of an outflow, but for IRAS F12397+3333, the map is consistent with a disk-like geometry or a spherical shell. The velocity-delay maps obtained by MEM provide a basis for the model selection in dynamical modeling method in our follow-up works. | 18 | 8 | 1808.00705 |
1808 | 1808.07580_arXiv.txt | Type IIb supernovae (SNe) are important candidates to understand mechanisms that drive the stripping of stripped-envelope (SE) supernova (SN) progenitors. While binary interactions and their high incidence are generally cited to favor them as Type IIb SN progenitors, this idea has not been tested using models covering a broad parameter space. In this paper, we use non-rotating single- and binary-star models at solar and low metallicities spanning a wide parameter space in primary mass, mass ratio, orbital period, and mass transfer efficiencies. We find that our single- and binary-star models contribute to roughly equal, however small, numbers of Type IIb SNe at solar metallicity. Binaries only dominate as progenitors at low metallicity. We also find that our models can account for less than half the observationally inferred rate for Type IIb SNe at solar metallicity, with computed rates $\lesssim4$\% of core-collapse (CC) SNe. On the other hand, our models can account for the rates currently indicated by observations at low metallicity, with computed rates as high as $15$\% of CC SNe. However, this requires low mass transfer efficiencies ($\lesssim 0.1$) to prevent most progenitors from entering contact. We suggest that the stellar wind mass-loss rates at solar metallicity used in our models are too high. Lower mass-loss rates would widen the parameter space for binary Type IIb SNe at solar metallicity by allowing stars that initiate mass transfer earlier in their evolution to reach CC without getting fully stripped. | \label{s:intro} CC SNe are explosions marking deaths of stars with zero-age main sequence (ZAMS) masses $\gtrsim 8 M_\sun$ \citep[see e.g.][for a recent review]{2009ARA&A..47...63S}. Depending on the absence or presence of hydrogen lines in the supernova spectrum, SNe are classified into Type I or Type II, respectively. The absence of hydrogen features in a Type I CC SN spectrum is attributed to a progenitor star that lacks its outer hydrogen layers. Type IIb SN progenitors exhibit `mild' stripping of their outer layers, initially exhibiting prominent hydrogen spectral features that weaken and disappear in the weeks following explosion. Type I CC (also known as Type Ibc) and Type IIb SNe are therefore also referred to as SE SNe. The mechanisms that drive the stripping and regimes in which they dominate are still open questions. The leading candidates are close binary interactions \citep[e.g.,][]{1992ApJ...391..246P,2010ApJ...725..940Y,2017ApJ...840...10Y}, stellar winds \citep[e.g.,][]{1993ApJ...411..823W,2012A&A...542A..29G,2013A&A...558A.131G, 2017MNRAS.470L.102S}, stellar rotation \citep[e.g.,][]{2012A&A...542A..29G,2013A&A...550L...7G,2013A&A...558A.131G}, and nuclear burning instabilities \citep[e.g.,][]{2011ApJ...741...33A,2015ApJ...811..117S}. Binary interactions were initially the favored channel to strip stars due to the high observed binarity of Wolf-Rayet (WR) stars \citep{1973IAUS...49..205K}. However, with spectroscopic UV observations indicating strong stellar winds, sufficient to strip stars \citep{1978A&A....63..103C}, they became the preferred channel. The trend is now appearing to be reversing, with binary interactions gaining traction as the preferred formation channel. This is due to a variety of pieces of evidence. First, clumping in stellar winds suggest that currently used mass-loss rates are too high; hot wind mass-loss rates are lower by a factor of $2 - 3$ than those typically used in stellar evolution calculations (\citealt{2014ARA&A..52..487S}; but also see \citealt{2012ApJ...751L..34V}). Second, recent observations indicate that massive stars are predominantly part of close binary systems \citep{2012Sci...337..444S,2014ApJS..213...34K,2017ApJS..230...15M}. Third, observed SE SN rates are too high to be explained solely by single-star evolution \citep{2011MNRAS.412.1522S}. Fourth, pre-SN masses estimated from light curves indicate lower ZAMS mass for SE SN progenitors than those that produce WR stars via the single-star channel \citep{2016MNRAS.457..328L,2019MNRAS.485.1559P}. Finally, in many cases circumstellar medium densities for SE SNe inferred from X-ray/radio observations indicate pre-SN mass-loss rates higher than WR winds \citep{2012ApJ...752...17W,2016ApJ...821...57D}, implying additional mass loss processes being active close to CC. Type IIb SNe are of particular interest in understanding formation channels to SE SNe because of a few reasons. First, they are the only class within the group that has several (five) identified progenitors\footnote{The progenitor of Type Ib SN iPTF13bvn was identified by \cite{2013ApJ...775L...7C} and confirmed by its disappearance by \cite{2016ApJ...825L..22F}. There is also a candidate for the progenitor of Type Ic SN 2017ein \citep{2018arXiv180301050V, 2018MNRAS.480.2072K}.}. % Second, there is evidence for the presence of a binary companion to the progenitor in three cases \citep{2014ApJ...790...17F,2014ApJ...793L..22F,2018ApJ...856...83R}. Finally, Type IIb SNe are quite abundant, accounting for $10-12\%$ of all CCSNe and $30-40$\% of all SE SNe \citep{2011MNRAS.412.1473L, 2011MNRAS.412.1522S,2017PASP..129e4201S}. SN 1993J is the prototypical Type IIb SN. A progenitor candidate was identified in ground-based pre-explosion images by \citet{1994AJ....107..662A}. They also found that its SED had a blue component which they attributed to either an OB association or a binary companion. Late-time observations of the region have provided stronger evidence for the presence of a putative companion star \citep{2004Natur.427..129M, 2014ApJ...790...17F}. SN 1993J is the best-studied Type IIb SN to date, in part due to its proximity, being the subject of several observational and theoretical investigations. Since 1993, putative progenitors of four more Type IIb SNe have been identified in pre-explosion images: SN 2008ax \citep{2008MNRAS.391L...5C,2015ApJ...811..147F}, SN 2011dh \citep{2011ApJ...739L..37M,2011ApJ...741L..28V,2013ApJ...762...74B}, SN 2013df \citep{2014AJ....147...37V,2015ApJ...807...35M}, and SN 2016gkg \citep{2017MNRAS.465.4650K,2017ApJ...836L..12T,2018Natur.554..497B}. There is also evidence for binary companions to the progenitors of SN 1993J \citep{2014ApJ...790...17F}, SN 2001ig \citep{2018ApJ...856...83R}, and SN 2011dh \citep{2014ApJ...793L..22F}. The Galactic supernova remnant, Cassiopeia A, is known to be the result of a Type IIb SN from light echo spectra \citep{2008Sci...320.1195K,2011ApJ...732....3R}. However, there is no companion even at deep limits \citep{2017arXiv171100055K,2018MNRAS.473.1633K}. With the discovery of more events, there were efforts, both observational and theoretical, to understand the population of Type IIb SNe. On the observational side, \cite{2010ApJ...711L..40C} studied a sample of Type IIb SNe and suggested that they can be further classified into two sub-types based on their radio SN shock velocities; compact and extended Type IIb SNe progenitors exhibit high and low velocities, respectively. However, this suggestion was quickly challenged by the discovery of SN 2011dh exhibiting both rapidly expanding radio shells \citep{2012ApJ...752...78S} and an extended yellow supergiant progenitor \citep{2014ApJ...793L..22F}. % \citet{2014ApJ...792....7F} suggested that some Type Ib/c SNe were spectroscopically more similar to Type IIb SNe. More recently, \citet{2016ApJ...827...90L} found a continuum in the signatures of Type IIb and Ib SN spectra. Most early theoretical investigations into progenitors of and evolutionary pathways to Type IIb focussed on SN 1993J \citep[e.g.,][]{1993Natur.364..509P, 1994ApJ...429..300W, 2004Natur.427..129M, 2009MNRAS.396.1699S}. \citet{2010ApJ...725..940Y} and \citet{2011MNRAS.414.2985D} studied progenitors of Type Ib/c SNe arising as a result of mass transfer in close binary systems and found that some of their Type Ib SN progenitors % exploded with small amounts of residual hydrogen. They suggested that these progenitors may be classified as Type IIb SNe if detected soon after explosion. \citet{2012A&A...538L...8G} and \cite{2013A&A...550L...7G} used solar-metallicity single-star, non-rotating and rotating models and identified their $20-25M_\sun$ rotating models as potential Type IIb SN progenitors. % \citet{2011A&A...528A.131C} performed the first parameter space search for single and binary progenitors of Type IIb SNe using detailed evolutionary calculations. However, they restricted their binary parameter space to initial primary masses $15 M_\sun$, initial secondary masses $10 M_\sun-15 M_\sun$, initial orbital periods $800 - 2100$ days, and solar metallicity. While they were able to analyze various evolutionary pathways to Type IIb SNe and their outcomes, they were unable to compute robust relative rates because of their limited parameter space coverage. As we show in this paper, the parameter space for binary Type IIb SNe varies significantly with initial primary mass. Another limitation is that they restrict their analysis to progenitors that explode with $0.1-0.5 M_\sun$ of residual hydrogen envelope. This excludes the group of more compact Type IIb SN progenitors suggested from analyses of Type IIb SN light curves \citep{2014MNRAS.445.1647M,2017ApJ...836L..12T,2018Natur.554..497B} and detailed non-LTE radiation hydrodynamical calculations \citep{2018A&A...612A..61D}. % Recently, \citet{2017ApJ...840...10Y} undertook a wide parameter space search for binary Type IIb and Ib SN progenitors using detailed evolutionary calculations, varying the initial primary-star mass from $10 - 18 M_\sun$, initial orbital period from $10 - 3000$ days, but keeping the initial mass ratio (= 0.9) and mass transfer efficiency (= 0.2) constant. They analyze Type IIb SN progenitors at two different metallicities, $Z =0.007$ and solar, and show that the parameter space for Type IIb SNe broadens significantly when lowering metallicity and that its effect is roughly analogous to lowering the wind mass-loss rate. However, their parameter space coverage precludes them from computing case C Type IIb SNe (mass transfer after core helium exhaustion), Type IIb SN relative rates, and statistical progenitor properties. Motivated by these gaps in theoretical analyses, in a two paper study, we investigate the progenitors (their evolutionary pathways, rates, and properties) of Type IIb SNe (henceforth referred to as SNe IIb) using a comprehensive parameter space coverage and statistical analysis. Our study yields a comprehensive database of full evolutionary histories of non-rotating single- and binary-star progenitors of SNe IIb at solar and sub-solar metallicities. This paper is dedicated to investigating the parameter space and evolutionary pathways to SNe IIb. In addition, our parameter space coverage allows us to compute theoretical SN IIb rates. The second paper will be dedicated to investigating the observable properties of SN IIb progenitors presented here. This paper is organized as follows. In Section \ref{s:models}, we provide a detailed description of our models. In Section \ref{s:channels}, we discuss key evolutionary channels, and governing physics, towards SNe IIb. In Section \ref{s:pspace}, we describe the parameter space for SN IIb progenitors at solar and sub-solar metallicities. In Section \ref{s:rates}, we provide theoretical rates for SNe IIb as a fraction of CC SNe. We discuss our results and conclude in Section \ref{s:conclusions}. We present numerical tests in the Appendix. | \label{s:conclusions} \subsection{Summary of Results} In this paper, we use non-rotating single- and binary-star models at solar and low metallicities covering a broad parameter space (see Table \ref{t:pspace}) to investigate evolutionary channels and parameter space for progenitors of SNe IIb. We find that metallicity and mass transfer efficiency play important roles in shaping the parameter space towards SNe IIb. Both of these effect the evolution of case EB binaries (binaries that initiate MT on the HG), where lower metallicities and mass transfer efficiencies are favorable for the production of SNe IIb. % Since the range of binary configurations that lead to case EB mass transfer is very wide, these binaries have a strong effect in shaping the parameter space, and therefore rates for SNe IIb. We use our models to compute theoretical SN IIb rates and compare them to observationally constrained values. We attempt to account for the uncertainties in the prior distributions of single- and binary-star birth properties and criteria used to define SN IIb progenitor structures by spanning possible values for them. This allows us to investigate to what extent our models are a priori able to explain observed SN IIb rates. We find that solar-metallicity single- and binary-star models, by themselves or together, account for less than half the observed SN IIb rate at solar metallicity. Our most optimistic rate is inconsistent within $2\sigma$ of the observationally inferred rates. Also, at high metallicity, singles and binaries contribute to roughly equal numbers (singles slightly more) of SNe IIb. At low metallicity, however, our models can account for the implied rate. Here, binaries are the sole producers of SNe IIb. The group of case EB binary SNe IIb (which in turn is courtesy of lower wind mass-loss rates at low metallicity) is responsible for this change. Further, even using our most restrictive criterion on a SN IIb progenitor structure, theoretical rates at low metallicity are still consistent with observed rates within the reported uncertainities (though these are quite large). If wind mass-loss rates are lower than those used in our models (which may be the case as discussed in Section \ref{ss:physics}), it might be instructive to compare observed rates at high metallicity to our low-metallicity models \citep{2017ApJ...840...10Y}. We find that rates for our low-metallicity and low mass transfer efficiency models are consistent with observed rates at high metallicity. This is true even after using our most restrictive definition of a SN IIb progenitor. However, this requires $f_{\rm bin} > 0.5$ and priors favoring equal mass binaries and favoring a flat orbital period distribution to be consistent within reported uncertainties for observationally inferred rates. We discuss implications of this in Section \ref{ss:takeaways}. \subsection{Implications of Assumed Physics} \label{ss:physics} Although we have tried to be conservative in our predicted upper and lower limits for rates, it is instructive to point out additional physical uncertainties in our models. Wind mass-loss rates have an important effect on single- and binary-star channels towards SNe IIb. % Stellar winds in our \mesa\ models are computed using different wind mass-loss prescriptions at various evolutionary phases. WR mass-loss rates have the strongest impact on evolutionary outcomes. However, these rates are also known to be quite uncertain, particularly for low mass WR stars. Most estimates, including those used in this work are expected to be overestimated due to clumping in the winds \citep{2014ARA&A..52..487S}. Lower wind mass-loss rates would expand the parameter space for binary SNe IIb \citep{2017ApJ...840...10Y,2019MNRAS.486.4451G}. In addition, though it is commonly assumed that mass-loss rates from red supergiants decrease with metallicity, some studies suggest that they are more or less metallicity independent \citep{2000A&A...354..125V, 2005A&A...438..273V, 2017MNRAS.465..403G}. If this is the case, we would expect low-metallicity single stars to produce SN IIb at a relative rate similar to that of high-metallicity single stars. Another important point is that we have not included stellar rotation. Rotation is thought to play an important role in massive star evolution, with very rapidly rotating stars even being predicted to evolve chemically homogeneously due to rotationally enhanced mixing \citep{1987A&A...178..159M}. Rotationally enhanced mass-loss rates can also reduce the minimum mass required for a single star to remove its hydrogen envelope \citep{2003A&A...404..975M}, thus increasing the rate of SN IIb progenitors from single stars. However, \citet{2013A&A...550L...7G} found that the initial mass range that produces SNe IIb does not change much for their rotating models with $v_{rot} \sim 250$ km s$^{-1}$. Further, observations show that the bulk of the massive star population rotates at lower velocities \citep{2013A&A...560A..29R,2015A&A...580A..92R}, and that stars rotating at faster rates can potentially be accounted for in terms of spun-up secondaries in post-mass transfer systems \citep{2013ApJ...764..166D}. Properly assessing the impact of rotation, along with the effect of tides and eccentricity (which is likely small due to the wide orbits considered here) would require future calculations that include these effects, but they are beyond the scope of the present study. \subsection{Key Takeaways and Future Directions} \label{ss:takeaways} Our results have the following implications for progenitor channels to SNe IIb. In order to account for the observationally inferred SN IIb rates, \begin{enumerate} \item solar-metallicity wind mass-loss rates need to be lower than those used in our models: lower wind mass-loss rates increase the contribution of binaries to SNe IIb, by populating the parameter space available via case EB mass transfer, while driving down the contribution from single stars. \item low mass transfer efficiencies are needed: this allows case EB accretors to thermally adjust to the transferred mass and avoid contact. \end{enumerate} We note that it is likely that even after accounting for robust mass-loss rates, the uncertainty in prior distributions of stellar birth properties and SN IIb definitions will still pose significant challenges for comparing theoretical models to observations. Therefore, in order to address the question of SNe IIb progrenitors we need four pieces of information: (1) accurate determinations of mass-loss rates for stripped stars as a function of metallicity, (2) constraints on structural properties of SN IIb progenitor stars, (3) robust distributions for single- and binary-star birth properties, and (4) on the observational side, SN IIb rates as a function of metallicity. An exploration of case A mass transfer towards SNe IIb at low metallicity and merger channels would also be worthwhile. In this paper we focussed on understanding the evolutionary channels and parameter space of SN IIb progenitors. In a follow-up paper we will focus on using our models to draw statistical comparisons to all three independent observational probes into SN IIb progenitors: (1) direct progenitor detections in archival images, (2) progenitor structure and explosion parameters from analyzing multi-band light-curves, and (3) circumstellar medium properties from X-ray/radio observations. The aim would be to confirm that our models are able to reproduce the full range of observational constraints for SNe IIb, that the results of such comparison is consistent with conclusions in this paper, and finally, to identify regions in the observational parameter space that are implied from our models but remain unprobed by observations to help guide observing strategies. The advent of high-cadence surveys like the ZTF \citep{2014htu..conf...27B, 2017NatAs...1E..71B}, ASSA-SN \citep{2017PASP..129j4502K}, DLT40 \citep{2017ApJ...848L..24V}, KAIT \citep{2001ASPC..246..121F}, and, at the turn of the decade, LSST \citep{2002SPIE.4836...10T,2008arXiv0805.2366I} have created the need for a comprehensive database of theoretical models to provide reliable progenitor characterization and in the case of binary progenitors, their companions. Conversely, the wealth of data that will be available will provide unprecedented statistics on progenitor properties, along with their dependence on properties such as host galaxy metallicity. Efforts such as the one undertaken here are therefore imperative in order to understand the physics of massive star evolution. | 18 | 8 | 1808.07580 |
1808 | 1808.09554_arXiv.txt | A massive planet in a protoplanetary disc will open a gap in the disc material which acts as a transition between Type I and Type II planetary migration. Type II migration is slower than Type I migration, however it is still desirable to slow down Type II migration to allow gas giant planets with semi-major axis in the range $5-10$AU to exist, similarly to our Solar system. We investigate a method of slowing down and reversing Type II migration by heating the outer gap edge due to incident radiation from the central star. Using an approximate vertically averaged heating method we find that Type II migration can be slowed or in extreme cases reversed if we assume near maximum allowed irradiation from the central star. Therefore, we believe this is a very interesting phenomenon that should be investigated in greater detail using three dimensional hydrodynamic and radiative transfer simulations. | Planets form in protoplanetary discs by accretion of gas and dust particles and will excite density waves in the disc material that transport angular momentum away from the planet. This angular momentum is deposited in the disc, exerting a torque on the disc material. If the angular momentum is deposited close to the planet \citep{Lin&Papaloizou1993,Bryden1999} and the resultant torque is stronger than the viscous diffusion torque of the disc material \citep{Lin&Papaloizou1979,Goldreich&Tremaine1980,Takeuchi1996,Crida2006} then a gap is formed around the location of the planet. These two conditions are known as the thermal and viscous criteria for gap opening. It has also been shown that small mass planets can open gaps at a greater distance from the planet, where the excited density wave shocks \citep{Goodman&Rafikov2001, Rafikov2002}. The interaction between the planet and the disc plays an important role in the evolution of the planet's orbital radius. There are two main regimes of planetary migration, Type I and Type II, which are differentiated between by the mass of the planet and as a result, the presence of a gap in the disc material. Type I migration occurs for planets that are not massive enough to open a gap in the disc, and migrate due to the Lindblad and corotation torques acting upon them \citep{Goldreich&Tremaine1980}. The corotation torque depends on the magnitude of the thermal and viscous diffusion in the disc \citep{Masset2001,Paardekooper2010,Paardekooper2011} and these can alter the speed of migration. Type II migration occurs when a planet is massive enough to open a gap in the disc. The low density area surrounding the planet slows down the Type I migration, and the planet now migrates with the gap on the viscous evolution timescale of the disc \citep{LinPapaloizou1986}. Hence Type I migration is considerably faster than Type II migration \citep{Ward1997}. Despite this reduction in migration rate, it has been found that the predicted timescale for Type II migration is shorter than the lifetime of the disc, suggesting that most planets undergoing Type II migration will migrate into and be absorbed by the central star \citep{Hasegawa&Ida2013}. Observational evidence dictates that gas giant planets are more common at orbital distances $R > 1 \textrm{AU}$ \citep{Mayor2011, Cassan2012,Fressin2013,Santerne2016}, which would not be the case if this predicted timescale was correct. Current models estimate that for giant planets to survive migration they must form at distances $R>20\textrm{AU}$ \citep{Coleman&Nelson2014}. Hence it is implied that something must be limiting the migration speed of a planet in the Type II migration regime \citep{Nelson2000}. The argument for classical Type II migration relies on the absence of material surrounding the planet, a result of the torque exerted on the disc by the massive planet clearing the area around it and forming a gap. This is clearly the case in one dimensional simulations, such as those in \cite{LinPapaloizou1986}, however extending this to higher dimensions this is found not to be the case \citep{Kanagawa2015,Hallam&Paardekooper2017}. It can easily be seen from two or three dimensional simulations that disc material can be present within the gap, and can cross the gap on horseshoe orbits. Hence we can only view classical Type II migration as an idealised case. The result is a deviation from the classical Type II migration rate, the viscous evolution timescale of the disc. The migration rate can now be significantly faster or slower than the viscous evolution timescale of the disc and is dependent on the planet and disc parameters \citep{Duffel2014}. \cite{Durmann&Kley2015} found that only for small disc masses $M_d/M_J < 0.2$ is the migration rate slower than the classical rate. Therefore, we investigate a new method of slowing the rate of planet migration in the Type II regime. We propose that radiation from the central star is incident on the outer edge of gap formed by the giant planet, visible as a result of the flaring of the disc. This incident radiation heats the outer disc, increasing the scale height in this region. This process is outlined in Figure \ref{Fig:Diagram}. As the one-sided torques scale with the aspect ratio as $h^{-3}$, the result of this may be a lowered contribution to the net torque on the planet from the planet's outer wake, which in turn causes the net torque on the planet to become more positive. As the rate of migration is proportional to the magnitude of the torque on the migrating planet and the direction given by the sign of the torque, this process could slow the planet's migration. Observational signatures of irradiated gap edges have been studied by \cite{Jang-Condell2013}. \begin{figure} \includegraphics[width=\columnwidth]{Fig1.png} \vspace*{-5mm} \caption{A diagram showing the process by which the outer gap edge is heated by the star. This is the method in which the torque contribution from the outer disc may be reduced and hence the net torque on the planet may become more positive.} \label{Fig:Diagram} \end{figure} This paper is arranged as follows. In Section \ref{Sec:Basic_Eq} we derive the relevant equations solved to simulate disc evolution. In Section \ref{Sec:NumSet} we discuss the code used and the numerical setup for our simulations. In Section \ref{Sec:Just} we justify the choice of parameters used in our simulations. In Section \ref{Sec:Inst} we discuss the onset of instabilities in our simulations and how we go about avoiding them. In Section \ref{Sec:Res} we present our results of heating the gap edge and the resultant dependence of the net torque on previously justified parameters. In Section \ref{Sec:Disc} we discuss our results in the context of prior work while justifying assumptions made and postulating any impact they may have. Finally, in Section \ref{Sec:Conc} we present our conclusions. | \label{Sec:Conc} We have investigated the heating of the outer gap edge by stellar radiation using a Gaussian distribution in the sound speed profile of a disc. Using this method it is possible to lower the contribution to the net torque on the planet from the planets outer wake, while leaving the contribution from the inner wake mostly unchanged. The goal was to address the problem that observational evidence dictates that giant planets are more common at orbits $R>1\textrm{AU}$ \citep{Mayor2011,Cassan2012,Fressin2013,Santerne2016} and that for a gas giant planet to survive throughout the discs lifetime it must begin Type II migration at an orbital radius $R>20\textrm{AU}$ \citep{Coleman&Nelson2014}. At this radius the core accretion timescale exceeds the lifetime of the disc, therefore there must be a process that limits Type II migration speed \citep{Nelson2000}. We found heating the gap edge via stellar radiation is a method of reducing the net torque on the planet and by extension the rate of inwards migration. We also found that for the extreme case of maximum to near maximum gap edge irradiation it is possible to return a positive net torque on the planet, which corresponds to outwards migration for the planet. We have investigated the impact on the net torque of modifying the following Gaussian parameters: \begin{itemize} \item $A$, the amplitude of the Gaussian. \item $\mathcal{W}$, the full width at half maximum of the Gaussian. \item $R_G$, the radial location of the Gaussian peak. \item The impact of asymetric Gaussians. \end{itemize} We found that the range of applicable Gaussians was limited by $\mathcal{W}$ constraints to avoid Rossby and Rayleigh unstable regimes. Our test case of $R_G = 1.20$ was studied extensively, and was found to provide very low magnitude positive torque for a small range of $A$ and $\mathcal{W}$, providing the potential for weak outward migration, or a large reduction in the rate of inwards migration. At larger radii, $R_G = 1.30 -1.40$, we found that there is a significantly larger parameter space for which outward migration is possible, both in $A$ and $\mathcal{W}$ and that the magnitude of the net torque is larger, so the planet is further from the edge of the outwards migration regime and has a higher rate of migration. We found that in general the height of the Gaussian must be at least $A = 1.75$ to achieve outwards migration and that increasing the $\mathcal{W}$ of the Gaussian lowers the net torque on the planet, a result of the Gaussian impacting the torque contribution from both the inner and outer wakes, rather than just the outer wake. We find that for a number of sample stellar classifications the predicted magnitude of heating at the gap edge corresponds to $A \approx 2.0$ at maximum or near maximum stellar irradiation with a weak radial dependence, $A \propto R_G^{1/8}$, which would make outwards migration a possibility. For more modest $A$, which could account for the presence of an inner disc, re-radiation heating the inner gap edge and self-shadowing of the outer disc, we still find significant effects on the net torque on the planet, meaning that even if outwards migration is not a possibility there is still the potential for reduction in Type II migration rate. We also investigate asymmetric Gaussian distributions, in which $\mathcal{W}$ is greater on the outer Gaussian edge, and at the limit of stability on the inner Gaussian edge. We found that increasing $\mathcal{W_\textrm{out}}$ can both increase the net torque on the planet at smaller radii ($R_G = 1.20$) and decrease the net torque at larger radii, such as the reduction at $R_G = 1.40$. This could potentially solve the problem of the small parameter space for outwards migration at $R_G = 1.20$, while the weak reduction in net torque for ranges $R_G = 1.30 - 1.35$ still leaves a large parameter space for outwards migration there. Overall we believe our results show that there is a significant possibility for the irradiation of the gap edge to severely slow the rate of Type II planetary migration. The results from our simplified approach can act as a good starting point for future three dimensional radiation-hydrodynamical simulations, which will reveal how important this effect is for allowing gas giant planets to remain at larger orbital radii. | 18 | 8 | 1808.09554 |
1808 | 1808.09762_arXiv.txt | {The Large Area Telescope aboard the \textit{Fermi} spacecraft has detected more than 200 $\gamma$-ray pulsars since its launch in 2008. By concurrently fitting standard geometric model light curves onto \textit{Fermi} and radio data, researchers have constrained the inclination and observer angles of a number of pulsars. At first this was done by comparing observed and modelled light curves by eye, and later via statistical approaches. We fit modelled light curves of 16 pulsars to radio and $\gamma$-ray data by optimising a custom test statistic that we have developed for combining light curves across the two wavebands, taking their disparate errors into account. We present geometrical constraints found using this process, and compare them with results found by eye or using other statistical methods.} \FullConference{5th Annual Conference on High Energy Astrophysics in Southern Africa\\ 4-6 October, 2017\\ University of the Witwatersrand (Wits), South Africa} \begin{document} | Pulsar observations have traditionally been dominated by data from the radio wavelength range, with sources in this band totalling 2 627 at the time of writing, according to the ATNF Pulsar Catalogue \cite{manchester2005}. This is in comparison with the fewer than ten sources that had been detected in the $\gamma$-ray domain prior to 2008 \cite{thompson2001}. The launch of the Large Area Telescope (LAT) aboard the \textit{Fermi} satellite that year marked the deployment of the first instrument sensitive enough to detect a significant number or $\gamma$-ray photons from pulsars. To date, it has discovered more than 205 new $\gamma$-ray pulsars at high sensitivity and resolution, ushering in what's been called a $\gamma$-ray pulsar revolution \cite{caraveo2014}. Multi-wavelength pulsar studies can now be done that aid our understanding of pulsar emission in several wavelength domains. In particular, using some goodness-of-fit test to compare observed and modelled pulsar light curves (LCs) in the radio and $\gamma$-ray bands, one can infer a pulsar's most likely physical configuration, in terms of its tilt angle ($\alpha$) and the observer angle ($\zeta$) (see e.g. \cite{venter2009}). Unfortunately, the available $\gamma$-ray and radio geometric models often produce very different $\alpha$ and $\zeta$ constraints when fit separately to observations. Various attempts have been made to find "compromise" ($\alpha,\zeta$) constraints, i.e.\ pairs of these parameters that lead to model LCs that adequately match observations in both bands. This endeavour has, however, been complicated by the disparity between the errors characterising the available radio and $\gamma$-ray pulsar observations. Due to this error discrepancy, any goodness-of-fit test that is dependent on the observational errors, such as Pearson's $\chi^2$ test, will deliver very different values of the test statistic in the radio and $\gamma$-ray wavebands. Adding the test statistics obtained for individual fits and then minimising the result typically leads to a good fit in one waveband and a relatively poor fit in the other. Previous studies have applied Pearson's $\chi^2$ goodness-of-fit test to compare modelled LCs to radio and $\gamma$-ray observations of millisecond pulsars (MSPs) \cite{pierbattista2015} and canonical pulsars \cite{johnson2014}. These studies attempted to circumvent the error-disparity problem by artificially inflating the relative uncertainties on the observed radio data to match those of the $\gamma$-ray observations before addition of the respective test statistics of the radio-only and $\gamma$-only LC fits. This produced LC fits in both wavebands that are invariably qualitatively better than those obtained without taking error disparity into account. Compromising the radio data is, however, not optimal for achieving the best possible LC fits, and deliver formally bad fits, the minimum value of the test statistic being orders of magnitude larger than the number of degrees of freedom. Furthermore, it has been suggested to scale the dynamic range of Pearson's $\chi^2$ test statistic to match in the radio and $\gamma$-ray bands before addition \cite{seyffert2016}. This method was recently applied to a sample of 11 MSPs and canonical pulsars, with mixed results \cite{bezuidenhout2017}. Recently, a bespoke test statistic called the Scaled-Flux Normalised $\chi^2$ test statistic has been put forward as an alternative to Pearson's $\chi^2$ test (Seyffert et al., 2018; in preparation). This test statistic partially negates the influence of significantly differing data errors, making it comparable between wavebands without the need for error inflation, scaling, or any other manipulation. In this work we apply the proposed test statistic to a sample of 16 pulsars previously fit by either of the previous studies (\cite{pierbattista2015},\cite{johnson2014}). We then compare the best-fit $(\alpha,\zeta)$\footnote{Note that pairs of constraints on $\alpha$ and $\zeta$ are reported throughout in this order and as measured in degrees.} pairs we find using the SFN test statistic to those found by the previous studies, as well as those obtained using by-eye comparison of modelled and observed LCs. For more details see Bezuidenhout et al. (2018; in preparation). | Our aim was to test the utility of the novel SFN test statistic put forward by Seyffert et al. (2018; in preparation) as it applies to dualband pulsar LC fitting. We believe, based on the correlations laid out in Table \ref{tab:correlations}, that the use of this test statistic results in LC fits that are more aligned to what one typically finds through by-eye fitting than is the case with fits found through manipulation of Pearson's $\chi^2$ test statistic. As such, fitting using the SFN test statistic may be useful for any further pulsar LC modelling efforts. Future work may focus on applying this fitting method to LCs modelled using a variety of different radio and $\gamma$-ray models, including the new dissipative $\gamma$-ray models (e.g. \cite{kalap2014}), in order to assess the merit of these models. This fitting procedure may also be extended to incorporate different wavebands such as x-rays, or even to multiband fitting of pulsars' spectra. | 18 | 8 | 1808.09762 |
1808 | 1808.08144_arXiv.txt | {The space density of the various classes of cataclysmic variables (CVs) could only be weakly constrained in the past. Reasons were the small number of objects in complete X-ray flux-limited samples and the difficulty to derive precise distances to CVs. The former limitation still exists. Here the impact of Gaia parallaxes and implied distances on the space density of X-ray selected complete, flux-limited samples is studied. The samples are described in the literature, those of non-magnetic CVs are based on ROSAT (RBS -- ROSAT Bright Survey \& NEP -- North Ecliptic Pole), that of the Intermediate Polars stems from Swift/BAT. All CVs appear to be rarer than previously thought, although the new values are all within the errors of past studies. Upper limits at 90\% confidence for the space densities of non-magnetic CVs are $\rho_{\rm RBS} < 1.1 \times 10^{-6}$\,\ppc\, and $\rho_{\rm RBS+NEP} < 5.1 \times 10^{-6}$\,\ppc\, for an assumed scale height of $h=260$\,pc and $\rho_{\rm IPs} < 1.3 \times 10^{-7}$\,\ppc\, for the long-period Intermediate Polars at a scale height of 120\,pc. Most of the distances to the IPs were under-estimated in the past. The upper limits to the space densities are only valid in the case where CVs do not have lower X-ray luminosities than the lowest-luminosity member of the sample. These results need consolidation by larger sample sizes, soon to be established through sensitive X-ray all-sky surveys to be performed with eROSITA on the Spektrum-X-Gamma mission. } | The space density of cataclysmic variable stars (CVs) is one of the bigger unknowns in the field. Being the outcome of binary star evolution through a common envelope and subsequent angular momentum loss this number is important for several parameters and processes that are relevant in binary evolution: the initial binary separation, the initial mass distribution, the common envelope efficiency, the angular momentum loss in the post-common envelope phase and that in the CV stage. It is the most relevant number to compare with binary population synthesis. The simple question is: How many are out there? The question is relevant for stellar evolution but has implications for models of the total energy output of the Milky Way. One of the contenders to explain the infamous Galactic Ridge X-ray Emission~\citep[GRXE,][]{worrall+82} are CVs and in particular magnetic CVs of the Intermediate Polar type (IPs). Through a deep Chandra pointing close to the galactic centre the apparently diffuse GRXE was largely resolved into point sources ~\citep{revnivtsev+09}. The composition however remained uncertain and was discussed thoroughly in the recent past and depends sensitively on the space density and the luminosity functions of the main source classes~\citep[e.g.][]{warwick14,nobukawa+16}. The second data release from the Gaia satellite opens the opportunity to re-assess the space density of X-ray selected cataclysmic variables. Past studies were hampered by small sample sizes and imprecisely determined distances. While the former limitation cannot be overcome presently the latter has essentially vanished. In this paper the non-magnetic CVs (dwarf nova and nova-like systems) and one class of the magnetic CVs, the IPs, are addressed. Both sub-classes have relatively well-understood X-ray spectra which results in a relatively clean selection of objects. The strongly magnetic CVs, the polars, will not be addressed here. Polars are special due to their more complex X-ray spectra with a thermal and a potentially strong soft component. ROSAT has uncovered many soft polars~\citep{beuermann+schwope94}, but all new discoveries made with XMM-Newton lack the pronounced soft component \citep[see e.g.][ and references therein]{webb+18}. It is therefore questionable if the observed sample of polars may be regarded as representative of the parent population. \begin{table*} \caption{Non-magnetic CVs found in the RBS. \#713 is no longer considered non-magnetic, whereas \#664 was previously classified as an IP. The luminosity is given in the ROSAT spectral band 0.1--2.4\,keV. Distances and luminosities from ~\citet{schwope+02} are listed in columns with tilde, all other values are from this work. $V_{\rm gen}$ was computed for a scale height of 200\,pc.} \label{t:nmcvs} \begin{tabular}{rlccrrrrrrr} RBS\# &Name & $\widetilde{D}$ & $\widetilde{\log L_X}$ & $f_{\rm X}$ & $p$ & $\Delta p$ & $\log L_X$ & $D$ & $D_{\rm max}$ & $V_{\rm gen}$ \\ & & [pc]& [s$^{-1}$] & [erg cm$^{-2}$ s$^{-1}$] & [mas] & [mas] & [s$^{-1}$] & [pc]& [pc]& [pc$^{3}$] \\ \hline 2 &EF Tuc & 500:& 32.0& 3.38E-12& 0.720& 0.023& 32.9& 1335& 1462& 2.09E+08\\ 22 &WW Cet & 100 & 31.0& 9.16E-12& 4.588& 0.047& 31.7& 216& 389& 3.28E+07\\ 280 &TT ARI & 135 & 31.1& 5.96E-12& 3.884& 0.070& 31.7& 256& 372& 4.20E+07\\ 288 &WX HYI & 265 & 31.6& 6.26E-12& 4.273& 0.029& 31.6& 232& 346& 3.21E+07\\ 372 &IQ Eri & 130 & 30.8& 3.76E-12& 5.186& 0.167& 31.2& 192& 222& 1.11E+07\\ 490 & & 33 & 29.6& 3.36E-12& 3.098& 0.126& 31.6& 320& 349& 3.51E+07\\ 512 &VW Hyi & 65 & 30.9& 1.59E-11& 18.531& 0.022& 30.7& 54& 128& 3.26E+06\\ 664 &TW Pic & -- & -- & 3.13E-12& 2.284& 0.022& 31.8& 432& 455& 7.99E+07\\ 694 &SU UMa & 280 & 32.1& 1.41E-11& 4.535& 0.029& 31.9& 219& 490& 9.25E+07\\ 710 &SW UMa & 140 & 30.9& 3.81E-12& 6.148& 0.080& 31.1& 162& 188& 9.09E+06\\ (713 &EI UMa) & -- & --& & & & & &&\\ 728 &BZ UMa & 110 & 30.9& 6.12E-12& 6.557& 0.064& 31.2& 152& 224& 1.39E+07\\ 1008 &T Leo & 100 & 31.1& 9.88E-12& 7.814& 0.069& 31.3& 128& 240& 1.33E+07\\ 1411 &RHS40 & 460 & 32.0& 3.47E-12& 2.919& 0.671& 31.7& 367& 407& 3.72E+07\\ 1900 &TY PsA & 300:& 31.7& 4.68E-12& 5.431& 0.065& 31.3& 183& 236& 1.26E+07\\ 1955 &V405 Peg & 30 & 29.5& 3.05E-12& 5.784& 0.062& 31.0& 172& 179& 8.06E+06\\ 1969 &CC Scl & 165 & 31.1& 4.23E-12& 4.904& 0.148& 31.3& 203& 249& 1.37E+07\\ \hline \end{tabular} \end{table*} | \label{s:results+discussion} The space densities of X-ray selected samples of magnetic and non-magnetic CVs was redetermined using recently published parallaxes and distances from Gaia-DR2. The results are summarized in Tab.~\ref{t:dens} and Figs.~\ref{f:lx_wei} and \ref{f:lxdist}. Fig.~\ref{f:lxdist} shows distributions of the original and the revised luminosities found for the RASS-CVs and the IPs, while Fig.~\ref{f:lx_wei} shows the weight per object (inverse of the generic volume) over its luminosity. Just as a reference, the published values of $\rho$ for the RBS, the RASS and the IP samples were $\sim1.5\times10^{-6}$\,\ppc, $4_{-2}^{+6} \times 10^{-6}$\,\ppc, and $1_{-0.5}^{+1} \times 10^{-6}$\,\ppc, for scale heights of 200\,pc, 260\,pc, and 120\,pc, respectively {\bf \citep{schwope+02,pretorius+knigge12,pretorius+mukai14}}. The comparison with the values listed in Tab.~\ref{t:dens} shows that all newly derived densities are smaller than published ones but still within the published errors. The statistical errors of $\rho$ of the non-magnetic CVs could be reduced very significantly thanks to precise Gaia data. The statistical error of $\rho$ for the IP sample is still large due to the shallow flux limit. For the RBS-CVs the space density is safely below $1.8 \times 10^{-6}$\,\ppc\ at 90\% confidence but could be smaller than $1.1 \times 10^{-6}$\,\ppc\ if the scale height would be as high as 260\,pc as assumed by \cite{pretorius+knigge12}. This result remains unchanged if the fluxes in the band $0.5-2.0$\,keV are used (third row in Tab.~\ref{t:dens}). The RBS-sample consists of long- and short period CVs. It thus appears possible that not all objects belong to the same galactic population. The use of just one scale to characterise the sample is likely an oversimplification. The RASS-CVs (RBS$+$NEP) are compatible with a significantly higher space density thanks to the lower flux limit of the NEP. The inclusion of just 4 CVs from the NEP-survey implies a space density a factor 4 to 7 larger than without those. Fig.~\ref{f:lx_wei} illustrates that at a given luminosity each of those CVs has a factor $\sim$10 higher weight than a corresponding RBS-CV. The whole sample is dominated by just one CV, the low-luminosity object EX Dra, $\log L_X = 29.8$\,\ergs, a rather unhealthy situation for the whole analysis. The pre-Gaia distances of the RASS-CVs were quite reliable, hence the median X-ray luminosity of the RASS sample remained unchanged at $\log L_X = 31.2$\,\ergs. The distribution of luminosities is less dispersed in the center but a bit more fuzzy at the outskirts (Fig.~\ref{t:dens}). The standard deviation of $\log L_X$ was 0.54\,dex and is now 0.57\,dex, omitting the highest and lowest values it was 0.46\,dex and is now 0.38\,dex. For the Intermediate Polars the first important thing to note is that all but one object, GK Per, have larger distances than previously assumed. Hence, they are more luminous then thought, the median luminosity is shifted from $\log L_X = 33.1$\,\ergs\ to $33.5$\,\ergs, the survey volume becomes larger and the space density conversely smaller. An upper limit to the space density of the IPs is $\rho < 1.3 \times 10^{-7}$\,\ppc\ at 90\% confidence, a significant reduction compared to the published value. The most likely value at $7.4_{-1.7}^{+4.8} \times 10^{-8}$\ppc\ is at 74\% of the published one. Distances to the IPs used by \cite{pretorius+mukai14} were either taken from the literature (3 trigonometric and 4 photometric parallaxes from the donor star) or were newly determined and based on WISE IR-magnitudes combined with the semi-empirical donor sequence by \cite{knigge06}. Not surprisingly, the mean distance ratio new/old is reasonably small, $d_{\rm rat} = 1.12$ for the three IPs which had a trigonometric parallax previously, among them GK Per with a Gaia distance smaller than published. The IPs with photometric parallaxes of the donor have $d_{\rm rat} = 1.33$ and those with estimated distances from WISE and the empirical donor sequence have a mean ratio $d_{\rm rat} = 1.74$. This leaves the two possibilities that either the IR donor sequence is somehow biased or that an additional emission component (dust, cyclotron radiation, free-free emission) mimics brighter secondaries. Otherwise the IP sample appears more homogeneous than the sample of non-magnetic CVs. There is not one object or a subgroup of objects that dominates the space density. However, given the rather high flux limit the number for $\rho$ derived here is valid only for the potentially rare objects with high luminosities. The putative class of low-luminosity IPs remains yet to be uncovered \citep[see e.g.~][]{worpel+18}. In this study an update on the space density of X-ray selected CVs was given. CVs appear to be rarer than previously thought. While the limitations due to uncertain distances are overcome thanks to Gaia major obstacles preventing further progress remain. These are the small sample sizes due to shallow flux limits of past X-ray surveys and the ignorance about the proper scale heights of the samples. It also appears likely that the existing samples are inhomogeneously composed as far as their scale height is concerned, they contain long-and short-period objects. Part of them lack determinations of their orbital period, which could be used to assign class membership, belonging to an older or younger population with corresponding scale height. The limitation given by the small sample sizes will hopefully soon be overcome as a result of the upcoming eROSITA all-sky surveys~\citep{merloni+12,schwope12} with an all-sky flux limit comparable to the ROSAT-NEP survey but with enlarged energy coverage, $0.3-10$\,keV, and better spatial resolution compared to ROSAT. Performing the survey is just the first step on a longer ladder which will involve spectroscopic identification, classification and detailed follow-up to determine orbital periods. | 18 | 8 | 1808.08144 |
1808 | 1808.08975_arXiv.txt | Wormhole solutions in a generalized hybrid metric-Palatini matter theory, given by a gravitational Lagrangian $f\left(R,\cal{R}\right)$, where $R$ is the metric Ricci scalar, and $\mathcal{R}$ is a Palatini scalar curvature defined in terms of an independent connection, and a matter Lagrangian, are found. The solutions are worked in the scalar-tensor representation of the theory, where the Palatini field is traded for two scalars, $\varphi$ and $\psi$, and the gravitational term $R$ is maintained. The main interest in the solutions found is that the matter field obeys the null energy condition (NEC) everywhere, including the throat and up to infinity, so that there is no need for exotic matter. The wormhole geometry with its flaring out at the throat is supported by the higher-order curvature terms, or equivalently, by the two fundamental scalar fields, which either way can be interpreted as a gravitational fluid. Thus, in this theory, in building a wormhole, it is possible to exchange the exoticity of matter by the exoticity of the gravitational sector. The specific wormhole displayed, built to obey the matter NEC from the throat to infinity, has three regions, namely, an interior region containing the throat, a thin shell of matter, and a vacuum Schwarzschild anti-de Sitter (AdS) exterior. For hybrid metric-Palatini matter theories this wormhole solution is the first where the NEC for the matter is verified for the entire spacetime keeping the solution under asymptotic control. The existence of this type of solutions is in line with the idea that traversable wormholes bore by additional fundamental gravitational fields, here disguised as scalar fields, can be found without exotic matter. Concomitantly, the somewhat concocted architecture needed to assemble a complete wormhole solution for the whole spacetime may imply that in this class of theories such solutions are scarce. | Within general relativity, wormholes were found as exact solutions connecting two different asymptotically flat regions of spacetime \cite{Morris:1988cz,Visser:1995cc} as well as two different asympotically de Sitter (dS) or anti-de Sitter (AdS) regions \cite{lemoslobooliveira}. The fundamental ingredient in wormhole physics is the existence of a throat satisfying a flaring-out condition. In general relativity this geometric condition entails the violation of the null energy condition (NEC). This NEC states that $T_{ab}\,k^a k^b \geq 0$, where $T_{ab}$ is the matter stress-energy tensor and $k^a$ is any null vector. Matter that violates the NEC is denoted as exotic matter. Wormhole solutions have been also found in other theories, see, e.g., \cite{Camera:1995,Nandi:1998,Bronnikov:2002rn,Camera:2003, Lobo:2007zb,Garattini:2007ff,Lobo:2008,Garattini:2009,Lobo:2010sb, Garattini:2011fs} and \cite{Lobo:2017oab} for reviews. In these works the NEC for the matter is also violated. However due to its nature, it is important and useful to minimize its usage. In fact, in the context of modified theories of gravity, it has been shown in principle that normal matter may thread the wormhole throat, and it is the higher-order curvature terms, which may be interpreted as a gravitational fluid, that support these nonstandard wormhole geometries. Indeed, in \cite{Lobo:2009ip} it was shown explicitly that in $f(R)$ theories wormhole throats can be theoretically constructed without the presence of exotic matter, in \cite{Garcia:2010xb,MontelongoGarcia:2010xd} nonminimal couplings were used to build such wormholes, and in \cite{Harko:2013yb} generic modified gravities were used also with that aim in mind, i.e., the wormhole throats are sustained by the fundamental fields presented in the modified gravity alone. This type of solutions were also found in Einstein-Gauss-Bonnet theory \cite{Bhawal:1992,Dotti:2007,Mehdizadeh:2015jra}, and are mentioned in brane world scenarios \cite{Lobo:2007}, in Brans-Dicke theories \cite{Anchordoqui:1997}, and in a hybrid metric-Palatini gravitational theory \cite{Capozziello:2012hr}. It is our aim to find wormhole solutions whose matter obeys the NEC not only at the throat but everywhere in a generalized hybrid metric-Palatini gravity, with action $f(R,{\cal R})$. The action $f(R,{\cal R})$ is well motivated. Indeed, a promising approach to modified gravity consists in having a hybrid metric-Palatini gravitational theory \cite{Harko:2011nh}, which consists in adding to the Einstein-Hilbert action $R$, a new term $f(\cal{R})$, where $\cal{R}$ is a curvature scalar defined in terms of an independent connection, and $f$ is some function of $\cal{R}$. In this approach, the metric and affine connection are regarded as independent degrees of freedom, and contrary to general relativity where the metric-affine, or Palatini, formalism coincides with the purely metric formalism, in an $R+f({\cal R})$ theory the two formalisms lead to different results \cite{Olmo:2011uz}. In the $R+f({\cal R})$ theory one retains through $R$ the positive results of general relativity, with further gravitational degrees of freedom being represented in the metric-affine $f({\cal R})$ component. One can further express the $R+ f({\cal R})$ theory in a dynamically equivalent scalar-tensor representation which simplifies the analysis. In this representation, besides wormhole solutions \cite{Capozziello:2012hr}, solar system tests and cosmological solutions have been analyzed and found \cite{Capozziello:2012ny,Capozziello:2012qt,Capozziello:2013yha, Capozziello:2013uya,Capozziello:2015lza}, see also \cite{HarkoLobo} for a review. A natural generalization to $R+f({\cal R})$ is to consider an $f(R,{\cal R})$ action, i.e., the gravitational action is taken to depend on a general function of both the metric and Palatini curvature scalars \cite{Tamanini:2013ltp}. One can also find the scalar-tensor representation of this generalized hybrid metric-Palatini gravity, which now has two scalar fields. Exact solutions were constructed representing an FLRW (Friedmann-Lema\^itre-Robsertson-Walker) universe in a generalized hybrid metric-Palatini theory \cite{Rosa:2017jld}. Among other relevant results, it was shown that it is possible to obtain exponentially expanding solutions for flat universes even when the cosmology is not purely vacuum. In addition, the junction conditions in this theory have been worked out \cite{rosalemos1}. In this work, our aim is to find static and spherically symmetric wormholes solutions in the generalized $f(R,{\cal R})$ hybrid metric-Palatini matter theory in which the matter satisfies the NEC everywhere, from the throat to infinity, so there is no need for exotic matter. This fills a gap in the literature as most of the work that has been done in this area has been aimed at finding solutions where the NEC is satisfied solely at the wormhole throat paying no attention to the other regions. As far as we are aware our work is the first where the NEC is verified for the entire spacetime. This paper is organized as follows. In Sec.~\ref{secII}, we consider the action and write out the gravitational field equations, both in the curvature and in the equivalent scalar-tensor representation. In Sec.~\ref{secIII}, we present the equations of motion for static and spherically symmetric wormholes solutions. In Sec.~\ref{secIV}, we impose specific choices for the metric redshift and shape functions and for the potential governing the scalar fields, and find solutions where the matter threading the wormhole satisfies the NEC in the whole spacetime. In Sec.~\ref{conclusion}, we conclude. | In this work, we found traversable asymptotically AdS wormhole solutions that obey the NEC everywhere in the generalized hybrid metric-Palatini gravity theory, so there is no need for exotic matter. The generalized hybrid theory consists in a gravitational action given by $f\left(R,\cal{R}\right)$, where $R$ is the metric Ricci scalar, and $\mathcal{R}\equiv\mathcal{R}^{ab}g_{ab}$ is the Palatini curvature defined in terms of an independent connection, to which a matter action is added. The gravitational action can be given in the scalar-tensor representation where the metric Ricci scalar $R$ is still present but now coupled to two scalar fields $\varphi$ and $\psi$. The equations of motion in this representation were obtained. The interior wormhole solution obtained verifies the NEC near and at the throat of the wormhole. The matching of the interior solution to an exterior vacuum solution yields a thin shell respecting also the NEC. Finding a combination of parameters that allow for the interior and the exterior solution to be matched without violating the NEC in the interior solution and at the thin shell is a problem that requires fine-tuning, but can be acomplished. We presented a specific combination of parameters with which it is possible to build the full wormhole solution and found that it is asymptotically AdS. Most of the work that has been done in hybrid metric-Palatini gravity theories has been aimed to find solutions where the NEC is satisfied solely at the throat of the wormhole. Within these theories our solution is the first where the NEC is verified for the entire spacetime. There are two main interesting conclusions. On one hand, the existence of these solutions is in agreement with the understanding that traversable wormholes supported by extra fundamental gravitational fields, here in the guise of scalar fields, can occur without the need of exotic matter. On the other hand, the rather contrived construction necessary to build a spacetime complete wormhole solution may indicate that there are not many such solutions around in this class of theories. \centerline{} \vskip 1cm | 18 | 8 | 1808.08975 |
1808 | 1808.08234_arXiv.txt | Recent hydrodynamical models of supernova remnants (SNRs) demonstrate that their evolution depends heavily on the inhomogeneities of the surrounding medium. As SNRs expand, their morphologies are influenced by the non-uniform and turbulent structure of their environments, as reflected in their radio continuum emission. In this paper, we measure the asymmetries of 96 SNRs in radio continuum images from three surveys of the Galactic plane and compare these results to the SNRs' radii, which we use as a proxy for their age. We find that larger (older) SNRs are more elliptical/elongated and more mirror asymmetric than smaller (younger) SNRs, though the latter vary in their degrees of asymmetry. This result suggests that SNR shells become more asymmetric as they sweep up the interstellar medium (ISM), as predicted in hydrodynamical models of SNRs expanding in a multi-phase or turbulent ISM. | \label{sec:intro} Synchrotron radiation is produced by electrons that are accelerated at the shocks of supernova remnants (SNRs) and interact with the local magnetic field (see e.g., \citealt{berezhko04}). This emission dominates at radio wavelengths, particularly at lower frequencies (e.g. 1.4~GHz) where thermal bremsstrahlung contributes less to the spectrum. Radio continuum surveys are the primary means by which new SNRs are identified (see e.g., \citealt{chomiuk09} and references therein). At present, 295 objects have been classified as SNRs in the Milky Way \citep{green17}, and 95\% of these are radio sources \citep{dubner15}. Historically, SNRs have been classified based on their non-thermal radio shells (e.g. \citealt{milne69}). After many SNRs were resolved at radio wavelengths, astronomers began to categorize them based on their morphologies (e.g., \citealt{green84}). The categories most commonly used are shell-type, composite, and mixed-morphology. Shell-type SNRs are those with limb-brightened radio shell emission (e.g., SN~1006: \citealt{winkler14}, G1.9$+$0.3: \citealt{dehorta14}) and 79\% of Galactic SNRs fall into this group \citep{dubner15}. Composite SNRs are ones that have both the shell and the center-filled emission from a pulsar wind nebula) (e.g., MSH 15-56: \citealt{dickel00}, Vela SNR: \citealt{dodson03}). Additionally, X-ray imaging led to another class of SNRs called mixed-morphology or thermal-composite \citep{rho98,lazendic06,shelton04} that have radio shells with center-filled X-rays (e.g., W44, IC443: \citealt{kawasaki05}). \begin{deluxetable*}{llcrcrrrc}[ht] \footnotesize \tablenum{1} \tablecaption{List of SNRs in Our Sample from THOR\label{tab:tableTHOR}} \tablehead{ \colhead{No.} & \colhead{Source\tablenotemark{a}} & \colhead{Alternate Names} & \colhead{Distance\tablenotemark{b}} & \colhead{Evidence of} & \colhead{$P_{1}/P_{0}$} & \colhead{$P_{2}/P_{0}$} & \colhead{$P_{3}/P_{0}$} & \colhead{References\tablenotemark{d}} \\ \colhead{} & \colhead{} & \colhead{} & \colhead{(kpc)} & \colhead{Explosion Type\tablenotemark{c}} & \colhead{($\times10^{-5}$)} & \colhead{($\times10^{-7}$)} & \colhead{($\times10^{-7}$)} & \colhead{}} \startdata 1 & G15.9$+$0.2 & & 8.5* & N & 84.9$_{-7.2}^{+9.0}$ & 8.50$_{-6.37}^{+6.70}$ & 4.90$_{-2.83}^{+2.46}$ & 1 \\ 2 & G16.7$+$0.1$^\dagger$ & & 10 & N & 3.01$\pm{1.3}$ & 18.7$_{-8.3}^{+7.3}$ & 2.79$_{-1.82}^{+1.84}$ & 2 \\ 3 & G18.1$-$0.1$^\dagger$ & & 6.4$\pm{0.2}$ & E & 95.9$_{-1.4}^{+1.5}$ & 326$\pm{9}$ & 5.25$_{-0.64}^{+0.51}$ & 3 \\ 4& G18.8$+$0.3$^\dagger$ & Kes~67 & 13.8$\pm{4.0}$ & E & 240$_{-2}^{+1}$ & 1119$_{-13}^{+12}$ & 149$\pm{2}$ & 3,4 \\ 5&G20.0$-$0.2 & & 11.2$\pm{0.3}$ & E, N & 27.3$\pm{0.7}$ & 43.9$_{-2.8}^{+3.3}$ & 2.71$_{-0.47}^{+0.57}$ & 3 \\ 6&G20.4$+$0.1 & & 7.8 & E, N & 74.9$\pm{2.9}$ & 259$_{-18}^{+19}$ & 33.4$_{-2.6}^{+2.4}$ & 5 \\ 7&G21.5$-$0.1 & & 8.5* & -- & 13.1$_{-2.7}^{+2.2}$ & 16.2$_{-11.7}^{+12.5}$ & 2.74$_{-1.94}^{+1.62}$ & -- \\ 8& G21.8$-$0.6$^\dagger$ & Kes~69 & 5.2$_{-0.3}^{+0.0}$ & E & 370$\pm{1}$ & 760$\pm{5}$ & 100$\pm{1}$ & 6, 7 \\ 9&G22.7$-$0.2 & & 4.4$\pm{0.4}$ & E & 36.4$_{-0.2}^{+0.3}$ & 162$\pm{2}$ & 6.03$_{-0.20}^{+0.17}$ & 8\\ 10&G23.3$-$0.3$^\dagger$ & W41 & 4.4$\pm{4}$ & E & 96.5$_{-0.3}^{+0.4}$ & 951$_{-4}^{+5}$ & 94.5$\pm{7}$ & 9, 10\\ 11 &G27.4$+$0.0$^\dagger$ & Kes~73 & 5.8$\pm{0.3}$ & N, A & 16.2$\pm{0.4}$ & 85.2$_{-3}^{+4}$ & 3.64$_{-0.32}^{+0.38}$ & 3, 11\\ 12&G28.6$-$0.1 & & 9.6$\pm{0.3}$ & E & 3.82$_{-0.23}^{+0.21}$ & 819$\pm{12}$ & 189$\pm{3}$ & 5\\ 13&G29.6$+$0.1$^\dagger$ & & 10.0$\pm{5.0}$ & N & 5.95$_{-1.12}^{+1.13}$ & 81.3$_{-15.9}^{+12.4}$ & 17.2$_{-3.4}^{+3.0}$ & 12 \\ 14& G31.9$+$0.0$^\dagger$ & 3C~391 & 7.1$\pm{0.4}$ & E, A & 63.9$\pm{0.3}$ & 240$_{-2}^{+3}$ & 6.54$_{-0.18}^{+0.17}$ & 13, 14\\ 15&G32.4$+$0.1$^\dagger$ & & 17 & E & 48.5$_{-5.4}^{+4.9}$ & 170$_{-34}^{+28}$ & 92.5$_{-14.2}^{+14.6}$ & 14, 15\\ 16 & G32.8$-$0.1$^\dagger$ & Kes~78 & 4.8 & E, N & 87.9$_{-1.6}^{+1.3}$ & 4005$_{-50}^{+53}$ & 94.1$_{-2.5}^{+3.0}$ & 16, 17\\ 17&G33.2$-$0.6$^\dagger$ & & 8.5* & -- & 180$_{-4}^{+3}$ & 106$\pm{8}$ & 0.34$_{-0.24}^{+0.25}$ & --\\ 18&G33.6$+$0.1$^\dagger$ & Kes~79 & 3.5$\pm{0.3}$ & N, A & 96.8$_{-0.9}^{+0.8}$ & 3.70$_{-0.50}^{+0.60}$ & 24.2$\pm{0.6}$ & 3,18,19\\ 19 & G34.7$-$0.4$^\dagger$ & W44 & 3.0$\pm{0.3}$ & E, N, A & 5.28$_{-0.03}^{+0.04}$ & 575$\pm{2}$ & 27.0$\pm{0.2}$ & 3,20,21 \\ 20 & G35.6$-$0.4 & & 3.6$\pm{0.4}$ & E & 5.21$_{-0.20}^{+0.23}$ & 538$\pm{9}$ & 5.28$_{-0.53}^{+0.47}$ & 22\\ 21 & G36.6$-$0.7 & & 8.5* & -- & 190$\pm{3}$ & 825$\pm{22}$ & 34.8$_{-2.8}^{+2.9}$ & -- \\ 22 & G49.2$-$0.7$^\dagger$ & W51C & 5.4$\pm{0.6}$ & E, A & 8.56$_{-0.09}^{+0.08}$ & 726$_{-3}^{+2}$ & 50.7$\pm{0.3}$ & 23, 24\\ \enddata \tablenotetext{a}{ $\dagger$ Denotes SNRs with evidence of interaction with a molecular cloud: G16.7$+$0.1: \citealt{green97} (G97); \citealt{reynosomangum00}; \citealt{hewitt08} (H08); \citealt{kilpatrick16} (K16); G18.1$-$0.1: \citealt{froebrich15} (F15); G18.8$+$0.3: \citealt{dubner04,tian07}; G21.8$-$0.6: G97; H08; \citealt{zhou09,hewitt09b}; F15; G23.3$-$0.3: \citealt{frail13} G27.4$+$0.0: F15, K16; G29.6$+$0.1: K16; G31.9$+$0.0: \citealt{frail96,reachrho99,reach02}, H08, F15, K16; G32.4$+$0.1: K16; G32.8$-$0.1: \citealt{koralesky98,zhouchen11}; F15; G33.2$-$0.6: F15, K16; G33.6$+$0.1: K16, \citealt{zhou16}; G34.7$-$0.4: \citealt{claussen97,seta98b,reach05}; H08; G49.2$-$0.7: G97; \citealt{koomoon97ii}; H08.} \tablenotetext{b}{ * Denotes a SNR with an assumed distance of 8.5 kpc (the International Astronomical Union recommended distance to the Galactic center) because the source does not have good constraints on its distance.} \tablenotetext{c}{Evidence of Explosion Type: N = Neutron star detection; E = Environment suggestive of core-collapse SNe (e.g., molecular cloud interaction, nearby H{\sc ii} regions); A = metal abundances from X-ray observations are consistent with core collapse SNe.} \tablenotetext{d}{References: (1) \citealt{reynolds06}; (2) \citealt{helfand03}; (3) \citealt{ranasingheleahy18}; (4) \citealt{tian07}; (5) \citealt{ranasinghe18}; (6) \citealt{zhouchen09}; (7) \citealt{leahytian08}; (8) \citealt{su14}; (9) \citealt{su15}; (10) \citealt{frail13}; (11) \citealt{vasisht97}; (12) \citealt{vasisht00}; (13) \citealt{ranasinghe17}; (14) \citealt{kilpatrick16}; (15) \citealt{yamaguchi04}; (16) \citealt{zhouchen11}; (17) \citealt{bamba16}; (18) \citealt{sato16}; (19) \citealt{auchettl14}; (20) \citealt{uchida12}; (21) \citealt{radhakrishnan72}; (22) \citealt{zhu13}; (23) \citealt{tianleahy13}; (24) \citealt{sasaki14}} \vspace{-10mm} \end{deluxetable*} \begin{deluxetable*}{llcrcrrrc}[ht] \footnotesize \tablenum{2} \tablecaption{List of SNRs in Our Sample from CGPS\label{tab:tableCGPS}} \tablehead{ \colhead{No.} & \colhead{Source\tablenotemark{a}} & \colhead{Alternate Names} & \colhead{Distance\tablenotemark{b}} & \colhead{Evidence of} & \colhead{$P_{1}/P_{0}$} & \colhead{$P_{2}/P_{0}$} & \colhead{$P_{3}/P_{0}$} & \colhead{References\tablenotemark{d}} \\ \colhead{} & \colhead{} & \colhead{} & \colhead{(kpc)} & \colhead{Explosion Type\tablenotemark{c}} & \colhead{($\times10^{-5}$)} & \colhead{($\times10^{-7}$)} & \colhead{($\times10^{-7}$)} & \colhead{}} \startdata 23 & G65.1$+$0.6$^\dagger$ & -- & 9.2$^{+0.4}_{-0.2}$ & N, E & 50.1$_{-1.0}^{+0.9}$ & 1748$^{+19}_{-18}$ & 79.2$^{+2}_{-3}$ & 1 \\ 24 & G67.7$+$1.8 & -- & 8.5* & A & 59.9$_{-1.2}^{+1.4}$ & 16.9$^{+2.7}_{-2.4}$ & 48.4$\pm2$ & 2 \\ 25 & G69.0$+$2.7 & CTB~80 & 1.5$^{+0.6}_{-0.4}$ & N & 6040$\pm38$ & 20060$^{+181}_{-186}$ & 2073$^{+24}_{-22}$ & 3, 4 \\ 26 & G69.7$+$1.0 & -- & 8.5* & -- & 11.7$\pm0.4$ & 67.6$\pm2.8$ & 4.16$^{+0.37}_{-0.40}$ & -- \\ 27 & G73.9$+$0.9$^\dagger$ & -- & 1.3$^{+0.7}_{-0.8}$ & E & 22.6$\pm0.2$ & 25.4$^{+8.4}_{-8.3}$ & 12.4$\pm0.3$ & 5 \\ 28 & G78.2$+$2.1 & $\gamma$~Cygni SNR & 2.0$^{+0.6}_{-0.3}$ & N & 77.1$\pm0.1$ & 1139$^{+1}_{-2}$ & 93.3$\pm0.2$ & 6, 7 \\ 29 & G84.2$-$0.8 & -- & 6.0$\pm$0.2 & -- & 780$_{-8}^{+7}$ & 5151$^{+6}_{-5}$ & 87.3$^{+2.8}_{-3.0}$ & 8 \\ 30 & G85.4$+$0.7 & -- & 3.5$\pm$1.0 & -- & 792$_{-8}^{+9}$ & 1429$\pm3$ & 88.9$^{+3.7}_{-4.2}$ & 9 \\ 31 & G93.7$-$0.2 & CTB~104A & 1.5$\pm$0.2 & -- & 329$_{-1}^{+1}$ & 230$\pm3$ & 221$\pm2$ & 10 \\ 32 & G94.0$+$1.0$^\dagger$ & 3C~434.1 & 4.5$\pm$1.5 & E & 126$\pm1$ & 489$\pm2$ & 5.08$\pm0.14$ & 11, 12 \\ 33 & G106.3$+$2.7 & -- & 0.8$^{+1.2}_{-0.1}$ & N, E & 50.0$_{-0.6}^{+0.7}$ & 6010$^{+4}_{-3}$ & 104$\pm2$ & 13 \\ 34 & G109.1$-$1.0$^{\dagger}$ & CTB~109 & 3.2$\pm$0.2 & N, E & 30.7$\pm0.1$ & 739$^{+3}_{-2}$ & 12.1$^{+0.2}_{-0.1}$ & 14 \\ 35 & G114.3$+$0.3 & -- & 0.7$^{+0.9}_{-0.0}$ & N & 53.1$\pm0.2$ & 128$\pm1$ & 1.66$^{+0.07}_{-0.06}$ & 15, 16 \\ 36 & G116.5$+$1.1 & -- & 1.6$\pm$0.6 & -- & 24.3$\pm0.3$ & 1458$\pm9$ & 12.6$^{+0.5}_{-0.4}$ & 15 \\ 37 & G116.9$+$0.2 & CTB~1 & 1.6$^{+1.5}_{-0.0}$ & A & 87.7$\pm0.4$ & 331$^{+3}_{-2}$ & 10.6$\pm0.3$ & 15, 17 \\ 38 & G120.1$+$1.4$^{\diamond}$ & Tycho & 2.4$^{+2.6}_{-0.9}$ & A, L & 8.00$\pm0.03$ & 0.09$\pm0.01$ & 0.65$\pm0.01$ & 18, 19 \\ 39 & G127.1$+$0.5$^{\dagger}$ & R~5 & 1.2$\pm$0.1 & E & 20.6$\pm0.1$ & 124$\pm1$ & 10.3$\pm0.2$ & 20 \\ 40 & G132.7$+$1.3$^{\dagger}$ & HB~3 & 2.2$\pm$0.2 & E, A & 308$_{-0.4}^{+0.5}$ & 583$\pm2$ & 74.0$\pm0.4$ & 17, 21 \\ 41 & G160.9$+$2.6 & HB~9 & 0.8$^{+1.0}_{-0.4}$ & -- & 47.7$\pm0.1$ & 364$\pm1$ & 0.33$\pm0.02$ & 22 \\ 42 & G166.0$+$4.3$^{\dagger}$ & VRO~42.05.01 & 4.5$\pm$1.5 & E & 1.13$_{-0.10}^{+0.09}$ & 745$^{+11}_{-9}$ & 63.8$^{+1.4}_{-1.4}$ & 23 \enddata \tablenotetext{a}{ $\diamond$ Denotes SNRs thought to be from Type Ia SNe. $\dagger$ Denotes SNRs with evidence of interaction with a molecular cloud: G65.1$+$0.6: F15; G73.9$+$0.9: \citealt{zdz16}; G94.0$+$1.0: \citealt{jeong13}; G109.1$-$1.0: \citealt{sasaki06}; G127.1$+$0.5: \citealt{zhou14}; G132.7$+$1.3: K16, \citealt{zhou16b}} \tablenotetext{b}{ * Denotes a SNR with an assumed distance of 8.5 kpc (the International Astronomical Union recommended distance to the Galactic center) because the source does not have good constraints on its distance.} \tablenotetext{c}{Evidence of Explosion Type: N = Neutron star detection; E = Environment suggestive of core-collapse SNe (e.g., molecular cloud interaction, nearby H{\sc ii} regions); A = metal abundances from X-ray observations; L = light echo spectrum.} \tablenotetext{d}{References: (1) \citealt{tian06}; (2) \citealt{hui09}; (3) \citealt{li05}; (4) \citealt{leahy12}; (5) \citealt{loz93}; (6) \citealt{leahy13}; (7) \citealt{hui15}; (8) \citealt{leahy12b}; (9) \citealt{jackson08}; (10) \citealt{uya02}; (11) \citealt{foster05}; (12) \citealt{jeong13}; (13) \citealt{kothes01}; (14) \citealt{kothes12}; (15) \citealt{yar04}; (16) \citealt{kulkarni93}; (17) \citealt{lazendic06}; (18) \citealt{tian11}; (19) \citealt{krause08}; (20) \citealt{leahy06}; (21) \citealt{routledge91}; (22) \citealt{leahy07}; (23) \citealt{landecker89}} \end{deluxetable*} \begin{deluxetable*}{llcrcrrrc}[ht] \footnotesize \tablenum{3} \tablecaption{List of SNRs in Our Sample from MOST\label{tab:tableMOST}} \tablehead{ \colhead{No.} & \colhead{Source\tablenotemark{a}} & \colhead{Alternate Names} & \colhead{Distance\tablenotemark{b}} & \colhead{Evidence of} & \colhead{$P_{1}/P_{0}$} & \colhead{$P_{2}/P_{0}$} & \colhead{$P_{3}/P_{0}$} & \colhead{References\tablenotemark{d}} \\ \colhead{} & \colhead{} & \colhead{} & \colhead{(kpc)} & \colhead{Explosion Type\tablenotemark{c}} & \colhead{($\times10^{-5}$)} & \colhead{($\times10^{-7}$)} & \colhead{($\times10^{-7}$)} & \colhead{}} \startdata 43 & G289.7$-$0.3 & -- & 8.5* & -- & 71.8$^{+0.4}_{-0.3}$ & 971$\pm5$ & 0.04$\pm0.01$ & \\ 44 & G290.1$-$0.8$^{\dagger}$ & MSH~11$-$61A & 7.0$\pm$1.0 & N, E, A & 23.8$\pm0.1$ & 230$\pm1$ & 9.40$\pm0.07$ & 1, 2, 3 \\ 45 & G294.1$-$0.0 & -- & 8.5* & -- & 788$^{+11}_{-10}$ & 6189$^{+134}_{-121}$ & 2376$^{+37}_{-42}$ & \\ 46 & G296.1$-$0.5 & -- & 3.0$\pm$1.0 & A & 523$\pm2$ & 2578$^{+16}_{-19}$ & 688$^{+5}_{-6}$ & 4, 5 \\ 47 & G296.8$-$0.3 & -- & 9.6$\pm$0.6 & -- & 1134$\pm3$ & 3630$\pm17$ & 373$\pm3$ & 6 \\ 48 & G298.6$-$0.0$^{\dagger}$ & -- & 8.5* & E & 57.4$\pm0.1$ & 1285$^{+4}_{-3}$ & 121$\pm1$ & 7 \\ 49 & G299.6$-$0.5 & -- & 8.5* & -- & 115$\pm1$ & 67.3$^{+3.0}_{-2.6}$ & 77.7$^{+1.3}_{-1.5}$ & \\ 50 & G301.4$-$1.0 & -- & 8.5* & -- & 79.6$^{+1.8}_{-2.0}$ & 2422$\pm1$ & 591$^{+14}_{-13}$ & \\ 51 & G302.3$+$0.7$^{\dagger}$ & -- & 8.5* & E & 48.2$^{+0.7}_{-0.6}$ & 946$\pm10$ & 567$^{+3}_{-4}$ & 8 \\ 52 & G304.6$+$0.1$^{\dagger}$ & Kes~17 & 9.7$^{+4.3}_{-1.7}$ & E, A & 30.1$\pm0.1$ & 137$\pm1$ & 6.71$\pm0.03$ & 9, 10 \\ 53 & G308.1$-$0.7 & -- & 8.5* & -- & 12.5$\pm0.4$ & 74.8$^{+3.0}_{-3.4}$ & 141$\pm2$ & \\ 54 & G308.8$-$0.1 & -- & 6.9$^{+8.1}_{-2.9}$ & N & 447$\pm1$ & 2612$^{+7}_{-6}$ & 221$\pm1$ & 11 \\ 55 & G309.2$-$0.6 & -- & 4.0$\pm$2.0 & -- & 4.93$\pm0.09$ & 2278$^{+6}_{-5}$ & 65.5$^{+3.3}_{-3.1}$ & 12 \\ 56 & G309.8$+$0.0 & -- & 8.5* & -- & 503$\pm2$ & 1237$^{+17}_{-16}$ & 1418$^{+7}_{-8}$ & -- \\ 57 & G310.6$-$0.3 & Kes~20B & 8.5* & -- & 228$\pm1$ & 1005$\pm6$ & 47.8$\pm0.6$ & -- \\ 58 & G310.8$-$0.4 & Kes~20A & 13.7 & E & 209$\pm1$ & 4222$^{+9}_{-7}$ & 509$\pm2$ & 13, 14 \\ 59 & G311.5$-$0.3$^{\dagger}$ & -- & 14.8 & E & 1.92$\pm0.02$ & 9.75$^{+0.14}_{-0.16}$ & 0.28$\pm0.02$ & 14 \\ 60 & G312.4$-$0.4$^{\dagger}$ & -- & 6.0$^{+8.0}_{-0.0}$ & E & 441$\pm1$ & 172$^{+4}_{-3}$ & 151$\pm1$ & 15 \\ 61 & G315.4$-$2.3$^{\diamond}$ & RCW~86 & 2.5$^{+0.3}_{-0.2}$ & E, A & 116$\pm6$ & 4379$\pm2$ & 2098$^{+7}_{-6}$ & 16, 17, 18 \\ 62 & G316.3$-$0.0 & MSH~14$-$57 & 7.2$\pm$0.6 & -- & 2.28$^{+0.05}_{-0.04}$ & 228$^{+1}_{-2}$ & 61.1$\pm0.3$ & 10 \\ 63 & G317.3$-$0.2 & -- & 8.5* & -- & 267$\pm1$ & 7315$\pm15$ & 781$\pm3$ & \\ 64 & G318.2$+$0.1 & -- & 8.5* & -- & 66.6$^{+0.9}_{-1.0}$ & 2537$^{+21}_{-23}$ & 120$^{+4}_{-3}$ & \\ 65 & G321.9$-$0.3 & -- & 6.5$^{+3.5}_{-1.0}$ & N & 289$\pm2$ & 1588$^{+18}_{-16}$ & 85.9$^{+1.9}_{-2.2}$ & 19 \\ 66 & G321.9$-$1.1 & -- & 8.5* & -- & 2051$^{+25}_{-23}$ & 7226$^{+122}_{-155}$ & 3639$\pm6$ & \\ 67 & G322.5$-$0.1 & -- & 8.5* & N & 50.0$^{+0.5}_{-0.6}$ & 257$^{+5}_{-4}$ & 112$\pm1$ & 20 \\ 68 & G323.5$+$0.1 & -- & 8.5* & -- & 21.4$^{+0.1}_{-0.2}$ & 460$\pm3$ & 51.6$^{+0.5}_{-0.4}$ & -- \\ 69 & G326.3$-$1.8 & MSH~15$-$56 & 4.1$\pm$0.7 & N & 35.6$\pm0.1$ & 114$\pm1$ & 0.60$\pm0.02$ & 16, 21 \\ 70 & G327.1$-$1.1 & -- & 8.5$\pm$0.5 & N & 45.3$\pm0.3$ & 126$^{+1}_{-2}$ & 11.8$\pm0.2$ & 22 \\ 71 & G327.4$+$0.4 & Kes~27 & 4.3$^{+1.1}_{-0.0}$ & N & 237$\pm1$ & 1194$^{+5}_{-4}$ & 113$\pm1$ & 23, 24 \\ 72 & G327.4$+$1.0 & -- & 8.5* & -- & 79.2$\pm0.4$ & 1040$^{+5}_{-4}$ & 240$\pm1$ & -- \\ \enddata \tablenotetext{a}{ $\diamond$ Denotes SNRs thought to be from Type Ia SNe. $\dagger$ Denotes SNRs with evidence of interaction with a molecular cloud: G290.1$-$0.8: \citealt{filipovic05}; G298.6$-$0.0: \citealt{acero16}; G302.3$+$0.7: \citealt{frail96} (F96); G304.6$+$0.1: F96; \citealt{hewitt09}; G311.5$-$0.3: \citealt{andersen11}; G312.4$-$0.4: F96; G332.4$+$0.1: F96; G332.4$-$0.4: F96; \citealt{paron06}; G337.0$-$0.1: F96; G337.8$-$0.1: \citealt{koralesky98,zhang15}; G346.6$-$0.2: \citealt{koralesky98,hewitt09,andersen11}; G348.5$+$0.1: F96, \citealt{reynoso00}} \tablenotetext{b}{ * Denotes a SNR with an assumed distance of 8.5 kpc (the International Astronomical Union recommended distance to the Galactic center) because the source does not have good constraints on its distance.} \tablenotetext{c}{Evidence of Explosion Type: N = Neutron star detection; E = Environment suggestive of core-collapse SNe (e.g., molecular cloud interaction, nearby H{\sc ii} regions); A = metal abundances from X-ray observations; L = light echo spectrum.} \tablenotetext{d}{References: (1) \citealt{reynoso06}; (2) \citealt{auchettl15}; (3) \citealt{kaspi97}; (4) \citealt{castro11}; (5) \citealt{longmore77}; (6) \citealt{gaensler98b}; (7) \citealt{bamba16b}; (8) \citealt{frail96}; (9) \citealt{washino16}; (10) \citealt{caswell75}; (11) \citealt{caswell92}; (12) \citealt{rakowski01}; (13) \citealt{reach06}; (14) \citealt{andersen11}; (15) \citealt{doherty03}; (16) \citealt{rosado96}; (17) \citealt{sollerman03}; (18) \citealt{williams11}; (19) \citealt{stewart93}; (20) \citealt{whiteoak96}; (21) \citealt{temim13}; (22) \citealt{sun99}; (23) \citealt{mcclure01}; (24) \citealt{chen08}; (25) \citealt{park09}; (26) \citealt{vink04}; (27) \citealt{reynoso04}; (28) \citealt{frank15}; (29) \citealt{kaspi96}; (30) \citealt{eger11}; (31) \citealt{rakowski06}; (32) \citealt{yamaguchi14}; (33) \citealt{kothes07}; (34) \citealt{giacani11}; (35) \citealt{yamaguchi12}; (36) \citealt{koralesky98}; (37) \citealt{tian12}; (38) \citealt{halpern10}; (39) \citealt{tian07b}; (40) \citealt{giacani09}} \end{deluxetable*} \begin{deluxetable*}{llcrcrrrc}[ht!] \footnotesize \tablenum{3} \tablecaption{List of SNRs in Our Sample from MOST (continued) \label{tab:tableMOST2}} \tablehead{ \colhead{No.} & \colhead{Source\tablenotemark{a}} & \colhead{Alternate Names} & \colhead{Distance\tablenotemark{b}} & \colhead{Evidence of} & \colhead{$P_{1}/P_{0}$} & \colhead{$P_{2}/P_{0}$} & \colhead{$P_{3}/P_{0}$} & \colhead{References\tablenotemark{d}} \\ \colhead{} & \colhead{} & \colhead{} & \colhead{(kpc)} & \colhead{Explosion Type\tablenotemark{c}} & \colhead{($\times10^{-5}$)} & \colhead{($\times10^{-7}$)} & \colhead{($\times10^{-7}$)} & \colhead{}} \startdata 73 & G329.7$+$0.4 & -- & 8.5* & E & 106$\pm1$ & 5.63$^{+0.05}_{-0.04}$ & 64.6$^{+0.8}_{-0.7}$ & -- \\ 74 & G330.2$+$1.0 & -- & 4.9$^{+5.0}_{-0.0}$ & N & 144$\pm1$ & 114$\pm1$ & 24.1$^{+0.2}_{-0.3}$ & 23, 25 \\ 75 & G332.0$+$0.2 & -- & 8.5* & -- & 0.17$\pm0.01$ & 770$\pm2$ & 60.6$\pm0.3$ & -- \\ 76 & G332.4$+$0.1$^{\dagger}$ & Kes~32 & 7.5$^{+3.5}_{-0.9}$ & E & 18.2$\pm0.1$ & 1658$\pm2$ & 142$\pm1$ & 26 \\ 77 & G332.4$-$0.4$^{\dagger}$ & RCW~103 & 3.3$^{+1.3}_{-0.2}$ & N, E, A & 20.2$^{+2.8}_{-3.2}$ & 12.8$\pm0.1$ & 4.29$\pm0.02$ & 27, 28 \\ 78 & G335.2$+$0.1 & -- & 1.8 & N & 19.8$^{+10.2}_{-12.0}$ & 43.7$\pm0.6$ & 91.1$\pm0.4$ & 29, 30 \\ 79 & G336.7$+$0.5 & -- & 8.5* & -- & 534$^{+62}_{-54}$ & 2092$\pm6$ & 241$\pm1$ & -- \\ 80 & G337.2$-$0.7$^{\diamond}$ & -- & 2.0$\pm$0.5 & A & 4.43$\pm0.02$ & 3.71$^{+0.07}_{-0.08}$ & 2.75$\pm0.03$ & 31, 32 \\ 81 & G337.3$+$1.0 & Kes~40 & 8.5* & -- & 2.15$\pm0.02$ & 165$\pm1$ & 34.7$^{+0.2}_{-0.1}$ & -- \\ 82 & G337.8$-$0.1$^{\dagger}$ & Kes~41 & 11.0 & E & 5.51$\pm0.02$ & 1017$\pm1$ & 20.4$^{+0.1}_{-0.1}$ & 33 \\ 83 & G340.4$+$0.4 & -- & 8.5* & -- & 93.0$\pm0.2$ & 188$\pm1$ & 68.4 $\pm0.4$ & -- \\ 84 & G340.6$+$0.3 & -- & 15.0 & -- & 39.1$\pm0.2$ & 126$\pm1$ & 11.6$\pm0.1$ & 33 \\ 85 & G341.9$-$0.3 & -- & 8.5* & -- & 204$\pm1$ & 385$\pm3$ & 24.3$^{+0.3}_{-0.4}$ & -- \\ 86 & G342.0$-$0.2 & -- & 8.5* & -- & 50.0$\pm0.2$ & 203$\pm2$ & 9.25$^{+0.17}_{-0.21}$ & -- \\ 87 & G342.1$+$0.9 & -- & 8.5* & -- & 13.0$\pm0.2$ & 218$^{+3}_{-2}$ & 2.98$\pm0.16$ & -- \\ 88 & G343.1$-$0.7 & -- & 8.5* & -- & 34.4$^{+0.2}_{-0.3}$ & 2504$\pm9$ & 516$\pm3$ & -- \\ 89 & G344.7$-$0.1$^{\diamond}$ & -- & 6.3$^{+7.7}_{-0.1}$ & A & 52.9$\pm0.2$ & 63.7$^{+0.6}_{-0.7}$ & 14.3$^{+0.1}_{-0.2}$ & 34, 35 \\ 90 & G346.6$-$0.2$^{\dagger}$ & -- & 11 & E & 4.48$\pm0.04$ & 26.1$\pm0.3$ & 5.38$^{+0.06}_{-0.07}$ & 36 \\ 91 & G348.7$+$0.3 & CTB~37B & 13.2$\pm$0.2 & N & 578$\pm1$ & 88.3$^{+3.0}_{-2.7}$ & 43.5$^{+0.3}_{-0.4}$ & 37, 38 \\ 92 & G351.2$+$0.1 & -- & 8.5* & -- & 41.4$\pm0.1$ & 333$\pm1$ & 19.7$\pm0.1$ & \\ 93 & G351.7$+$0.8 & -- & 13.2$\pm$0.5 & -- & 131$\pm1$ & 261$\pm2$ & 2.88$\pm0.14$ & 39 \\ 94 & G351.9$-$0.9 & -- & 8.5* & -- & 270$\pm1$ & 533$\pm5$ & 92.2$\pm1.1$ & \\ 95 & G352.7$-$0.1$^{\diamond}$ & -- & 7.5$^{+0.9}_{-0.7}$ & A & 0.07$\pm0.01$ & 218$\pm1$ & 36.1$\pm0.1$ & 40, 41 \\ 96 & G354.8$-$0.8 & -- & 8.5* & -- & 11.7$\pm0.3$ & 3750$\pm2$ & 520$\pm3$ & \\ \enddata \tablenotetext{a}{ $\diamond$ Denotes SNRs thought to be from Type Ia SNe. $\dagger$ Denotes SNRs with evidence of interaction with a molecular cloud: G290.1$-$0.8: \citealt{filipovic05}; G298.6$-$0.0: \citealt{acero16}; G302.3$+$0.7: \citealt{frail96} (F96); G304.6$+$0.1: F96; \citealt{hewitt09}; G311.5$-$0.3: \citealt{andersen11}; G312.4$-$0.4: F96; G332.4$+$0.1: F96; G332.4$-$0.4: F96; \citealt{paron06}; G337.0$-$0.1: F96; G337.8$-$0.1: \citealt{koralesky98,zhang15}; G346.6$-$0.2: \citealt{koralesky98,hewitt09,andersen11}; G348.5$+$0.1: F96, \citealt{reynoso00}} \tablenotetext{b}{ * Denotes a SNR with an assumed distance of 8.5 kpc (the International Astronomical Union recommended distance to the Galactic center) because the source does not have good constraints on its distance.} \tablenotetext{c}{Evidence of Explosion Type: N = Neutron star detection; E = Environment suggestive of core-collapse SNe (e.g., molecular cloud interaction, nearby H{\sc ii} regions); A = metal abundances from X-ray observations; L = light echo spectrum.} \tablenotetext{d}{References: (1) \citealt{reynoso06}; (2) \citealt{auchettl15}; (3) \citealt{kaspi97}; (4) \citealt{castro11}; (5) \citealt{longmore77}; (6) \citealt{gaensler98b}; (7) \citealt{bamba16b}; (8) \citealt{frail96}; (9) \citealt{washino16}; (10) \citealt{caswell75}; (11) \citealt{caswell92}; (12) \citealt{rakowski01}; (13) \citealt{reach06}; (14) \citealt{andersen11}; (15) \citealt{doherty03}; (16) \citealt{rosado96}; (17) \citealt{sollerman03}; (18) \citealt{williams11}; (19) \citealt{stewart93}; (20) \citealt{whiteoak96}; (21) \citealt{temim13}; (22) \citealt{sun99}; (23) \citealt{mcclure01}; (24) \citealt{chen08}; (25) \citealt{park09}; (26) \citealt{vink04}; (27) \citealt{reynoso04}; (28) \citealt{frank15}; (29) \citealt{kaspi96}; (30) \citealt{eger11}; (31) \citealt{rakowski06}; (32) \citealt{yamaguchi14}; (33) \citealt{kothes07}; (34) \citealt{giacani11}; (35) \citealt{yamaguchi12}; (36) \citealt{koralesky98}; (37) \citealt{tian12}; (38) \citealt{halpern10}; (39) \citealt{tian07b}; (40) \citealt{giacani09}; (41) \citealt{sezer14}} \end{deluxetable*} Barrel-shaped or bilateral SNRs, a subgroup of shell-type SNRs, are characterized by an axisymmetric morphology with two bright limbs. \cite{gaensler98} analyzed a sample of bilateral SNRs at radio frequencies and showed that their axes tended to be aligned with the Galactic plane. More recently, \cite{west16} investigated all Milky Way bilateral SNRs, and they showed that a simple model of SNRs expanding into an ambient Galactic magnetic field could reproduce the observed radio morphologies. To quantify the complex and varied morphologies of SNRs, \citealt{lopez09bkgnd} developed and applied several mathematical tools. Using the power-ratio method (PRM), \cite{lopez09,lopez11} showed that the thermal X-ray emission of Type Ia SNRs is more symmetric and circular than that of core-collapse SNRs. Subsequently, \cite{peters13} extended this approach to infrared images of SNRs and found similar results as in the X-ray. Recently, \cite{tyler17} used the PRM to compare the SNR soft X-ray morphologies to neutron star velocities and showed that the neutron stars are moving opposite to the bulk of the SN ejecta in many sources. For a detailed summary of these results and those of other groups, see \cite{lopez18}. In this paper, we investigate the radio morphologies of SNRs in the Milky Way to examine how asymmetries evolve with size and age. SNRs are observable for \hbox{$\sim10^{4}-10^{5}$}~years at radio wavelengths \citep{sarbadhicary17}, and their morphologies are shaped by interactions with the surrounding medium (e.g., \citealt{zhang18}) and by the magnetic field (e.g., \citealt{orlando07,west17}). This paper is structured as follows. In Section~\ref{sec:data}, we describe the radio data and introduce our sample of Galactic SNRs. Section~\ref{sec:methods} outlines the power-ratio method which we employ to measure the asymmetries of the sources. Finally, Section~\ref{sec:results} presents our results and discusses the implications regarding SNR evolution. | \label{sec:results} In Figure \ref{fig:P2}, we plot the dipole power ratio ($P_{1}/P_{0}$; left), the quadrupole power ratio ($P_{2}/P_{0}$; middle), and the octupole power ratio ($P_{3}/P_{0}$; right) versus radius. To convert the radii to parsecs, we adopt the distances listed in Tables~\ref{tab:tableTHOR}--\ref{tab:tableMOST} and the angular extents given in \cite{green17}. \begin{deluxetable}{ccc} \tablenum{4}\tablecolumns{3} \tablecaption{Median Power-Ratio Results by Radius\label{tab:results}} \tablehead{\colhead{Power-Ratio} & \colhead{Radius $\lesssim$10~pc} & \colhead{Radius $\gtrsim$10~pc}} \startdata \cutinhead{All 96 SNRs} $P_{1}/P_{0}$ & $2.0\times10^{-4}$ & $6.7\times10^{-4}$ \\ $P_{2}/P_{0}$ & $8.1\times10^{-6}$ & $7.5\times10^{-5}$ \\ $P_{3}/P_{0}$ & $5.3\times10^{-7}$ & $6.8\times10^{-6}$ \\ \cutinhead{63 SNRs with Measured Distances} $P_{1}/P_{0}$ & $2.0\times10^{-4}$ & $5.1\times10^{-4}$ \\ $P_{2}/P_{0}$ & $8.5\times10^{-6}$ & $7.4\times10^{-5}$ \\ $P_{3}/P_{0}$ & $6.5\times10^{-7}$ & $6.5\times10^{-6}$ \\ \enddata \end{deluxetable} The sample spans a wide range of power-ratio values, and the median power ratios of the SNRs by radius are listed in Table ~\ref{tab:results}. Generally, SNRs with radii $\lesssim$10~pc have smaller power ratios than those with radii $\gtrsim$10~pc, and this trend is most pronounced in $P_{2}/P_{0}$. We note that 33 of the 96 SNRs have no distance measurements to date (the triangles in Figure~\ref{fig:P2}), and we have assumed a distance of 8.5~kpc to those sources. However, if those SNRs are located at closer distances, their power ratios would increase (see Section~\ref{sec:methods}). Thus, we also list the median power ratios of only the 63 SNRs with measured distances in Table~\ref{tab:results}. We find that this sub-sample gives similar results, with the smaller SNRs giving lower power ratios, indicative of less asymmetries than the larger SNRs. Comparing the SNRs associated with Type Ia versus those from core-collapse explosions, G337.2$-$0.7 and Tycho and have among the lowest power-ratio values among both samples. However, the other Type Ia SNRs have near the median or greater power ratios than the CC SNRs. These results suggest that radio continuum morphology is not a reflection of explosion type, in contrast to X-ray and infrared morphologies where the two classes have distinct symmetries \citep{lopez09,lopez11,peters13}. Gven that G337.2$-$0.7 and Tycho have the smallest and third-smallest radii, respectively, of the 96 SNRs in the sample, their symmetric morphologies in the radio may simply reflect their young age and small size. The radii $R_{\rm s}$ of the SNRs are a rough proxy of age $t$, since $R_{\rm s} \propto t^{m}$, where $m$ is the expansion parameter, and the shock velocity $v_{\rm s}$ is given by $v_{\rm s} = mR_{\rm s}/t$. The value of $m$ depends on the evolutionary stage of the SNR. During free expansion, $m \sim 1$, and as the shock begins to decelerate, $m \sim 0.6-0.8$ \citep{chevalier82a,chevalier82b}. Once the shock has swept-up a mass $M_{\rm sw}$ that is comparable to the mass of the ejecta $M_{\rm ej}$, then the SNR enters the Sedov-Taylor (ST) phase, when $m = 0.4$ \citep{sedov59}. Subsequently, $m$ decreases from $m = 0.33$ in the pressure-driven snowplow stage (e.g., \citealt{blondin98}) to $m = 0.25$ in the momentum-conserving stage \citep{cioffi88}. \begin{deluxetable}{lrrrrc} \tablenum{5} \tablecaption{Ages of the THOR SNRs\label{tab:ages}} \tablehead{ \colhead{Source} & \colhead{$R_{\rm s}$\tablenotemark{a}} & \colhead{$n_{0}$} & \colhead{$M_{\rm SW}$} & \colhead{$t_{\rm kyr}$} & \colhead{References\tablenotemark{d}} \\ \colhead{} & \colhead{(pc)} & \colhead{(cm$^{-3}$)} & \colhead{($M_{\sun}$)} & \colhead{(kyr)} & \colhead{}} \startdata G15.9$+$0.2 & 7.4$\pm$3.5 & 0.7 & 41 & 2.9\tablenotemark{d} & 1 \\ G16.7$+$0.1 & 5.8$\pm$2.3 & 1.0\tablenotemark{b} & 28 & 1.5 & -- \\ G18.1$-$0.1 & 7.4$\pm$0.2 & 0.6 & 35 & 4.4\tablenotemark{d} & 2 \\ G18.8$+$0.3 & 28.1$\pm$8.1 & 1.0\tablenotemark{b} & 3211 & 76\tablenotemark{b} & -- \\ G20.0$-$0.2 & 16.3$\pm$0.4 & 1.0\tablenotemark{b} & 627 & 19\tablenotemark{b} & -- \\ G20.4$+$0.1 & 9.1$\pm$4.7 & 1.0\tablenotemark{b} & 109 & 4.5\tablenotemark{b} & -- \\ G21.5$-$0.1 & 6.2$\pm$2.9 & 1.0\tablenotemark{b} & 34 & 1.7\tablenotemark{b} & -- \\ G21.8$-$0.6 & 15.1$^{+0.0}_{-0.9}$ & 1.0\tablenotemark{b} & 498 & 16\tablenotemark{b} & -- \\ G22.7$-$0.2 & 16.6$\pm$1.5 & 1.0\tablenotemark{b} & 662 & 20\tablenotemark{b} & -- \\ G23.3$-$0.3 & 17.3$\pm$15.7 & 4.0 & 2997 & 45 & 3 \\ G27.4$+$0.0 & 3.4$\pm$0.2 & 0.6 & 3.4 & 1.1\tablenotemark{d} & 4 \\ G28.6$-$0.1 & 15.4$\pm$0.5 & 0.2 & 106 & 7.5\tablenotemark{d} & 5 \\ G29.6$+$0.1 & 7.3$\pm$3.7 & 1.0\tablenotemark{b} & 40 & 2.6\tablenotemark{c} & -- \\ G31.9$+$0.0 & 6.2$\pm$0.3 & 2.0 & 69 & 3.5\tablenotemark{d} & 6 \\ G32.4$+$0.1 & 14.8$\pm$3.5 & 1.0\tablenotemark{b} & 469 & 15 & -- \\ G32.8$-$0.1 & 11.9$\pm$9.9 & 0.1 & 24 & 4.2\tablenotemark{d} & 7 \\ G33.2$-$0.6 & 22.3$\pm$10.5 & 1.0\tablenotemark{b} & 1605 & 43 & -- \\ G33.6$+$0.1 & 5.1$\pm$0.4 & 0.4 & 7.7 & 3.0\tablenotemark{d} & 8 \\ G34.7$-$0.4 & 13.5$\pm$1.4 & 5.0 & 1780 & 10 & 9 \\ G35.6$-$0.4 & 6.8$\pm$0.8 & 1.0\tablenotemark{b} & 46 & 2.2 & -- \\ G36.6$-$0.7 & 30.9$\pm$14.5 & 1.0\tablenotemark{b} & 4269 & 96 & -- \\ G49.2$-$0.7 & 23.6$\pm$2.6 & 0.1 & 190 & 18\tablenotemark{d} & 10 \enddata \tablenotetext{a}{The error bars on $R_{\rm s}$ reflect the uncertainties in the source distances from Table~\ref{tab:tableTHOR}. If no uncertainties in distance are given in the literature (or if we have assumed a distance of 8.5~kpc), we assign a distance uncertainty of 4~kpc.} \tablenotetext{b}{Assumed $n_{0} = 1.0$~cm$^{-3}$ as no constraints on density were found in the literature \citep{fer01}.} \tablenotetext{c}{Assumed $E_{51} = 1.0$ as no constraints on explosion energy were found in the literature.} \tablenotetext{d}{We have scaled the age estimates from the references to the distances listed in Table~\ref{tab:tableTHOR} and the SNR radii in this table.} \tablenotetext{e}{References: (1) \citealt{reynolds06}; (2) \citealt{leahy14}; (3) \citealt{castro13}; (4) \citealt{kumar14}; (5) \citealt{bamba01}; (6) \citealt{chen04}; (7) \citealt{zhouchen11}; (8) \citealt{sun04}; (9) \citealt{reach05}; (10) \citealt{koo95}} \end{deluxetable} SNR ages are typically derived by assuming that SNRs are in the ST phase of their evolution, the stage that most radio-bright SNRs are thought to be observed \citep{berk86,berezhko04}. In this case, ages are determined based on the observed shock velocity by the expression $t = 2 R_{\rm s} / 5 v_{\rm s}$. Alternatively, given an estimate of the mass density of the ISM $\rho_{\rm o}$ and the explosion energy $E$, SNR ages can be derived using the ST solution \citep{sedov59}: \begin{equation} R_{\rm s}=1.15\bigg(\frac{E}{\rho_{\rm o}}\bigg)^{1/5}t^{2/5}=5.0~E_{51}^{1/5} n_{\rm o}^{-1/5} t_{\rm kyr}^{2/5}~\textrm{pc}. \label{eq:sedov} \end{equation} \noindent On the right-hand side of the equation, $E_{51}$ is the explosion energy in units of $10^{51}~{\rm erg}$ and $t_{\rm kyr}$ is the SNR age in kyr. To demonstrate the large uncertainties in age and the challenge of comparing those values to the power ratios, we compiled the age estimates of the THOR SNRs found in the literature (listed in Table~\ref{tab:ages}). When available (e.g., for G15.9$+$0.2, G32.8$-$0.1, G33.6$+$0.1), these values are the dynamical ages derived from the shock radius $R_{\rm s}$ and velocity $v_{\rm s}$, using the relation $t = 2 R_{\rm s} / 5 v_{\rm s}$. For those SNRs without constraints on $v_{\rm s}$, we scale the ages estimated from the ST solution in the references (see Table~\ref{tab:ages}) to the distances in Table~\ref{tab:tableTHOR} and radii in Table~\ref{tab:ages}. For those SNRs without any age estimates in the literature, we adopt the ISM densities $n_{\rm o}$ in Table~\ref{tab:ages} and assume explosion energies of $E_{51} = 1$ to estimate $t_{\rm kyr}$. To evaluate the validity of the assumption that all of the SNRs are in the ST phase, we calculated the mass swept-up $M_{\rm SW}$ by their forward shocks: \hbox{$M_{\rm SW} = \frac{4}{3} \pi R_{\rm s}^3 \times 1.4 m_{\rm H} n_{\rm o}$}, where $m_{\rm H}$ is the mass of hydrogen. The resulting $M_{\rm SW}$ for each SNR is listed in Table~\ref{tab:ages}, along with the adopted shock radii $R_{\rm s}$ and ISM densities $n_{\rm o}$ to derive $M_{\rm SW}$. For two SNRs (G27.4$+$0.0 and G33.6$+$0.1), $M_{\rm SW} < 10 M_{\sun}$, and thus their forward shocks may have swept up less than their ejecta masses $M_{\rm ej}$. This result would indicate that they may not have reached the ST phase yet, so their age estimates $t_{\rm kyr}$ in Table~\ref{tab:ages} are upper limits. Five SNRs (G18.8$+$0.3, G23.3$-$0.3, G33.2$-$0.6, G34.7$-$0.4, G36.6$-$0.7) have $M_{\rm SW} > 10^{3} M_{\sun}$ and thus may have transitioned past the ST phase, which occurs at a time $t_{\rm tr} = 2.9\times10^{4} E_{51}^{4/17} n_{\rm o}^{-9/17}~{\rm years}$ when the shock has swept up $M_{\rm SW} \approx 10^{3} E_{51}^{15/17} n_{\rm o}^{-14/17}~M_{\sun}$ \citep{blondin98}. In these cases, the age estimates $t_{\rm age}$ should be interpreted as lower-limits. In Figure~\ref{fig:age}, we plot $P_{2}/P_{0}$ (left panel) and $P_{3}/P_{0}$ (right panel) versus age $t_{\rm kyr}$ of the THOR sample. We find a weak trend that younger SNRs ($\lesssim$3~kyr old) have lower $P_{2}/P_{0}$ and $P_{3}/P_{0}$ than the older SNRs ($\gtrsim$3~kyr old), consistent with the results shown in Figure~\ref{fig:P2}. These findings suggest that SNRs' forward shocks are initially more symmetric, and their expansion into an inhomogeneous medium increases the asymmetries with time. The large dispersion in the power-ratio values in the small/young SNRs indicates that the objects may begin with different degrees of asymmetry as well, possibly reflecting their explosion geometries or the inhomogeneous environments immediately surrounding the SN. \begin{figure*} \includegraphics[width=\textwidth]{pages.pdf} \caption{Plots of $P_{2}/P_{0}$ (left) and $P_{3}/P_{0}$ (right) versus age $t_{\rm kyr}$ (listed in Table~\ref{tab:ages}), assuming the SNRs are in the Sedov-Taylor phase. Error bars on $t_{\rm kyr}$ are represent the uncertainty in the ambient density $n_{\rm o}$, which we conservatively assume to be one order of magnitude.} \label{fig:age} \end{figure*} One reason that the plots in Figure~\ref{fig:age} show less correlation compared to the power ratios versus radius in Figure~\ref{fig:P2} may be due to the large uncertainty in the SNR ages. Specifically, the explosion energies and the ISM densities are not well constrained in many cases. For many SNRs, we assumed $n_{\rm o} = 1$~cm$^{-3}$ (as an approximation of the density of the warm neutral medium: \citealt{fer01}) and $E_{51} = 1$ due to a lack of any observational constraints. Realistically, these parameters can range from $E_{51} \sim$ 0.1--10 (e.g., \citealt{sukhbold16}) and $n_{\rm o} \sim 10^{-3}$--10$^{2}$~cm$^{-3}$. In particular, $n_{\rm o}$ spans several orders of magnitude and depends on the environment of the SNR. If the SNR is expanding into a progenitor's wind bubble, then $n_{\rm o} \sim10^{-2}$--10$^{-3}$ (e.g., RCW~86: \citealt{broersen14}); if the SNR is interacting with a molecular cloud, then $n_{\rm o} \sim 10$--100~cm$^{-3}$ (e.g., gamma-ray bright SNRs: \citealt{castro10}). Generally, $n_{\rm o}$ can be estimated from modeling of X-ray or gamma-ray observations \citep{castro13}. Thus, some age estimates could be off by a factor of $\sim$10 or more, shifting their placement in Figure~\ref{fig:age}. Given that the uncertainty in $n_{\rm o}$ is responsible for a large uncertainty in the ages $t_{\rm kyr}$, we compute the error bars of $t_{\rm kyr}$ by calculating the ages for an order-of-magnitude smaller and larger $n_{\rm o}$. Consequently, the horizontal error bars in Figure~\ref{fig:age} are conservative estimates. Our results are consistent with recent three-dimensional, hydrodynamical simulations of SNRs expanding into an inhomogeneous medium \citep{kim15,martizzi15,walch15,zhang18}. In these works, the authors follow the evolution of SNRs in a multi-phase or turbulent ISM. They find that the SNRs in an inhomogeneous medium become progressively more asymmetric compared to those in a homogeneous ISM. For example, \cite{martizzi15} showed that the blast wave travels faster in the inhomogeneous medium case, particularly in areas of low-density channels around the SNR. \cite{zhang18} found that the mean ambient density is the primary factor influencing SNRs' evolution and that a smoother (lower Mach number) turbulent structure leads to faster, more asymmetric expansion\footnote{We note that \cite{martizzi15} and \cite{zhang18} had different results on the impact of turbulent structure on SNR expansion and morphology. \cite{martizzi15} showed that SNRs expand faster in a more turbulent medium, while \cite{zhang18} found that smoother turbulent structure leads to faster shock expansion. \cite{zhang18} attributes the disparity to how each study models turbulence: \cite{martizzi15} used a lognormal density distribution, whereas \cite{zhang18} adopted an initial Gaussian velocity perturbation in a uniform medium that grows with time.}. Thus, the increase in SNR asymmetries with radius (Figure~\ref{fig:P2}) and with age (Figure~\ref{fig:age}) reflect the inhomogeneous, turbulent structure of the ISM. We note that \cite{lopez09,lopez11} found no age evolution in the power-ratios derived from X-ray images of SNRs. However, the soft X-rays trace thermal emission from SNR ejecta, so SNR X-ray morphologies may reflect explosion asymmetries. Furthermore, these previous studies considered only young sources, with ages $t_{\rm kyr}\lesssim$10, whereas our sample spans a wider range, with $t_{\rm kyr}\sim$1--100. Further investigation is necessary to determine if size/age evolution is unique to the SNRs' radio continuum morphologies. To compute ages, we assumed a uniform ambient density, though our results indicate that the SNRs are expanding into inhomogeneous environments (since a homogeneous ISM would lead to no size/age evolution in the asymmetries). \cite{zhang18} showed that the mean ambient density is the dominant factor in determining shock expansion with time, though inhomogeneities are important in shaping SNRs overall. Thus, the assumption of a single $n_{\rm o}$ may be sufficient. However, given the large uncertainties in the SNR ages, the radii are the best observable indicator of SNR evolutionary stage. In this work, we have analyzed nearly one-third of the SNR population of the Milky Way. In the future, application and comparison to extragalactic SNRs may reveal differences in the turbulent structure of nearby galaxies. While the fractal nature of the Milky Way ISM \citep{elmegreen96} is also observed in nearby galaxies (e.g., in the Small Magellanic Cloud: \citealt{stan99}), the medium can differ substantially, e.g., in porosity \citep{bagetakos11} or in molecular gas velocity dispersion \cite{sun18}. Ultimately, comparison of SNRs' morphological evolution with simulations (e.g., \citealt{kim15,martizzi15,walch15,zhang18}) may place new constraints on the ISM properties of the Milky Way and nearby galaxies. | 18 | 8 | 1808.08234 |
1808 | 1808.05536_arXiv.txt | {} {With the aim of performing a suitable comparison of the internal process of galactic bars with respect to the external effect of interactions on driving gas toward the inner most region of the galaxies, we explored and compared the efficiency of both mechanisms on central nuclear activity in optically selected active galactic nuclei (AGNs) in spiral galaxies.} {We selected homogeneous samples of barred AGNs and active objects residing in pair systems, derived from the Sloan Digital Sky Survey (SDSS). In order to carry out a reliable comparison of both samples (AGNs in barred hosts in isolation and in galaxy pairs), we selected spiral AGN galaxies with similar distributions of redshift, magnitude, stellar mass, color and stellar age population from both catalogs. With the goal of providing an appropriate quantification of the influence of strong bars and interactions on nuclear activity, we also constructed a suitable control sample of unbarred spiral AGNs without a companion and with similar host properties to the other two samples.} {We found that barred optically selected AGNs show an excess of nuclear activity (as derived from the $Lum[OIII]$) and accretion rate onto a central black hole ($\cal R$) with respect to AGNs in pairs. In addition, both samples show an excess of high values of $Lum[OIII]$ and $\cal R$ with respect to unbarred AGNs in the control sample. We also found that the fractions of AGNs with powerful nuclear activity and high accretion rates increase toward more massive hosts with bluer colors and younger stellar populations. Moreover, AGNs with bars exhibit a higher fraction of galaxies with powerful $Lum[OIII]$ and efficient $\cal R$ with respect to AGN galaxies inhabiting pair systems, in bins of different galaxy properties. Regarding AGNs belonging to pair systems, we found that the central nuclear activity is remarkably dependent on the galaxy pair companion features. The $Lum[OIII]$ for AGNs in pairs is clearly enhanced when the galaxy companion exhibits a bright and more massive host with high metallicity, blue color, efficient star formation activity and young stellar population. The results of this work reveal an important capacity of both mechanisms, bars and interactions, to transport material towards the galaxy central regions. In this context, it should also be noted that the internal process of the bar is more efficient at improving the central nuclear activity in AGN objects than that corresponding to the external mechanism of the galaxy-galaxy interactions.} {} | The most accepted hypothesis surrounding the origin or active galactic nuclei (AGNs) proposes that they arise from accretion of material onto a central massive black hole triggering nuclear activity. It is widely known that the main fueling mechanisms of the central engine in galaxies are related to dynamical perturbations transporting gas to the inner most central regions \citep{LB69,Rees84}. In this sense, several authors agree that galactic bars and galaxy mergers/interactions are usually considered the two principal processes for torquing material to the centers of active galaxies (e.g., Mihos \& Hernquist 1996; Combes 2003; Alonso et al. 2007, Di Matteo et al. 2005; Alonso et al. 2013, 2014). Bars play a fundamental role in the dynamical evolution of their host galaxies and can also affect several properties of spirals on relatively short timescales. In this context, bar perturbations can modify star formation activity, stellar populations, colors, chemical composition and even galactic structure \citep{atha83,buta96,comb93,cheu13,martin95,robi17,vera16}, promoting evolution of their host galaxies (Ellison et al. 2011a, Zhou et al. 2015). Moreover, the gas infall produced by bars toward the innermost regions of galaxies is a mechanism that may efficiently trigger nuclear activity in the central zone of AGN galaxies \citep{cor03,comb93,ju18}. Different studies based on numerical simulations show a loss in galaxy angular momentum produced by interactions between gas clouds and the edges of the bar, driving a flow of material toward the central regions of barred galaxies \citep{SBF90}. \\ Following this line, \cite{pet18} analyzed different mass models of spiral galaxies under the influence of tidal interactions finding no strong correlation between bar length or pattern speed and the interaction strength. However, these authors show that interactions slightly accelerate bar formation in some models. On the other hand, there is a slower disk-dominated rotation-curve model likely due to interactions of gas clumps. In agreement with this point, \cite{zana18} carried out a study of external versus internal bursts of bar formation, finding a strong dependence on the mass of the disk.\\ Furthermore, the ``bars within bars'' scenario states that an external bar transports material over distances of a few parsecs. In this region an internal secondary bar produces gravitational instability in the accumulated gas, enabling flow toward regions near the massive black hole \citep{SBF89}. Supporting this theory, several authors using different observational techniques have observed secondary bars inside strong external bars \citep{emse01,malk98,laine02,caro02,mac2000}. In this context, in a recent study, \cite{du17} suggest that these short bars, generally embedded in large-scale bars, are an important mechanism for driving gas inflow, feeding the central black hole. Nevertheless, this mechanism becomes unstable, and inner bars are destroyed when the black hole mass grows to $\sim0.1\%$ of the total stellar mass. This event slows down, or even stops, the growth of a central black hole. There is clear observational evidence that bars enhance central nuclear activity compared to non-barred spiral galaxies: \cite{oh12} found that bars produce an increment in the central activity of blue galaxies with low black hole masses from a sample of barred late-type AGNs. Furthermore, Alonso et al. (2013, hereafter A13) show that isolated barred AGN galaxies display a higher fraction of powerful nuclear activity, in comparison with a suitable control sample of unbarred AGNs with similar distributions of redshift, magnitude, morphology and local environment. They also found that barred AGN galaxies show an excess of objects with high accretion rates in comparison to unbarred ones. In addition, from the analysis of barred AGN spiral galaxies inhabiting groups and clusters, Alonso et al. (2014) found that the increment of nuclear activity produced by bar perturbations is also notable in barred active galaxies located in higher-density environments. Recently, \cite{gal15}, using a sample of disk galaxies from Sloan Digital Sky Survey (SDSS) and Galaxy Zoo 2, found a higher fraction of barred AGNs than of star-forming barred galaxies, although the central black hole accretion rate shows no dependence on the presence of a bar. \cite{cheu15} deepened this study towards high redshifts concluding that large-scale bars cannot be considered the dominant fueling mechanism for supermassive black hole growth; see also \cite{gou17}. Galaxy interactions can be an effective mechanism to modify different host galaxy properties, mainly by triggering star formation \citep{alo06,alo12,barton,Lam03, Lam12,kenni,mesa}, and also affect the galaxy stellar mass function \citep{gama}. The presence of a close galaxy companion drives a clear enhancement in galaxy morphological asymmetries, and this effect is statistically significant up to projected separations of at least 50 $\kpc$ \citep{patton16}. The physical processes behind galaxy-galaxy interactions have been explained by theoretical and numerical analyses (e.g., Martinet 1995; Toomre \& Toomre 1972; Barnes \& Hernquist 1992, 1996; Mihos \& Hernquist 1996), showing that collisional disruption, material dissipation, and gas inflows are produced by the tidal torques generated during near encounters. In addition to feeding star formation, these material inflows could also feed a central black hole and increase nuclear activity (Sanders et al. 1988). The performance of this process depends on the gas reservoir and the particular internal characteristics of galaxies involved in the interaction. Based on observational evidence, the connection between mergers/interactions and nuclear activity is fairly well accepted. In this sense, several studies have found clear clues of interactions in luminous quasar hosts (e.g., Canalizo \& Stockton 2001; Bennert et al. 2008; Urrutia, Lacy \& Becker 2008; Ramos Almeida et al. 2011; Bessiere et al. 2012; Urrutia et al. 2012). In addition, different analyses have found a clear increase of the nuclear activity in less-luminous AGNs with tidal interaction features or distorted morphologies with respect to non-interacting AGN galaxies (e.g., Koss et al. 2010, 2012; Ellison et al. 2011a, 2013; Silverman et al. 2011; Sabater et al. 2013). In particular, Alonso et. al (2007) performed a statistical analysis of nuclear activity comparing AGN galaxies in 1607 close pair systems to AGNs without companions. They found that the nuclear activity of active galaxies with strong interaction features is significantly larger than for AGNs in an isolated environment. The accretion rate also shows that AGNs in merging pairs are actively feeding their central black holes. \cite{mesa} constructed a sample of spiral galaxy pairs from SDSS, and classified them according to the spiral arms' rotation pattern, detecting an increment in the nuclear activity in systems whose spiral arms rotate in opposing directions. More recently, Sabater et al. (2015) found that galaxy interactions affect AGN activity in an indirect way, by influencing the central gas supply. In this sense, these studies provide obvious clues about AGN fueling and its link with galactic mergers and interactions. In addition, \cite{bar17} asserts that the enhancements in specific star formation rates (SFRs) are positively correlated with enhanced AGN luminosity, suggesting that both values are mutually triggered by the merger events, the latter being significantly greater than the former. An interesting approach is to study and compare the role of bars and galaxy interactions in feeding central black holes. Motivated by this, we analyze the effect of the internal process of bars in comparison with the external mechanism of galaxy interactions on the central nuclear activity in AGN galaxies. For this purpose, using data from the Sloan Digital Sky Survey, we obtain large and homogeneous samples of barred AGN galaxies and AGNs in pair systems required to derive a direct and consistent comparison of these two mechanisms. The conclusions of these studies will allow us to expand on what is currently known about the governing mechanism of the induction of radial gas inflow to galactic centers. This paper is structured as follows. Section 2 describes the procedure used to construct the samples of barred AGN galaxies, AGNs in pairs, and control AGN galaxies. In Section 3, we study the effects of bars and galaxy interactions on nuclear activity and the relation with the host galaxy properties. Section 4 explores the role of the galaxy pair companion in feeding central black holes, and in Sect. 5 we summarize our main conclusions. The cosmology adopted here is $\Omega = 0.3$, $\Omega_{\lambda} = 0.7$, and $H_0 = 100~ \kms \rm Mpc$. | We have performed a comparative analysis of the effect of bars and mergers/interactions on the central nuclear activity of spiral AGN galaxies. We acknowledge that different AGN selection criteria can lead to the selection of galaxies with different SFR activity. Low-luminosity radio-selected AGN are significantly biased towards low SFR values compared to IR selected galaxies in the SDSS. Ellison et al. (2016) suggest the dominance of mergers in IR selected AGNs, a lower merger incidence amongst optically selected AGNs, and that secular fuelling dominates low-excitation radio galaxies (LERGs). Our study is based on homogeneous samples of optically selected AGNs with strong bars, in relative isolation or in pair systems. In order to carry out a suitable comparison of the effects of bars and mergers/interactions, we selected AGN spiral galaxies in both samples with similar redshift, $g-$band apparent magnitude, absolute $r-$band magnitude, stellar mass, color, and stellar age population distributions. To obtain an appropriate quantification of the effect of the two processes (bars and interactions) on the nuclear activity, we also constructed a suitable control sample of unbarred spiral AGN galaxies, without a pair companion, with similar host properties to those of the other samples. The main results and conclusions of our analysis are summarized as follows. We found that barred active galaxies show an excess of nuclear activity compared to AGN galaxies in pair systems. Moreover, both samples show an excess of high $Lum[OIII]$ values with respect to unbarred spiral active galaxies in the control sample. We also analyzed the accretion strength onto a central black hole for AGN host galaxies in the different samples. From this study, we conclude that barred active galaxies have an excess of objects with high accretion rate values with respect to AGN hosts inhabiting pair systems. Furthermore, active galaxies in both samples exhibit higher $\cal R$ values than the control active galaxies. We studied the fraction of powerful AGN galaxies ($Lum[OIII] > 10^{6.4} L_{\odot}$) as a function of stellar mass, color, and stellar age population. We found that the number of active galaxies with efficient central nuclear activity increases when selecting host galaxies with larger stellar mass, bluer colors, and younger stellar populations. From this analysis, we also show that the fraction of $Lum[OIII] > 10^{6.4} L_{\odot}$ for the AGNs in pairs is slightly higher than that of the AGNs in the control sample, throughout the whole range of $log(M^*)$, $(M_u-M_r)$ and $D_n(4000)$. Concurrently, we found that barred active objects show a moderately higher fraction of powerful AGNs with respect to the other AGN samples, in bins of different galaxy properties. We also explored the fraction of AGN galaxies with high accretion rate ($\cal R$$>-0.6$) as a function of the host galaxy properties. We found that the fraction of active galaxies with $\cal R$$>-0.6$ increases towards more massive hosts with bluer colors and younger stellar populations. From this analysis we also show that barred AGN host objects exhibit a slightly higher fraction of efficient accretion rate with respect to the other samples, in bins of different galaxy properties. We also found that AGN inhabiting pair systems display higher $\cal R$$>-0.6$ values in comparison with active galaxies in the control sample. Furthermore, we found that the smallest black holes exhibit a higher fraction of $\cal R$$>-0.6$, and, in this context, barred AGN galaxies show a moderate excess of efficient accretion rate with respect to active galaxies in the other samples, throughout the whole range of the $log(M_{BH})$. For AGNs inhabiting pair systems, we found that the nuclear activity is remarkably dependent on the galaxy pair companion properties. In this context, we show that the nuclear activity of an AGN companion presents a noticeable increment when its galaxy host is brighter and more massive, with higher metallicity, bluer colors, more efficient star formation activity and a young stellar population. We also explored the fraction of AGNs in pair systems with higher OIII luminosity as a function of $log(M^*)$, considering galaxy companions with different host features separately. We found that the fraction of $Lum[OIII] > 10^{6.4} L_{\odot}$, for AGNs belonging to pair systems, is similar to that of barred active galaxies when galaxy pair companions present bright and massive hosts with high metallicity, blue colors, efficient star formation activity and a young stellar population. Conversely, the fraction of powerful nuclear activity is similar for AGNs in pairs and for AGNs in the control sample, when the galaxy pair companions present low masses, luminosities and metallicities, red colors, less star formation activity and an old stellar population. The results found in this work suggest an important effect of both bars and interactions in driving radial gas inflows towards the innermost regions of galaxies. These mechanisms produce an enhancement in nuclear activity and accretion rate of central black holes in spiral active nuclei galaxies. It should also be noted that the internal process of the bar perturbation presents a more effective transport of the gas flow to the central zones in comparison with the external mechanism of the mergers and interactions. Furthermore, the impact of the AGNs belonging to pair systems on the central nuclear activity is noticeably influenced by the galaxy pair companion properties. | 18 | 8 | 1808.05536 |
1808 | 1808.10735_arXiv.txt | The LOFAR radio telescope creates Petabytes of data per year. This data is important for many scientific projects. The data needs to be efficiently processed within the timespan of these projects in order to maximize the scientific impact. We present a workflow orchestration system that integrates LOFAR processing with a distributed computing platform. The system is named Automated Grid-enabled LOFAR Workflows (AGLOW). AGLOW makes it fast and easy to develop, test and deploy complex LOFAR workflows, and to accelerate them on a distributed cluster architecture. AGLOW provides a significant reduction in time for setting up complex workflows: typically, from months to days. We lay out two case studies that process the data from the LOFAR Surveys Key Science Project. We have implemented these into the AGLOW environment. We also describe the capabilities of AGLOW, paving the way for use by other LOFAR science cases. In the future, AGLOW will automatically produce multiple science products from a single dataset, serving several of the LOFAR Key Science Projects. Target: IEEE International Conference On e-Science Keywords: Workflow management software, Radio Astronomy, Distributed Computing, Big Data applications | \label{sec:intro} Data sets in radio astronomy have increased 1000-fold over the past decade\cite{sabater_datasize}. It is no longer feasible to move, store and process these data sizes at university clusters, nor to process these data manually. LOFAR, the Low-Frequency Array\cite{LOFAR} is a modern and powerful radio telescope that creates more than 5 Petabytes of data per year. At present, the majority of LOFAR time is allocated to several Key Science Projects (KSPs)\cite{lotss}. These projects need to process hundreds or thousands of observations. Typical observations produce approximately 14 TB of archived data. Obtaining high fidelity images from this data requires complex processing steps. To manage and automate the data processing, workflow management software is needed. This software needs to accelerate LOFAR processing on a High Throughput Computing (HTC) cluster while ensuring it is easy to prototype, test, and integrate future algorithms and pipelines. To automate LOFAR data processing, we have worked with the LOFAR Surveys KSP (SKSP). Together, we designed a software suite that integrates LOFAR software\cite{cookbook} with the Dutch grid infrastructure\cite{dutchinfra}. This software, based on Apache Airflow\footnote{\url{https://airflow.apache.org/}}, makes it easy to add future science cases, extend and modify pipelines, include data quality checks, and rapidly prototype complex pipelines. For the SKSP use cases, AGLOW achieves a significant reduction in development time: from months to days, allowing researchers to concentrate on data analysis rather than management of processing. Additionally, and perhaps more importantly, the software versions and repositories used are defined within the workflow. This makes reproducibility an integral part of the AGLOW software. Finally, the software is built to leverage an HTC cluster by seamlessly submitting the processing jobs through the cluster's job submission system\cite{glite}. The work presented here builds on our previous work parallelizing single LOFAR jobs\cite{mechev} on a distributed environment. The majority of processing was done at SURFsara at the Amsterdam Science Park\cite{SurfSara}, which is one of the sites used by the LOFAR Long Term Archive (LTA)\footnote{\url{https://lta.lofar.eu}}. Ongoing efforts include scheduling and processing data at clusters in Pozna\`{n} in Poland and J\"{u}lich in Germany. \textbf{Contributions:} The main features of the AGLOW software are the following: \begin{itemize} \item Integration of the Grid middleware with Apache Airflow, allowing us to dynamically define, create, submit and monitor jobs on the Dutch national e-infrastructure. \item Integration of the LOFAR (LTA) utilities in Airflow, facilitating pipeline developers to automate staging (moving from tape to disk) and retrieval of LOFAR data. \item Integration of the SURFsara storage with Airflow, making LOFAR pipelines aware of the storage layer available at the Dutch national e-infrastructure. \item Ease of creating simple software blocks, with which users can integrate and test their pipelines. \item Storing all software versions and script repositories as part of the workflow to make LOFAR processing reproducible and portable. \end{itemize} \textbf{Outline:} The organization of this manuscript is as follows: We provide background on data processing in radio astronomy and why LOFAR science requires complex workflows and cover workflow management algorithms and capabilities (section \ref{sec:background}). We discuss related work in workflow management (section \ref{sec:related}). In section \ref{sec:AGLOW}, we introduce our software and two use cases. Both of our use cases require acceleration at an HTC cluster and automation by a workflow orchestration software. We follow these examples with details on the integration between LOFAR software, LOFAR data and the resources at SURFsara in Amsterdam in section \ref{sec:extending}. Finally, we discuss our results (sect. \ref{sec:results}) and look ahead to the demands of future LOFAR projects and upcoming telescopes in section \ref{sec:conclusions}. | \label{sec:conclusions} In this work, we have detailed a comprehensive workflow management software for processing radio astronomy data on a distributed infrastructure. We leverage an industry standard workflow management software, Airflow. Using its capabilities, we make it possible to build, test, automate and deploy LOFAR pipelines on short timescales, generally from months to days. With the flexibility of Airflow's Python and Bash operators, users can design their own workflows, as well as co-ordinate more complex science cases. In this way, AGLOW facilitates reproducible processing of scientific data. In the future, AGLOW will support additional LOFAR science cases including Long Baselines and Spectroscopy. In this article, we have described our implementation of the data processing pipelines used by the LOFAR Surveys Key Science Project. Future work includes further de-coupling of the Grid-setup and pipeline logic. We will do this by creating `sub-dags' (details in \ref{sec:Subdags}) for each type of LOFAR jobs. Using these sub-dags will reduce the complexity of scientific workflows while also making the code even more reusable and thus easier to maintain and upgrade. Efforts to integrate processing at the other two LTA sites, (J\"{u}lich and Pozna\'{n}) have already started with `prefactor' runs being performed on J\"{u}lich using a modified version of the SKSP workflow. The software also currently works on the Eagle cluster at Pozna\'{n}. Combining the J\"{u}lich and SURFsara workflows will be done in the future so that AGLOW can track and start processing at multiple clusters. Finally, AGLOW can be used as a `LOFAR As A Service' model. In this model, users only provide an observation ID and processing parameters and receive the final results upon job completion. This model will build upon previous success offering LOFAR processing to users without login to the \texttt{GINA} cluster \cite{oonk_prep}. This previous work was already useful for studying radio absorption in Cassiopeia A \cite{maria} and a `data-to-images' service will be valuable to the whole LOFAR community. Our experience with automating LOFAR scientific workflows on a distributed architecture will be valuable when setting up data processing for future Radio Telescopes such as the Square Kilometer Array \cite{ska} . | 18 | 8 | 1808.10735 |
1808 | 1808.09232_arXiv.txt | \rxj ~($z=0.451$) is one of the most luminous X-ray galaxy clusters, which hosts a prominent cool core and exhibits a signature of a major merger. We present the first {\it direct} observational evidence for sub-sonic nature of sloshing motion of the cool core. We find that a residual X-ray image from the {\it Chandra X-ray Observatory} after removing the global emission shows a clear dipolar pattern characteristic of gas sloshing, whereas we find no significant residual in the Sunyaev-Zel'dovich effect (SZE) image from the Atacama Large Millimeter/submillimeter Array (ALMA). We estimate the equation of state of perturbations in the gas from the X-ray and SZE residual images. The inferred velocity is $420^{+310}_{-420}$\,km\,s$^{-1}$, which is much lower than the adiabatic sound speed of the intracluster medium in the core. We thus conclude that the perturbation is nearly isobaric, and gas sloshing motion is consistent with being in pressure equilibrium. Next, we report evidence for gas stripping of an infalling subcluster, which likely shock-heats gas to high temperature well in excess of 20\,keV. Using mass distribution inferred from strong lensing images of the {\it Hubble Space Telescope} (\HST), we find that the mass peak is located away from the peak position of stripped gas with statistical significance of $> 5\sigma$. Unlike for the gas sloshing, the velocity inferred from the equation of state of the excess hot gas is comparable to the adiabatic sound speed expected for the 20\,keV intracluster medium. All of the results support that the southeast substructure is created by a merger. On the other hand, the positional offset between the mass and the gas limits the self-interaction cross section of dark matter to be less than $3.7~h^{-1}~{\rm cm^2~g^{-1}}$ (95\%~CL). | Galaxy clusters are the largest gravitationally-bound and virialized objects in the universe. They are located at the knots of filaments in the large-scale structure and provide us with unique cosmological information. Galaxy clusters are also dynamically young and are continuously growing through mergers between smaller clusters. Such merger activities induce various high-energy phenomena in the hot and optically thin plasma, i.e., the intracluster medium (ICM), in the gravitational potential well of clusters. A large amount of observational and numerical studies have suggested that mergers lead to shock-heating of the gas \citep[e.g.,][]{Ricker01, Markevitch02, Takizawa05, Takizawa06, Bourdin13}, production of cold fronts \citep[e.g.,][]{Markevitch01, Ascasibar06, ZuHone10, Roediger11, Blanton11, Ueda17, Hitomi18d}, ram-pressure stripping of the gas from infalling galaxies \citep[e.g.,][]{David04, Roediger07, Sasaki16}, non-equilibrium ionization of the ICM \citep[e.g.,][]{Akahori08, Akahori10, Akahori12, Inoue16}, and (re-) acceleration of relativistic particles \citep[e.g.,][]{Feretti12, Akamatsu13b, Brunetti14, van_Weeren17}. The entire picture of merging processes is however still far from clear; for example, heating mechanisms and dynamics of the ICM are under debate. Observational studies in multi-wavelength are therefore crucial for understanding the physics of galaxy cluster mergers. \rxj ~is one of the most luminous X-ray galaxy clusters and is located at the redshift of $z = 0.451$. It was thought to be a relaxed cluster when it was discovered in the {\it ROSAT} all sky survey \citep{Schindler97}. \cite{Komatsu99} made the first measurements of the Sunyaev-Zel'dovich effect \citep[SZE:][]{Sunyaev72} toward this cluster with the James Clerk Maxwell Telescope (JCMT) at 350\,GHz as well as with the 45\,m Nobeyama Radio Telescope at 21 and 43\,GHz. A higher angular resolution observation of the SZE was performed by \cite{Komatsu01} using the Nobeyama Bolometer Array (NOBA) and they found a prominent substructure which has no counterpart in the soft X-ray image by {\it ROSAT}. The presence of the substructure has been confirmed by \Chandra ~and \XMM ~\citep[e.g.,][]{Allen02, Gitti04} as well as by more recent SZE measurements \citep[][]{Mason10, Korngut11, Plagge13, Adam14, Kitayama16}. \cite{Allen02} measured the mean temperature of the ICM to be over 10\,keV, which is relatively high compared to other typical clusters. \cite{Kitayama04} and \cite{Ota08} found a very hot ($> 20$\,keV) component of the ICM in this cluster. In addition, the radial profile and spatial distribution of the ICM temperature indicate that the temperature drops to $\sim 6$\,keV toward the cluster center so that the cool core is formed \citep[e.g.,][]{Allen02, Ota08, Kreisch16}. A disturbed morphology is furthermore supported by the radio synchrotron observations \citep[e.g.,][]{Ferrari11} and the gravitational lensing maps \citep[e.g.,][]{Kohlinger14}. The total mass of \rxj ~within $r_{200}$ is estimated to be $\sim 1.5 \times 10^{15}\,h^{-1}$\,\MO ~by using the weak-lensing analysis, where $r_{200}$, the radius within which the mean mass density is 200 times the critical density of the universe, is $1.85\,h^{-1}$\,Mpc \citep{Lu10} for this galaxy cluster\footnote{They adopted the Hubble constant of 70\,km\,s$^{-1}$\,Mpc$^{-1}$.}. It is considered that the hot component is most likely associated with a past major merger event \citep[e.g.,][]{Johnson12, Kreisch16}, although its specific nature, such as geometry and dynamics of the collision, is still unclear. Recently, \cite{Kitayama16} presented the SZE image observed by Atacama Large Millimeter/submillimeter Array (ALMA) with angular resolution of 5\arcsec. Such high angular resolution enables us to remove the emission of the central active galactic nucleus (AGN) and to reconstruct an accurate SZE map. \cite{Kitayama16} found that the shape of the SZE is elongated toward the southeast and the peak position of the SZE is located at 11\arcsec ~southeast from the central AGN. The ALMA high-resolution SZE image motivates us to directly compare it with high-quality data in other wavelengths by \Chandra ~and {\it Hubble Space Telescope} (\HST). In this paper, we investigate the merger phenomena and merger history in \rxj ~by combining the data of \Chandra, ALMA, and \HST. We adopt $\Omega_{\rm m}=0.3$ and $\Omega_{\rm \Lambda}=0.7$. We use the dimensionless Hubble constant ($h \equiv H_{0}/100\,$km\,s$^{-1}$\,Mpc$^{-1}$); given controversial results on the value of $h$ \citep[e.g.,][]{Planck16, Riess16}, we do not fix it unless stated otherwise. In this cosmology, an angular size of $1''$ corresponds to a physical scale of 4.04\,$h^{-1}$\,kpc at the redshift $z=0.451$. Unless stated otherwise, quoted errors correspond to 1$\sigma$. | We have studied \rxj, one of the well-known major merging clusters, combining the high angular resolution, multi-wavelength data taken by \Chandra, ALMA, and \HST. The conclusions of this paper are summarized as follows. \begin{itemize} \item The residual image of the X-ray surface brightness shows a clear dipolar pattern in the cluster center, whereas, we find no excess SZE signal in the central region in the SZE residual image. The dipolar pattern indicates that a fraction of the gas in the cool core is disturbed by gas sloshing. We have estimated the equation of state of the perturbation in gas using the X-ray and SZE residual images. We find that the inferred velocity is $420^{+310}_{-420}$\,km\,s$^{-1}$, which is much lower than the adiabatic sound speed of the 10\,keV ICM inside the core; thus, the perturbation is consistent with being isobaric. This is the first direct evidence of sub-sonic nature of gas sloshing motion. \item Both X-ray and SZE residual images show an excess component in the southeast substructure. We find that the peak of excess hot ($\sim 30$\,keV) ICM is located at $27''$ ($109\,h^{-1}$\,kpc) to the southeast from the cluster center. This region is faint in X-rays but bright in the SZE. The X-ray inferred thermal pressure of the excess component is nearly constant among the regions where the SZE signal is prominent. The velocity inferred from the equation of state of the excess hot ICM is $1970 \pm 150$\,km\,s$^{-1}$, which is comparable to the adiabatic sound speed of the 20\,keV ICM. This result supports a picture that the perturbation in the southeast is generated by shock. \item The mass distribution of \rxj ~obtained with SL is reproduced well by the two dark halo components. Their mass peaks are in good agreement with the positions of the BCG and the 2nd BCG. The mass peak of the main cluster matches the X-ray centroid, while the mass peak of infalling subcluster is offset from the substructure in X-rays and the SZE. This indicates that the gas originated in the subcluster is stripped by ram-pressure of the main cluster and shock heated during an on-going major merger. \item In our scenario, this major merger is likely in the first passage. \rxj ~is therefore an exceptional cluster in which the excess hot gas, on-going major merger, and the sloshing cool core coexist within the central $150$\,$h^{-1}$\,kpc. \item We have also constrained the self-interaction cross section of DM. The resulting upper limit of the cross section is $\sigma_{\rm DM}/m < 3.7\,h^{-1}$\,cm$^{2}$\,g$^{-1}$ (95\% CL). \end{itemize} | 18 | 8 | 1808.09232 |
1808 | 1808.00063_arXiv.txt | The rotational evolution of cool stars is governed by magnetised stellar winds which slow the stellar rotation during their main sequence lifetimes. Magnetic variability is commonly observed in Sun-like stars, and the changing strength and topology of the global field is expected to affect the torque exerted by the stellar wind. We present three different methods for computing the angular momentum loss in the solar wind. Two are based on MHD simulations from \cite{finley2018dipquadoct}, with one using the open flux measured in the solar wind, and the other using remotely-observed surface magnetograms. Both methods agree in the variation of the solar torque seen through the solar cycle and show a $30-40\%$ decrease from cycle 23 to 24. The two methods calculate different average values, $2.9\times 10^{30}$ erg (open flux) and $0.35\times 10^{30}$ erg (surface field). This discrepancy results from the already well-known difficulty with reconciling the magnetograms with observed open flux, which is currently not understood, leading to an inability to discriminate between these two calculated torques. The third method is based on the observed spin-rates of Sun-like stars, which decrease with age, directly probing the average angular momentum loss. This method gives $6.2\times 10^{30}$ erg for the solar torque, larger than the other methods. This may be indicative of further variability in the solar torque on timescales much longer than the magnetic cycle. We discuss the implications for applying the formula to other Sun-like stars, where only surface field measurements are available, and where the magnetic variations are ill-constrained. | Angular momentum loss through stellar winds explains the rotational evolution of low mass stars ($M_*\leq1.3M_{\odot}$) on the main sequence. These stars are shown to have outer convection zones \citep{marcy1984observations, donati2006large, morin2008stable, donati2008magnetic, petit2008toroidal, morgenthaler2011direct, gregory2012can, reiners2012observations, folsom2016evolution, folsom2017evolution}, which are able to support magnetic fields through the interplay of rotation and convection, forming a dynamo \citep{brun2017magnetism}. The magnetic field generation of such dynamos is linked with rotation \citep{browning2008simulations, reiners2009evidence, reiners2010volume, vidotto2014stellar, see2015energy, shulyak2017strong}, such that a faster rotator will, in general, produce a larger field strength. Stellar winds are found to be more effective at slowing rotation in the presence of these large scale magnetic field \citep{weber1967angular, mestel1968magnetic, keppens2000stellar, matt2012magnetic, garraffo2015dependence, reville2015effect}. Therefore, the relation of stellar rotation, magnetism and angular momentum loss leads to the convergence of rotation periods at late ages \citep{skumanich1972time, soderblom1983rotational, barnes2003rotational, barnes2010simple, delorme2011stellar, van2013fast, bouvier2014angular}. Observations of the rotation rates of stars at different ages, and our knowledge of stellar structure, also give us direct constraints on the total external torque on the star. This value is independent from any knowledge of the physical mechanism for that angular momentum loss, but it probes only a long-time average torque (i.e., only on timescales smaller than the spin-down time, which can be in the range of tens to hundreds of Myr for main sequence stars). With the increasing number of accurate rotation period measurements available to compare with model results \citep[e.g.][]{agueros2011factory, mcquillan2013measuring, nunez2015linking, rebull2016rotation, covey2016rapidly, douglas2017poking, agueros2017setting}, we are able to examine the physical mechanisms of stellar wind braking in greater detail \citep{irwin2009ages, bouvier2014angular}. A variety of spin evolution models have been developed to date \citep[e.g.][]{gallet2013improved, van2013fast, gallet2015improved, johnstone2015stellar, matt2015mass, amard2016rotating, sadeghi2017semi,see2018open}, which relate basic stellar properties; mass, radius, rotation period, field strength and mass loss rate, with results from analytic or numerical models for the spin down torque applied to the star, and the subsequent redistribution of internal angular momentum. Stellar mass and radius remain essentially constant throughout the main sequence. However, in addition to the long-time secular changes of the magnetic field due to rotation, magnetic activity is also observed to vary significantly over timescales of years to decades \citep{baliunas1995chromospheric, azizi2017survey}. This is routinely observed for the Sun which is known to have a magnetic activity cycle \citep{babcock1961topology, wilcox1972annual, willson1991sun, guedel1997x, gudel2007sun, schrijver2008global}, moving from an activity maximum through minimum and back to maximum in roughly 11 years. The Sun's cyclic behaviour is apparent in changes to the large scale magnetic field \citep{derosa2012solar}, which significantly modifies the solar wind structure and outflow properties \citep{smith1995ulysses, mccomas2000solar, wang2000long, tokumaru2010solar}. Activity cycles on other stars are quantified using activity proxies such as the long term monitoring of Ca II $HK$ emission \citep{baliunas1995chromospheric, egeland2017mount}, observed lightcurve modulation due to star spots \citep{lockwood2007patterns}, X-ray activity \citep{hempelmann1996coronal} and more recently Zeeman Doppler Imaging, ZDI \citep{semel1989zeeman, donati1989zeeman, brown1991zeeman, donati1997zeeman}. The mass loss rate of the Sun is shown to vary with the magnetic cycle \citep{mccomas2013weakest} and is fundamentally connected with magnetic activity \citep{cranmer2007self}. This behaviour is expected to be similar for other low mass stars. Previous theoretical studies have shown the variation in angular momentum loss over magnetic cycles \citep{pinto2011coupling, garraffo2015dependence, reville2015solar, alvarado2016simulating, reville2017global}. However they require costly MHD simulations which attempt to simultaneously fit the mass loss rate and magnetic field strengths for single epochs. In contrast, by utilsing stellar wind braking formulations from \cite{reville2015effect}, \cite{finley2017dipquad}, \cite{pantolmos2017magnetic} and \cite{finley2018dipquadoct}, hereafter FM18, which can easily predict the torque for any known mass loss rate and magnetic field strength/geometry. This allows, for the first time, a more continuous calculation of the angular momentum loss rate. Using the multitude of current observations of the Sun (this work), and multi-epoch studies of other stars from the ZDI community (Paper II), we can now evaluate the variation of stellar wind torques over decadal timescales. We briefly reiterate the angular momentum loss prescriptions from FM18 in Section 2, collate solar observations in Section 3, and implement them in Section 4 to produce the most up-to-date determination of the solar braking torque, using methods based on the surface magnetogram data obtained from SOHO/MDI and SDO/HMI, and evaluating the open magnetic flux from the \textit{Ulysses} and the Advanced Composition Explorer (ACE) spacecrafts, along with an estimate based on the rotational behaviour of other Sun-like stars. Section 5 then discusses our result and addresses the observed discrepancy between surface field and open flux methods, along with the differences between our torque value and the derived long-time average result.% Previous theoretical studies have shown the variation in angular momentum loss over magnetic cycles \citep{pinto2011coupling, garraffo2015dependence, reville2015solar, alvarado2016simulating, reville2017global}. However they require costly MHD simulations which attempt to simultaneously fit the mass loss rate and magnetic field strengths for single epochs. By contrast, using the stellar wind braking formulations from \cite{reville2015effect}, \cite{finley2017dipquad}, \cite{pantolmos2017magnetic} and \cite{finley2018dipquadoct}, hereafter FM18, which can easily predict the torque for any known mass loss rate and magnetic field strength/geometry, without need for new simulations. This allows, for the first time, a more continuous calculation of the angular momentum loss rate. Using the multitude of current observations of the Sun (this work), and multi-epoch studies of other stars from the ZDI community (Paper II), we can now evaluate the variation of stellar wind torques over decadal timescales. We briefly reiterate the angular momentum loss prescriptions from FM18 in Section 2, collate solar observations in Section 3, and implement them in Section 4 to produce the most up-to-date determination of the solar braking torque, using methods based on the surface magnetogram data obtained from SOHO/MDI and SDO/HMI, and evaluating the open magnetic flux from the \textit{Ulysses} and the Advanced Composition Explorer (ACE) spacecrafts, along with an estimate based on the rotational behaviour of other Sun-like stars. Section 5 then discusses our result and addresses the observed discrepancy between surface field and open flux methods, along with the differences between our torque value and the derived long-time average result. | In this work we have utilised the wealth of current solar observations and the semi-analytic results from FM18 to produce an estimate of the current solar wind torque. This is compared with spin evolution calculations and shown to be a factor of 2-3 smaller than expected. Two angular momentum loss prescriptions from FM18 are implemented using observed surface field strengths from the SOHO and SDO spacecrafts, along with mass loss rates and open flux measurements from \textit{Ulysses} and ACE spacecrafts. The methods are found to produce average torques which either differ due to the amount of unsigned open flux the FM18 model produces from a given surface field observation, or the potential over-prediction of the open flux from spacecraft measurements. Assuming the open flux measurements from the in-situ spacecrafts are valid, we predict the solar wind torque has a present-day value of $2.3\times 10^{30}$erg, averaged over the last 22 years. The observation that the spin rates of Sun-like stars converge toward a single track that depends on age also allows us to derive equation (\ref{eq_omegaconverged}), describing how this spin-down depends on the torque and stellar properties. Then using solar parameters in equation (\ref{eq_tausunrot}) predicts that the long-time averaged torque should be $6.2\times 10^{30}$erg. Comparing this estimate of the torque from observed spin-evolution to the present-day torques predicted by the dynamical models gives additional insights. Differences in the average present-day torques to the spin-evolution torques, could be due to, (a) variability on a longer timescale than probed by the present-day variability presented here (but less than a spin-down time), (b) errors in using the dynamical models inferring present-day torque, or (c) that stars spin-down significantly different than Skumanich at ages of a few to several Gyr. We need additional information to discriminate between these possibilities. The required variability of (a) suggests we should observe stars like the Sun that are on average significantly more active (i.e., that they have larger torques) such that the average is correct. From the dynamical models, (b), uncertainties remain in the wind acceleration and effects of non-axisymmetric field components which both require further study to disentangle. Observationally, (c) requires more period-mass-ages for old stars to confirm or refute the \cite{van2016weakened} hypothesis. For other Sun-like stars, measurements of their unsigned open flux and mass loss rates are not readily available. Instead we rely on surface magnetic field measurements, which are gained through Zeeman Doppler Imaging and Doppler Broadening techniques. Using the FM18 formula, predictions of the angular momentum loss rates for these stars based on their surface measurements may be smaller than in reality. Future models should be refined to better match the wealth of solar data available, such models should be able to open the correct amount of flux from a given surface magnetic field observation and continue to remain general for application to other Sun-like and low-mass stars. | 18 | 8 | 1808.00063 |
1808 | 1808.09973_arXiv.txt | We study the components of cool and warm/hot gas in the circumgalactic medium (CGM) of simulated galaxies and address the relative production of OVI by photoionization versus collisional ionization, as a function of halo mass, redshift, and distance from the galaxy halo center. This is done utilizing two different suites of zoom-in hydro-cosmological simulations, VELA (6 halos; $z>1$) and NIHAO (18 halos; to $z=0$), which provide a broad theoretical basis because they use different codes and physical recipes for star formation and feedback. In all halos studied in this work, we find that collisional ionization by thermal electrons dominates at high redshift, while photoionization of cool or warm gas by the metagalactic radiation takes over near $z\sim2$. In halos of $\sim 10^{12}M_{\odot}$ and above, collisions become important again at $z<0.5$, while photoionization remains significant down to $z=0$ for less massive halos. In halos with $M_{\textrm v}>3\times10^{11}~M_{\odot}$, at $z\sim 0$ most of the photoionized OVI is in a warm, not cool, gas phase ($T\lesssim 3\times 10^5$~K). We also find that collisions are dominant in the central regions of halos, while photoionization is more significant at the outskirts, around $R_{\textrm v}$, even in massive halos. This too may be explained by the presence of warm gas or, in lower mass halos, by cool gas inflows. | \label{sec:1} Gas in the circumgalactic medium (CGM) is a key ingredient in galaxy evolution. Gas from the intergalactic medium (IGM) streams into the centers of DM halos and feeds star formation. Stellar feedback heats up cold gas and produces metallicity-enhanced warm/hot gas ($T>10^{4.5}$~K) that outflows to the CGM. Theory and simulations agree that halos with $M>M_{\textrm crit}\sim 3 \times 10^{11}~M_{\odot}$ develop a warm/hot CGM through both shock heating at $\sim R_{\textrm v}$ and outflows \citep{DekelBirnboim06}. In these systems cold gas in narrow inflowing streams can penetrate through warm/hot CGM only at $z>2$. In low mass halos ($M<M_{\textrm crit}$), however, the CGM is always dominated by cold ($T~<10^{3.8}$~K) and cool gas (10$^{3.8}$~K~$<T<10^{4.5}$~K) \citep{BirnboimDekel03,Keres05,DekelBirnboim06,Keres09,Correa17,Fielding17}. This implies that halos of mass $\sim10^{12}$ and above should have a significant fraction of their baryons in a warm-hot CGM, which may serve as a gas reservoir for star formation at later times \citep{Faerman17}. The understanding of the formation, evolution and properties of the CGM is thus mandatory for the study of the evolution of galaxies. \smallskip \\ Even if massive, this warm/hot gas component is difficult to detect as its density is low and its emission diffuse. In addition, most of the radiation is at high energies, in the UV and X-ray ranges, which are hard to observe \citep[e.g.,][]{Crain10}. The most common technique used to study the diffuse CGM is to look for absorption lines in background bright sources (e.g., QSO). This method has been useful for the detection of multiphase (cool and hot) gas in the CGM of galaxies, but it requires a close projection between the studied galaxy and the background source. Low temperature and high density gas has been observed through lines from gas in lower ionization states or from Lyman Limit Systems (LLS) and Damped Lyman-Alpha systems (DLAs). Warm/hot diffuse gas can only be studied by observations of lines from highly ionized metals. Each observation gives only a single measurement of CGM properties thus a large number of observations are required to get a general picture of the CGM surrounding galaxies. Many useful ion absorption lines probing CGM temperatures fall in the UV range; as a consequence, it is easier to study properties of the CGM in high-redshift galaxies where the rest-frame UV absorption lines have been redshifted to the optical range \citep[e.g.,][]{CowieSongalia95,Steidel10,Rudie13,Lehner14}. This is the case for OVI lines that for low redshift galaxies ($z<2.0$) fall in the UV spectral range. The most relevant observations in low-redshift CGMs have utilized the Cosmic Origins Spectrograph (COS) on the Hubble Space Telescope (HST) \citep[e.g.,][]{ThomChen08,Prochaska11}. In particular, the COS-Halos project provides measurements of highly ionized atoms, including OVI, for the CGMs of a large sample of $\sim L_*$ galaxies \citep{Tumlinson13}.\smallskip \\ The interpretation of the composition of the CGM based on a single metal ion is not straightforward. The determination of the distribution of mass, temperature, and density of the gas requires understanding of the ionization mechanism at work. Oxygen is collisionally ionized (CI, hereafter) to OVI between 10$^5$~K~$<T<10^6$~K, nearly independent of density. On the other hand, at lower densities and lower temperatures it may be produced by photoionization (PI, hereafter), once a hard external radiation field is present. Thus, OVI in the CGM can be a tracer of either collisionally-ionized warm/hot gas or photo-ionized cool gas \citep{Savage02,Peeples14,Werk14,Stern16,Faerman17}. The former could have originated, for example, from feedback-driven galactic outflows \citep{Mathews17} or from virial shock heating, and the latter could arise, e.g., from cold gas inflows \citep{Stern18} or from cooling of the warm-hot phase \citep{Bordoloi17}. Observers dispute which is the gas phase that dominates the CGM. On one hand, X-ray observations of higly ionized oxygen emission lines \citep{Henley10,HenleyShelton10} and also absorption lines within the spectrum of extremely bright QSOs \citep[e.g][]{Nicastro02,Rasmussen03}, suggest the presence of extended massive warm/hot coronae dominated by CI high ions (e.g. OVII/OVIII). Furthermore, extended X-ray line and continuum emission observed in low-$z$ galaxies \citep{AndersonBregman11,Bogdan13,Anderson15} points towards that direction. On the other hand, a detailed examination of the CGM absorption line profiles \citep{Tripp08,ThomChen08,Werk14,Stocke13} have shown a high covering factor of photoionized gas clouds in the CGM of low-$z$ star-forming galaxies. Also, COS-Halos found that it is possible that a large fraction of the CGM gas mass is in a cool low ionization state. This cool gas appears to be in dense knots that are not in hydrostatic equilibrium with their surroundings \citep{Werk14}. How such gas is supported is not well understood. A combined scenario of both a warm/hot and a cool CGM has been also proposed by \citet{Stocke14} and \citet{Pachat16}, who postulated that cool PI clouds can be embedded in a hotter more massive CI diffuse gas, based on OVI and Ly$\alpha$ absorption line detections. \smallskip \\ Hydrodynamical simulations have been used to study the CGM properties, the OVI distribution, and to attempt to break the degeneracy concerning its production mechanism. Some simulations obtained an OVI column density (N$_{\textrm OVI}$, hereafter) that is lower by an order of magnitude from the observed value \citep[e.g.,][]{Hummels13,Suresh15,RocaFabrega16,Ford16,Oppenheimer16}. Other simulations reveal better, although still not full, agreement with the observations \citep[e.g.,][]{Gutcke17,Suresh17,Nelson18}. Some of these works also found a clear dependence of the N$_{\textrm OVI}$ on the star formation rate and the $M_{200}$ of the host galaxy, and on its environment \citep{Oppenheimer16, Liang16}. Feedback efficiency showed to play a key role in modifying the shape and normalization of the N$_{\textrm OVI}$ profile, due to its effect on the metal production \citep{Rahmati16,Oppenheimer16}. Also it has been found that the total OVI mass and distribution depends on the feedback strength \citep[e.g. ][]{Liang16}. Feedback strength determines the amount of warm/hot gas that is produced, and how much that generates OVI through CI \citep{Fielding17}. Simulations have also indicated that in low-mass halos the OVI is produced by PI more than in high-mass halos \citep[e.g.][, at $z=0$]{Gutcke17}. Recent works also agree on that in halos with masses larger than 10$^{12}M_{\odot}$, at low-$z$, OVI is mainly produced through collisional ionization \citep{Liang18}. These results agree with the theoretical framework proposed by \citet{BirnboimDekel03} where higher mass halos are able to generate a warm/hot CGM through virial shock heating while lower mass systems are dominated by cold/cool flows. \smallskip \\ In this work we use two simulation suites to study the CGM as a function of redshift, halo mass, and distance from the halo center. In Sec.~\ref{sec:2} we describe the simulation suites. In Sec.~\ref{sec:4} and ~\ref{sec:4.7} we analyze properties of the CGM gas in simulations and we study the OVI distribution as function of halo mass, radius and redshift. The summary and conclusions are presented in Sec.~\ref{sec:8}. \smallskip \\ | \label{sec:8} In this work we study properties of CGM gas in two sets of simulations obtained using different codes and stellar feedback recipes, the VELA and the NIHAO suites. We focus our analysis on the mass-redshift dependence of the ionization mechanism that produces the OVI observed in the CGM. We note that collisional ionization of OVI occurs in gas at $T \sim 3 \times 10^5$~K (see Fig.\ref{fig:0_1}). Photoionized OVI can originate both in cool or warm gas, as shown by the black contours in Fig.~\ref{fig:10}. \smallskip \\ The main results are: \begin{itemize} \item {\bf Redshift dependence:} The OVI ionization mechanism depends strongly on redshift. At high-$z$ ($z>2-3$) collisions dominate due to a low flux of the ionizing UV-background and low metallicity of cool gas. At lower redshifts photoionization becomes more important as the UV-background flux peaks at $z\sim~2$, and cool/warm gas inflows from the IGM are enhanced by metals from outflows. At lower redshift, collisional ionization becomes important again, especially for high mass systems. The increase in the CI OVI fraction at low-$z$ is a consequence of both a decrease in the metagalactic UV-background intensity and the lower efficiency of inflows which bring fresh cool gas into the warm/hot gas halos of massive galaxies. Systems with $10^{11}~M_{\odot} < M < 10^{12}~M_{\odot}$, develop a warm CGM and at $z \sim 0$, OVI is created both by collisions and photoionization of the warm gas. Low mass systems ($M < 10^{11}~M_{\odot}$) show a different behavior at low-$z$: they do not develop a hot CGM and in addition cool gas at low-$z$ is already enhanced by metals, so PI remains dominant. \item {\bf Mass dependence:} We confirm that the dominance of an ionization mechanism also depends on the halo mass. This result was previously found by \citet{Gutcke17} when analysing halos at $z=0$, and in this work we extend this analysis to higher redshift. We find that low mass halos, which do not develop a warm/hot corona ($M < 10^{11}~M_{\odot}$), are dominated by photoionization from $z=1-1.5$ down to $z=0$. In systems with higher halo mass, gas is heated by accretion shocks and a warm/hot CGM is formed. The warm/hot gas corona is able to slow down or stop the accretion of fresh cool/warm gas and also keeps hot metallic outflows in the CGM. As a consequence, higher mass halos show higher dominance of collisional ionization in the CGM at low-$z$. \item {\bf Radial dependence:} We observe a radial dependence on the dominant OVI ionization mechanism. In general, CI is more significant in the central regions while PI is more dominant in the CGM outskirts. The PI is marginally dominant in the outskirts even in high-mass systems at low-$z$. This result may be explained by the existence of warm gas with $T<3 \times 10^5$~K, or the presence of low-density cool gas inflows that are able to push the virial shock to lower radii. Penetration of inflowing gas will depend on the accretion rate and also on the outflow strength. \item {\bf Feedback dependence:} Although the general behavior of the dominant ionization mechanism as function of redshift and halo mass is similar in VELA and NIHAO, the radial distribution, location in redshift of the peak photoionization fraction and the absolute values of the OVI collisional ionization fraction will strongly depend on feedback recipes used in simulations. The OVI collisional ionization fraction profile, both in redshift and radius, could be used in the future as a new theoretical/observational constraint on the feedback recipes used in simulations. \item {\bf Origin of the CGM gas:} We have studied the origin of CGM gas in the NIHAO SPH simulations and we have found that each gas phase shows different origins. Interestingly, warm gas can be produced either by cooling of CGM hot gas, by SN heating of disk cold gas, or by inflowing shock heated IGM gas. Cool gas on the other hand can be fresh gas infalling from the CGM in clumps and through filaments, gas from the disk that was heated up by stellar feedback and later on cooled down again, or gas that was heated by the virial shock and cooled down from the IGM/CGM warm/hot gas for the first time. \end{itemize} | 18 | 8 | 1808.09973 |
1808 | 1808.07643_arXiv.txt | We report the observations of the solar chromosphere from a newly commissioned solar telescope at the incursion site near Pangong Tso lake in Merak (Leh/Ladakh). This new H$_{\alpha}$ telescope at the Merak site is identical to the Kodaikanal H$_{\alpha}$ telescope. The telescope is installed in the month of August, 2017 at the Merak site. A 20-cm doublet lens with additional re-imaging optics makes the telescope. A Lyot filter with 0.5~\AA~passband isolates the Balmer line of the hydrogen spectra to make the observations of the solar chromosphere. The observations made in H$_{\alpha}$ wavelength delineates the magnetic field directions at the sunspot and the quiet regions. A CCD detector records the images of the chromosphere with a pixel resolution of 0.27$^{\prime\prime}$ and covers 9.2$^{\prime}$ field-of-view. This telescope has a good guiding system that keeps the FOV in the intended position. We report the development of control software for tuning the filter unit, control detector system, observations, and calibration of the data to make it useful for the scientific community. Some preliminary results obtained from the Merak H$_{\alpha}$ telescope are also presented. This high altitude facility is a timely addition to regularly available H$_{\alpha}$ images around the globe. | In 1889, George Ellery Hale has built his own spectrohelioscope to make a map of the sun in particular wavelength (Hale and Ellerman, 1960). Since then, several observatories around the world have built slit based spectroheliographs which have entrance and exit slits. The exit slit decides the pass-band of the spectrohelioscope. Later, many observatories started using the filter based spectral isolators to image the sun, particularly the solar chromosphere. Using these instruments, observatories have made full-disk Ca-K and H$_{\alpha}$ images recorded in the photographic plates (Srivastava, Ambastha and Bhatnagar, 1991; Zirin, Ligget, and Patterson, 1982). With the advancement in the detector technology, many have started utilizing the electronic based detectors which are more sensitive to the light and obtained more uniform images (Verma et al. 1997; Denker et al. 1999; Otruba and Potzi, 2003). The H$_{\alpha}$ data shows the conditions of the chromosphere during the flare events. This is the best wavelength region to observe the solar flare ribbons from ground based observatories. The observations made in H$_{\alpha}$ wavelength shows the fibrilar structures which delineates the magnetic fields (Woodard and Chae, 1999). The filament eruption, pre-flare brightening, rotation of prominences before the eruption, dynamics of spicules and many more interesting features could be seen in this wavelength. Mass motions inside the filaments and rotations of the prominences are studied in several cases during its eruption (Zirker {\it et~al.} 1998; Lin {\it et~al.} 2003; Labrosse {\it et~al.} (2010); Mackey {\it et~al}. (2010)). With the high-resolution and long hours of observations of the filaments/prominences it is easy to study their properties using H$_{\alpha}$ observations leading up to the eruption. During the flare, at the chromospheric level two ribbons are seen which move away from each other at certain speed (Vemareddy, Maurya, and Ambastha (2012)). The identification of newly brightened pixels can be done using H$_{\alpha}$ wavelength. These newly brightened pixels along with the corresponding locations in the magnetograms can be used to compute the magnetic flux swept-up by the ribbons as they move apart in the chromosphere (Qiu {\it et~al} (2004)) and can be used to infer the reconnection rate inside the reconnecting region. The details of various other science cases that can be studied with the H$_{\alpha}$ data is reported in Ravindra {\it et al.} (2016). Kodaikanal observatory has a history of making the full-disk solar observations in white-light, Ca-K and in H$_{\alpha}$ wavelengths. The white-light and Ca-K observations have started in 1904 and 1905 respectively. The H$_{\alpha}$ observations have started in 1912. Both the Ca-K and H$_{\alpha}$ observations discontinued in 2007 because the production of photographic plates/films has stopped. Though, the white-light observations using photographic plates continued till today (Ravindra, et al. 2013). Apart from Kodaikanal, the solar observations were also carried out at Aryabhatta Research Institute of Observational Sciences (ARIES) and at Udaipur Solar Observatory (USO) using 15-cm objective, starting from 1972. To continue the long term full-disk observations of the sun in H-alpha, Indian Institute of Astrophysics (IIA) has worked out the main optical design and configuration of the telescope. Based on the design, two identical telescopes and Lyot filters were manufactured by Nanjing Institute of Astronomical Optics and Technology (NIAOT), National Astronomical Observatories (NAO), Chinese Academy of Sciences (CAS), China. Both the telescopes were assembled and tested in NIAOT for its performance. Out of two telescopes, one of the telescopes was installed at the Kodaikanal Observatory in October, 2014 (Ravindra et al. (2016)). The second one was installed initially at Center for Research and Education in Science and Technology (CREST) campus at Hosakote by our IIA team. We have carried out in-house development of the software control and tracking system of the telescope. After completion of successful testing at CREST Hosakote, telescope was taken to Merak for final installation, which is also the proposed site for the 2-m class National Large Solar Telescope (NLST; Hasan, 2012). In this paper, we report the installation of the second telescope, development of telescope control software, technique used for guiding the telescope and filter unit. We present some of the representative results in the paper and the paper concludeswith the summary of the telescope installation at Merak. | Indian Institute of Astrophysics (IIA) has planned and designed the H$_{\alpha}$ telescope. Two identical telescopes were fabricated at NIAOT, china and tested for its performance. The NIAOT has made the 2 Lyot filters whose passbands are 0.4 and 0.5~\AA. A team from NIAOT has installed one of the telescopes at the Kodaikanal Solar Observatory in 2014 to continue the observations in H$_{\alpha}$ wavelength which was initiated in 1912 at Kodaikanal. The Lyot filter installed at Kodaikanal has a passband of 0.4~\AA. The second one is installed at the Merak site by IIA team, near the Lake shore of Pangong Tso, where the IIA is planning to setup a large 2-m class telescope dedicated for solar observations. The H$_{\alpha}$ telescope is installed at Merak site in August 30, 2017. Since, then it is making the observations of the sun whenever the sky is clear. The passband of the Lyot filter installed at Merak is 0.5~\AA. The chromospheric images obtained from the telescope shows fine features in H$_{\alpha}$ wavelength and showed several surges in many active regions. The H$_{\alpha}$ telescope installed at Merak site provides another addition to the existing solar facilities at Kodaikanal. The Lyot filter has the facility to tune it to any position on the H$_{\alpha}$ line profile. A stepper motor attached to the filter unit can be tuned to any point on the line profile with an accuracy of 10~m\AA. The tuning takes about 1~s to move the central passband from one position to the other on the line profile. This property of the filter can be used to reconstruct the line profile and also can be used to find the Doppler shift in the line profile position. Using this, it is possible to make dopplergrams at the chromospheric heights (Dhara, Ravindra and Banyal 2016; Chae, Park and Park 2006; Joshi et~al. 2013). In future, we will produce such dopplergrams and make it available for scientific use. At present, a smaller size CCD camera is used. This covers only 9.2$^{\prime}$ FOV. In future, we would like to replace it with a large format CCD to cover the entire 32$^{\prime}$ FOV. The Merak H$_{\alpha}$ telescope is a new addition to the existing facilities at Indian subcontinent. This trans Himalayan region largely remains unaffected by regular monsoon and thus enable useful and uninterrupted observations of the sun. | 18 | 8 | 1808.07643 |
1808 | 1808.08059_arXiv.txt | Massive planets form within the lifetime of protoplanetary disks and therefore they are subject to orbital migration due to planet-disk interactions. When the first planet reaches the inner edge of the disk its migration stops and consequently the second planet ends up locked in resonance with the first one. We detail how the resonant trapping works comparing semi-analytical formulae and numerical simulations. We restrict to the case of two equal-mass coplanar planets trapped in first order resonances but the method can be easily generalised. We first describe the family of resonant stable equilibrium points (zero-amplitude libration orbits) using series expansions up to different orders in eccentricity as well as a non-expanded Hamiltonian. Then we show that during convergent migration the planets evolve along these families of equilibrium points. Eccentricity damping from the disk leads to a final equilibrium configuration that we predict precisely analytically. The fact that observed multi-exoplanetary systems are rarely seen in resonances suggests that in most cases the resonant configurations achieved by migration become unstable after the removal of the protoplanetary disk. Here we probe the stability of the resonances as a function of planetary mass. For this purpose, we fictitiously increase the masses of resonant planets, adiabatically maintaining the low-amplitude libration regime until instability occurs. We discuss two hypotheses for the instability, that of a low-order secondary resonance of the libration frequency with a fast synodic frequency of the system, and that of minimal approach distance between planets. We show that secondary resonances do not seem to impact resonant systems at low-amplitude of libration. Resonant systems are more stable than non-resonant ones for a given minimal distance at close encounters, but we show that the latter nevertheless play the decisive role in the destabilisation of resonant pairs. We show evidence that, as the planetary mass increases and the minimal distance between planets gets smaller in terms of mutual Hill radius, the region of stability around the resonance center shrinks, until the equilibrium point itself becomes unstable. | Super-Earths (SE) are planets with a mass between 1 and $\sim$20 Earth masses or a radius between 1 and $\sim$4 Earth radii, and so-far discovered with orbital period typically shorter than $\sim 100$ days. They are estimated to orbit 30 -- 50\% of Sun-like stars (\cite{2011arXiv1109.2497M}; \cite{2012ApJS..201...15H}; \cite{2013ApJ...766...81F}; \cite{2013PNAS..11019273P}) and multi-planetary systems are not rare. The fact that close-in SE systems form so frequently around stars, but not always (for instance not in the Solar System) is an interesting constraint on planetary formation models. It is generally expected that SEs form (mostly) within the lifetime of the protoplanetary disk of gas and therefore, regardless of whether they form in the inner or outer part of the disk, they should undergo radial migration towards the central star, as a result of planet-disk interactions (\cite{2015A&A...578A..36O}; \cite{2017MNRAS.470.1750I}). Migration brings the SEs to the inner edge of the disk, where inward migration stops \cite{2006ApJ...642..478M}. For this reason, the SEs are captured into mutual mean motion resonances, where the ratios of orbital periods are equal to the ratios of integer numbers. This is observed in all simulations (e.g.{ \cite{2007ApJ...654.1110T}; \cite{2008A&A...482..677C}; \cite{2008A&A...478..929M}, and the aforementioned \cite{2015A&A...578A..36O} and \cite{2017MNRAS.470.1750I}). Due to this renewed interest in resonant captures, in the first part of this paper we revisit the problem of capture in first order resonances of two equal-mass coplanar planets in convergent migration, using a semi-analytical approach and numerical simulations. In Section \ref{sec:AnalyticalSetup} we compute analytically the locus of equilibrium points of first-order resonances, where both the resonant and secular oscillations of the planetary orbits have a null amplitude. Our calculations are developed for unexpanded Hamiltonians, which allows to follow the dynamics up to arbitrarily large eccentricities (e.g. \cite{2006MNRAS.365.1160B}, \cite{2006CeMDA..94..411M}). We compare the results with those obtained with first and second order expansions of the Hamiltonian in the eccentricity, showing qualitative and quantitative disagreements. The quantitative accuracy of our results is validated with simulations in which planets are forced to migrate towards each-other, without any eccentricity damping. These simulations have to follow the loci of the equilibrium points, and show perfect agreement with the unexpanded model. Moreover, we calculate the two frequencies of libration around the equilibrium points, therefore obtaining a complete understanding of the system; we again check the validity of the analytical calculations against numerical simulations in which the amplitudes of resonant and secular librations are slightly excited and the frequencies of oscillation of the semi major axis and the eccentricity are measured. In Section \ref{sec:ConvergentInwardMigration} we introduce the eccentricity damping exerted by the disk onto the planets. This leads to a final equilibrium configuration where convergent migration stops. The analytic calculation of the equilibrium eccentricities and semi major axes ratio is presented in the Appendix. We check against numerical simulations the validity of these analytical predictions, showing excellent agreement. Despite resonant capture is typical of migration simulations, the observed SEs systems show little preference for near-integer period ratios and their orbital separations are usually much wider than those characterising planets in resonant chains. \cite{2017MNRAS.470.1750I} showed that this observation is not inconsistent with the migration/resonant trapping paradigm. In fact, simulations show that, after the removal of the disk of gas, the resonant planetary systems often become unstable. \cite{2017MNRAS.470.1750I} showed that the observations are very well reproduced if the fraction of the resonant systems that eventually become unstable exceeds 90\%. The reasons for these instabilities, however, are unexplained. \cite{2012Icar..221..624M} studied numerically the stability of resonant multi-planetary systems for high-integer first-order mean motion resonances. They built the desired resonant configuration by simulating the Type-I migration phase in a protoplanetary disk of gas; then they slowly depleted the disk. They observed that there is a critical number of planets above which the resonant systems go naturally unstable, with a crossing time comparable to that of non-resonant systems, and studied how this number changes with the planetary masses and resonant order. In other words, they demonstrated that, given the planetary masses, there is a limit number of planets that can form a stable resonant chain or, given the number of planets, there is a limit mass for stability. The reason of the instability, however, was not discussed. Thus, in the second part of this paper we address why resonant planets become unstable if they are too massive. We focus here on a system of two coplanar planets and study the stability of the resonant center as a function of the planets' masses (assumed to be equal for simplicity). In a subsequent work, we will generalise this study to more populated resonant chains. Again, we follow a double approach: analytic and numeric. In Section \ref{sec:MassIncrease} we start from a pair of small-mass planets deep in resonance and we slowly increase their masses. The mass growth preserves, by the adiabatic principle, the original small libration amplitude. In this way we can explore the stability of the resonance center up the threshold mass for instability. At the same time, we detail how one can follow analytically the evolution of the system to a good approximation up to high value of the planetary masses. To understand why the planets ultimately become unstable, we compare the numerical evolution of the system with an analytically computed map of secondary resonances (resonances between the libration frequencies or between a libration frequency and one of the short-periodic harmonic) as well as a map of minimum approach distance between the planets, finding that one of them matches well the instability limit observed in the numerical simulations. We summarise our results in the final Section \ref{sec:Conclusions}. | \label{sec:Conclusions} In this work, we investigated the dynamics of resonant planetary system, from the capture in mean motion resonance via convergent migration in a protoplanetary disk, to the stability of systems with low-amplitude libration of the resonant angles. We treat the simple case of the planar three-body problem, with two equally massive planets. We present the analytical techniques needed to describe the system in Section \ref{sec:AnalyticalSetup}. There, we develop the theory for unexpanded Hamiltonians and find semi-analytically the equilibrium points; we validate numerically the analytical calculations, showing perfect agreement. We compare these with equilibrium points resulting from low-order expansions in the eccentricities, showing that the latter they do not capture qualitatively or quantitatively the simulations. Since we are interested in the dynamics in the region of the phase space around the equilibrium points, we calculate the frequencies of librations in the regime of vanishing amplitude of libration, and check again the results with numerical simulations. In Section \ref{sec:ConvergentInwardMigration} we describe the forces which result from the interactions between the planets and a disk of gas, which is used in the numerical simulations in order to capture the planets in first-order mean motion resonance. These interactions include a damping in the eccentricity and a torque which results in an inward Type-I migration. To ensure convergent migration and resonant trapping, a planetary trap (\cite{2006ApJ...642..478M}) is implemented at the edge of the disk of gas. These dissipative forces are implemented in our code using simple analytical formul\ae\ which simulate the disk-planet interactions of real hydro-dynamical simulations (e.g.\ \cite{2006A&A...450..833C}). We present in the Appendix \ref{subsec:EqInResonantCapture} an analytical description of the capture in mean motion resonance following a general approach, and derive formul\ae\ to calculate analytically the final equilibrium configuration. We compare our formul\ae\ with similar ones from previous works, and validate our results with numerical simulations, showing perfect agreement. In Section \ref{sec:MassIncrease} we investigate the stability of resonant systems at low amplitude of libration. We describe our numerical experiments where we fictitiously increase the planetary mass to follow the low-amplitude regime until the onset of instability. At the same time, we detail how one can follow analytically the evolution of the system to a good approximation up to high value of the planetary masses. We test against two possible reasons for instability. The first is that of a secondary low-order resonance between the frequency of libration of a resonant angle (which grows with the planetary mass, while maintaining adiabatically the same amplitude around the equilibrium point) and the frequency of the fast synodic angle $\LAMBDA_1-\LAMBDA_2$ (which is constant with the planetary mass), where $\LAMBDA_i$ is the mean longitude of a planet. We construct a map of the frequency of libration of the resonant angles as a function of the planetary mass and the eccentricity, and compare the calculated values with the frequency of the synodic angle. We see that some systems become unstable before reaching the 2-1 resonance between the libration and the synodic frequency, while others cross it unaffected. We therefore conclude that these secondary resonances do not play a significant role in the instability of pairs of planets in first order mean motion resonance at low amplitude of libration. The second hypothesis is that of instability due to close encounters between planets, inspired by the mutual Hill radius stability criterion \cite{1993Icar..106..247G}. In this case we build a map of the minimal distance reached by the planets in their orbital configuration as a function of the mass and the eccentricity. We see that resonant planetary systems are more stable than those with randomly chosen orbital parameters, as they can reach a minimal distance that is smaller then the critical distance $d_{crit} = 2 \sqrt{3} r_H$, where $r_H$ is the mutual Hill radius, which is the usual critical distance below which two planets go unstable (cfr.\ \cite{1993Icar..106..247G}, \cite{2017Icar..293...52O}). We find nonetheless that there is a critical distance after which the system goes unstable, which is a fraction of the usual $d_{crit}$. We see that for systems with bigger amplitude of libration of the resonant angles this critical distance approaches more and more the usual $d_{crit}$. This indicates that the region of stability around the equilibrium point shrinks as the mass increases, until the point itself becomes unstable and the system exits the resonance. \newcommand{\aap}{A\&A} \newcommand{\aj}{AJ} \newcommand{\apj}{ApJ} \newcommand{\apjs}{ApJS} \newcommand{\mnras}{MNRAS} \newcommand{\icarus}{ICARUS} \newcommand{\apss}{APSS} | 18 | 8 | 1808.08059 |
1808 | 1808.06758_arXiv.txt | {We review the physics of intergalactic electromagnetic cascades in the presence of the extragalactic magnetic field (EGMF). Various regimes of intergalactic electromagnetic cascades are considered depending on the number of cascade generations, the value of the cascade electron deflection angle, and the relations between the EGMF coherence length, typical cascade $\gamma$-ray mean free path, and electron energy loss length. We also review contemporary constraints on the EGMF parameters and explore the sensitivity of various $\gamma$-ray instruments to the EGMF parameters.} | \label{sec:intro} Primary $\gamma$-rays from extragalactic sources may be absorbed on extragalactic background light (EBL) \cite{nikishov_1962,gould_schreder_1967} and cosmic microwave background (CMB) photons \cite{jelley_1966} through the $\gamma\gamma \rightarrow e^{+}e^{-}$ pair production (PP) process. For sufficiently distant sources with redshift $z>$0.1, the optical depth $\tau$ of the PP process exceeds unity for the primary energy $E_{0}>$1 TeV. Therefore, above some energy $E(\tau=1)$ (usually called ``the gamma-ray horizon'' and defined by the so-called Fazio-Stecker relation \cite{fazio_stecker_1970}) primary spectra of extragalactic $\gamma$-ray sources are strongly distorted. This effect, evidently, creates appreciable obstacles for the direct study of distant very high energy (VHE, $E>$100 GeV) $\gamma$-ray emitters. On the other hand, $\gamma$-ray absorption could be an asset (e.g. \cite{neronov_semikoz_2007}, \cite{neronov_semikoz_2009}). Indeed, secondary electrons and positrons (in what follows they are called simply ``electrons'', unless it is necessary to distinguish $e^{+}$ and $e^{-}$) produce secondary (cascade) $\gamma$-rays through inverse Compton (IC) scattering. These cascade $\gamma$-rays are observable; their spectral shape may help constrain the shape of the primary $\gamma$-ray spectrum in the optically thick ($\tau>$1) energy range. Moreover, observable spectral, angular, and temporal distributions of cascade $\gamma$-rays are sensitive to parameters of the extragalactic magnetic field (EGMF). Thus, these parameters, such as the EGMF characteristic strength $B$ and coherence length $\lambda$, could be constrained using observations of extragalactic $\gamma$-ray sources. This case exemplifies how intergalactic electromagnetic (EM) cascades may be a valuable tool of astroparticle physics. Deep understanding of the underlying physics is important in order to use this tool effectively. EM cascades in the magnetized Universe may reveal a number of qualitatively different regimes depending on the following basic parameters: the primary energy, the distance from the source to the observer, the strength and coherence length of the EGMF, etc. The most general constraints on the EGMF parameters in voids of the large scale structure (LSS) that define the ``EGMF parameter window'' are provided in section~\ref{sec:window}. In section~\ref{sec:regimes} we discuss various regimes of intergalactic EM cascades. In section~\ref{sec:egmf} we briefly review some contemporary constraints on the EGMF parameters $(B,\lambda)$. Section~\ref{sec:sensitivity} is devoted to a study of sensitivity of various $\gamma$-ray instruments to the EGMF parameters. Some other case studies performed by us were described in \cite{dzhatdoev_et_al_2017}. Finally, we provide conclusions in section~\ref{sec:conclusions}. In this work we assume $z$=0.186 and the EBL model of \cite{gilmore_et_al_2012} unless stated otherwise. This EBL model is consistent with contemporary constraints (e.g. \cite{korochkin_rubtsov_2018}). We made use of the ROOT analysis framework \cite{brun_rademakers_1997} and the MINUIT routine \cite{james_roos_1975}. | \label{sec:conclusions} The strength and structure of the EGMF are currently poorly constrained. EM cascades are a promising tool to probe the intergalactic medium. However, while doing so, one has to remember that EM cascades in the magnetized Universe may display a widely varying behaviour according to the regime in operation. The most important regimes were discussed above. We have performed a detailed sensitivity study of various $\gamma$-ray instruments to the EGMF parameters and found that by far the best constraints on the EGMF would be obtained using a combination of space-based and ground-based instruments. Finally, we have identified a new promising technique to measure extremely low ($B<$0.1 fG) EGMF, namely, the time projection chamber approach to $\gamma$-ray astronomy. | 18 | 8 | 1808.06758 |
1808 | 1808.05345.txt | We have derived new physical quantities for several High-Mass X-ray Binaries (HMXBs) with supergiant (SG) companions through their cyclotron lines. The parameters are: the terminal velocity of the wind, the mass loss rate of the donor, the effective temperature and the magnetic fields. These parameters influence significantly the improvement of the model of accretion. In spite of the variety of their observational properties, the corresponding magnetic field is around $B\sim 10^{12}$ G. This result can be constrained by the effects on stellar evolution. %%As a result, the difference in time scales of variation of accretion rate and different type of companion %%will be distinguished between the magnetized neutron stars. In addition, we have performed a segmentation in the parameter space of donors intended for several SG-HMXB listed in our sample set. %%As a result, associated with accretion from a stellar wind The parameter space can be categorized into five regimes depending on the possibility of disk formation associated with accretion from the stellar wind. This can give a quantitative clarification of the observed variability and the properties of these objects. %%We demonstrate that only once the system enters the inner edge of accretion regime, it can emit powerful X-rays %%and be detected as a SG-HMXB. %In addition, their wind velocities can also be used to identify them observationally. We show that, when these systems come into the direct accretion region, systems with corresponding parameters can emit X-rays. | The detection of Cyclotron Resonance Scattering Features (CRSFs) in spectra of many %that emitted by accreting neutron stars (NSs) with high magnetic field (B$\geq 10^{12}$ G) provides valuable insights into the physics of emitting regions and the evolution of these systems. They form due resonant scattering processes with electrons, protons, and other ions in the plane and perpendicular to the magnetic field (Voges et al. 1982; Wilson et al. 2008). As a result, the cyclotron line features provide the only direct estimate of the magnetic field strength of NSs in X-ray binary systems. In High-Mass X-ray Binaries (HMXBs), an NS accretes matter from a companion star via stellar wind. The accreted matter is channeled along field lines of the strong magnetic field of the NS onto the magnetic poles. X-ray emission from the NS is produced in regions around the magnetic poles. It is noteworthy to mention here that most observed cyclotron lines have been detected above 10 keV and are interpreted as electron features, with inferred magnetic fields $B\sim 10^{12}$ G (Heindl et al. 2001). The combined effects of poor statistics, photoelectric absorption and the lack of evidence for a remnant accretion disk have made these energy sources elusive. %In addition a cutoff below 2-3 keV would remain undetected %in the available spectra of some sources. Most efforts to calculate theoretical cyclotron lines have been performed in a line-forming region with a constant temperature and density of an electron-proton plasma permeated by a uniform magnetic field (Wheaton et al. 1979; Orlandini et al. 1999; Yamamoto et al. 2011). %Nishimura (2003) demonstrated that the variations on the %strength of a magnetic field affect the profile of the cyclotron %line. Nishimura (2005) calculated cyclotron lines assuming a strong variation in field strength with distance from an emission region. However, no model generating such high flux and high temperature at a layer deeper than absorbing heavy atoms has been proposed. According to recent studies, several pulsars show changes in luminosity dependence in the cyclotron resonance energy. The first aim of this paper is to derive magnetic field strengths, which is crucial for these systems, %One of the interesting properties of this %class is the observed correlation between the orbital period and the %spin period of the NS (Corbet 1986, 2004), and obtain clues about the evolution of HMXBs, which can be understood in terms of the conservative evolution of normal massive binary systems. The second aim of this study is to derive unknown parameters of HMXBs without uncertainty in the strength of the NS magnetic field. With robust data on the NS magnetic field, combined with spin period ($P_{\rm{spin}}$) and orbital period ($P_{\rm{orb}}$), we can fix several hitherto-unknown parameters, such as wind velocity and wind mass loss rate. These parameters influence significantly the model of wind-fed binary systems and can constrain the effects of binary evolution (Taani et al. 2012a,b; Taani \& Khasawneh 2017; Dai et al. 2017; Taani et al. 2019; Karino et al. 2019). From this standpoint, with observations of NS magnetic fields, we could constrain the end products of HMXBs, such as an NS-NS merger, which is considered to be one of the most powerful gravitational wave sources and also the most probable site for heavy element creation (Postnov \& Yungelson 20016; Taani 2015). In the next section, we introduce the recent results of NS magnetic field given by CRSF observations. In Section 3, we discuss the method to obtain hitherto-unknown binary parameters from robust data on the NS magnetic field in SG-HMXBs. In Section 4, we discuss our findings. The last section is devoted to conclusions. | %{\bf %THIS SECTION WILL BE MODIFIED BASED ON THE CONTENTS IN THE DISCUSSION PART.} %The work summarizes the known data regarding the cyclotron sources and the fundamental cyclotron line and field estimates of the sources. %We have studied various distribution of magnetic field HMXBs which generated by the cyclotron lines. We also %investigated the formation and evolution of our observed sample, spin period, orbital period and magnetic field %quantities. The following conclusions and implications are obtained: We have derived new physical quantities for several HMXBs with supergiant companions through their cyclotron lines. These parameters are: the terminal velocity of the wind, mass loss rate of the donor, magnetic field, effective temperature and corresponding luminosity. Furthermore, for all systems, our analysis (direct accretion condition shown by shaded region, and the solution of wind equation shown by solid curves) suggests that the wind velocity should be systematically slow. %These parameters greatly influence the evolutionary tracks, and can be constrained the effects on the stellar evolution. By adopting the accretion regime model by Bozzo et al. (2008), we have explored the parameter space in different regimes based on the intrinsic variabilities of mass accretion rate and wind velocity. This will allow us to describe an evolutionary path for several SG-HMXBs in these diagrams. Different regimes are sufficient to spatially separate bright X-ray sources, and can be probed through the magnetic field-wind velocity. As a result, persistent SG HMXBs within the shaded region can be observed through the direct accretion regime. This interpretation is based on its emission of accretion in high-energy X-rays. % %, the parameter space can be categorized %into: supersonic inhibition regime, subsonic inhibition %regime, supersonic propeller regime, subsonic propeller %regime, and direct accretion regime in order to describe the intrinsic variabilities of the mass accretion rate and wind velocity of these systems. %In addition, to describe an evolutionary path for several SG-HMXBs in these diagrams. %Different regimes are sufficient to spatially separate the bright X-ray sources, and can be probed through the magnetic field-wind velocity. %As a result, the persistent SG HMXBs within the shaded region can be observed through the direct accretion regime. %This interpretation is based on its emission of accretion in high-energy X-rays %%The X-ray emission of such a system is powered by the accretion created in the figures. It is seen that the wind velocity causes a significant effect on the results of their x-ray features and it can be used to determine the ejection mechanism. Consequently, when the wind velocity is slow, the accretion disk can be formed even in systems with large orbital period. This will allow to better characterize the HMXB of both types, SG and Be, hosting NS, by deriving accurate properties of these compact binaries. %will reveal information about the evolutionary history of different types of binaries. %From the updated measured HMXB cyclotron lines, we have derived the magnetic field strength of these systems. are all concentrated around From the updated measurement of HMXB cyclotron lines, the derived magnetic fields given by CRSF data are all concentrated around $\sim 10^{12} \rm{G}$. However, the fundamental energy during X-ray observation, spin and other physical parameters property diverges and varies. The existence of a high magnetic field has the potential to control their formation and evolution. %Thus, the magnetic field measure is of paramount importance to %having a thorough insight into the physics of the emitting region, and %provide us with constraints that help us improve our understanding %of binary evolution. %As a result, the cyclotron lines are of %As has been correctly pointed out, several sources show variation in wind velocity in the range of about 5 - 23 $\times 10^7 \rm{cm \, s}^{-1}$ %As has been correctly pointed out, several sources show variation in wind velocity in the range of about 5 - 23 $\times 10^7 \rm{cm \, s}^{-1}$ Our model cannot constrain the accretion mechanism for fast-spinning NSs (P$_{spin} \leq 40$ s) with a short orbital period (P$_{orb} \leq 10$ d), like in LMC X-4, Cen X-3 and OAO1657 (see Fig. 4). In LMC X-4 and Cen X-3, these two binary systems are extremely tight systems. Thus, the accretion mechanism may not be approximated by spherical wind, because in such tight systems the concentrated asymmetric wind or RLOF accretion should be considered. While for OAO1657, the donor of this system is a further evolved star with a long orbit, and can be observed as a Wolf-Rayet star with stellar wind mass loss rate during its evolution. %our model cannot constrain any further evolution for such binaries %our model no longer valid %There is no direct one-to-one relation between the %time variability at the smaller scales and the one induced by %the inhomogeneous wind. Matter can pile up for quite long %in intermediate regions (e.g. the shock or the corotation radius) %before an instability which makes accretion possible is %triggered. https://arxiv.org/pdf/1711.08709.pdf %As has been correctly pointed out, several sources show %variation in fundamental energy and hence x-ray observation. Since the decrease of the surface magnetic %field strength might be due to temporary or permanent burying of the field due to the %mass and angular momentum modes transported to the NS. %%\textbf{ %We have found the $\sim 10^{-8} M_{\odot} \rm{s}^{-1}$ is the critical value to let the disk to be %formed in the wind-fed X-ray binary systems. As a result, if the wind mass-loss rate decreased, %this would lead to the disruption of the disk and disappearance of the X-ray emission. %This supports the reliability of current mass-loss rate predictions. %%} %{\bf (This paragraph may be beyond the coverage of this paper.)} %cyclotron lines. The variation can be because during one rotation %phase different parts of the mound are probed, and also due to change %in local field structure due to accretion dynamics, e. g. change in %accretion rate, as is seen for sources like Her X-1 and V0332+53. %The difference in time scales of variation of E$_{cyc}$ and %luminosity will distinguish between magnetized accreting NSs. %However more fitting, statistical methods and representative sources with absorption %features in their X-ray spectra are needed in the future. This would help us to study the cyclotron %lines and to show the dependence of the magnetic field on %dynamical processes and to test possible variations in the mass %ratio distribution as tracers of different star formation %mechanisms. Finally, however the currently available data for CRSFs are not sufficiently accurate or numerous to allow precise analysis. One would hope that the results of this work will be improved with data from Suzaku, \emph{INTEGRAL}, $\emph{eRosita}$ and $\emph{HXMT}$, which will provide significant increases in the observational sensitivity of some cyclotron sources. %will constitute a base %of further studies on the observational properties of some cyclotron sources detected by Suzaku, %\emph{INTEGRAL}, $\emph{eRosita}$ and $\emph{HXMT}$ satellites and other sources. %(Merloni et al. %%2012; Doroshenko et al. 2014). %Recently, observed examples of faint HMXBs are in- %creasing. The existence of the faint persistent sources in %wind-fed systems would be a collateral evidence of the %fast-faint mode of wind accretion. However, in the actual %binary systems, one dimensional computation may be in- %sufficient. In the near future, two or three dimensional %modeling including orbital motion would figure out the %true nature of HMXBs %https://arxiv.org/pdf/1403.7600.pdf %However, J18027, , 4U1907 and Vela X-1 %(C) supersonic propeller regime (ra > rm > rco), %The propeller regime is defined when the accretion radius %of the disk is larger than the magnetic radius. %In addition, the subsonic %propeller regime becomes more clear as the strength of the %propeller increases. As the strength increases there is a sharp %decrease of the accretion rate to the star. | 18 | 8 | 1808.05345 |
1808 | 1808.08435_arXiv.txt | We have carried out interferometric observations of cyanopolyynes, HC$_{3}$N, HC$_{5}$N, and HC$_{7}$N, in the 36 GHz band toward the G28.28$-$0.36 high-mass star-forming region using the Karl G. Jansky Very Large Array (VLA) Ka-band receiver. The spatial distributions of HC$_{3}$N and HC$_{5}$N are obtained. HC$_{5}$N emission is coincident with a 450 $\mu$m dust continuum emission and this clump with a diameter of $\sim 0.2$ pc is located at the east position from the 6.7 GHz methanol maser by $\sim 0.15$ pc. HC$_{7}$N is tentatively detected toward the clump. The HC$_{3}$N : HC$_{5}$N : HC$_{7}$N column density ratios are estimated at 1.0 : $\sim 0.3$ : $\sim 0.2$ at an HC$_{7}$N peak position. We discuss possible natures of the 450 $\mu$m continuum clump associated with the cyanopolyynes. The 450 $\mu$m continuum clump seems to contain deeply embedded low- or intermediate-mass protostellar cores, and the most possible formation mechanism of the cyanopolyynes is the warm carbon chain chemistry (WCCC) mechanism. In addition, HC$_{3}$N and compact HC$_{5}$N emission is detected at the edge of the 4.5 $\mu$m emission, which possibly implies that such emission is the shock origin. | \label{sec:intro} Cyanopolyynes (HC$_{2n+1}$N, $n=1,2,3,...$) are one of the representative carbon-chain species. In low-mass star-forming regions, carbon-chain molecules are known as early-type species; they are abundant in young starless cores and deficient in star-forming cores \citep[e.g.,][]{1992ApJ...392..551S, 2009ApJ...699..585H}. In contrast to the general picture, cyanoacetylene (HC$_{3}$N), the shortest member of cyanopolyynes, is detected from various regions such as infrared dark clouds \citep[IRDCs; e.g.,][]{2012ApJ...756...60S}, molecular outflows \citep{1997ApJ...487L..93B}, protoplanetary disks \citep{2015Natur.520..198O, 2018ApJ...857...69B}, and comets \citep[e.g.,][]{2011ARA&A..49..471M} and it is interesting to trace cyanopolyyne chemistry for better understanding of the molecular evolution during star/planet formation process. Cyanopolyynes attract astrobiological as well as astrochemical interests. Since they contain the nitrile bond (--C$\equiv$N), cyanopolyynes have been suggested as possible intermediates in the synthesis of simple amino acids \citep[e.g.,][]{2017A&A...605A..57F, 2018arXiv180409210C}. Saturated complex organic molecules (COMs), consisting of more than six atoms with rich hydrogen atoms, are abundant around protostars. Such chemistry is known as hot core in high-mass star-forming regions and hot corino in low-mass star-forming regions. In addition to hot corino, around a few low-mass protostars, carbon-chain molecules are formed from CH$_{4}$ evaporated from dust grains, which is known as warm carbon chain chemistry \citep[WCCC; e.g.,][]{2013ChRv..113.8981S}. Progress in observational studies of carbon-chain molecules in high-mass star-forming regions has been slower, compared to low-mass star-forming regions. Regarding hot cores, HC$_{5}$N has been detected in chemically rich sources, Orion KL \citep{2013A&A...559A..51E} and Sgr B2 \citep{2013A&A...559A..47B}, while only a tentative detection of HC$_{7}$N in Orion KL was reported \citep{2015A&A...581A..71F}. \citet{2009MNRAS.394..221C} performed a chemical network simulation and suggested that cyanopolyynes could be formed in a hot core from C$_{2}$H$_{2}$ evaporated from grain mantles. Motivated by the chemical network simulation, \citet{2014MNRAS.443.2252G} carried out survey observations of HC$_{5}$N toward 79 hot cores associated with the 6.7 GHz methanol masers and reported its detection in 35 sources. However, the association with the maser is questionable, because they used a large beam (0.95\arcmin) and a low-excitation energy line ($J=12-11$; $E_{\rm u}/k = 10.0$ K), which can be excited even in dark clouds. \citet{2017ApJ...844...68T} carried out observations of long cyanopolyynes (HC$_{5}$N and HC$_{7}$N) toward four massive young stellar objects, where \citet{2014MNRAS.443.2252G} had reported the HC$_{5}$N detection, using the Green Bank 100-m and the Nobeyama 45-m radio telescopes, and detected high-excitation energy lines ($E_{\rm u}/k \approx 100$ K) of HC$_{5}$N. The detection of such lines means that HC$_{5}$N exists at least in the warm gas, not in cold molecular clouds ($T_{\rm {kin}} \simeq 10$ K). \citet{2018arXiv180405205T} found that the G28.28$-$0.36 high-mass star-forming region is a particular cyanopolyyne-rich source with less COMs compared with other sources. Hence, G28.28$-$0.36 is considered to be a good target region to study the cyanopolyyne chemistry around massive young stellar objects (MYSOs). Using the Nobeyama 45-m radio telescope, \citet{2016ApJ...830..106T} investigated the main formation mechanism of HC$_{3}$N in G28.28$-$0.36 from its $^{13}$C isotopic fractionation. The reaction of ``C$_{2}$H$_{2}$ + CN" was proposed as the main formation pathway of HC$_{3}$N, which is consistent with the chemical network simulation conducted by \citet{2009MNRAS.394..221C} and the WCCC model \citep{2008ApJ...681...1385}. \begin{figure}[ht!] \plotone{ir.eps} \caption{Spitzer's IRAC 3.6 $\mu$m image toward G28.28$-$0.36. The open circle and cross indicate the 6.7 GHz methanol maser \citep{2009ApJ...702.1615C} and ultracompact \ion{H}{2} (UC\ion{H}{2}) region \citep{2009A&A...501..539U}, respectively. \label{fig:IR}} \end{figure} In this paper, we carried out interferometric observations of cyanopolyynes (HC$_{3}$N, HC$_{5}$N, and HC$_{7}$N) toward the G28.28$-$0.36 high-mass star-forming region ($d = 3$ kpc) with the Karl G. Jansky Very Large Array (VLA). Figure \ref{fig:IR} shows the Spitzer IRAC 3.6 $\mu$m image\footnote{http://sha.ipac.caltech.edu/applications/Spitzer/SHA/} toward the region. G28.28$-$0.36 is classified as an Extended Green Object (EGO) source \citep{2008AJ....136.2391C} from the Spitzer Galactic Legacy Infrared Mid-Plane Survey Extraordinaire \citep[GLIMPSE;][]{2003PASP..115..953B}. In Figure \ref{fig:IR}, the open circle and cross indicate the 6.7 GHz methanol maser \citep{2009ApJ...702.1615C} and ultracompact \ion{H}{2} (UC\ion{H}{2}) region \citep{2009A&A...501..539U}, respectively. The 6.7 GHz maser is considered to give us the exact position of MYSOs \citep{2013MNRAS.431.1752U}. A UC\ion{H}{2} region seems to heat the environment. As shown in Figure \ref{fig:IR}, the ring structure around the UC\ion{H}{2} region is suggestive of expanding motion and on-going massive star formation. We describe the observational details and data analyses in Section \ref{sec:obs}. The resultant images and spectra of cyanopolyynes are presented in Section \ref{sec:res}. We compare the spatial distributions of cyanopolyynes with the infrared images and discuss possible formation mechanisms in Section \ref{sec:dis}. | \label{sec:con} We have carried out interferometric observations of cyanopolyynes (HC$_{3}$N, HC$_{5}$N, and HC$_{7}$N) toward the G28.28$-$0.36 high-mass star-forming region with the VLA Ka-band. We obtained the moment zero images of HC$_{3}$N and HC$_{5}$N and tentatively detected HC$_{7}$N. The spatial distributions of HC$_{3}$N and HC$_{5}$N are consistent with the 450 $\mu$m dust continuum clump, i.e., the Cyanopolyyne-rich clump. The HC$_{3}$N : HC$_{5}$N : HC$_{7}$N column density ratios are estimated at 1.0 : $\sim 0.3$ : $\sim 0.2$ at Position A. The Cyanopolyyne-rich clump seems to contain deeply embedded low- or intermediate-mass protostellar core(s). The most probable formation mechanism of the cyanopolyynes at the Cyanopolyyne-rich clump is the WCCC mechanism. We possibly found the HC$_{3}$N and HC$_{5}$N emission in the shock region. | 18 | 8 | 1808.08435 |
1808 | 1808.01208_arXiv.txt | We examine the case of a random isotropic velocity field, in which one of the velocity components (the ``radial" component, with magnitude $v_z$) can be measured easily, while measurement of the velocity perpendicular to this component (the ``transverse" component, with magnitude $v_T$) is more difficult and requires long-time monitoring. Particularly important examples are the motion of galaxies at cosmological distances and the interpretation of Gaia data on the proper motion of stars in globular clusters and dwarf galaxies. We address two questions: what is the probability distribution of $v_T$ for a given $v_z$, and for what choice of $v_z$ is the expected value of $v_T$ maximized? We show that, for a given $v_z$, the probability that $v_T$ exceeds some value $v_0$ is $p(v_T \ge v_0 | v_z) = {p_z(\sqrt{v_0^2 + v_z^2}})/{p_z(v_z)}$, where $p_z(v_z)$ is the probability distribution of $v_z$. The expected value of $v_T$ is maximized by choosing $v_z$ as large as possible whenever $\ln p_z(\sqrt{v_z})$ has a positive second derivative, and by taking $v_z$ as small as possible when this second derivative is negative. | Measuring radial velocities in astronomy through Doppler shifts in spectral lines is easy. Measuring transverse velocities through shifts in angular position is more challenging. Redshift measurements allowed Vesto Slipher to determine the radial velocities of distant galaxies more than 100 years ago, while the transverse motion of galaxies at cosmological distances has never been measured.\footnote{The transverse velocities of the much closer LMC and SMC have both been measured using the Hubble Space Telescope (Kallivayalil et al. 2006; Kallivayalil, van der Marel, \& Alcock 2006).} However, the latter situation may soon change (Darling, Truebenbach, \& Payne 2018). As noted previously by Sandage (1962) and Loeb (1998), precision redshift measurements taken over a significant time span would allow for a ``real time" measurement of the evolution of the Hubble parameter; there have been attempts to measure this effect using \ion{H}{1} 21 cm absorption line redshifts (Darling 2012). Similarly, precision astrometry might soon allow for the measurement of galactic proper motions (and hence, galaxy transverse velocities) in real time (Peebles et al. 2001; Nusser, Branchini, \& Davis 2012; Quercellini 2012; Darling \& Truebenbach 2018; Darling, et al. 2018). [The measurement of tranvserse velocities of distant galaxies using microlensing was considered by Grieger, Kayser, \& Refsdal (1986) and Gould (1995), while Hamden, et al. (2010) explored the possibility of using perspective rotation in clusters]. The possibility of making transverse velocity measurements with Gaia is discussed in detail by Nusser et al. (2012), while Darling et al. (2018) examine the potential for ngVLA to measure these transverse velocities. This leads to an obvious question: given the opportunity to monitor a limited set of galaxies with known radial velocities $v_z$, which of these are most likely to have the largest transverse velocities $v_T$? (Here $v_z$ corresponds to the peculiar line-of-sight velocity, with the contribution from the Hubble expansion subtracted off.) This question is perhaps less relevant for an all-sky survey such as Gaia, but other instruments such as ngVLA might monitor a limited sample of distant galaxies. Given a wide dispersion in the magnitudes of the total velocity, $v$, one might naturally assume that the galaxies with the largest radial velocities would also tend to have the largest transverse velocities. However, if $v$ is narrowly distributed around a single value, then the largest radial velocities would correspond to the {\it smallest} transverse velocities. This is easily seen for the case where $v$ is identically the same for all of the objects in question; in this case $v_T = \sqrt{v^2 - v_z^2}$ is maximized when $v_z=0$. Which of these two arguments is correct? Both of them are relevant. As we will see, the largest transverse velocities can correspond to either the largest or the smallest values of the radial velocity, $v_z$, depending on the distribution of $v_z$. This paper addresses the following questions: given an isotropic random velocity field, along with a known distribution of radial velocities $v_z$, what is the corresponding distribution of transverse velocities, $v_T$, and for what choice of $v_z$ is the expected value of $v_T$ maximized? While this discussion is motivated within the context of galaxy velocities, these questions are quite general, and it seems likely that our results would be applicable to other areas of astronomy as well, such as star clusters. We address these questions mathematically in the next section, and briefly discuss our results in Section 3. | We have derived an expression for the distribution of the transverse velocity, $v_T$, for a given fixed value of the radial velocity, $v_z$, valid for any isotropic velocity distribution (or indeed, for any isotropic vector field) in equation (\ref{CDF}). Our results indicate that the expected value for $v_T$ can be maximized by choosing the largest possible value of $v_z$ if $v_z$ has a strongly super-Gaussian distribution, and for the smallest possible value of $v_z$ if the distribution is strongly sub-Gaussian, where these terms are defined in the previous section. We now circle back to the question which originally motivated this investigation: what about the peculiar velocity field of galaxies? While current observations are beginning to probe this distribution (e.g., Tully, et al. 2013; Springob, et al. 2014; Tully, Courtois, \& Sorce 2016), the data are still too noisy to provide a precise estimate of $p_z(v_z)$. The uncertainties in the measured peculiar velocities are typically of order the velocities themselves at cosmological distances (Watkins \& Feldman 2015). However, this problem can be mitigated by binning the velocity data. Using the catalog of Tully, et al. (2013), Sorce (2015) derived a bias-corrected distribution for $v_z$ which is consistent with a Gaussian distribution. This is precisely the unique distribution for which the value of $v_T$ is insensitive to the value of $v_z$. It is also consistent with the theoretical model of Sheth and Diaferio (2001), which predicts a form for $p_z(v_z)$ that looks Gaussian at small $v_z$. However, their model also predicts an exponential distribution for $p_z(v_z)$ at large $v_z$. For the exponential distribution, we expect that $v_T$ will be largest when $v_z$ is maximized. This suggests that if one were monitoring a limited set of galaxies over a long time span, efforts should be concentrated on those with the largest radial pecular velocities. Future data sets to which these results might be applied include measurements of radial peculiar velocities from distance calibrators such as SN Ia (Riess 1999) or gravitational wave sources (Chen, Fishbach, \& Holz 2018) or from the kinetic Sunyaev-Zel'dovich effect (Akrami, et al. 2018). The derivations presented here can also be used as a constraint on models of the peculiar velocity field of galaxies in the standard $\Lambda$CDM cosmology, such as those in Sheth \& Diaferio (2001). In addition, our derivations can be applied to new astrometric data from the Gaia satellite on the proper motion of stars in globular clusters and dwarf galaxies (Helmi, et al. 2018) in an attempt to constrain their mass distribution (Milone, et al. 2018) or rotation (Bianchini, et al. 2018) as well as the possible existence of an intermediate black hole at their center (e.g., Kiziltan, et al. 2017). | 18 | 8 | 1808.01208 |
1808 | 1808.07797_arXiv.txt | {} {In this article we aim to show how the gradient of the thermal millimetre continuum spectrum, as emitted from the quiet solar atmosphere, may be used as a diagnostic for the optical thickness regime at the centre of the observing frequency band.} {We show the theoretical derivation of the gradient of the millimetre continuum for both logarithmic- and linear-scale spectra. We compare this expression with the empirical relationship between the slope of the millimetre continuum spectrum and the plasma optical thickness computed from both isothermal and multi-thermal two-dimensional cylindrical radiative transfer models.} {It is found that the logarithmic-scale spectral gradient provides a clear diagnostic for the optical thickness regime for both isothermal and multi-thermal plasmas, provided that a suitable correction is made for a non-constant gaunt factor over the frequency band. For the use of observers we present values for this correction at all ALMA bands and at a wide range of electron temperatures.} {We find that the spectral gradient can be used to find (a) whether the source is fully optically thin, (b) the optical thickness of the source if it lies within the transitional regime between optically thin and thick plasma ($\tau \approx 10^{-1} - 10^{1}$), or (c) whether the source is fully optically thick for an isothermal plasma. A multi-thermal plasma will act the same as an isothermal plasma for case (a), however, the transitional regime will only extend from $\tau \approx 10^{-1} - 10^{0}$. Above $\tau=1$ the slope of the continuum will depend increasingly on the temperature gradient, as well as the optical thickness, reducing the reliability of the diagnostic.} | 18 | 8 | 1808.07797 |
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1808 | 1808.00381_arXiv.txt | { Giant planets open gaps in their protoplanetary and subsequently suffer so-called type~II migration. Schematically, planets are thought to be tightly locked within their surrounding disks, and forced to follow the viscous advection of gas onto the central star. This fundamental principle however has recently been questioned, as migrating planets were shown to decouple from the gas' radial drift. } { In this framework, we question whether the traditionally used linear scaling of migration rate of a giant planet with the disk's viscosity still holds. Additionally, we assess the role of orbit-crossing material as part of the decoupling mechanism. } { We have performed 2D ($r,\theta$) numerical simulations of point-mass planets embedded in locally isothermal $\alpha$-disks in steady-state accretion, with various values of $\alpha$. Arbitrary planetary accretion rates were used as a means to diminish or nullify orbit-crossing flows. } { We confirm that the migration rate of a gap-opening planet is indeed proportional to the disk's viscosity, but is not equal to the gas drift speed in the unperturbed disk. We show that the role of gap-crossing flows is in fact negligible.} { From these observations, we propose a new paradigm for type~II migration\,: a giant planet feels a torque from the disk that promotes its migration, while the gap profile relative to the planet is restored on a viscous timescale, thus limiting the planet migration rate to be proportional to the disk's viscosity. Hence, in disks with low viscosity in the planet region, type~II migration should still be very slow. } | Planetary migration is a key ingredient to understand the architecture of planetary systems. This radial displacement of planets is due to their gravitational interaction with the protoplanetary disk. These disks surround most young stars, and have a lifetime of a few million years. Planetary migration leads to significant changes in the semi-major axis of all planets \citep[see][for a recent review]{baruteau+2014} and carves the structure of planetary systems. Migration of planets has been extensively studied in recent decades. Small mass planets, for which the response of the disk can be considered linear, do not perturb the density profile of the disk, and are in a regime called type~I migration. Giant planets, however, are massive enough to modify the disk radial density profile. They deplete the region around their orbit and create a gap \citep{lin-papaloizou1986a}, separating the inner from the outer disk. Once a gap is open, the planet is repelled inwards by the outer disk and outwards by the inner disk. The position of the planet within the gap adjusts so that the torques from the inner and out disks cancel out. However, as the disk spreads viscously and the gas accretes onto the central star, the gap, as well as the embedded planet that carved it, are carried with it. This is the classical scheme of the so-called type~II migration \citep{lin-papaloizou1986b}, responsible for inward motion of giant planets. In this scheme, the planet does not migrate with respect to the gas, but together with the gas, and acts as a gas-proof barrier between the parts of the disk. This standard scheme of type~II migration, where a planet follows exactly the viscous accretion speed of the gas has been questioned by several works. \citet{quillen+2004} note that if the inertia of the planet is much larger than that of the gas originally present in the gap, the disk has a hard time moving the planet, and the migration is slower than the viscous speed. \citet{crida-morbidelli2007} add that the corotation torque, exerted on the planet by the gas still present in the gap, may play a role, especially in regions where the background density profile is steep (which promotes a high corotation torque). This could slightly decouple the planet from the gas evolution, but this process relies on a non-empty gap, hence it could be seen as a situation where perfect type~II migration is not expected anyway. Furthermore, \citet{hasegawa-ida2013} remark that the assumption that a planet be locked in its gap had no solid physical ground. \citet{lubow-dangelo2006} and \citet{duffell+2014} show that, in simulations, gas is able to cross the gap during planetary migration. \citet{duermann-kley2015} (DK15 hereafter) explicitly question the idea that the planet stays in equilibrium in the middle of the gap. They suggest that when the gas reaches an equilibrium gap profile, such that the torques from the viscosity, the pressure, and the planetary gravity balance on each side \citep{crida+2006}, the planet does not necessarily feel a zero torque. Hence it moves to a different position inside the gap. The motion of the planet then forces the gas to re-adjust the equilibrium profile, by passing through the planet's orbit from the inner to the outer part of the disk. Because this transfer of gas is due to the planet-gas interaction and not to the disk's viscous spreading, the evolution of the planet becomes unlocked from the disk's viscous evolution. The author of DK15 find that the migration speed is independent of the disk's drift speed and mainly depends on parameters such as the mass of the disk or that of the planet. Furthermore, \cite{duermann-kley2017} (DK17 hereafter) showed that planetary accretion, in some cases, is able to cut the gas flow across the planet's orbit. In these cases the gap acts as a barrier between the inner and the outer disks and classical type~II migration regime could be reestablished. However the authors noticed that even in these cases the migration rate can differ from classical type~II migration rate. Both studies (DK15 and DK17) considered a classical $\alpha$-viscosity disk and focused on the dependence of migration speed on parameters like the disk mass and the planetary mass for a fixed viscosity value. These studies show that gas is crossing the gap and therefore the migration speed seems to be independent on the disk's drift. Although the authors already provided evidence that migration speeds depend on viscosity, a direct comparison of this dependence with the fundamental assumptions of classical type~II remains to be conducted. This is precisely the aim of the present paper. In fact, though it is now admitted that giant planets' migration does not follow classical type~II migration, it is very important to check whether some scaling of migration rate with viscosity is still preserved. Precisely, the existence of giant planets at orbits larger than $1$ a.u. in semi-major axis (so called warm Jupiters) is difficult to explain in the usual paradigm of viscously accreting disks unless considering low viscosity and confirming that migration speed scales with viscosity. Low viscosity is admitted in the central part of the disks, the so-called dead-zone, where magneto-rotational instability \citep{BalbusHawley91} does not operate. Moreover, recent studies \citep{BaiStone2013,Bai2016,Suzuki+2016} have shown that the disk's structure can be very different from that of a viscously accreting disk. Mass accretion onto the star could be ensured by magnetically driven winds providing angular momentum removal. These disks would have very low viscosity. It is beyond the purpose of this paper to model this kind of disks, however these studies motivate the investigation of the migration of giant planets in low viscosity disks. We will also investigate the role of gap crossing flows on migration. The paper is organized as follows. In \Cref{sec:setup}, we present the physical model and the numerical scheme, while the setup of initial conditions is reported in the appendix. In \Cref{sec:viscosity}, we study the scaling of migration speed with viscosity by performing numerical simulations of Jupiter-mass planets in disks with controlled inflow rates and various viscosities, but same gas surface density. In \Cref{sec:accretion}, the accretion of gas by the planets is modeled by removing, with various efficiencies, the gas entering the Hill sphere\,; it allows us to find the influence of planetary accretion on migration, and to quantify how important is cutting the gas flow across the gap. Finally, our findings are put summarized in \Cref{sec:conclu}, where we propose a new, consistent paradigm to explain giant planet migration. | \label{sec:conclu} To summarize, we have found in \Cref{sec:viscosity} that although type~II migration speed is proportional to the gas' viscosity, it is not driven by the radial inward drift of the gas. In particular, we confirmed that the giant planet can actually migrate faster than the gas drifts, and even in a stationary disk. We concluded that what drives type~II migration is the imbalance between the torques felt by the planet from the inner and the outer disk, as pointed out by DK15. However, the width and shape of the gap is not directly linked to the viscosity, especially at low $\nu$ where the pressure effects are dominant \citep{crida+2006}\,; hence, we do not expect this torque imbalance to be proportional to $\nu$, in contrast with the observed migration speed. In \Cref{sec:accretion}, we have seen that gap-crossing flows are actually negligible at low viscosity, and that cutting this small gas flow with planetary accretion hardly impacts the migration speed. Thus, the planet migrates faster than the disk drift, even when no gas is exchanged between the inner and outer disks. Gap-crossing flows can not be responsible for the observed fast migration, in contrast with the of case type~III migration \citep{masset-papaloizou2003}. These two results allow to us draw a new, consistent picture of type~II migration. As a giant planet forms, it opens a gap by perturbing the gas profile with the gravitational torque it exerts. The gas reaches a new equilibrium profile on each side of the gap. Nonetheless, the planet inside its gap feels a non-zero torque, because the inner and the outer torques have no reason to balance out \citep[as recently studied by][]{kanagawa+2018}. Thus, the planet has to migrate inwards. As it does so, some gas may cross the gap from the separatrix of the HSR, although this is not enough to restore the initial gap profile in the frame of the planet if the viscosity is low ad the gap is wide (regardless of whether the planet accretes or not). Therefore, the density distribution has to adapt to the new position of the planet\footnote{i.e. the bump in the surface density produced at the inner edge of the gap has to be redistributed over the inner disk while the outer disk has to spread down to the new gap's outer edge.}, and this is done over a viscous time. Once the gas is again at equilibrium with the planet, the planet is not in an equilibrium inside the gap anymore, and we are back to the initial situation. In this scheme, the planet may well migrate faster than the gas drifts, because it is pushed by a torque that has no connection with the drift of an unperturbed disk. But because the gap-crossing flow is negligible, the planet must migrate at a rate proportional to viscosity, otherwise it would pile gas up in the inner disk and leave a depleted outer disk behind, eventually halting its migration. Although the final result (migration speed of gap-opening planets is proportional to the viscosity) is in line with the standard picture of type~II migration, this new scheme is conceptually revolutionary in our understanding of this phenomenon, and allows us to reconcile all the puzzling observations that have been made recently, questioning the standard picture. Additionally, it confirms that even if some gas may cross the gap, a giant planet in a low viscosity disk should migrate slowly. In this frame, the abundance of warm Jupiters, who did not migrate all the way towards their star, may suggest that most protoplanetary disks have a low effective viscosity in the planet-forming region. | 18 | 8 | 1808.00381 |
1808 | 1808.09743_arXiv.txt | { Our velocity relative to the {cosmic microwave background} (CMB) generates a dipole from the CMB monopole, which was accurately measured by COBE. {The} relative velocity also modulates and aberrates the CMB fluctuations, generating a {small} signature of statistical isotropy violation in the covariance matrix. This signature was first measured by Planck 2013. Galaxy surveys are similarly affected by a Doppler boost. The dipole generated from the number count monopole has been extensively discussed, and measured (at very low accuracy) in the NVSS and TGSS radio continuum surveys. {For the first time,} we present an analysis of the Doppler imprint on the number count fluctuations, using {the} bipolar spherical harmonic formalism to quantify these effects. Next-generation wide-area surveys with a high redshift range are needed to detect the small Doppler signature in number count fluctuations. We show that radio continuum surveys with {the SKA should enable a detection at $\gtrsim 3 \sigma$ in Phase 2, with marginal detection possible in Phase 1.}} \begin{document} | \label{sec:intro} Our motion relative to the cosmic microwave background (CMB) rest frame gives rise to a Doppler boost effect on the radiation. {The maximum anisotropy detected is at the level $\sim 10^{-3}$, from the dipole.} This large amplitude indicates that we are moving relative to the CMB rest frame -- the {one} in which the dipole vanishes. Assuming that the dipole is entirely due to Doppler boosting of the CMB monopole, the velocity was measured from the COBE experiment as $v=369 \pm 0.9\,$km/s towards galactic coordinates $(l,b)=(263.99^{\circ}\pm 0.14^{\circ}, 48.26^{\circ}\pm 0.03^{\circ})$~\citep{Kogut1993,Hinshaw2009}. A Doppler boost generates not only a dipole {anisotropy} from the CMB monopole, but also affects the primordial fluctuations of the CMB {by generating a distinctive statistical isotropy violation signature of aberration and modulation.}. The effects are small -- a $\sim 10^{-3}$ correction to $\sim 10^{-5}$ fluctuations. There are two contributions~\citep{Kosowsky2011,Amendola2011, Notari2012, Planck2014}: \begin{itemize} \item Doppler modulation amplifies the fluctuations in the boost direction, and suppresses them in the {opposite} direction {(the same effect generates the kinematic dipole from the CMB monopole)} \item {Doppler aberration deflects} the arrival direction of photons towards the boost direction, distorting the anisotropy pattern. \end{itemize} These imprints on the CMB fluctuations were not detectable with COBE and WMAP, but Planck's high angular resolution and sensitivity and low noise levels allowed the first measurement of this velocity signature in 2013~\citep{Planck2014} (confirmed in Planck 2015~\citep{2016A&A...594A..16P}). Our motion also generates a dipole asymmetry in observed number of clusters which has been studied in \citep{Chluba2005}. Our motion affects not only the CMB but also galaxy surveys. According to the standard model, the rest frames of the CMB and of matter should coincide. This provides a critical test of statistical isotropy, as first pointed out by~\citep{Ellis1984}. Radio continuum surveys have been used as an alternative independent probe to extract the cosmic kinematic dipole. Pioneering tests were performed by~\cite{Baleisis1998}, but the detection of a kinematic dipole in NVSS radio number counts was only confirmed later by~\citep{Blake2002}. Subsequent analyses revealed a $>2\sigma$ discrepancy between the predicted and measured amplitudes of the velocity~\citep{Singal2011,Gibelyou2012,Rubart2013,Rubart2014,Tiwari2014,Tiwari2015,Tiwari2016,Colin2017,Bengaly2018}. The nonlinear effects of local large-scale structure and the allowance for flux density and calibration errors have not resolved the discrepancy~\citep{Rubart2014,Bengaly2018}. Current radio surveys do not have sufficient number density, sensitivity and resolution for a robust determination of the dipole amplitude and direction from the Doppler effect on the monopole. The SKA is forecast to measure this amplitude and direction with sufficient accuracy to detect significant deviations from theory~\citep{Schwarz2015,SKA}. The measurement of the Doppler boost vector using the method described here is less susceptible to local structure bias. The unprecedented angular resolution and sensitivity of the SKA continuum survey will facilitate the detection of the much weaker Doppler effect on the number count fluctuations. This has not previously been investigated. Here we present for the first time a forecast for such a detection in future SKA radio continuum surveys. | \label{sec:discussion} {There is a potential `contamination' of the unboosted map by galaxy peculiar velocities and by the lensing magnification effect on number counts. For a continuum survey, the peculiar velocities (and therefore RSD) are nearly averaged out by projection of number counts onto the unit sphere giving rise to a few percent change in power spectra.} Lensing magnification affects number counts through a change in solid angle and through magnification bias $s(z)$ (see Fig.~\ref{fig:nz_bz_sz}), which can move sources into or out of the map. This effect potentially leads to a bias on the estimate of the Doppler aberration effect, and in principle, the number count map should be de-lensed before applying the Doppler boost estimator. For SKA1, the lensing effect on continuum number counts is below the signal to noise, as shown in~\cite{Alonso2015} (see their Fig.~17), and can be safely neglected. For SKA2, the effect is above signal to noise unless the error on the magnification bias is $\gtrsim 20\%$. We computed the signal to noise forecast for detection of the Doppler boost, including both lensing and RSD effects. In summary, fluctuations in number count maps are affected by our motion with respect to the CMB rest frame. For the first time, we have proposed a method to extract the Doppler boost, {using the statistically anisotropic signature induced by modulation and aberration of fluctuations in number counts} from radio continuum surveys. Our forecasts indicate that SKA2 will be able to {detect} the Doppler field at a level that matches Planck at $\ell \sim 1000$, and improves on Planck if higher multipoles are accessed: \begin{eqnarray} \mbox{SKA2 detection:} \quad \gtrsim 3\sigma \quad \mbox{for}~~\ell_{\rm max} \gtrsim 1000\,. \end{eqnarray} For SKA in Phase 1, a marginal detection may be possible, as shown in Fig.~\ref{fig:10ujy}. Such a detection would provide an important further test of the Cosmological Principle, which dictates consistency between the CMB and the matter distribution. \newpage \subsection* | 18 | 8 | 1808.09743 |
1808 | 1808.05271_arXiv.txt | We present the first results from the GeMS/GSAOI Galactic Globular Cluster Survey (G4CS) of the Milky-Way globular clusters (GCs) NGC\,3201 and NGC\,2298.~Using the \textit{Gemini South Adaptive Optics Imager} (GSAOI), in tandem with the \textit{Gemini Multi-conjugate adaptive optics System} (GeMS) on the 8.1-meter Gemini-South telescope, we collected deep near-IR observations of both clusters, resolving their constituent stellar populations down to $K_s\simeq21$ Vega mag.~Point spread function (PSF) photometry was performed on the data using spatially-variable PSFs to generate $JHK_{s}$ photometric catalogues for both clusters.~These catalogues were combined with \textit{Hubble Space Telescope} (HST) data to augment the photometric wavelength coverage, yielding catalogues that span the near-ultraviolet (UV) to near-infrared (near-IR).~We then applied $0.14$ mas/year accurate proper-motion cleaning, differential-reddening corrections and chose to anchor our isochrones using the lower main-sequence knee (MSK) and the main-sequence turn-off (MSTO) prior to age determination.~As a result of the data quality, we found that the $K_{s}$ vs.~F606W~$\!-K_{s}$ and F336W vs.~F336W~$\!-K_{s}$ color-magnitude diagrams (CMDs) were the most diagnostically powerful.~We used these two color combinations to derive the stellar-population ages, distances and reddening values for both clusters.~Following isochrone-fitting using three different isochrone sets, we derived best-fit absolute ages of $12.2\pm0.5$ Gyr and $13.2\pm0.4$ Gyr for NGC\,3201 and NGC\,2298, respectively.~This was done using a weighted average over the two aforementioned color combinations, following a pseudo-$\chi^2$ determination of the best-fit isochrone set.~Our derived parameters are in good agreement with recent age determinations of the two clusters, with our constraints on the ages being or ranking among the most statistically robust. | \label{sec:intro} \subsection{Complexity of GC Stellar Populations} The study of resolved globular clusters (GCs) provides unique leverage on some of the most pressing questions in astronomy, ranging from our understanding of stellar structure and evolution, to the details of star formation and chemical enrichment processes, to the formation histories of galaxies and the beginnings of structure formation in the universe \citep[e.g.][and references therein]{gra04, ren13, ren17, kru14, van17}.~The {\it Hubble Space Telescope UV Legacy Survey of Galactic Globular Clusters} \citep{pio15} and the {\it Hubble Space Telescope ACS Survey of Galactic Globular Clusters} \citep{sar07} obtained deep near-UV\footnote{F275W[UV] and F336W[$U$] WFC3 filters} and optical\footnote{F438W[$B$] WFC3, and F606W[$V$] and F814W[$I$] ACS filters} photometry for about 60 Milky Way GCs, enabling the creation of high-quality panchromatic color-magnitude diagrams (CMDs).~Analogous to the theoretical stellar evolution-tracing Hertzsprung-Russell (HR) diagrams, deep CMDs can be used to test theories of stellar evolution and star cluster formation.~Individual GCs constitute the best available approximations to simple stellar populations, being composed of coeval stars with virtually the same chemical composition.~As such, they frequently serve as calibrators for stellar population synthesis models which can then be used to study distant, unresolved stellar systems.~The deep CMDs that arose from the Hubble Space Telescope (HST) surveys have produced a range of important benchmark measurements and several unexpected findings.~For instance, it was found that a number of Galactic GCs harbour multiple stellar populations \citep[e.g.][]{mil08, pio15}, a result with far reaching consequences for our understanding of their formation \citep[e.g.][]{ren15} and with close ties to star formation physics.~The second prominent result was the determination of GC relative ages with sufficient precision to directly probe the formation and evolution of the galactic halo \citep[e.g.][]{mar09, dot10, dot11, lea13}.~These results suggest stark differences between the early formation history of the inner and outer halo, providing further evidence that the evolution of the outer halo involved the disruption and accretion of dwarf satellite galaxies \citep[e.g.][]{sea78, fre02, hel08, ive12}. \subsection{The Role of Near-Infrared Observations} Near-infrared (near-IR) CMDs offer significant advantages when compared to pure UV-optical CMDs, such as:~{\it i}) near-IR colors are much less affected by differential reddening than optical colors and~{\it ii}) near-IR isochrones exhibit a very characteristic S-shaped main sequence (MS) which can be used for accurate, {\it absolute} age determinations.~This characteristic shape is attributed to low-mass MS stars, which show a well-defined (faint-ward) bending, due to collision-induced and roto-translational absorption of hydrogen molecules (H$_2$-H$_2$ and other molecular pairs, such as: H$_2$-He, H$_2$-CH$_4$, etc.)~at near-IR wavelengths \citep[also referred to as the MS ``knee'' or MSK;][]{bor01, bor02, gus09, fro10, ric12}.~The resultant difference in color and luminosity between the MS turn-off (MSTO) point and the lower MSK depends predominantly on metallicity, while the luminosity and color of the MSK are virtually independent of age at any given metallicity \citep{cal09}.~Additionally, because the convective motions in the stellar atmospheres are mostly adiabatic at low stellar masses \citep{sau08}, the MSK is anchored in a regime of well understood stellar structure and evolution.~Therefore, by exploiting the metallicity-dependence-dominated MSTO-MSK distance, the age-independence of the MSK at fixed metallicity and its well understood physics, the observed MSTO-MSK color and luminosity difference form a robust {\it absolute} theoretical reference frame which allows for the measurement of absolute GC ages with uncertainties of about 1~Gyr.~This remarkable property of near-IR CMDs was recently exploited by \cite{mas16}, \cite{cor16}, and \cite{sara16} using various isochrone fitting techniques.~Finally, because near-IR CMDs contain the two aforementioned MS reference points, i.e.\ the MSTO and MSK, the age information extracted from their analysis becomes independent of cluster distance and screen-type reddening effects. \subsection{The GeMS/GSAOI Galactic Globular Cluster Survey} In this paper we describe the first results from the {\it GeMS/GSAOI Galactic Globular Cluster Survey} (G4CS).~The G4CS is conducting a homogeneous, deep near-IR CMD study of Milky Way GCs covering a broad spectral energy distribution (SED) baseline in order to: {\it i}) investigate the morphology of near-IR CMDs, particularly at low stellar masses, {\it ii}) determine the internal consistency and calibration of model CMD predictions for various near-IR+optical+near-UV filter combinations, thereby testing the influence of variations in the molecular absorption bands on isochrones using information from publicly-available high-resolution spectroscopy \citep[e.g.][]{car09, car10},~{\it iii}) derive absolute GC ages with accuracies of about 1~Gyr,~and {\it iv}) quantify and characterize binary fractions and blue-straggler star (BSS) populations. The G4CS target list was compiled through selecting GCs included in the recent HST surveys of Galactic GCs, to ensure that deep near-UV+optical photometry is available \citep{sar07,pio15}, with the additional restriction that the GCs lay no further than 23~kpc to ensure that the MSK can be detected in a reasonable amount of time with GeMS/GSAOI at the 8.1-meter Gemini-South telescope.~In order to study as broad a parameter space as possible, the target sample covers a range in GC ages ($t\!\simeq\!8.0\!-\!13.5$~Gyr), metallicities ($[{\rm Z/H}]\!\simeq\!-2.37$ to $-0.32$~dex), and masses (${\cal M}_{GC}\!\simeq\!10^{3.75}\!-\!10^{6.05}M_{\odot}$, see also Fig.~\ref{fig:sample}). \begin{figure}[t!] \centering \includegraphics[width=\linewidth]{Figure1.pdf} \caption{The age-metallicity relationship for our full G4CS target sample with available data from \cite{dot10}.~The two GCs from this study are circled and correspondingly labeled.~The symbol size scales with the $V$-band integrated luminosity; larger symbols correspond to brighter GCs.~Note that the whole sample covers the transition region between the flat and steep age-metallicity relation.~Utilizing our excellent age/metallicity resolution, we aim to clearly resolve the transition between these two GC sub-populations.} \label{fig:sample} \end{figure} This paper presents the results for the first two GCs observed from the sample: NGC\,2298 and NGC\,3201.~These two GCs were selected, in part, due to little indication of multiple stellar populations in their near-UV/optical CMDs \citep{pio15}, making them two of the few GCs that may most closely resemble simple stellar populations.~Both clusters have HBs that could be influenced by the ``2nd parameter'' effect and NGC\,3201 may be about a Gyr younger than the oldest GCs \citep{mil14}.~This implies that NGC\,3201 is a candidate for having been formed in a dwarf galaxy and later accreted into the Milky Way halo \citep{mac04}.~Indeed, evidence of extra-tidal stars has been found around both clusters \citep{balb11,ang16}.~Both clusters also experience relatively high reddening ($E_{B-V} \simeq 0.2$\,mag), making it advantageous to study them in the near-IR.~Additionally, NGC\,3201 may be considered a control cluster since it has been used by both \cite{cal09} and \cite{bon10} to demonstrate the MSTO-MSK method using VLT MAD and HST, respectively.~Therefore, being able to compare our analysis to these earlier results will allow us test the efficacy of our methods.~The properties of the two GCs studied thus far are presented in Table~\ref{tab:gcprop}. \begin{deluxetable}{lrr} \tablecaption{Fundamental GC parameters. \label{tab:gcprop}} \tablehead{ \colhead{Parameter} & \colhead{NGC 3201} & \colhead{NGC 2298} } \startdata Distance [kpc] & 4.9 & 10.8\\ $(m-M)_V$ & 14.2 & 15.6\\ $E_{B-V}$ & 0.24 & 0.14\\ $r_h$ [pc] & 4.4 & 3.1\\ $M_{V}$ & $-7.45$ & $-6.31$\\ $[{\rm Fe/H}]$ & $-1.59$ & $-1.92$\\ \enddata \tablecomments{All values are taken from the Milky Way GC McMaster Catalog \cite{har10}. Note that the distance moduli and absolute magnitudes are not corrected for reddening, while the distance has been corrected.} \end{deluxetable} \begin{figure*}[t] \centering \includegraphics[width=0.49\linewidth]{Figure2a.pdf} \includegraphics[width=0.49\linewidth]{Figure2b.pdf} \caption{HST/WFC3 stacked UV+optical (F275W+F336W+F438W) greyscale images of NGC\,3201 ({\it left}) and NGC\,2298 ({\it right}) from \citet{sot17}.~The red squares represent the area covered by our GeMS/GSAOI observations.} \label{fig:hstfields} \end{figure*} This paper is organized as follows: the observations and data reduction are presented in Section~\ref{sec:obs}, in Section~\ref{sec:anls} we describe the various experiments undertaken to obtain the best PSF photometry given the characteristics of our observations, describe zero-point calibration, proper-motion cleaning and differential reddening corrections, Section~\ref{sec:results} describes our methodology and results of isochrone-fitting and, finally, Section~\ref{sec:sum} concludes with a summary of the work. | \label{sec:sum} The multi-conjugate adaptive optics system, near-IR imager combination, GeMS/GSAOI, mounted on the 8.1-m Gemini-South telescope was used to observe two Milky Way globular clusters, NGC\,3201 and NGC\,2298, reaching a depth of $K_s\!\simeq\!21$ Vega mag for both targets.~Spatially variable PSFs were created for both clusters using \texttt{DAOPHOT-IV} in order to perform PSF-fitting photometry, yielding high-quality photometric results.~The resulting photometric catalogues were combined with HST-near-UV data from \cite{pio15} and HST-optical data from \cite{sar07} to create panchromatic stellar photometry libraries for both clusters spanning near-UV to near-IR wavelengths.~Zero-point calibrations, proper motion cleaning and differential reddening corrections were applied to said catalogues to provide clean, precise photometry and to ensure that they contained only high-probability GC member stars.~Absolute age determinations were then performed, utilizing a characteristic scale length defined as the distance from MS-saddle to MSTO for both clusters.~The lower MSK was recovered in both clusters, for the first time in NGC\,2298, highlighting the diagnostic power of AO-supported near-IR data.~Three different isochrone sets (\texttt{DSED}, \texttt{VR} and \texttt{BaSTI}) were used in combination with two CMDs created using UV-near-IR (F336W vs.~F336W$-K_{s}$) and optical-near-IR ($K_s$ vs.~F606W$-K_{s}$) filter combinations to derive values of absolute age, distance and reddening, adopting values of [Fe/H] and [$\alpha$/Fe] derived via high-resolution spectroscopy of member stars.~The most internally consistent results were determined using the \texttt{DSED} isochrone-based measurements, yielding an age, reddening, and distance determination for NGC\,3201 of $12.2\pm0.5$ Gyr, $E_{B-V}=0.25\pm0.01$ mag, and $5.1\pm0.1$ kpc, respectively.~The \texttt{DSED} isochrones were also found to be the most internally consistent set for NGC\,2298, yielding age, reddening and distance determinations of $13.2\pm0.4$ Gyr, $E_{B-V}=0.20\pm0.01$ mag, and $10.6\pm0.2$ kpc, respectively.~We find very good agreement between literature values and our derived parameters, and show that our measurements are among the most statistically robust constraints determined thus far.~New observations for the G4CS are planned in order to apply this method to globular clusters with a wider range of metallicities and formation histories. | 18 | 8 | 1808.05271 |
1808 | 1808.09575_arXiv.txt | We observed two full orbital phase curves of the transiting brown dwarf KELT-1b, at \three and \fouralt, using the Spitzer Space Telescope. Combined with previous eclipse data from \cite{beatty2014}, we strongly detect KELT-1b's phase variation as a single sinusoid in both bands, with amplitudes of $964\pm36$\,ppm at \three and $979\pm54$\,ppm at \fouralt, and confirm the secondary eclipse depths measured by \cite{beatty2014}. We also measure noticeable Eastward hotspot offsets of $28.4\pm3.5$ degrees at \three and $18.6\pm5.2$ degrees at \fouralt. Both the day-night temperature contrasts and the hotspot offsets we measure are in line with the trends seen in hot Jupiters \citep[e.g.][]{crossfield2015}, though we disagree with the recent suggestion of an offset trend by \cite{zhang2018}. Using an ensemble analysis of Spitzer phase curves, we argue that nightside clouds are playing a noticeable role in modulating the thermal emission from these objects, based on: 1) the lack of a clear trend in phase offsets with equilibrium temperature, 2) the sharp day-night transitions required to have non-negative intensity maps, which also resolves the inversion issues raised by \cite{keating2017}, 3) the fact that all the nightsides of these objects appear to be at roughly the same temperature of 1000\,K, while the dayside temperatures increase linearly with equilibrium temperature, and 4) the trajectories of these objects on a Spitzer color-magnitude diagram, which suggest colors only explainable via nightside clouds. | Orbital phase curve observations are one of the few ways in which we can directly investigate the global climates of exoplanets. This is particularly important for strongly irradiated planets such as hot Jupiters, since there can be temperature differences of over one thousand degrees between their day- and nightsides. This is believed to drive noticeable atmospheric composition changes between the two hemispheres, to say nothing of radically altering the vertical temperature structure as a function of longitude. One critically important component of the day-to-night changes in hot Jupiters is the possible formation and clearing of clouds on their night- and daysides. Though the possible role of clouds in exoplanet atmospheres has been appreciated for quite some time \citep{burrows1997,marley1999}, much of the 3D modeling of hot Jupiter atmospheres has assumed they are cloud free \citep[e.g.,][]{showman2008,kataria2016}. As a practical matter, this is due to the Herculean task of constructing accurate 3D global circulation models (GCMs) that properly deal with ``just'' dynamics and radiative transfer \citep{showman2008}. The modeling effort to add self-consistent cloud physics, including condensation processes and size distributions, that link to the established radiative and dynamics codes is just getting underway \citep[e.g.][]{lee2016}. As a result, the results of Spitzer phase curve observations are usually contextualized using a framework of competing ``thermal-only'' effects within hot Jupiters' atmospheres. For example, there is now a well-established trend that hot Jupiters with higher zero-albedo complete heat redistribution equilibrium temperatures (i.e., higher stellar irradiation) also show a higher temperature contrast between their day- and nightsides \citep{perezbecker2013}. Both the early theoretical work that predicted this trend \citep{showman2002} and more recent GCM analyses \citep{komacek2016} explain this using differences in the atmospheric radiative timescales and the atmospheric advective \citep{showman2002} or drag \citep{komacek2016} timescales. Put another way, the temperature difference between the day- and nightside of a hot Jupiter is determined by the balance of how fast the atmosphere cools and how fast it moves heat to the nightside. In all these analyses it has been made clear that the inclusion of clouds has the potential to strongly affect the results of the simulations, and recently more effort has been devoted to incorporating clouds into GCMs. In particular, \cite{parmentier2013}, \cite{lee2016}, and \cite{macdonald2017} have all found that 3D or 2D atmospheric models of HD 209458b that include cloud physics do a better job of replicating that planet's emission spectrum than cloud free models. Recently, \cite{powell2018} described a general atmosphere model that couples dynamics, radiative transfer, and cloud physics -- and which predicts that hot Jupiters should generally possess a nightside cloud deck. Observationally, the presence of high altitude clouds along planetary terminators has been evident in transmission spectroscopy measurements for some time \citep[e.g.][]{kreidberg2014}, but the signatures of clouds in emission measurements have been more difficult to see. This is because the daysides of hot Jupiters -- which provide us with our best emission spectra -- are believed to be mostly cloud-free \citep{parmentier2016}, though \cite{beatty2017a} recently inferred nightside TiO condensation on Kepler-13Ab based on that planet's dayside emission. The consideration of global, day and night, cloud coverage has largely been driven by the availability of red-optical phase curve data from Kepler. Initially, the phases curves of some individual planets showed clear Westward hotspot offsets that seem to strongly indicate clouds \citep[e.g.][]{demory2013}. \cite{parmentier2016} performed an ensemble analysis of Kepler phase curve results and found that clouds were generally required to explain not only the reflection signals themselves, but also how the amplitude and offsets of the planetary phase curves changed with temperature in the Kepler bandpass. One notable result from the Kepler phase curve observations is the apparently variable cloud cover on HAT-P-7b \citep{armstrong2016}. Four years of Kepler data show that not only does the location of maximum flux shift by up to 80 degrees over hundreds of days, but the shape of the phase curve itself is also variable on the same time scale. \cite{armstrong2016} interpreted this as changes in the weather on HAT-P-7b as the planetary cloud cover changed both its extent and its location. Interestingly, their toy model to explain the observations required that the observable thermal emission from HAT-P-7b was being altered by the variable clouds -- and not just the reflected light signal. In the infrared, \cite{mendonca2018} reanalyzed \cite{stevenson2017}'s \three and \four Spitzer phase curves of WASP-43b using a toy-model for clouds, by approximating their presence as a constant additional atmospheric opacity on the planetary nightside. \cite{mendonca2018} found that their cloudy simulations agreed more closely with the low observed nightside thermal emission, as the inclusion of clouds caused the modeled nightside flux to be significantly lower than predicted in cloud-free atmosphere. The \cite{armstrong2016} and \cite{mendonca2018} results indicate that even at the thermal infrared wavelengths probed by Spitzer, we should be considering how clouds modulate the thermal emission of hot Jupiters. This was recently, again, evident in \cite{dang2018}'s single \four phase curve of CoRoT-2b, which displayed a Westward hotspot offset, and was taken to be evidence for clouds affecting the dayside thermal emission of CoRoT-2b. Additionally, if the long-term cloud variation that \cite{armstrong2016} saw in HAT-P-7b is representative of the entire population of hot Jupiters, then there is the distinct possibility that Spitzer phase curve results are observing the combination of short-term weather effects on top of the equilibrium climates of hot Jupiters. \begin{figure*}[t] \vskip 0.00in \includegraphics[width=1.05\linewidth,clip]{RawPlot.pdf} \vskip -0.0in \caption{The raw photometry we used for our analysis combined new observations covering an entire orbit at \three and \four with \three and \four eclipse photometry previously analyzed in \cite{beatty2014}. Both data sets display correlations between the measured intensity and the x- and y-pixel position of the stellar centroid, which are typical features of \three and \four Spitzer photometry.} \label{rawplot} \end{figure*} To investigate the role of thermal-only effects versus clouds in hot Jupiters, we therefore observed Spitzer phase curves of the transiting brown dwarf KELT-1b \citep{siverd2012}. KELT-1b is a 27.23\mj\ object, with a radius of 1.116\rj. This is a mild, but significant, radius inflation compared to brown dwarf model predictions at the KELT-1 system age of 1.65\,Gyr \citep{siverd2012}. In isolation and in the field, we would expect KELT-1b to have an effective temperature of $\sim850$\,K due to its internal heat \citep{saumon2008}, but KELT-1b is on a short, 1.27 day, orbit around a 6500\,K host star. Previous observations of KELT-1b's dayside emission thus show it to be considerably hotter than the expectation from internal heat alone, at 3200\,K, and identical to a field M5 spectrum \citep{beatty2017}. The broadband thermal dayside emission from KELT-1b has previously been observed by Spitzer at \three and \four \citep{beatty2014}, $K$ \citep{croll2015}, and $z^\prime$ \citep{siverd2012} eclipses. \cite{beatty2017} also observed an $\mathrm{R}\approx50$ $H$-band eclipse spectrum. The relatively high mass of KELT-1b gives it a surface gravity approximately 22 times higher than a typical hot Jupiter. This high gravity could potentially change the atmospheric dynamics of KELT-1b, but since KELT-1b also receives the same level of external irradiation as a hot Jupiter, it can serve as a direct test of atmospheric dynamics in this regime. | \begin{enumerate} \item The available set of \three and \four phase offset measurements -- including KELT-1b -- are consistent with the observed planets having a constant phase offset of 14 deg. for all planetary equilibrium temperatures, though with a high scatter. This conflicts with cloudless GCM predictions that cooler planets should show large ($\sim70$\,deg.) offsets \citep{zhang2018}. \item The low disk-integrated nightside fluxes measured for KELT-1b and other hot Jupiters require that the underlying latitudinally-average atmospheric intensity map show relatively sharp transition from a hot dayside to a nearly constant and cooler nightside, such that disk-integrated nightside observations cannot ``see'' the hot dayside atmosphere. This requirement solves the problem pointed out by \cite{keating2017} that many hot Jupiters appear to have negative nightside intensities -- if one assumes a sinusoidal intensity map. \item The day- and nightside brightness temperatures for all the planets at \three and \four show two remarkable trends (Figure \ref{daynight}). First, the dayside brightness temperatures show a clear linear trend as a function of planetary equilibrium temperature. Second, the nightside brightness temperatures in both bands are consistent with all the planets having constant, $\sim1000$\,K, nightsides. This provides a new way to view the well-known trend \citep{showman2002,perezbecker2013} towards higher day-night temperature contrasts at higher equilibrium temperatures (Figure \ref{contrasts}). Namely, that this trend is not primarily the result of changing heat redistribution or radiative and advective timescales in the planetary atmospheres \citep[e.g.,][]{perezbecker2013}, but rather arises because the daysides of hot Jupiters simply become hotter under increasing stellar irradiation, while the nightsides of all the hot Jupiters possess clouds that show nearly uniform $\sim1000$\,K emission. Note that nearly simultaneously with the submission of this manuscript \cite{keating2018} also described these day- and nightside temperature correlations and also suggested that they are evidence for nightside clouds. \item Using Gaia DR2 parallaxes for KELT-1b and the other planets we can trace the phase evolution of their atmospheres on a color-magnitude diagram (Figure \ref{combinedcmd}). These trajectory plots suggest that most of the planets have nightside colors that are explainable by the presence of clouds \citep{saumon2008}, and not as easily explained by the presence of atmospheric \methane\ \citep{cooper2006}. \end{enumerate} In principle, it is possible to explain most of these observational results using cloudless atmosphere models, by specifically adjusting the balance between the atmospheric thermal timescales \citep{komacek2016}. However, the agreement of all these trends over such a wide range of parameter space would require a physical balancing act too coincidental to be easily believable, so we therefore prefer the theory that this is the result of nightside clouds. As mentioned in the Introduction, modelers of hot Jupiter atmospheres have long expected that clouds should play a noticeable role in setting the thermal emission properties of these planets. However, in the past we have lacked observations of sufficient quantity and precision to justify the Herculean task of constructing GCMs that include the effects of clouds and cloud formation. The arguments above indicate, however, that we have reached a point in our observations of hot Jupiters that cloudless GCM predictions can no longer accurately capture the trends that we see in the data, as particularly evidenced by the apparently constant phase offsets (Figure \ref{offsets}) and nightside temperatures of all the hot Jupiters (Figure \ref{daynight}). The wide range of possible cloud properties will make modeling their effect on planetary emission a difficult task without detailed observational results. Perhaps the most direct line of attack for this problem would be to observe the nightside emission spectra themselves. For hot Jupiters, though the signature of the nightside emission is present in spectroscopic transit observations, the typical transmission spectroscopy signal should be much higher than the nightside emission signal in nearly all cases. KELT-1b is a notable exception to this since as a brown dwarf its transmission spectroscopy signal will be close to zero. One way to directly observe nightside emission spectra would therefore be to target transiting hot brown dwarfs. More generally, for the hot Jupiters one will need spectroscopic phase curve observations to identify the planetary emission levels in transit. Though time intensive, spectroscopic phase curve observations do have the added benefit of allowing us to directly see how clouds change the spectra of individual planets as they switch from clear to overcast skies, which provide invaluable input into modeling the cloud formation processes on these planets. Another potential method to observe the cloud dynamics on hot Jupiters would be to conduct repeat observations of the same targets. As noted in Section 4.1, nearly all of the planets that have Spitzer phase curve observations have only a single measurement in each \textsc{irac} channel. Repeated broadband phase curve observations could therefore allow us to determine the time-variability in hot Jupiter atmospheres, which presumably would be driven by changing clouds (i.e., weather). As described in Section 4.1, this would allow us to determine if the lack of a trend in the measured phase offsets seen in Figure \ref{offsets} is actually the result of weather-driven scatter in the observations. An alternate way to gauge the presence of weather on hot Jupiters would be to take repeated secondary eclipse spectra. The toy model proposed by \cite{armstrong2016} to explain the variability in HAT-P-7b's optical phase curve requires substantial cloud-driven modulation of the dayside thermal emission, which might be detectable in precise eclipse spectra collected at separated times. Finally, more indirect methods may also allow us to constrain some of the cloud properties in hot Jupiters. \cite{beatty2017a} suggested that determining the point at which atmospheric cold traps become effective could constrain the average cloud particle size on planetary nightsides. Precise eclipse mapping of a planetary dayside \citep[e.g.,][]{rauscher2018} would reveal both the longitudinal and latitudinal surface intensity maps for the dayside, potentially resolving the degeneracies in inverting phase curves to global intensity maps described in Section 4.2. \ | 18 | 8 | 1808.09575 |
1808 | 1808.00454_arXiv.txt | We present a multiwavelength study of 28 Galactic massive star-forming \HII regions. For 17 of these regions, we present new distance measurements based on Gaia DR2 parallaxes. By fitting a multicomponent dust, blackbody, and power-law continuum model to the 3.6~\micron\ through 10~mm spectral energy distributions, we find that ${\sim}34\%$ of Lyman continuum photons emitted by massive stars are absorbed by dust before contributing to the ionization of \HII regions, while ${\sim}68\%$ of the stellar bolometric luminosity is absorbed and reprocessed by dust in the \HII regions and surrounding photodissociation regions. The most luminous, infrared-bright regions that fully sample the upper stellar initial mass function (ionizing photon rates $N_C \ge 10^{50}~{\rm s}^{-1}$ and dust-processed $L_{\rm TIR}\ge 10^{6.8}$~\Lsun) have on average higher percentages of absorbed Lyman continuum photons ($\sim$51\%) and reprocessed starlight ($\sim$82\%) compared to less luminous regions. Luminous \HII regions show lower average PAH fractions than less luminous regions, implying that the strong radiation fields from early-type massive stars are efficient at destroying PAH molecules. On average, the monochromatic luminosities at 8, 24, and 70~\micron\ combined carry 94\% of the dust-reprocessed $L_{\rm TIR}$. $L_{70}$ captures ${\sim}52\%$ of $L_{\rm TIR}$, and is therefore the preferred choice to infer the bolometric luminosity of dusty star-forming regions. We calibrate SFRs based on $L_{24}$ and $L_{70}$ against the Lyman continuum photon rates of the massive stars in each region. Standard extragalactic calibrations of monochromatic SFRs based on population synthesis models are generally consistent with our values. | Dust plays several prominent roles in the physics of the interstellar medium (ISM) and star formation. Dust absorbs ultraviolet (UV) radiation emitted by young stars. This absorbed UV radiation is re-emitted at infrared (IR) wavelengths, cooling the dust and also cooling the gas via collisions with dust grains \citep{Draine78, Dwek86, Hollenbach+97}. Studies of nearby star-forming galaxies suggest that, on average, nearly half of emitting starlight is reprocessed by dust \citep{Draine03, Tielens+05}, and the thermal infrared (IR) emission from dust grains dominates the 10--100 \micron\xspace spectral energy distributions (SEDs) of galaxies. Dust is an important tracer of star formation activity and provides an indirect measure of the star formation rate (SFR) in external galaxies. The bolometric, thermal IR luminosity ($L_{\rm TIR}$) is one of the most reliable tracers of dust-obscured star formation \citep{Kennicutt98, PerezGonzalez+06, Kennicutt+09, Kennicutt+12}. \HII regions are locations of recent, active star formation, where massive OB stars emit ionizing UV photons that can interact with the surrounding dust or escape into ISM. \HII regions are generally composed of multiple different components, including the central ionizing cluster(s) of OB stars, a surrounding photodissociation region (PDR), and the remnants of the giant molecular cloud from which the star cluster(s) formed. \HII regions are frequently seen in close proximity to one another, sometimes so much so that their components overlap. It is therefore common to regard \HII regions more generally as star-forming regions within a galaxy. In general, star formation in the Milky Way cannot be studied using the same observational techniques as external galaxies \citep{Chomiuk+Povich11}. Sight-lines through the Galactic disk suffer very high extinction, so SFR diagnostics that depend upon optical/UV observational tracers (in particular, H$\alpha$) cannot be applied. Distances to Galactic \HII regions are often highly uncertain, and confusion arises from multiple star forming regions overlapping along a given line of sight. However, mid- and far-IR SFR tracers can be applied to Galactic and extragalactic regions \citep{Calzetti+07,Calzetti+10,Li+10,Li+13,Stephens+14,V+E13,VEH16}, and thermal radio continuum observations can easily resolve individual Galactic regions and be used as a substitute for recombination-line diagnostics to count ionizing photons \citep{Paladini+03}. The Milky Way offers the unique opportunity to study individual massive star forming regions (MSFRs) resolved over sub-parsec distance scales, where the associated young stellar populations can be directly observed. The Massive Young Star-Forming Complex Study in Infrared and X-ray \citep[MYStIX;][]{Feigelson+13, Broos+13} has characterized hundreds of OB stars in $\sim$20 young Galactic star-forming regions within 4 kpc, and $\sim$100 more obscured OB stars in the MYStIX point-source catalog have recently been found by \citet{Povich+17}. The MSFRs Omnibus X-ray Catalog \citep[MOXC;][]{Townsley+14} produced X-ray point-source catalogs and diffuse emission maps from archival {\it Chandra X-ray Observatory} data on seven MYStIX MSFRs and four additional Galactic MSFRs out to 7 kpc (plus 30 Doradus in the Large Magellanic Cloud). To better understand the interplay between massive stars and IR/radio nebular tracers of star formation, we have conducted a study of 29 Galactic MSFRs, with 21 drawn from the MYStIX and MOXC surveys, and nine additional, prominent regions that have similar high-resolution X-ray through mid-IR archival data available. We construct SEDs by performing aperture photometry on data from the {\it Spitzer Space Telescope}, the {\it Midcourse Space Experiment (MSX)}, the {\it Infrared Astronomical Satellite (IRAS)}, the {\it Herschel Space Observatory}, and the \planck satellite. We then fit a multi-component \citet{Draine+Li07} dust, blackbody, and power-law continuum model to the mid-IR through radio SEDs for each region to measure $L_{\rm TIR}$, constrain dust properties, and search for evidence of supernova contamination in the radio continuum. We use MYStIX point-source database of X-ray and IR-detected OB stars, along with supplementary lists of massive stars from the literature for non-MYStIX targets, to predict the ionizing photon rate injected into each region and to calculate the fraction of emitted luminosity that is reprocessed by dust. This paper is organized as follows: Section~\ref{section:observations} describes the data sources used in this paper. Section~\ref{section:SED_model_description} describes our SED modeling procedure, while Section~\ref{section:SED_fit_results} summarizes the trends observed in the resulting fits. In Section~\ref{section:reprocessing} we discuss the relationship between the MSFRs and their ionizing stellar clusters. In Section~\ref{section:lum_and_SFRs} we discuss commonly used SFR indicators that rely on monochromatic luminosities, and investigate the differences in predicted SFRs these indicators yield when applied to our sample of MSFRs. | We have presented a comprehensive study of the globally integrated IR-radio emission of 28 Galactic MSFRs. We fit the 3.6~\micron--10~mm spectral SEDs constructed from aperture photometry on \spitzer, \msx, \iras, and \herschel images plus \planck extended sources with models consisting of one or two \citet{Draine+Li07} dust components, one cold blackbody component, and a power-law continuum. From our SED model fits and adopted distances to each MSFR we derive the total IR luminosity $L_{\rm TIR}$ and ionizing photon rate $N_C^{\prime}$ required to maintain each radio \HII region. Our sampled MSFRs span three orders of magnitude in luminosity, ranging over $10^4~L_{\sun}\la L_{\rm TIR}\la 2\times 10^7~L_{\sun}$ in dust-reprocessed total infrared luminosity and $3\times 10^{47}~\rm{s}^{-1} \la N_C^{\prime}\la 5\times 10^{50}~\rm{s}^{-1}$ in ionizing photon rate required to maintain the observed radio \HII regions. Modeling the IR+radio SED simultaneously offers considerable advantages over studying either the IR or radio emission alone. The free-free continuum is negligible at shorter mid-IR wavelengths. Although the true ionized gas spectrum departs from a pure power law at short wavelengths, this is unlikely to significant impact our results (e.g., the power law continuum contributes at most a few percent of the total flux at 3.6~\micron\; see Figure~\ref{figure:SED_examples}). However, the incorporation of the Br$\alpha$ emission-line flux at 4.5 \micron, constrained by the radio spectrum, has enabled an improved (although still not perfect) fit to the Spitzer [4.5] mid-IR band compared to models based on dust emission alone \citep{Stephens+14}. We searched the literature to compile lists of the known massive stellar population in each MSFR to estimate the stellar bolometric luminosity ($L_{\star}$) and emitted Lyman continuum photon rate ($N_C$). We balance the ``energy budget'' in each MSFR in terms of the ratios $L_{\rm TIR}/L_{\star}$ and $N_C^{\prime}/N_C$. In 10/28 MSFRs the emergent dust-processed luminosity in the SED exceeds the bolometric luminosity input by the cataloged stars, leading us to conclude that the census of the massive stellar population is incomplete. Our main results are summarized as follows: \begin{enumerate} \item A significant fraction ($f_{C,\rm{abs}}$) of Lyman continuum photons emitted by massive stars is absorbed by dust before contributing to the ionization of \HII regions. This absorption increases with bolometric luminosity; $f_{C,\rm{abs}}=34\%$ averaged across the 14 MSFRs for which it could be calculated and increases to $51\%$ averaged over the 4 most luminous MSFRs in our sample, which average $L_{\rm TIR}=10^7$~\Lsun\ each (Table~\ref{table:subgroup_mean}). This empirical result agrees well with the theoretical predictions of \citet{McKee+Williams97}, who calculated that the dust opacity in giant \HII regions increases with ionizing photon luminosity, reaching an average $\langle f_{C,\rm{abs}}\rangle=0.46$ for Galactic radio \HII regions with $N_{C}^{\prime}>1.5\times 10^{50}$~s$^{-1}$. \item We calculate an average PAH fraction from our dust models and find that it is systematically higher in regions that are powered by a single O6-type star or later, with lower PAH fractions observed in regions will fully populated upper IMFs. This radiation fields in these lower-luminosity \HII regions are relatively weak, and inefficient at destroying PAH molecules. \item We calibrate SFRs based on the monochromatic luminosities $L_{24}$ and $L_{70}$ from our SED models against the Lyman continuum photon rates of the cataloged massive stars in each region. We find that standard extragalactic calibrations of monochromatic SFRs based on population synthesis models are generally consistent with our values, although there is large variation among the 28 individual MSFRs in our sample. Our results are consistent with the \citet{Calzetti+07} 24~\micron\ calibration, and an extrapolation of the \citet{Li+13} 70~\micron\ SFR to the smaller size scales of the Galactic regions is broadly consistent our SFRs. \item The preferred monochromatic luminosity for measuring obscured SFRs is $L_{70}$, which captures, on average, $52\%$ of $L_{\rm TIR}$ in our regions, a result that is in excellent agreement with comparable extragalactic studies \citep[e.g.,][]{Calzetti+10}. \end{enumerate} SFR studies using Galactic radio \HII regions have long included corrections for Lyman continuum photons lost to dust absorption \citep{SBM78,Inoue+01,Murray+Rahman10,Lee+12}. Such corrections are typically not incorporated into extragalactic calibrations, as most H$\alpha$ emission observed on galaxy-wide scales originates from regions with negligible dust. Other SFR tracers, such as integrated UV emission, that do not rely on Lyman continuum photon rates avoid this issue entirely. However, dust absorption becomes significant for spatially-resolved studies of obscured star formation. While current, widely-used calibrations of obscured SFRs account for Lyman continuum photons that escape into the diffuse ISM by using a combination of recombination lines and IR broadband emission \citep[e.g][]{Calzetti+07,Kennicutt+09}, these calibrations could be biased toward low SFRs at the smallest spatial scales and/or highest dust obscurations. The IR and radio SFR calibrations presented in this work are preferred for application to Milky Way studies over the analogous extragalactic calibrations, given the orders-of-magnitude differences in timescales, physical sizes, and luminosities separating whole galaxies from individual Galactic star-forming regions. Systematics due to heating of dust by older stellar populations are most pronounced for the total IR or $L_{70}$ SFR tracers. The \citet{Calzetti+07} $L_{24}$ calibration, which was based on individual IR-bright knots in nearby galaxies, is most consistent with our $L_{24}$ calibration, and both appear to give reasonable results when applied to Galactic regions with sufficiently high IR luminosities \citep[see][whose sample of Galactic star-forming regions overlaps with the low-luminosity end of our sample]{VEH16}. Even within the Milky Way, our calibrations would likely break down when applied to star-forming clouds that are either too low-mass or too early in their evolution to have formed massive stars ionizing radio \HII regions \citep{V+E13,Povich+16}. Thermal radio continuum has been relied upon over the past four decades to measure the total ionizing photon rate of the Milky Way and hence the Galactic SFR \citep[][and references therein]{Chomiuk+Povich11,Kennicutt+12}. We have demonstrated, across nearly three orders of magnitude in luminosity, that the average ionizing photon rate required to maintain the ionization of radio \HII regions is only one-third of the Lyman continuum photon rate emitted by the massive stellar content of these regions. It is therefore important to account for both the escape of Lyman continuum photons from compact radio \HII regions and their absorption by dust within the \HII regions to derive accurate SFRs or simply to infer the ionizing stellar populations within radio \HII regions. For example, the work by \citet{Murray+Rahman10} to measure the Galactic SFR using free-free emission measured by the {\it Wilkinson Microwave Anisotropy Probe ({\it WMAP})} used the calculations of \citet{McKee+Williams97} to correct their Lyman continuum photon rates for dust absorption. This absorption correction seems appropriate, and because {\it WMAP} measured free-free emission across Galactic scales the escape of ionizing photons would be negligible. For smaller, less-luminous \HII regions such as those studied by \citet{VEH16}, neither absorption nor escape of Lyman continuum photons can be safely neglected. Our comparisons of Galactic and extragalactic SFR calibrations required that we assume a standard conversion of Lyman continuum photon rates to absolute SFR based on population synthesis models. While it is encouraging to see convergence between the IR nebular SFR tracers in the Galactic and extragalactic cases, \citet{Chomiuk+Povich11} warned that the assumed star formation timescale in this conversion is likely to be too long by a factor of a few compared to the actual duration of star formation in individual, IR-bright regions, hence such calibrations likely underestimate the true absolute SFRs. In future work we will measure SFRs directly from the spatially resolved low- and intermediate-mass stellar populations to provide a more direct, empirical SFR calibration for the IR and radio nebular tracers. | 18 | 8 | 1808.00454 |
1808 | 1808.07873_arXiv.txt | We present lightcurves and derive periods and amplitudes for a subset of 38 near earth objects (NEOs) observed at 4.5 \micron\ with the IRAC camera on the the \Sp\ Space Telescope, many of them having no previously reported rotation periods. This subset was chosen from about 1800 IRAC NEO observations as having obvious periodicity and significant amplitude. For objects where the period observed did not sample the full rotational period, we derived lower limits to these parameters based on sinusoidal fits. Lightcurve durations ranged from 42 to 544 minutes, with derived periods from 16 to 400 minutes. We discuss the effects of lightcurve variations on the thermal modeling used to derive diameters and albedos from \Sp\ photometry. We find that both diameters and albedos derived from the lightcurve maxima and minima agree with our previously published results, even for extreme objects, showing the conservative nature of the thermal model uncertainties. We also evaluate the NEO rotation rates, sizes, and their cohesive strengths. | Near Earth Objects (NEOs) are small Solar System bodies whose orbits bring them close to the Earth's orbit. NEOs are compositional and dynamical tracers from elsewhere in the Solar System. The study of NEOs allows us to probe environmental conditions throughout the Solar System and the history of our planetary system, and provides a template for analyzing the evolution of planetary disks around other stars. NEOs are the parent bodies of meteorites, one of our key sources of detailed knowledge about the development of the Solar System, and so studies of NEOs are essential for understanding the origins and evolution of our Solar System and others. As of 2018 June there are over 18,000 known NEOs. Roughly 2000 new NEOs are being discovered each year, primarly by the Catalina Sky Survey \citep{leonard17} and Pan-STARRS \citep{veres15}, and the rate will significantly increase when LSST begins operations \citep{veres17}. However, little is known about most NEOs after their discovery, beyond their orbits and optical magnitudes. The size of objects that pass close to Earth can be measured with radar, for example using the Arecibo or Goldstone facilities. Over 750 NEOs have been observed\footnote{https://echo.jpl.nasa.gov/asteroids/index.html}, at a rate of $\sim$75 -- 100 objects per year during the past three years. This rate cannot be easily scaled up, however, and is not keeping pace with the rate of new NEO discoveries. Optical or near-IR spectra of NEOs can determine the surface properties and allow their taxonomic classification \citep{bus99,bus02a,bus02b,demeo09}. However, currently less than 2\% of the NEOs in the JPL Small-Body Database\footnote{https://ssd.jpl.nasa.gov/sbdb\_query.cgi} have assigned taxonomic types. Small NEOs are especially difficult to characterize: for example, \citet{perna18} recently conducted a 30-night GTO program at the NTT and obtained spectra of 147 NEOs, focusing on smaller (<300m) objects. With 24 usable nights, they were able to observe $\sim$ 6 objects per night on this moderately-sized telescope. It would take a major effort on large telescopes to increase the fraction of spectrally-classified objects. The IRAC instrument\citep{2004ApJS..154...10F} on the \Sp\ Space Telescope \citep{2004ApJS..154....1W} is a powerful NEO characterization system. NEOs typically have daytime temperatures $\sim$250 K, hence their thermal emission at 4.5 \micron\ is almost always significantly larger than their reflected light at that wavelength. We can therefore use a thermal model using the optical and IR fluxes to derive NEO properties, including diameters and albedos (see \citealt{2010AJ....140..770T,2016AJ....152..172T}). Measuring the size distribution, albedos, and compositions for a large fraction of all known NEOs will allow us to understand the scientific, exploration, and civil-defense-related properties of the NEO population. After an initial pilot study to verify our observing techniques and analysis methods with the \Sp\ data \citep{2008ApJ...683L.199T}, our team has conducted three major surveys of NEOs with \Sp/IRAC in the Warm/Beyond Mission phases: the ExploreNEOs program \citep{2010AJ....140..770T}, the NEO Survey \citep{2016AJ....152..172T}, and the NEO Legacy Survey \citep{Trilling2017}. As of 2018 March, \Sp\ has completed a total of over 1800 NEO observations, with an expected total of over 2100 observations by the time that the NEO Legacy program has completed in early 2019. Our initial NEO survey results are summarized in \citet{2010AJ....140..770T,2016AJ....152..172T} and \citet{2011AJ....141...75H}. Since then we have examined the albedo distribution and related them to taxonomic classifications \citep{2011Icar..212..158T}, performed a physical characterization of NEOs in our sample \citep{2014Icar..228..217T}, and examined the physical properties of subsets of the sample, including low-$\Delta\nu$ NEOs \citep{2011AJ....141..109M} and dormant short-period comets\citep{2015AJ....150..106M}. We examined individual objects more closely, such as in our discovery of cometary activity associated with the NEO Don Quixote \citep{2014ApJ...781...25M}. We have also performed additional observations on specific NEOs of interest, including the small (<10 m) NEOs 2009~BD \citep{2014ApJ...786..148M} and 2011~MD \citep{2014ApJ...789L..22M}, and the Hayabusa-2 mission target 162173~Ryugu \citep{2017A&A...599A.103M}. One part of our \Sp\ observations of 162173~Ryugu consisted of repeated integrations during its full period to obtain an IR lightcurve to help to constrain the object's shape and size. This led us to conclude that we could perhaps extract similar lightcurves for objects in the survey programs, which were designed only to obtain a single flux measurement from the mosaic image averaging over all of the exposures in the observation. We found that our predicted NEO fluxes were fairly conservative in many cases, and that we could detect most of the NEOs in the individual IRAC exposures. The Wide-field Infrared Survey Explorer \citep[\WISE;][]{wright10} has similarly used infrared observations to characterize a large sample of main-belt asteroids and NEOs. This Explorer-class mission obtained images in four broad infrared bands at 3.4, 4.6, 12 and 22 \micron. \WISE\ conducted its 4-band survey of the sky starting in 2010 January, and after the cryogen was depleted later that year, it continued to operate with its 3.4 and 4.6 \micron\ bands until 2011 February. The spacecraft was reactivated in 2013 December as {\it NEOWISE} \citep{mainzer14} and has since been conducting a sky survey in the 3.4 and 4.6 \micron\ bands to focus on NEO discovery and characterization, using a thermal modeling technique similar to what we have employed with \Sp\ as described above. Over its lifetime, {\it NEOWISE} has observed over 860 NEOs\footnote{https://neowise.ipac.caltech.edu/} and published their estimated diameters and albedos \citep[e.g.,][]{masiero17}. The \WISE\ data can also be used to derive lightcurves of asteroids \citep[e.g.,][]{sonnett15}. However, the cadence is quite different; the \WISE\ survey typically provides repeated observations separated by 3 hr over a 1.5 day period, making it useful for sampling periodicities on the order of 1 -- 2 days. The \Sp\ data samples cadences from a few minutes to hours, making it ideal for small and fast-rotating NEOs, and complementary to the data that \WISE\ provides. Also, since \Sp\ has a larger primary mirror, and it can track the observatory to follow the apparent motion of the NEO, we can integrate for longer periods on each NEO and therefore are more sensitive, detecting objects at the level of a few $\mu$Jy. In this paper we present the results from an analysis of a sample of the available \Sp\ lightcurve data. Section \ref{ObsRed} describes the observations and the reduction techniques. Section \ref{periods} describes the analysis techniques used to derive periods and amplitudes of the lightcurves and presents those results. Section \ref{NEATM} discusses the effects of rotation-induced brightness variability on the thermal modeling results. | We have presented a sample of 38 NEO lightcurves obtained from data taken as part of the ExploreNEOs, NEO Survey, and NEO Legacy \Sp\ programs. We derived periods and amplitudes based on Lomb-Scargle or Plavchan fits for 10 objects where we appear to have complete sampling of the periods, and also present lower limits for another 28 objects based on sine fits to lightcurves shorter than or about equal to one period. Six lightcurves were fit with both periodogram and sine fits and found to have consistent periods. Enabled by the sensitivity and stability of \Sp/IRAC, the NEO surveys have observed thousands of objects where lightcurves can be extracted and periods and amplitudes can be determined or constrained by the data. Because of \Sp's current position in its orbit, it can observe NEOs that are not currently accessible by earth-based observatories. With the 4.5 \micron\ data, we can also estimate the diameter and measure albedos of the NEOs using the same observations. By analyzing the full database as we have done for this small sample, we will be able to extract lightcurves for hundreds of NEOs and determine or set limits on their periods and amplitudes. | 18 | 8 | 1808.07873 |
1808 | 1808.03088_arXiv.txt | {We investigate the reliability of standard N-body simulations by modelling of the well-known Hernquist halo with the help of \texttt{GADGET-2} code (which uses the tree algorithm to calculate the gravitational force) and \texttt{ph4} code (which uses the direct summation). Comparing the results, we find that the core formation in the halo center (which is conventionally considered as the first sign of numerical effects, to be specific, of the collisional relaxation) has nothing to do with the collisional relaxation, being defined by the properties of the tree algorithm. This result casts doubts on the universally adopted criteria of the simulation reliability in the halo center. Though we use a halo model, which is theoretically proved to be stationary and stable, a sort of numerical 'violent relaxation' occurs. Its properties suggest that this effect is highly likely responsible for the central cusp formation in cosmological modelling of the large-scale structure, and then the 'core-cusp problem' is no more than a technical problem of N-body simulations.} | N-body simulations of the origin of present-day cosmic objects from initial small perturbations is not only popular, but almost inevitable method of studying the large-scale structure formation in the Universe: the task is highly nonlinear and complex, which makes numerical methods to be the only direct approach to the problem. Simulating the structure formation under various cosmological assumptions and comparing the results with observations one may check the applicability of the cosmological models. Probably, the strongest result that has been derived from N-body cosmological simulations is the core-cusp problem (see \citep{corecuspreview} for a review). Simulations of the $\Lambda$CDM Universe always predict that dark matter (hereafter DM) halos have either a high density core or a very steep density profile in the center \citep{gao2008, dutton2014} (for the sake of simplicity, this dense central region is named 'a cusp', despite of its shape). However, abundant observations of galaxies favor a respectively large and low-density cores in their central regions: the dark matter mass there is much lower than the cusp should contain \citep{mamon2011, walker2011}. This discrepancy is known for many years at the galaxy scale, but some recent observations disfavor cusps in the galaxy clusters as well \citep{clustercores}. The core-cusp problem may indicate that the cold dark matter model is incorrect. However, only irreproachably reliable estimations of the simulation accuracy and convergence may permit to make so strong physical conclusions. N-body simulations may suffer from several numerical effects; we mention only the major ones. Collisional relaxation of the test particles is an unphysical effect, if one models dark matter, which is believed to be collisionless \citep{binney2002}. In the case of a spherical halo, the process may be characterized by the relaxation time \citep[eqn. 1.32]{bt} \begin{equation} \tau_r =\dfrac{N(r)}{8\ln\Lambda}\cdot\tau_d \label{21a1} \end{equation} where $\ln\Lambda =\ln (b_{max}/b_{min})$ is the Coulomb logarithm ($b_{max}$ and $b_{min}$ are the characteristic maximum and minimum values of the impact parameter), $N(r)$ is the number of test particles inside radius $r$, $\tau_d=(4\pi G\bar\rho(r)/3)^{-1/2}$ is the characteristic dynamical time of the system at radius $r$, $\bar\rho(r)$ is the average density inside $r$. It is noteworthily that the potential softening used in the N-body simulations allows to avoid close collisions of the particles, but scarcely affects the collisional relaxation time: $b_{min}$ is of the order of the softening radius, but $\tau_r$ depends on it only logarithmically. We will use $\ln\Lambda =\ln (3 r_s/\epsilon)$, where $\epsilon$ is the softening radius, like in \citep{17}. The softening of the gravitational potential is an another source of computation errors. If $\epsilon$ is too small, close two-body collisions on large angles occur. If the value of $\epsilon$ is too large, it smooths away the high density peaks and introduces a bias in the force computation. The optimal choice of softening has been a subject of many studies, see, e.g. \citep{merritt96,2000MNRAS.314..475A,knebe2000,dehnen01,zhan06}. The third potential source of undesirable numerical effects is the gravitational interaction computations. The direct summation is exact, but very resource-intensive algorithm, since it requires to calculate $N-1$ partial forces per particle in a system containing $N$ particles. It means that we need $\sim N^2$ partial force calculations for all the system. Many N-body codes use a hierarchical tree algorithm \citep{tree1,tree2}, which is significantly more economical. The idea of the method is to group distant particles into larger cells and calculate their gravity as a single multipole force. A recursive subdivision of space is used to achieve the required spatial adaptivity. For instance, the algorithm of the \texttt{GADGET} code \citep{springel2005} divides the space in cubic cells. Each cell is repeatedly subdivided into eight daughter ones until the ratio between the distance to the cell and the size of the cell exceeds the parameter specified by the user. The \texttt{GADGET} code that we consider in this paper splits the cells until the following cell-opening criterion is satisfied \begin{equation} \label{21b1} \frac{G M_{cell}}{r^2}\left(\frac{l}{r}\right)^2 \le \fac |\vec a|, \end{equation} where $M_{cell}$ and $l$ are the cell mass and extension, $r$ is the distance to the cell, $|\vec a|$ is the value of the total acceleration obtained in the last timestep, and $\fac$ is a tolerance parameter (see \citep[equation (18)]{springel2005}, where $\fac$ is denoted by $\alpha$). The interaction with the closest particles is calculated by the direct summation. Requiring only $O(N\ln N)$ partial force calculations for all the system, the tree algorithm is, generally speaking, much faster than the direct summation. The price to pay is that the result represents only an approximation to the true force. Finally, figure 4 in \citep{springel2005} gives a highly visual illustration of significant numerical phenomena appearing during the temporal integration of particle orbits. Thus, several unphysical effects inevitably occur in the simulations, and their correct estimation is necessary for an appropriate interpretation of simulation results. However, the present state-of-art of N-body simulation tests (especially, of the DM behavior) can hardly be named adequate. The commonly-used criterion of the convergence of N-body simulations in the halo center is solely the density profile stability \citep{power2003}: these simulations show that the central cusp (close to $\rho\propto r^{-1}$) is formed quite rapidly ($t<\tau_r$). The shape of the cusp and its stability is insensitive to the simulation parameters (see, e.g. \citep{hayashi2003}). Considering the temporal dependence of the overdensity inside various radii $r$ from the halo center (i.e., the ratio of the average density inside radius $r$ from the halo center to the average universe density), \citep{power2003} find that a core forms in the center no earlier, than $t(r)= 1.7 \tau_r(r)$. We will denote this time by $\tau_p(r)= 1.7 \tau_r(r)$. The authors assume that cusp universality and stability implies the negligibility of numerical effects, and the core formation is the first sign of the collisional relaxation. Therefore, \citep{power2003} suggest that N-body simulation results are trustable until $t(r)= 1.7 \tau_r(r)$. The reasons why the collision influence can be neglected at almost two relaxation times are not quite clear. Moreover, \citep{hayashi2003} and \citep{klypin2013} report that the cusp is stable much longer, probably, up to tens of relaxation times. The weak point of this reasoning is that the profile stability by itself does not guarantee the absence of the collisional influence. Indeed, if the collisions are already significant, but the particles mainly scatter on small angles in the collisions (the latter is true for N-body simulations), the system can be described by the Fokker-Planck equation. This equation may have a stationary solution \citep{evans1997, 13, 18}, and if it is stable, it works as a sort of attractor: the collisions tend to transform an initial distribution into the stationary one. Then the attractor profile should survive for much longer than $\tau_r$, since the Fokker-Planck diffusion is self-compensated in this case. The profile corresponding to the stationary solution should also be quite universal, since its shape is defined by the the Fokker-Planck coefficients (i.e., by the potential of the particle interaction) that are similar for various N-body packages. However, the profile universality and stability have nothing to do with the simulation veracity: it is already created by test particle collisions, that is, by a purely numerical effect. It seems to be no coincidence that the Fokker-Planck equation really has a stationary solution that is similar to $\rho\propto r^{-1}$ in the halo center \citep{evans1997, 13} and close to the Einasto profile with $n\sim 6$ at $r\sim 10^{-1} r_s$ \citep{18}. Another way to test N-body simulations is to model an analytical solution and compare the results with theoretical predictions. Simulations of an isolated spherically symmetric halo allow to go beyond the density profile stability and consider the full array of dynamical parameters of the particles. \citep{17} modelled a Hernquist halo and found that all integrals of motion characterizing individual particles experience strong unphysical variations along the whole halo, revealing an effective interaction between the test bodies. Moreover, the simulations show that the cusp stability is really provided with the particle collisions, as we described in the previous paragraph: intense upward and downward Fokker-Planck streams of particles in the cusp region occur, and the cusp is stable because they compensate each other. This result suggests that the cusps in cosmological N-body simulations may also be a consequence of numerical effects. However, the paper \citep{17} has not answered to several important questions. Though significant unphysical effects were found, the immediate causes of them, as well as their dependence from the N-body code parameters, remained unclarified. The fact that the variations of the integrals of motion of individual particles were significant at very large radii, where the influence of the collisions and potential softening was certainly negligible, implied that the integral variations there are most likely due to the potential calculating algorithm. In order to clarify this point, we perform a new simulation of the Hernquist halo, following the way described in \citep{17}. However, contrary to that work, we evolve the same initial conditions for the halo using the \texttt{GADGET-2} code with various parameter settings. Besides, we follow the system evolution in a 4-th order Hermite code, which uses the direct summation algorithm to calculate the gravitation force, and thus is free from the tree algorithm drawbacks. Comparing the results, we may elucidate the cause of the unphysical effects and clarify their dependence from the code parameters. In Sect.~\ref{21sec2} we describe the codes we used and the simulation setup, in Sect.~\ref{21sec3} we present the methods we use to treat the data, in Sect.~\ref{21sec4}, we present the results of the simulations and discuss them. Finally, in Sect.~\ref{21sec5} we briefly summarize the obtained results. | otherwise any discrepancy with observations may turn out to be just a result of an unaccounted numerical effect. | 18 | 8 | 1808.03088 |
1808 | 1808.08165_arXiv.txt | We present a detailed light curve analysis of RR Lyrae variables at multiple wavelengths using Fourier decomposition method. The time-series data for RR Lyrae variables in the Galactic bulge and the Magellanic Clouds are taken from the Optical Gravitational Lensing Experiment survey while the infrared light curves are compiled from the literature. We also analyse the multiband theoretical light curves that are generated from the stellar pulsation models of RR Lyrae stars for a wide range of metal-abundances. We find that the theoretical light curve parameters with different metal abundances are consistent with observed parameters in most period bins at both optical and infrared wavelengths. The theoretical and observed Fourier amplitude parameters decrease with increase in wavelength while the Fourier phase parameters increase with wavelength at a given period. We use absolute magnitudes for a subset of theoretical models that fit the observed optical RR Lyrae light curves in the Large Magellanic Cloud to estimate a distance modulus, $\mu_\textrm{LMC}=18.51\pm0.07$, independent of the metallicity. We also use Fourier analysis to study the period-color and amplitude-color relations for RR Lyrae stars in the Magellanic Clouds using optical data and find that the slope of period-color relation at minimum light is very shallow or flat and becomes increasingly significant at the maximum light for RRab stars. We also find that the metallicity dependence of the period-color relations increases as we go from minimum to maximum light, suggesting that the mean light results are indeed an average of the various pulsational phases. We summarize that the average variation in these relations is consistent between theory and observations and supports the theory of the interaction of the stellar photosphere and the hydrogen ionization front. | RR Lyrae stars are low-mass, helium-burning, horizontal branch stars, exhibiting periodic light curves with a pulsation period of $\sim0.2-1.0$~day and amplitude variation of $\le$~2 mag \citep{preston1964, kolenberg2012}. These radially pulsating variables are valuable tracers of old and metal-poor stellar populations and provide extragalactic distance estimates with well-defined period-luminosity relations ($PLR$s), especially in the near-infrared bands \citep{longmore1986, bono2001, catelan2004, sollima2006, muraveva2015, braga2015, neeley2015}. Additionally, a narrow range of the intrinsic colors of RR Lyrae stars also serves as a useful reddening indicator \citep{sturch1966, mateo1995, guldenschuh2005,wagner-kaiser2017}. RR Lyrae and Cepheid variables are both excellent probes for the understanding of the theory of stellar pulsation and evolution. A comparison of their observed light curves and pulsation properties with theory provides very useful constraints for the stellar pulsation models \citep{simon1985,marconi2013,marconi2017,bhardwaj2017}. Light curve structures of Cepheid and RR Lyrae variables were first studied using the Fourier decomposition method by \citet{simon1981} and \citet{simon1982}, respectively. The variation of lower-order Fourier parameters with period for 70 field RR Lyrae stars was discussed by \citet{simon1982}. They found that fundamental-mode (RRab) and first-overtone (RRc) mode RR Lyrae stars can be easily differentiated on the Fourier amplitude plane. \citet{petersen1984} carried out a Fourier analysis of RR Lyrae stars in the $\omega$ Centauri and found evidence of Cepheid-like progressions \citep{simon1981} in RRab stars in the period range 0.5 day to 1.5 days. \citet{simon1985} compared the light curves of RR Lyrae stars with those from hydrodynamical models using Fourier decomposition and while they found consistent results between theoretical and observed data for the RRc stars, there was a discrepancy in the Fourier phase parameters for RRab stars. \citet{bono1996} computed pulsation models of RR Lyrae stars with different chemical compositions and found that the bolometric amplitudes increase for RRab stars and decrease for RRc stars with an increase in the metal abundance. Using a smooth grid of models covering a range of stellar masses, luminosities and metallicities, \citet{bono2001} provided theoretical constraints on the $PLR$ for RR Lyrae stars in the $K$-band. Most of these theoretical studies focussed on the comparison with observed pulsation properties \citep{marconi2011,bono2011,bono2016,marconi2016}, while others carried out a model fitting of the light curves e.g. \citet{simon1985,wood1997,bono2002,natale2008,marconi2010,marconi2013,marconi2013b,marconi2017} for Cepheids and \citet{kovacs1998,bono2000,castellani2002,difabrizio2002,marconi2005,marconi2007} for RR Lyrae stars. However, this work makes use of the modern time-series data and the most recent stellar pulsation models of RR Lyrae stars from Marconi et al. (2015) to provide an extensive comparison of theoretical and observed light curves of RR Lyrae variables. Fourier analysis of RR Lyrae stars has also been employed to obtain photometric metallicities using the empirical light curve structure and metallicity relations \citep{kovacs1995}. \citet{jurcsik1996} investigated the best-fit relations between $[Fe/H]$ and the Fourier parameters and found a linear relation among $[Fe/H]$, period and the Fourier phase parameter ($\phi_{31}$). \citet{morgan2007} found a $[Fe/H]-\phi_{31}-P$ relation for RRc stars in the globular clusters analogous to that found by \citet{jurcsik1996} for RRab stars. The correlation of Fourier phase parameters with metallicity for RR Lyrae stars has been revisited by several authors, for example, \citet{smolec2005, nemec2013, skowron2016, ngeow2017} and references therein. These empirical relations allow studies of morphology, structure and the metallicity of the Galaxy and the Magellanic Clouds \citep{haschke2012,deb2014,deb2015,hajdu2015,skowron2016}. Additionally, light-curve data of RR Lyrae stars can also be used to study period-color (PC) and amplitude-color (AC) relations as a function of pulsation phase to understand the interaction of the stellar photosphere and the hydrogen ionization front \citep{simon1993, kanbur2005, bhardwaj2014, ngeow2017}. Recent near-infrared observations of RR Lyrae stars provide evidence of a tight $PLR$ at longer wavelengths \citep{sollima2006, principe2006, coppola2012, muraveva2015, braga2015, neeley2017} that can be used to measure the value of the Hubble constant up to a few percent precision \citep{beaton2016}. However, metallicity contribution to $PLR$s for RR Lyrae stars is not well-constrained as the observed $P-L-[Fe/H]$ relations differ from those obtained using theoretical pulsation models \citep[for example,][]{muraveva2015,marconi2015}. A detailed light curve analysis of RR Lyrae stars can provide insights into the metallicity effects on the light curve structure and subsequently on the mean light PLRs. Further, a comparison of the observed light curves of RR Lyrae stars with theoretical models is crucial for understanding important constraints for the stellar pulsation codes. This analysis has been recently carried out for Cepheid variables by \citet{bhardwaj2017} and we extend this work for RR Lyrae stars in the present analysis. \begin{table} \caption{A summary of 410 RR Lyrae models (274 RRab and 136 RRc) used in the present analysis. The last two columns denote the number of RRab/RRc models with a unique combination of (Z,Y,$\dfrac{M}{M_{\odot}}$,$\log\dfrac{L}{L_{\odot}}$). ${T_e}$ indicates the range of effective temperature for each of these combinations.} \label{tab:Data} \centering \begin{tabular}{c c c c c c c} \hline \hline Z & Y & $\dfrac{M}{M_{\odot}}$ & $\log\dfrac{L}{L_{\odot}}$ & ${T_e}$ (K) & RRab & RRc \\ [0.5ex] \hline \hline & & 0.51 & 1.69 & 5700$-$6800 & 7 & 3 \\[-1ex] & & 0.51 & 1.78 & 5600$-$6800 & 7 & 3 \\[-1ex] \raisebox{1ex}{0.02} & \raisebox{1ex}{0.27} & 0.54 & 1.49 & 6000$-$7100 & 5 & 3\\[-1ex] & & 0.54 & 1.94 & 5200$-$6600 & 8 & $-$ \\[1ex] & & 0.55 & 1.62 & 5900$-$7100 & 12 & 4 \\[-1ex] & & 0.55 & 1.72 & 5800$-$7000 & 12 & 10 \\[-1ex] & & 0.56 & 1.60 & 5900$-$7100 & 11 & 4 \\[-1ex] \raisebox{1ex}{0.008} & \raisebox{1ex}{0.256} & 0.56 & 1.70 & 5800$-$7000 & 10 & 9 \\[-1ex] & & 0.57 & 1.58 & 6000$-$7100 & 5 & 3\\[-1ex] & & 0.57 & 2.02 & 5400$-$6680 & 6 & $-$\\[1ex] & & 0.53 & 1.81 & 5700$-$6800 & 8 & 4 \\[-1ex] & & 0.55 & 1.71 & 6000$-$7000 & 10 & 8 \\[-1ex] & & 0.55 & 1.81 & 5700$-$6900 & 13 & $-$ \\[-1ex] & & 0.56 & 1.65 & 6000$-$7100 & 10 & 4 \\[-1ex] \raisebox{0ex}{0.004} & \raisebox{0ex}{0.25} & 0.56 & 1.75 & 5800$-$7000 & 10 & 8 \\[-1ex] & & 0.57 & 1.63 & 6000$-$7100 & 10 & 4\\[-1ex] & & 0.57 & 1.73 & 5900$-$7000 & 10 & 8\\[-1ex] & & 0.59 & 1.61 & 6000$-$7200 & 10 & 3\\[-1ex] & & 0.59 & 2.02 & 5700$-$6700 & 7 & $-$\\[1ex] & & 0.58 & 1.87 & 5900$-$6900 & 7 & 4 \\[-1ex] \raisebox{0ex}{0.001} & \raisebox{0ex}{0.245} & 0.64 & 1.67 & 6000$-$7200 & 9 & 4 \\[-1ex] & & 0.64 & 1.99 & 5700$-$6800 & 10 & 5\\[1ex] & & 0.6 & 1.89 & 5700$-$6900 & 9 & 7 \\[-1ex] \raisebox{0ex}{0.0006} & \raisebox{0ex}{0.245} & 0.67 & 1.69 & 6000$-$7200 & 9 & 6 \\[-1ex] & & 0.67 & 2.01 & 5800$-$6800 & 9 & 6\\[1ex] & & 0.65 & 1.92 & 5800$-$6900 & 6 & 7 \\[-1ex] \raisebox{0ex}{0.0003} & \raisebox{0ex}{0.245} & 0.716 & 1.72 & 6000$-$7200 & 8 & 7 \\[-1ex] & & 0.716 & 1.99 & 5700$-$6900 & 11 & 3\\[1ex] & & 0.72 & 1.96 & 5800$-$6900 & 7 & 2 \\[-1ex] \raisebox{0ex}{0.0001} & \raisebox{0ex}{0.245} & 0.8 & 1.76 & 6000$-$7200 & 8 & 7\\[-1ex] & & 0.8 & 1.97 & 5800$-$6700 & 10 & $-$\\[1ex] \hline \end{tabular} \end{table} The structure of this paper is as follows: Section \ref{sec:data} describes the theoretical models and the observational data used in this analysis. The Fourier decomposition technique is briefly discussed in Section \ref{sec:fourier} and we study the variation of amplitude and Fourier parameters with wavelength, period and metallicity and present the comparison of theoretical and observed light curve parameters. We discuss the extinction corrected PC, AC and PC-metallicity relations for RR Lyrae stars in Section \ref{sec:PCAC}. Finally, we summarise the results of this study in Section \ref{sec:results}. \begin{table*} \caption{The observed light curve data used in the present analysis with the number of stars available in each dataset. N$_\textrm{RRab}$, N$_\textrm{RRc}$ and N$_\textrm{T2C}$ refer to the number of RRab, RRc and type II Cepheid variables.} \centering \begin{tabular}{c c c c c c c} \hline\hline & Band & N$_\textrm{RRab}$ & N$_\textrm{RRc}$ & Reference & N$_\textrm{T2C}$ & Reference\\ \hline \hline & V & 11736 & 5603 & & 517 & \\[-1ex] \raisebox{1ex}{Bulge} & & & & \raisebox{1ex}{\citet{soszynski2014}} & & \raisebox{1ex}{\citet{soszynski2017}}\\[-1ex] & I & 25929 & 10261 & & 873 & \\[1ex] & V & 25762 & 8758 & & 202 & \\[-1ex] \raisebox{1ex}{LMC} & & & & \raisebox{1ex}{\citet{soszynski2016}} & & \raisebox{1ex}{\citet{soszynski2008}}\\[-1ex] & I & 26209 & 8893 & & 203 & \\[1ex] & V & 4769 & 739 & & 42 & \\[-1ex] \raisebox{1ex}{SMC} & & & & \raisebox{1ex}{\citet{soszynski2016}} & & \raisebox{1ex}{\citet{soszynski2010}}\\[-1ex] & I & 4825 & 746 & & 43 & \\[1ex] & 3.6 $\mu$m & 14 & $-$ & & $-$ & \\[-1ex] \raisebox{1ex}{M4} & & & & \raisebox{1ex}{\citet{neeley2015}} & & \raisebox{1ex}{$-$}\\[-1ex] & 4.5 $\mu$m & 14 & $-$ & & $-$ & \\[1ex] \hline \end{tabular} \label{tab:Observed_Data} \end{table*} \begin{table*} \caption{The light curve parameters of RR Lyrae stars from the theoretical models. The columns provide the chemical composition (Z and Y), filter ($\lambda$), stellar mass, mode of pulsation, effective temperature ($T_e$), logarithmic luminosity, logarithmic period, amplitude ($A$), mean magnitude ($m_0$), Fourier amplitude ($R_{21}$,$R_{31}$) and phase ($\phi_{21}$,$\phi_{31}$) parameters and the mean radius.} \centering \scalebox{0.95}{ \begin{tabular}{c c c c c c c c c c c c c c c} \hline\hline Z & Y & $\lambda$ & $\frac{M}{M_{\odot}}$ & Mode & $T_e$ & $\log\frac{L}{L_{\odot}}$ & $\log(P)$ & A & $m_0$ & $R_{21}$ & $R_{31}$ & $\phi_{21}$ & $\phi_{31}$ & $\log\frac{R}{R_{\odot}}$\\ [0.5ex] \hline \hline 0.0200& 0.270& U& 0.51& FU& 6800& 1.69& -0.193& 1.258& 1.056& 0.407& 0.152& 2.838& 5.627& 0.702\\ 0.0200& 0.270& U& 0.51& FU& 6700& 1.69& -0.171& 1.661& 1.123& 0.494& 0.176& 3.041& 6.006& 0.715\\ 0.0200& 0.270& U& 0.51& FU& 6500& 1.69& -0.127& 1.596& 1.192& 0.474& 0.089& 3.463& 0.594& 0.741\\ 0.0200& 0.270& U& 0.51& FU& 6200& 1.69& -0.058& 1.044& 1.293& 0.458& 0.131& 4.462& 1.993& 0.782\\ 0.0200& 0.270& U& 0.51& FU& 6500& 1.69& -0.127& 1.596& 1.192& 0.474& 0.089& 3.463& 0.594& 0.741\\ 0.0200& 0.270& U& 0.51& FU& 5900& 1.69& 0.012& 0.433& 1.431& 0.871& 0.205& 4.722& 1.338& 0.826\\ 0.0200& 0.270& U& 0.51& FU& 5700& 1.69& 0.064& 0.519& 1.592& 0.448& 0.224& 5.107& 1.510& 0.856\\ 0.0200& 0.270& U& 0.54& FU& 6800& 1.49& -0.379& 1.703& 1.586& 0.476& 0.283& 2.807& 5.510& 0.603\\ 0.0200& 0.270& U& 0.54& FU& 6700& 1.49& -0.356& 1.481& 1.604& 0.449& 0.243& 2.848& 5.662& 0.615\\ 0.0200& 0.270& U& 0.54& FU& 6500& 1.49& -0.313& 1.169& 1.653& 0.388& 0.196& 3.028& 6.215& 0.641\\ \hline \end{tabular}} \begin{tablenotes} \small \item \textbf{Notes:} This table is available entirely in a machine-readable form in the online journal as supporting information. \end{tablenotes} \label{tab:Models} \end{table*} | \label{sec:results} We have carried out a detailed light curve analysis for the largest available dataset of RR Lyrae stars in the Bulge, LMC and SMC from the OGLE-IV survey using the Fourier decomposition technique and compared the results with the most recent stellar pulsation models of RR Lyrae stars from \citet{marconi2015}. The models show a decrease in amplitude with an increase in wavelength, except for a few period ranges where $U_{amp}<B_{amp}$, depending on the effective temperature of the RR Lyrae star. This is consistent with observations, albeit the mean amplitudes from the models are slightly higher than those from observations $-$ an increase in the mixing length parameter can cause a decrease in the pulsation amplitudes \citep{criscienzo2004b}. Also, the uncertainties on the assumed convective efficiency affect the pulsation amplitudes of the theoretical light curves \citep{fiorentino2007}. An investigation of the variation of Fourier parameters with mass predicts a decrease in Fourier amplitude parameters and an increase in the Fourier phase parameters with an increase in mass for a given period range, especially in the $K$-band. The availability of the NIR RR Lyrae data in the near future would be useful for providing constraints on the M-L combinations of the RR Lyrae stars. The variation of Fourier parameters with wavelength presents a decrease in amplitude parameters and an increase in phase parameters with an increase in wavelength, for a given period. The scatter in the amplitude parameters decreases as we go from optical to infrared, given that the metallicity effects are less at longer wavelengths. The observed Fourier parameters of RR Lyrae stars are in reasonable agreement with those obtained from the models $-$ with a better consistency in the infrared bands. For the long-period range $0 < \log(P) < 0.2$, models show marginal inconsistency in phase parameters at optical wavelengths. We found a subset of 25 RRab stars from the LMC with I-band light curves that match well with models. These subset of models were used to obtain an average distance modulus to LMC of $18.51\pm0.07$ mag, which is in good agreement with published results. We study the period-color and amplitude-color relations at minimum and maximum light to understand the interaction of the stellar photosphere with the hydrogen ionisation front. While the PC$_{\mathrm{min}}$ slope is nearly-flat for Bulge RRab stars and consistent with previous results, it has a small but significant positive/negative slope for LMC/SMC RRab stars. However, the change in slope from minimum to maximum light is significant and thereby, the theory of the interaction of stellar photosphere and hydrogen ionisation front is consistent. Unlike their fundamental mode counterparts, the RRc stars show a smaller difference in PC slope from min/max $-$ this is consistent with the theory of the HIF$-$stellar photosphere interaction because these stars are hotter and so the HIF and stellar photosphere engagement occurs at a much higher temperature in a range where Saha ionization equilibrium is more sensitive to temperature. Using the photometric median metallicities, we find Z=0.001 for Bulge, Z=0.0006 for LMC and Z=0.0003 for SMC and use these models for comparison with the observations. The models predict a flat PC$_{\mathrm{min}}$ at minimum light for Z=0.02, Z=0.001, Z=0.0006 and Z=0.0003. It is interesting to note that Z=0.0003 predicts a slightly negative slope for PC$_{\mathrm{min}}$, similar to that observed in SMC. We, therefore, suggest that PC$_{\mathrm{min}}$ may be used as a constraint for models. The metallicity dependence of the PC relations increases as we go from minimum to maximum light. At maximum light, the PC relation slope increases from Bulge to LMC to SMC. The results of PC and AC relations in both theory and observations are found to be consistent with the previous works and the theory of the interaction of stellar photosphere and hydrogen ionization front. The multi-wavelength light curve analysis of fundamental-mode RR Lyrae stars has been carried out extensively for the first-time in the present analysis using both theoretical models and observed light curves. Although our results suggest an overall consistency of the models with the observations, there are cases of discrepancies such as the higher amplitudes at optical bands and the sensitivity of convection towards the redder edge of the instability strip that need further investigation. A smoother grid of models along with variation in the mixing length, viscosity etc. should result in a better agreement between the models and observations. The results of this work can provide stringent constraints for the theoretical stellar pulsation codes that incorporate static atmosphere models to generate RR Lyrae light curves at multiple wavelengths. | 18 | 8 | 1808.08165 |
1808 | 1808.04382_arXiv.txt | We present observations of supernova (SN)~2017ens, discovered by the ATLAS survey and identified as a hot blue object through the GREAT program. The redshift $z=0.1086$ implies a peak brightness of $M_{g}=-21.1$ mag, placing the object within the regime of superluminous supernovae. We observe a dramatic spectral evolution, from initially being blue and featureless, to later developing features similar to those of the broadlined Type Ic SN~1998bw, and finally showing $\sim2000$\,km\,s$^{-1}$ wide H$\alpha$ and H$\beta$ emission. Relatively narrow Balmer emission (reminiscent of a SN~IIn) is present at all times. We also detect coronal lines, indicative of a dense circumstellar medium. We constrain the progenitor wind velocity to $\sim50$--60\,km\,s$^{-1}$ based on P-Cygni profiles, which is far slower than those present in Wolf-Rayet stars. This may suggest that the progenitor passed through a luminous blue variable phase, or that the wind is instead from a binary companion red supergiant star. At late times we see the $\sim2000$\,km\,s$^{-1}$ wide H$\alpha$ emission persisting at high luminosity ($\sim3\times10^{40}$\,erg\,s$^{-1}$) for at least 100 day, perhaps indicative of additional mass loss at high velocities that could have been ejected by a pulsational pair instability. | \label{sec:intro} Type Ic supernovae (SNe) arise from the core collapse of a massive star that has lost its hydrogen and helium layers prior to exploding, through either strong stellar winds or interaction with a binary companion \citep[e.g.,][]{1997ARA&A..35..309F, 2017hsn..book..195G}. Their light curves are powered by the radioactive decay of $^{56}$Ni that is produced in the SN explosion. Related to these events, but with luminosities up to 100 times higher, are the Type I superluminous SNe (SLSNe~I; see \citealt{2012Sci...337..927G, 2018ApJ...854..175I, 2018SSRv..214...59M} for reviews of observations and models). SLSNe exhibit spectral similarities to SNe~Ic \citep{2010ApJ...724L..16P}, but their luminosities are such that they cannot be powered solely by radioactive decay \citep{2011Natur.474..487Q}. The nature of the additional energy source remains unknown, with suggestions ranging from a central engine \citep{2010ApJ...717..245K,2010ApJ...719L.204W} to interaction with a massive H and He-free circumstellar medium \citep[CSM;][]{2011ApJ...729L...6C}. Some SNe~Ib/Ic have been observed to develop relatively narrow ($\sim 500-1000$\,km\,s$^{-1}$) emission lines of hydrogen in their spectra; examples include SNe~Ib 2014C and 2004dk \citep{2015ApJ...815..120M, 2018arXiv180307051M} and SNe~Ic 2001em and 2017dio \citep{2017hsn..book..195G, 2018ApJ...854L..14K}. This has been interpreted as evidence that for at least some H-poor SNe, the fast ejecta are colliding with H-rich material relatively far from the star. This late-time interaction has also been observed in some SLSNe~Ic which show H$\alpha$ emission at $+70$ to $+250$\,d after their peak brightness \citep{2015ApJ...814..108Y,2017ApJ...848....6Y}. In this Letter we report on the discovery of an unusual SN with our GREAT ({\bf GR}OND-{\bf e}PESSTO-{\bf AT}LAS; \citealt{2008PASP..120..405G, 2015A&A...579A..40S, 2018PASP..130f4505T}) Survey. We introduce this program here, which is designed to rapidly identify hot, blue transients, with the specific goal of finding very young SLSNe in faint galaxies \citep{2017ATel10510....1C}. SN~2017ens (ATLAS17gqa) was discovered by the ATLAS survey on 2017 June 5 (UT dates are used herein), located at (J2000) $\alpha=12^h04^m09^s.37$, $\delta=-01^\circ55'52.2''$. Prompted by the high blackbody temperature of $21,000\pm3000$\,K that we measured with our GREAT data on 2017 June 8 \citep{2017ATel10478....1C}, we began an intensive spectroscopic and photometric follow-up campaign (Sec.\,\ref{sec:obs}). The adopted redshift of SN~2017ens, $z=0.1086$ (Sec.\,\ref{sec:analysis_result_coronal}), implies an absolute magnitude of $M_{g}=-21.1$ at peak, and thus a luminosity consistent with a SLSN \citep{2012Sci...337..927G}. In Sec.\,\ref{sec:analysis_result} we present the spectral evolution of SN~2017ens, which began to show $\sim2000$\,km\,s$^{-1}$ wide H$\alpha$ and H$\beta$ emission after $+163$\,d (phases are corrected for time dilation and are relative to the GROND $r$-band maximum on MJD = 57,924.011). We compare the spectral properties of SN~2017ens to those of other SLSNe and broadlined SNe~Ic (SNe~Ic-BL), and also present the detections of rarely seen coronal lines. The bolometric light curve and modeling results are described in Sec.\,\ref{sec:LC_model}. Finally, in Sec.\,\ref{sec:dis} we discuss plausible scenarios that may explain the spectral evolution and luminosity of SN~2017ens. We adopt a cosmology of $H_{0}=72~\mathrm{km~s^{-1}~Mpc^{-1}}$, $\Omega_\Lambda=0.73$, and $\Omega_{m}=0.27$. The foreground reddening toward SN~2017ens is $A_{V}=0.058$\,mag \citep{2011ApJ...737..103S}, and we assume that host-galaxy extinction is negligible because no Na~I~D absorption is visible in the SN spectrum. | \label{sec:dis} One important clue to interpreting the possible powering mechanisms behind SN~2017ens is that we measured the H-rich material to have a velocity of $\sim50$--60\,\kms\ from the blueshifted absorption of the narrow P-Cygni profiles. This wind velocity is far slower than those present in Wolf-Rayet star winds. If this wind is from the progenitor, it could come from a massive H-rich progenitor (such as a luminous blue variable) that explosively ejected its H envelope shortly before the SN explosion. Alternatively, this wind could come from a pulsational pair-instability SN with a slow and long-term stable wind \citep{2017ApJ...836..244W}. It is also possible that SN~2017ens exploded as a SN~Ic-BL inside a patchy, H-rich CSM from a binary companion; the expanding ejecta interact with the bulk of the CSM at later times, as has been suggested for SN~2017dio \citep{2018ApJ...854L..14K}. Alternatively, as proposed for ASASSN-15no \citep{2018MNRAS.476..261B}, a dense inner CSM may have hidden the SN features at early times, before they become briefly visible as the CSM was swept up by the ejecta. At late times they could have again been masked by an increasingly strong interaction component. A special CSM geometry (e.g., doughnut shape) is also probable, and we see the SN~Ic-BL along a certain viewing angle. In the case of a binary companion, the wind of $\sim50$--60\,\kms\ and mass-loss rate of $5\times10^{-4}\,M_\odot\,\mathrm{yr^{-1}}$ are consistent with a red supergiant \citep{2017MNRAS.465..403G}, albeit at the more extreme end, which can be explained by the companion having gained mass from the SN progenitor during an earlier accretion phase. If so, this may suggest that the progenitor of SN~2017ens lost its H and He layers through interaction with a binary companion. We must also consider the apparent $\sim2000$\,\kms\ material given its high luminosity. If this is associated with mass loss from the progenitor, and the line width is not from electron scattering as seen in many SNe~IIn, then the material is moving much faster than the winds of H-rich stars (or the CSM of SNe~IIn). It is difficult to imagine how this could be produced by anything other than a sudden ejection of the H envelope, shortly before the SN explosion. In fact, the luminosity of the $\sim2000$\,km\,s$^{-1}$ wide component of H$\alpha$ is comparable to that seen in SN~1995N \citep{2002ApJ...572..350F} ($\sim2.3\times10^{40}$\,erg\,s$^{-1}$), and it may be too large to be coming solely from swept-up material. A pulsational pair-instability explosion is at least qualitatively consistent with an outburst that can unbind the H envelope shortly before an SN explosion. This scenario is also consistent with the measured low-metallicity environment. The unique spectroscopic evolution of SN~2017ens together with its high luminosity poses challenges to all currently known SN scenarios. While detailed modeling can help elucidate the nature of this transient, ongoing surveys for SLSNe such as GREAT will find more such peculiar transients. With a larger sample and high-cadence follow-up spectroscopy, we will be able to further understand the nature of SN 2017ens-like objects and the role of interaction in SLSNe. | 18 | 8 | 1808.04382 |
1808 | 1808.08215_arXiv.txt | In 1964 (Solar Cycle 20; SC 20), Patrick McIntosh began creating hand-drawn synoptic maps of solar magnetic features, based on H$\alpha$ images. These synoptic maps were unique in that they traced magnetic polarity inversion lines, and connected widely separated filaments, fibril patterns, and plage corridors to reveal the large-scale organization of the solar magnetic field. Coronal hole boundaries were later added to the maps, which were produced, more or less continuously, into 2009 (i.e., the start of SC 24). The result was a record of $\sim45$ years ($\sim570$ Carrington rotations), or nearly four complete solar cycles of synoptic maps. We are currently scanning, digitizing and archiving these maps, with the final, searchable versions publicly available at NOAA's National Centers for Environmental Information. In this paper we present preliminary scientific studies using the archived maps from SC 23. We show the global evolution of closed magnetic structures (e.g., sunspots, plage, and filaments) in relation to open magnetic structures (e.g., coronal holes), and examine how both relate to the shifting patterns of large-scale positive and negative polarity regions. | The solar magnetic field is constantly changing, driven by the dynamo below and driving in turn a field that permeates the heliosphere. Concentrated magnetic flux is generated in the Sun's interior and emerges through its surface, e.g., as sunspots. Ongoing diffusion and transport by solar-surface flows results in a shifting pattern of positive and negative magnetic polarity that is an evolving boundary on the global magnetic field. Because the hot corona results in an expanding solar wind, both ``closed'' and ``open'' magnetic fields extend upwards from this boundary. Closed-field regions include sunspots, plage, and, if magnetic shear/twist is concentrated at the polarity inversion line (PIL), filaments. Of particular note on a global scale are polar crown filaments, which may extend nearly $360^\circ$ around the sun at high latitudes. Open-field regions manifest as unipolar coronal holes, which, depending on solar-cycle phase, may appear predominantly at the poles or as isolated structures at lower latitudes. \begin{figure}[t] \begin{center} \includegraphics[width=5in]{level0_1928.pdf} \caption{Example of original, hand-drawn McA synoptic solar map. Magnetic polarity is indicated by +/-; PILs are dashed lines with filaments indicated by extensions; coronal hole boundaries are indicated by hashed lines; plage by light dots, and sunspots by darker dots.} \label{figorigmap} \end{center} \end{figure} \begin{figure}[b!] \begin{center} \includegraphics[width=5in]{level3_1928.pdf} \caption{Example of colorized, digitized McA synoptic solar map (for the same CROT as Fig. \ref{figorigmap}). Magnetic features are identified with a distinct color, as described in the legend.} \label{figcolormap} \end{center} \end{figure} | The unique power of the McIntosh archive is its capability for simultaneously representing closed and open magnetic structures over a range of time scales. The completion of the full McA digitization will provide the community with a comprehensive resource for addressing key questions including: How do active longitudes vary within and between solar cycles, for both closed and open magnetic features? Where are closed and open magnetic features rooted (as evidenced by rotation rate), and how does this depend on solar cycle phase, feature lifetime, and latitude? How does the evolution of open and closed magnetic features relate to surface flows on solar-cycle time scales (e.g., torsional oscillations, \cite[Howe et al., 2013]{howeetal_13})? Answering any or all of these questions has important implications for our understanding of the solar dynamo, and for our interpretation of periodic variations of Earth's space environment and upper atmosphere. | 18 | 8 | 1808.08215 |
1808 | 1808.00618_arXiv.txt | The study of extended, cold dust envelopes surrounding R Coronae Borealis (RCB) stars began with their discovery by IRAS. RCB stars are carbon-rich supergiants characterized by their extreme hydrogen deficiency and for their irregular and spectacular declines in brightness (up to 9 mags). We have analyzed new and archival {\it Spitzer} Space Telescope and {\it Herschel} Space Observatory data of the envelopes of seven RCB stars to examine the morphology and investigate the origin of these dusty shells. {\it Herschel}, in particular, has revealed the first ever bow shock associated with an RCB star with its observations of SU~Tauri. These data have allowed the assembly of the most comprehensive spectral energy distributions (SEDs) of these stars with multi--wavelength data from the ultraviolet to the submillimeter. Radiative transfer modeling of the SEDs implies that the RCB stars in this sample are surrounded by an inner warm (up to 1,200 K) and an outer cold (up to 200 K) envelope. The outer shells are suggested to contain up to 10$^{-3}$ M$_\odot$ of dust and have existed for up to $10^5$ yr depending on the expansion rate of the dust. This age limit indicates that these structures have most likely been formed during the RCB phase. | \label{sec:intro} R Coronae Borealis (RCB) stars provide an excellent opportunity to understand more about the advanced stages of stellar evolution \citep{1996PASP..108..225C,2012JAVSO..40..539C}. They form a rare class of hydrogen--poor, carbon--rich supergiants. Two formation scenarios have been proposed for their origin: the single degenerate final helium-shell flash (FF) model and the double degenerate (DD) white dwarf (WD) merger model \citep{1996ApJ...456..750I, 2002MNRAS.333..121S}. The latter involves the merger of a CO and a He WD \citep{1984ApJ...277..355W}, while the former takes the hot evolved central star of a planetary nebula (PN) and turns it into a cool supergiant \citep{Fujimoto:1977lr, 1979sss..meet..155R}. The trademark behavior of RCB stars is their spectacular and irregular declines in brightness. These declines can take an RCB star up to 9 magnitudes fainter than its peak brightness, and are caused by the formation of discrete, thick clouds of carbon dust along the line of sight \citep{1935AN....254..151L, 1939ApJ....90..294O,1996PASP..108..225C}. All RCB stars show an infrared excess due to the presence of warm circumstellar material (CSM, \citealp[and references therein]{Feast:1997lr,2012JAVSO..40..539C}). Further, some RCB stars have been found to have cold, extended nebulosity \citep[e.g.,][]{Schaefer:1986lq,Walker:1985rr,1986ASSL..128..407W,2011MNRAS.414.1195B, 2011ApJ...743...44C}. The origins of this CSM material as well as the progenitors of the central RCB stars still remain shrouded in mystery. One important difference between the RCB stars formed in the two scenarios is that in the FF model, they would be surrounded by a fossil, neutral hydrogen-rich (HI-rich) PN shell \citep{Walker:1985rr, Gillett:1986cr, 1990MNRAS.247...91L, 1999ApJ...517L.143C, 2011ApJ...743...44C}. Three stars (Sakurai's Object, V605 Aquilae, and FG Sagittae) have been observed to undergo FF outbursts that transformed them from hot evolved stars into cool giants with spectroscopic properties similar to RCB stars \citep{1997AJ....114.2679C,1998ApJS..114..133G,1998A&A...332..651A,Asplund:1999bh,Asplund:2000qy,2006ApJ...646L..69C}. These FF stars are all surrounded by PNe which are still ionized. However, the cooler RCB central stars are no longer able to provide the needed ionizing radiation so the atoms in the shell have recombined. The velocity of the fossil PN shell would be similar to its ejection velocity, $\sim$ 20 -- 30 km s$^{-1}$. In the DD scenario, the stars may have had PN phases but they would have occurred so long ago, $\sim$10$^9$ years, that no structure resembling a fossil envelope would remain when the two WDs finally merge to form an RCB star. These shells could be material lost during the WD merger event, itself. This would have happened much more recently, $\lesssim$10$^4$ years ago, and would imply these structures are much less massive than previously estimated \citep{Gillett:1986cr, 2011ApJ...743...44C}. A third explanation for the observed shells is that they could have formed during the RCB phase. RCB stars are thought to produce dust at a rate of 10$^{-7}$ to 10$^{-6}$ M$_\odot$ yr$^{-1}$ \citep{2012JAVSO..40..539C}. \citet{2013AJ....146...23C} have found that newly forming clouds are propelled away from the central star at speeds up to 400 km s$^{-1}$. This also could result in the observed envelopes on a timescale of about 10$^4$ years. We are now in an era where high spatial resolution and high sensitivity far-IR (FIR), submillimeter (sub-mm), and even radio observations exist of RCB stars in order to study their cold CSM material. We present unpublished {\it Spitzer} and {\it Herschel} observations of the RCB$/$HdC stars: MV Sagittarii (MV~Sgr), R Coronae Borealis (R~CrB), RY Sagittarii (RY~Sgr), SU Tauri (SU~Tau), UW Centauri (UW~Cen), V854 Centuari (V854~Cen), V Coronae Australis (V~CrA), and HD~173409. We have constructed multi--wavelength datasets ranging from the ultraviolet (UV) to sub-mm in order to better determine the mass, size, and morphology of the diffuse material surrounding these RCB stars. | \citet{2015AJ....150...14M} presented and explored three possible interpretations for the origins of the diffuse, dusty nebulosity that surrounds some RCB stars: (1) they are fossil planetary nebulae (PNe); (2) they are remnant material from the merger of a CO and a He white dwarf binary, (3) they have been constructed from dust ejection events during the current RCB phase. We will now examine the results presented here in the context of these three scenarios. The results of the MOCASSIN models are presented in Table 11. The MCRT modeling of these SEDs suggests the existence of two discrete, concentric spherical shells around each of our sample RCB stars. The construction of these shells during the current RCB phase is critically tied to the number of dust puffs produced, the expansion velocity of the dust puffs, and the lifetime of RCB stars. It has been suggested that during a decline a single puff contains $\sim$10$^{-8}$ M$_{\sun}$ of dust \citep{1992ApJ...397..652C,2011ApJ...743...44C}. Then, $\sim$10$^{-7}$ M$_{\sun}$ of dust would form per year if a dust puff forms somewhere around the star every 50 days. Thus, an RCB inner envelope would be produced in about 10 years and it is unlikely that any of the inner shells are the remnant material of a WD merger or fossil PN. The origin of the outer shells is of greater uncertainty. The data seem to suggest that at some point dust formation ceased and then restarted, or that the inner and outer shells have different origins. For example, the inner shell could be from the RCB phase and the outer shell could be a fossil PN shell or remnant material from WD merger. A knowledge of the hydrogen abundance in these shells would help determine whether they are fossil PN shells or not. If the envelopes are fossil PNe then they should be H--rich. H I measurements at 21--cm of R~CrB put lower limits on any hydrogen in its dust shell \citep{2015AJ....150...14M}. Assuming R~CrB is a typical RCB star, then it is unlikely its dust shell is a fossil PN. Further, recent modeling of the merger rates of WD binaries by \citet{2015ApJ...809..184K} found that typically the merger will not take place for at least 500 Myr after both stars become WD. This is reinforced by the discovery that the nearby system WD 1242-105 is a binary white dwarf expected to merge in 740 Myr \citep{2015AJ....149..176D}. After these lengths of time, it is very unlikely that any PN material would still be around an RCB star. Hydrodynamic modeling of the material that remains following a WD merger suggests that these envelopes would contain M$_{\rm Dust} \leq 10^{-6}$ M$_\odot$ \citep{2015AJ....150...14M}. The mean mass of the outer envelopes in this sample is $10^{-3}$ M$_\odot$. The least massive envelope (V854~Cen) is implied to contain $2.60 \times 10^{-5}$ M$_\odot$ of dust. This is still an order of magnitude higher than predicted for remnant material from a WD merger. The velocity of the expanding dust has been estimated as being tens to hundreds of km s$^{-1}$. Slower expansion velocities have been suggested by \citet[and references therein]{2011ApJ...739...37G}. Faster outward movement is suggested by the He I $\lambda$10830 line, which suggests that the dust is rapidly accelerated up to 400 km s$^{-1}$ \citep{1992ApJ...397..652C,2003ApJ...595..412C,2013AJ....146...23C}. The outer envelopes in our sample have implied outer radii that range from 10$^{18}$ cm to 10$^{19}$ cm (see Table 11). Material at these distances represents the oldest material to be shed by RCB stars. Dust moving with slower velocities, 20 km s$^{-1}$, would take about $10^4$ to $10^5$ years to reach anywhere from 10$^{18}$ cm to 10$^{19}$ cm, respectively. These times drop by an order of magnitude if the dust velocities agree more with the results of the He I $\lambda$10830 analysis. These timescales are both much longer than we have known about the RCB phenomenon \citep{1797RSPT...87..133P}. We have compared the observations of the RCB stars to the hydrogen--deficient carbon (HdC) stars and stars that have been observed to undergo a final flash (FF). HdC stars are essentially spectroscopic twins of RCB stars. HdC stars, however, do not experience decline events and lack any IR excess. The HdC star HD~173409 was observed with both PACS and SPIRE on {\it Herschel}. No emission associated with HD~173409 was detected in any of the {\it Herschel} observations. The SED for this star also shows no evidence for any IR excess when fit by a single 7000 K blackbody. Recently, one HdC star, HD~175893, was found to have an IR excess from analysis of WISE colors and could either represent a missing link between the two classes of objects or an RCB star going through an extended period of low dust formation \citep{2012A&A...539A..51T}. The results of our sample were compared to the FF star, V605~Aql, and the findings of \citet{2013ApJ...771..130C}. \citet{2013ApJ...771..130C} presented the SED for V605~Aql, which indicates the presence of $\sim$10$^{-3}$ M$_\odot$ of dust associated with its 1919 ejecta. This is on a similar level to the dust masses derived from our MOCASSIN modeling for the outer shells. In this scenario, these envelopes would have been created in the recent past. However, the rapid evolution in the effective temperature of V605~Aql from 5,000 K to 95,00 K in around 80 years \citep{2006ApJ...646L..69C} has not been found in any RCB star. The {\it Herschel} observations of SU~Tau with the PACS and SPIRE instruments have led to the discovery of a bow shock like structure. This is the first known RCB star to exhibit this type of feature, which represents interactions between the SU~Tau CSM and the local interstellar medium (ISM). The bow shock extends between 30$''$ to 50$''$ from the central position of SU~Tau with a brighter feature in the southeast possibly indicating a location where more material is beginning to pile up. RCB stars are among the most uncommon and bizarre objects discovered in the Universe. However, they provide the opportunity to greatly advance our knowledge in areas such as stellar evolution and stellar chemistry. Additional examination of these objects, especially at 21--cm, is needed to determine the origin of the cold, diffuse CSM seen around the RCB stars. | 18 | 8 | 1808.00618 |
1808 | 1808.09739_arXiv.txt | {Signal artefacts due to Radio Frequency Interference (RFI) are a common nuisance in radio astronomy. Conventionally, the RFI-affected data are tagged by an expert data analyst in order to warrant data quality. In view of the increasing data rates obtained with interferometric radio telescope arrays, automatic data filtering procedures are mandatory. Here, we present results from the implementation of a RFI-detecting recurrent neural network (RNN) employing long-short term memory (LSTM) cells. For the training of the algorithm, a discrete model was used that distinguishes RFI and non-RFI data, respectively, based on the amplitude information from radio interferometric observations with the GMRT at $610\, \mathrm{MHz}$. The performance of the RNN is evaluated by analyzing a confusion matrix. The true positive and true negative rates of the network are $\approx 99.9\,\%$ and $\approx 97.9\,\%$, respectively. However, the overall efficiency of the network is $\approx 30\%$ due to the fact that a large amount non-RFI data are classified as being contaminated by RFI. Matthews correlation coefficient is ~0.42 suggesting that a still more refined training model is required.} | In Sec.\ref{perf}, we show that the RNN reaches an accuracy of $\approx 98\%$ after sufficient training. However, this seemingly high accuracy is due to the large number of data in the TN category, see Eq.\ref{confmatr}. When studying the PPV and the FDR, a weakness of the chosen method becomes apparent which lowers its overall efficiency. A large amount of data (FP category, $415179$ data points), which are actually non-RFI, are classified as being RFI resulting in a PPV of $\approx 18\,\%$. This also becomes evident when taking the MCC into account, which is $\approx 0.42$, meaning the classification is not random with respect to the data, but the correlation is not strong either. When calculating the F1 score, the overall precision of the network is $\approx 0.31$. An improvement of the efficiency of the method can be expected from the following refinements: \begin{itemize} \item \textbf{Data usage:} In this study, we used only the amplitude information in the data. However, the amplitude differences with respect to the channles, baselines, and times steps should also be used, adding four more axes to the training data cube. In addition, the phase (spatial) information in the data could be further utilized. \item \textbf{Model complexity:} The discrete amplitude-based model to distinguish RFI and non-RFI may be adjusted to cope with more complex signal shapes and strength patterns. \item \textbf{Network architecture:} The network training could be extended to consider the time step, polarization and baseline sequences instead of the channel sequence only. Thus, the amount of data in the FP category will be reduced. By also adding the information on the image-level, it is possible to combine the RNN with the advantages of a CNN which would give information of prominent features in an image, giving a hierarchy of dominant features like RFI. This would result in a change of the architecture into a recurrent convolutional neural network (RCNN) \end{itemize} The results of this study mark an encouraging milestone and path towards a highly dynamical RFI filter meeting the challenges of future radio antenna arrays. | 18 | 8 | 1808.09739 |
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1808 | 1808.07809_arXiv.txt | Fast Radio Bursts (FRBs) are a class of short-duration transients at radio wavelengths with inferred astrophysical origin. The prototypical FRB is a broadband signal that occurs over the extent of the receiver frequency range, is narrow in time, and is highly dispersed, following a $\nu^{-2}$ relation. However, some FRBs appear band-limited, and show apparent scintillation, complex frequency-dependent structure, or multi-component pulse shapes. While there is sufficient evidence that FRBs are indeed astrophysical, their one-off nature necessitates extra scrutiny when reporting a detection as \emph{bona~fide} and not a false positive. Currently, there is no formal validation framework for FRBs, rather a set of community practices. In this article, we discuss potential sources of false positives, and suggest a framework in which FRB-like events can be evaluated as real or otherwise. We present examples of false-positive events in data from the Arecibo, LOFAR, and Nanshan telescopes, which while FRB-like, are found to be due to instrumental variations, noise, and radio-frequency interference. Differentiating these false-positive detections from astrophysical events requires knowledge and tests beyond thresholded single-pulse detection. We discuss post-detection analyses, verification tests, and data sets which should be provided when reporting an FRB detection. | \label{sec:intro} \begin{figure*} \centering \begin{subfigure}[t]{0.45\textwidth} \centering\captionsetup{width=.95\linewidth} \includegraphics[width=1.0\textwidth]{figures/D20161204_buf23_Beam0.pdf} \caption{} \label{fig:beam0_dynamic_spec} \end{subfigure} \begin{subfigure}[t]{0.45\textwidth} \centering\captionsetup{width=.95\linewidth} \includegraphics[width=1.0\textwidth]{figures/D20161204_buf4_Beam5.pdf} \caption{} \label{fig:beam5_dynamic_spec} \end{subfigure} \caption{ Dynamic spectrum (top) and \gls{snr}-maximized de-dispersed time series (bottom) of an FRB-like event that was detected simultaneously in Beam 0 and 5 of the ALFA receiver on December 4, 2016. The dynamic spectrum has been bandpass normalized. (a) Detected FRB-like event in Beam 0 of ALFA. The characteristic dip before and after the event is due to zero-DM removal which is part of the ALFABURST RFI exciser. The strong, narrow band source at 1480~MHz around 100~ms is due to a local RFI source. (b) Same event detected in Beam 5. } \label{fig:dynamic_spec} \end{figure*} The astrophysical origin of \glspl{frb} has been a mystery since they were first reported \citep{2007Sci...318..777L}. Though, the detection of multiple \glspl{frb} in \cite{2013Sci...341...53T} put forth a convincing case for their astrophysical nature and, subsequent detections have re-enforced this case. The \gls{dm} associated with the reported events indicate they occur well beyond our Galaxy, possibly at cosmological distances. Given their extragalactic distance, observed fluxes suggest extremely energetic progenitors; however, their emission mechanism remains unknown. The consensus that \gls{frb} are astrophysical events developed from detections with multiple telescopes, over multiple wavelengths, using different receiver systems. \glspl{frb} are difficult to detect as they require high-gain telescopes which typically have a small field of view, such that despite many thousands of observing hours, only a few dozen have been reported as of this writing \citep{2016PASA...33...45P}\footnote{http://frbcat.org/}. The prototypical \gls{frb} is broad-band across the observable bandwidth of the receivers used. The pulse is narrow-in-time---on the order of a few milliseconds in width---and highly dispersed, exhibiting a $\nu^{-2}$ relation with frequency. Despite follow-up observations in the direction of FRB events, only FRB\,~121102 has been shown to repeat \citep{2016Natur.531..202S}. Several reported detections deviate from this prototypical form, exhibiting complex frequency-dependent structure and are possibly band-limited. The pulse width also varies, due to either the emission process or propagation effects such as scattering from the \gls{ism} and \gls{igm}. These rare events are detected by automated, high-performance software pipelines that extensively search a broad range of trial \glspl{dm} and pulse widths \citep{Barsdell2012, 2015MNRAS.452.1254K, Bannister2017, Chime2018}. Each de-dispersed time series is then thresholded -- any peaks above a minimum \gls{snr} are reported as candidates. The number of candidates is usually overwhelming due to \gls{rfi} and system gain variations. Initially, candidates were reviewed manually; however, with the amount of data acquired in recent surveys, it has become a significant time effort to do so. As such, and as our understanding of the expected signal properties has grown, so too has the use of machine-learning-based classifiers of candidate events \citep[e.g.][]{Wagstaff2016, 2018MNRAS.474.3847F, 2018MNRAS.478.1209F, Connor2018}. As FRBs are rare and appear not to repeat (with the notable exception of FRB\,121102), being able to confidently verify them is an important issue. There is significant \gls{rfi} detectable at all radio observatories, and there are known anthropogenic sources that appear FRB-like \citep{2015MNRAS.451.3933P}. Given the significant number of \gls{dm} trials and high-time resolution of the spectra in a typical survey, there are a large number of false-positives (type-I errors) that pass the automated post-processing detection thresholds. This is by design, as we would like to severely limit the potential for false-negatives (type-II errors) in our detection pipelines by accepting a number of false-positives during automated searches and manually discarding them later. But, given the large sampling of the parameter search space it can be difficult to differentiate between a true, astrophysical \gls{frb} (true-positive) and a `terrestrial \gls{frb}' (false-positive) due to \gls{rfi}, systematics, or other local effects. As the survey time increases, the likelihood of detecting such a false-positive will increase. In this paper, we present a list of verification criteria to counter false detection of `terrestrial \glspl{frb}', that can be applied \emph{post facto} to recorded data. We firstly discuss examples of FRB-like sources, which after investigation, were shown to be non-astrophysical (\S~\ref{sec:false-pos}). Using these and previously reported \glspl{frb}, we develop a set of criteria to test detections against (\S~\ref{sec:verify_crit}). We then apply these criteria to some of the previously reported detections (\S~\ref{sec:appln_to_previous_detections}) as examples of usage, then discuss future observational methods to reduce ambiguity (\S~\ref{sec:future_methods}). We conclude with suggestions as to what data, in addition to the detected dynamic spectrum, should be reported to provide a robust statement about \gls{frb} detections in general. | From an observational point of view, \glspl{frb} are unique astronomical sources as, so far, no follow-up observations have been able to verify a source using a different telescope or observing frequency, except in the case of FRB\,121102 which is known to repeat. Thus, it is necessary to provide evidence for an astrophysical origin when reporting a detection. In Section \ref{sec:verify_crit} we have presented a set of criteria which can be used to verify an \gls{frb} detection relative to a prototypical \gls{frb} model and observational data. We have shown which of the criteria are more essential to verification than others. For example, the strongest evidence for an event being of astrophysical origin is a multi-site detection. In combination these criteria become stronger evidence for the origin of a detected \gls{frb}. Effort should be made to test these criteria using the available data when reporting a detection. The set of criteria and methods for applying them laid out here should ideally be applied to all new FRB detections. Publishing a table similar to Figure \ref{fig:heat_map} provides comprehensive information for assessing new detections. As \gls{frb} surveys continue to increase in sensitivity, sampled parameter space, and time, there will be an increase in the number of false-positive detections, even as rejection models are improved. Differentiating between true, astrophysical \glspl{frb} and false-positive events will become more difficult. Further, it could be that most astrophysical \glspl{frb} are not prototypical as it is defined now, such as in the case of FRB\,121102. Which indicate that some of the criteria must be relaxed, making the differentiation more difficult still. This should warrant the drive towards interferometric and multi-site observations, complex-voltage data recorders, sub-band pulse searches, and standard reporting of the observing system. Automation of some of the verification tests can help to formalize criteria into metrics. Machine learning-based models are routinely employed to detect and classify candidate detections. \cite{Zhang2018} used simulated pulses to train a deep neural network detection model. \cite{Connor2018} built a classifier model based on the dynamic spectrum, pulse profile, and DM-trial space of a detection. This model was similarly trained with simulated pulses. UTMOST \gls{frb} detections \citep{2018MNRAS.478.1209F} are automated through the use of a classifier model which takes into account multi-beam detections to reduce the number of false-positive detections due to local \gls{rfi}. The ALFABURST system \citep{2018MNRAS.474.3847F} which detected the event in Section \ref{sec:D20161204} uses a similar classifier model. Inherent to these models is a definition of a prototype pulse. They provide the initial assessment of the detection, which can then be followed by testing the more complex criteria which still require expert analysis. Detection reporting incurs an economic cost in observing time, resources, and human effort. Though, there is a scientific cost to delayed reporting as follow-up observations could detect further pulses or multi-wavelength observations could reveal an unknown counterpart to the radio pulse. An initial detection will not be able to be verified against many of the criteria presented here. But basic tests such as \gls{snr}, dispersion relation, and telescope state can be automated to determine an initial detection `importance' in a VO-Event trigger \citep{2017arXiv171008155P}. This importance factor can then be adjusted accordingly as more criteria are tested. Reporting of false-positive events, even if the source is not explained, helps to improve the robustness of search pipelines against systematics and \gls{rfi}. Reporting these events also helps to improve the case for \glspl{frb} being of astrophysical origin, just as the explanation for Perytons \citep{2015MNRAS.451.3933P} removed doubt about detections using Parkes. An attempt should be made to make the raw data, either unprocessed filterbank or complex-voltage data, available to be used in independent verification and testing on search pipelines. Jupyter notebooks of the verification tests and terrestrial \glspl{frb} are hosted on our public git repository\footnote{https://github.com/griffinfoster/terrestrial-frb-letter}. | 18 | 8 | 1808.07809 |
1808 | 1808.05384_arXiv.txt | Gamma-ray bursts (GRBs) are exceptionally bright electromagnetic events occurring daily on the sky. The prompt emission is dominated by X-/$\gamma$-rays. Since their discovery over 50~years ago, GRBs are primarily studied through spectral and temporal measurements. The properties of the emission jets and underlying processes are not well understood. A promising way forward is the development of missions capable of characterising the linear polarisation of the high-energy emission. For this reason, the SPHiNX mission has been developed for a small-satellite platform. The polarisation properties of incident high-energy radiation (50--600~keV) are determined by reconstructing Compton scattering interactions in a segmented array of plastic and Gd$_3$Al$_2$Ga$_3$O$_{12}$(Ce) (GAGG(Ce)) scintillators. During a two-year mission, $\sim$200 GRBs will be observed, with $\sim$50 yielding measurements where the polarisation fraction is determined with a relative error $\leq$10\%. This is a significant improvement compared to contemporary missions. This performance, combined with the ability to reconstruct GRB localisation and spectral properties, will allow discrimination between leading classes of emission models. | Gamma-ray bursts (GRBs) are the brightest electromagnetic events in the universe, occurring randomly on the sky approximately daily~\cite{paciesas1999,paciesas2012}. The emission is characterised by two epochs - the prompt phase, lasting for seconds to minutes and dominated by X-/$\gamma$-rays, and the afterglow, lasting for days and emitting at lower energies. The afterglow is relatively well understood~\cite{Zhang&Kumar}. Many aspects of the prompt phase remain unknown, but the fireball model~\cite{Meszaros06} is a generally accepted scenario, where the GRB is formed during the collapse of a massive object into a black hole. Recent observations of gravitational waves~\cite{GW} in coincidence with a short GRB (typical prompt duration $<$2~s) have shown that the progenitor is a merger of two neutron stars. Neutron star-black hole mergers can also result in short GRBs. For long GRBs (typical prompt duration up to 100~s), the progenitor is instead connected to broad line type 1c supernovae~\cite{Cano}. Understanding the GRB prompt emission mechanism holds the key to using GRBs as probes of the early universe and of extreme physics such as relativistic jets, relativistic magneto-hydrodynamics, aberration of light, relativistic shock waves, and Lorentz invariance violation~\cite{context1,context2}. During the process of collapse and merger, two highly relativistic jets of plasma are emitted along the rotational axis of the black hole. A GRB is observed if one of the jets is directed to earth. Energy dissipation within the outflow (e.g. internal shocks or magnetic reconnection) gives rise to the prompt emission (keV--MeV). The prompt energy spectrum is featureless and is often well described by a smoothly broken power-law with a break at a peak energy, $E_p$, typically occurring in the range from a few tens to several hundred keV. The Band function is often used, smoothly connecting a low-energy power-law, with spectral index $\alpha$ (the photon flux $N_E \propto E^\alpha$, where $E$ is the energy) to a high-energy power-law with spectral index $\beta$~\cite{Band}. There is a large spread in the measured values of $\alpha$, $\beta$ and $E_p$ \cite{paciesas2012}, with a typical GRB having $\alpha=-1$, $\beta=-2.5$ and $E_p=$ 200~keV. As the jet interacts with the surrounding medium it is decelerated producing a forward external shock, the emission from which forms the afterglow. GRB detectors typically measure the energy distribution (spectra) and the arrival time (light-curve) of the GRB photons. Even though large samples of bursts have been observed, the properties of the jets and the underlying emission process remain poorly understood. The study of GRB jets is inherently difficult since images cannot be produced due to the large observing distance. Measurements of the linear polarisation properties of the detected photons address this problem. Linear polarisation is described using two parameters: $(i)$ the polarisation fraction (PF, \%) describing the magnitude of beam polarisation; and, $(ii)$ the polarisation angle (PA, degrees) which defines the orientation of the electric field vector of the incident photon beam relative to, e.g., celestial north. Despite the scientific value of GRB polarimetry, there is a lack of reliable observational data~\cite{mcconnell}. Between 2010 and 2012, the GAP polarimeter on-board the IKAROS spacecraft measured the polarisation of three bright GRBs in the 70-300~keV energy band~\cite{GAP1,GAP2}. The measurements indicated high PF values, and for the brightest burst, polarisation parameters were determined in two time bins revealing a 90$^\circ$ change in PA. The measurements had weak statistical significance and additional observations with more sensitive missions are required. POLAR is a gamma-ray polarimeter mission (50--500~keV) launched with the Chinese Tiangong-2 space station in 2016~\cite{polar}. After $\sim$6~months of operations, data-taking ceased due to instrument malfunction. Polarisation data is expected for some of the $\sim$50 observed GRBs~\cite{Merlin}. The AstroSat mission was launched in 2015. The CZTI instrument is a general purpose coded aperture spectrometer for X-ray observations which can be used for polarimetry. During the first year of operations, 47 GRBs were detected and polarisation parameters determined for 11 of the brightest bursts~\cite{astrosatGRB}. High PF values were observed for the majority of GRBs, albeit with relatively large uncertainties, as discussed in Section~\ref{sec:comp}. Observations have also been reported from instruments not designed for polarimetry, e.g. the INTEGRAL~\cite{integralGRB1,integralGRB2} and RHESSI~\cite{rhessiGRB} missions. Since no polarimetric calibration was performed prior to observations, it is difficult to ascertain the reliability of the reported results~\cite{dispute1,dispute2,dispute3}. This paper describes the Satellite Polarimeter for High eNergy X-rays (SPHiNX) -- a satellite-borne instrument for hard X-ray polarimetric studies of GRBs. The instrument design is optimised for polarimetry and with a large field-of-view, $\sim$120$^\circ$ opening angle, and collecting area, $\sim$800~cm$^2$, a large sample of $\sim$200 GRBs will be provided during the two year mission. The light-curve and spectral shape will be determined for all GRBs. Polarisation parameters will be reconstructed in the energy range 50-600~keV for $\sim$50 GRBs. The arrival time of X-rays will be determined with an absolute timing accuracy of 1~ms to allow correlation with other missions. These instrument characteristics will allow discrimination between different classes of GRB emission models. This paper is organised as follows. In Section~\ref{sec:motivation}, the scientific motivation for the mission is presented. An overview of the mission parameters and constraints is presented in Section~\ref{sec:mission}. The instrument design is described in Section~\ref{sec:design} and the operational and calibration strategy is outlined in Section~\ref{sec:ops+calib}. The simulated instrument characteristics and scientific performance are discussed in Section~\ref{sec:performance}. An outlook is presented in Section~\ref{sec:outlook}. | 18 | 8 | 1808.05384 |
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1808 | 1808.03293.txt | The physics of planet formation is investigated using a population synthesis approach. We develop a simple model for planetary growth including pebble and gas accretion, and orbital migration in an evolving protoplanetary disk. The model is run for a population of 2000 stars with a range of disk masses and radii, and initial protoplanet orbits. The resulting planetary distribution is compared with the observed population of extrasolar planets, and the model parameters are improved iteratively using a particle swarm optimization scheme. The characteristics of the final planetary systems are mainly controlled by the pebble isolation mass, which is the mass of a planet that perturbs nearby gas enough to halt the inward flux of drifting pebbles and stop growth. The pebble isolation mass increases with orbital distance so that giant planet cores can only form in the outer disk. Giants migrate inwards, populating a wide range of final orbital distances. The best model fits have large initial protoplanet masses, short pebble drift timescales, low disk viscosities, and short atmospheric cooling times, all of which promote rapid growth. The model successfully reproduces the observed frequency and distribution of giant planets and brown dwarfs. The fit for super Earths is poorer for single-planet systems, but improves steadily when more protoplanets are included. Although the study was designed to match the extrasolar planet distribution, analogs of the Solar System form in 1--2\% of systems that contain at least 4 protoplanets. | \label{sec:intro} Planetary systems exhibit a remarkable degree of diversity in terms of the number of observed planets per system, the masses of these planets, their densities, and their orbits \citep{winn:2015}. Presumably, all these systems formed as a result of a common set of physical processes acting on protoplanetary disks that were qualitatively similar, at least initially. Thus, the observed diversity demands an explanation. Differences between planetary systems could have arisen for several reasons including differences in the initial disks, different environmental conditions, or chance events during planet formation. The relative importance of these factors is unclear, however. Numerical models for the formation of planetary systems can help address this question, and may tell us why a particular system looks the way it does. Unfortunately, our current understanding of planet formation is far from complete, and this limits the usefulness of these models. For example, the main physical process driving the evolution of gaseous protoplanetary disks remains uncertain. Conventional ``$\alpha$-disk'' models are difficult to reconcile with observations \citep{rafikov:2017}. Recent work has focussed on several evolution mechanisms including disk winds, hydrodynamic turbulence, and planetary torques \citep{suzuki:2016, bai:2016, hartmann:2018, fung:2017}, although the relative importance of these remains unclear. As a result, the dynamics of small, solid particles is also somewhat unclear. This in turn affects our picture of the growth, radial mixing, and histories of dust grains and pebble-size particles \citep{ciesla:2007, birnstiel:2010, bitsch:2018}. Another major area of uncertainty is the mechanism by which pebbles are converted into asteroid-size planetesimals, as well as the time scale and efficiency of this conversion \citep{cuzzi:2010, johansen:2011, birnstiel:2016}. We cannot say with certainty whether the main building blocks of solid planets and the cores of gas-giant planets are pebbles, planetesimals, or both, nor precisely how the building blocks change with time and location in a disk. Finally, it is apparent that planetary orbits can evolve during and after planet formation \citep{ward:1997, trilling:1998}. However, the range of circumstances in which orbital migration is important is still being investigated, and the extent of migration in a typical planetary system is unclear. Population synthesis models provide one approach to reducing these uncertainties. These studies adopt relatively simple models for disk evolution, planetary growth, and planetary migration. The main uncertainties in the physics of planet formation are encapsulated in model parameters. Typically, models are run many times, and the results compared with the Solar System or the observed extrasolar planet distribution to constrain the parameter values. \citet{ida:2004} conducted one of the earliest population synthesis studies using an analytic model for the growth of planets accreting planetesimals and gas in a weakly viscous disk. The simulations included ``type-II'' migration in which the orbits of planets massive enough to open a gap migrate with the viscous evolution of the disk \citep{ward:1997}. The authors followed the evolution of a large number of single-planet systems, with planets starting at random locations in disks with a range of disk masses. The simulations produced populations of terrestrial and gas-giant planets somewhat similar to the observed distribution with a dearth of planets at intermediate masses inside a few AU. At larger distances, planets of all masses formed. The study also produced a large population of giant planets orbiting close to the star, which is not seen by observational surveys. We note that the model used by \citet{ida:2004} was essentially deterministic, so that the variety of final planetary systems arose due to different initial conditions. In a follow-up study, \citet{ida:2008} added ``type-I'' migration in which planets that are too small to open a gap in the disk migrate due to linear tidal torques \citep{ward:1997}. These simulations yielded a population of small planets at a range of distances, and a pile up of planets at the inner edge of the disk, in common with the earlier study. However, planets always migrated to the inner edge of the disk before they could become gas giants unless type-I migration rates were artificially reduced by at least an order of magnitude. With reduced migration rates, planets that grew at early times were often lost, but planets growing at later times could survive, and some became gas giants. \citet{mordasini:2009a,mordasini:2009b} used a somewhat more detailed semi-analytic model to study the growth of single-planet systems using a range of disk conditions. Both type-I and type-II migration were included. The simulations yielded a diverse population of planets including giants, provided that type-I migration rates were reduced by 1--3 orders of magnitude. The authors also determined which subset of their sample (mainly giant planets) would be observable using existing Doppler radial velocity surveys. Comparing the mass and orbital distributions of this subset with the observed population, the authors found an acceptable match, again provided that type-I rates were greatly reduced. \citet{ida:2010} carried out a study using multiple planets per system. Dynamical interactions between planets were neglected apart from the possibility that planets undergoing converging migration could be trapped in mean-motion resonances. However, planets in the same system competed for resources, so that the growth of one planet was affected by the presence of its neighbors. The authors ignored gas accretion, focussing instead on the growth and evolution of solid planets. They found that small planets in the inner disk commonly captured one another in resonances, subsequently migrating together through the disk. The end result was a population of low-mass planets on short-period orbits qualitatively similar to the population of ``super-Earths'' found by the Kepler mission. \citet{alibert:2013} also considered multiple planets per system, but included dynamical interactions between the largest planets using an $N$-body integrator. Thus, their model added the effects of chance events during the evolution due to the chaotic orbital interactions between planets. The study compared the outcome of ensembles of simulations using single-planet and 10-planet systems. The authors found that the population of giant planets produced in each case was quite similar, suggesting that the formation of giant planets is not greatly affected by neighbors or chance events. However, the population of low-mass planets in each case was very different. The simulations with multiple planets generated a large population of short-period Earth and sub-Earth mass planets that was absent from the single planet case. \citet{ndugu:2018} looked at planet formation in disks subject to heating from other stars in their birth cluster, thus adding an element of environmental diversity to planet formation. The authors carried out a population synthesis of single-planet systems for planets accreting pebbles and gas, and subject to migration. The study found that planet formation via pebble accretion is sensitive to the local cluster temperature, primarily because it affects the disk scale height and pebble accretion efficiency in the outer disk. The authors concluded that the efficiency of giant planet formation, and its dependence on stellar metallicity, could be strongly affected by disk heating due to nearby stars. Thanks to observational surveys, we now have a large dataset of extrasolar systems that can be used as tests for population synthesis studies of planet formation. These surveys have found a broad variety of systems, most of which are very different than the Solar System. Planets with periods less than 100 days and masses between that of Earth and Neptune seem to be very common \citep{fulton:2017}, although there is no analog in the Solar System. Giant planets are relatively uncommon by comparison \citep{cumming:2008}. However, the observed planetary distribution is subject to a number of observational biases. For example, the Doppler radial velocity method tends to find planets that impart large reflex motions on their star, which favors the discovery of massive, short-period planets. Most discoveries are announced after a planet has completed at least one orbital period, which limits the maximum period to values comparable to the length of the survey. The transit method is even more biased towards discovering short-period planets, and tends to find planets with large radii, which favors planets with large mass, low density or both. Various studies have tried to determine the underlying planetary distribution taking into account observational biases. These studies often estimate the fraction of stars with planets in a certain range of orbital periods and a range of masses or radii. Some other studies estimate the number of planets per star with certain planetary characteristics. We will make use of several estimates of each kind here. In this paper, we adopt a population synthesis approach to examine the physics of planet formation. In common with many earlier studies, we consider the possibility that much of the observed diversity of planetary systems is due to differences in initial conditions. Unlike most existing studies, we focus on the growth of planets by accreting pebbles rather than planetesimals, since recent work suggests pebble accretion plays an important role in planet formation, especially the formation of giant planets \citep{lambrechts:2012}. Like some earlier works, we quantify how well the synthetic planetary distribution produced by our model matches observations, taking into account observational biases. However, we go further by iterating this procedure using an optimization scheme to successively refine the parameter values for the model. In this way we can determine the best parameter values that match the available observations, and thus help to constrain the physics of planet formation. The rest of the paper is organized as follows. Section~2 describes the planet formation model we use. Section~3 gives details of the observational constraints used to test the model. Section~4 describes how the model parameters are optimized iteratively by comparing the model output with observations. In Section~5 we look at the results for single planet systems, while Section~6 describes the results for multi-planet systems. Section~7 examines Solar System analogs formed by the best-fit models. Section~8 contains a discussion, and the main results are summarized in Section~9. % % PLANET FORMATION MODEL % | In this paper we have developed a simple model for planet formation with 8 free parameters that describe some of the main uncertainties in the physics of protoplanetary disk evolution, planetary growth, and planet orbital migration. We ran this model for sets of 2000 stars with protoplanetary disk masses and radii compatible with observed distributions, and random initial protoplanet orbital distances. The planets produced by the model were compared with 7 published studies that have estimated the orbital and mass/radius distributions of extrasolar planets and brown dwarfs. The model parameters were then improved iteratively using a particle swarm optimization scheme in order to find the best-fit set of parameters. The planet formation model considers a small number of protoplanets embedded in an evolving protoplanetary disk. Planets accrete solid mass in the form of pebbles that drift inwards through the disk due to gas drag. Objects that reach a certain size---the pebble isolation mass---perturb nearby gas sufficiently to halt the inward flux of pebbles. Planets above this mass can accrete gas from the disk at a rate that depends on the planetary mass, the cooling timescale for the planet's atmosphere, and the inward flux of gas. Low-mass planets undergo inward type-I migration due to tidal interactions with the disk. Planets massive enough to open a gap in the disk inwards migrate more slowly at a parameterized rate. Converging migrating planets can be trapped in their 2:1 mean-motion resonance. For both single-planet and multi-planet systems, the closest fit to the observations occurs for large initial protoplanet masses, low disk viscosities, short pebble drift lifetimes, and short cooling times for planetary atmospheres when gas accretion begins. The first three of these promote rapid growth of solid planets, while the last promotes rapid gas accretion. All four of these parameter values tend to lie close to the maximum or minimum permitted values. The best fit solutions find that planets reach the pebble isolation mass at a small fraction of the mass needed to open a gap in the gas disk. The best fits also indicate that pebbles undergo collisional disruption at speeds around 1 m/s, and that gas giants can accrete at most about half of the inward flux of gas through the disk. Overall, the best-fit models do a good job of matching the observed frequency and orbital distribution of giant planets and the observed abundance of brown dwarfs. The best fits provide a poorer match for the observed population of super Earths on short-period orbits. This is especially true for objects with radii 2--3 times that of Earth, which are substantially under represented in the model compared to observations. A few fits provide a better match to the observed super Earth population but at the expense of greatly underestimating the number of giant planets and brown dwarfs. We find that adding more protoplanets per system in the model improves the fit substantially compared to single-planet systems. This is significant because single-planet systems are often used in population synthesis studies. The best fit solutions continue to improve as more planets are added up to 8 per system, which is the maximum number that matches our constraint that neighboring objects begin exterior to their 2:1 mean-motion resonance. The main effect of adding more protoplanets is to increase the number of super Earths that form relative to the number of giant planets, which tends to improve the fit to the observed planetary distribution. Pebble isolation appears to be the key process that determines the characteristics of planetary systems produced by the model used here. Typically, planets do not start accreting gas until they reach the pebble isolation mass, since heat released by infalling pebbles prevents the atmosphere from cooling and contracting so that gas can flow inwards \citep{lambrechts:2014}. In our model, the pebble isolation mass increases with distance from the star since it is proportional to the mass needed to open a gap in the gas disk, which also increases with distance \citep{rafikov:2002}. Protoplanets also tend to grow faster closer to the star because pebble accretion is more effective. A combination of these factors determines where gas giant planets can form. In the inner disk, most planets quickly grow to the pebble isolation mass. In principle, gas accretion can begin at this point. However, the planetary masses are small, so gas accretion is extremely slow. These planets never accrete more than a small amount of gas as a result. In the outer disk, the pebble isolation mass is much larger. Any planets that reach this size can undergo rapid gas accretion and become giant planets. However, in most disks, the growth of solid planets in the outer disk is too slow to reach the pebble isolation mass before the disk disperses, especially since the pebble flux dwindles rapidly at later times. Only when the disk mass is large can solid planets in the outer disk grow fast enough to reach pebble isolation and undergo rapid gas accretion to form giant planets. Giants migrate inwards during and after their formation, although the degree of migration varies depending on the planetary mass, the time, and the nature of the disk. Thus, the final orbits of giant planets tend to occupy a wide range of distances extending to the inner edge of the disk. The solutions favored by the optimization scheme lead to rapid early growth of solid planets due especially to the small values of the turbulence strength $\alpha$ and the pebble flux decay timescale. Partly this is because rapid growth increases the chance that giant planet cores can form early enough to undergo the prolonged stage of slow gas accretion before runaway gas accretion begins. At least as important, however, is the fact that the gas surface density declines over time, and this means that the pebble isolation mass decreases as well. In order to form a giant planet, a protoplanet must exist in a part of this disk where the pebble isolation mass is large enough for runaway gas accretion to occur, and the object must grow fast enough to reach the pebble isolation mass before it has decreased too far. In this study, we find very low values for the pebble isolation mass. For a typical gas surface density of 100 g/cm$^2$ at 5 AU, the isolation mass is around an Earth mass for the best-fit single-planet case, and 2--3 times smaller for runs with 4 or more planets. By contrast, \citet{lambrechts:2014} found a pebble isolation mass of roughly 20 Earth masses at 5 AU. We note however that these authors used a high disk viscosity with $\alpha=0.006$ compared to $\alpha\sim10^{-5}$ for our best fit models. For viscosities this low, analytic calculations and hydrodynamical simulations both suggest a planet with a mass of about 10 Earth masses would open a deep gap at 5 AU \citep{rafikov:2002, duffell:2013}. The planetary mass needed perturb the gas enough to stop inward drifting pebbles is likely to be at least a few times smaller than the gap-opening mass \citep{lambrechts:2014}. Thus, the very small pebble isolation masses found here may not be too unreasonable. In the planet formation model used here, solid planets grow entirely due to pebble accretion rather than by accreting planetesimals or merging with other planets. Runaway and oligarchic growth of planetary embryos, which feature prominently in the standard picture of planet formation \citep{wetherill:1993, kokubo:1998} do not occur here. One consequence is that terrestrial planets form rapidly, within the lifetime of the gas disk. Earth is thought to have formed in $\sim 10^8$ years \citep{jacobson:2014}, much longer than the lifetime of a typical disk \citep{haisch:2001}, with much of its growth involving occasional giant impacts between planets after the gas dispersed. This discrepancy can be reconciled by treating the small, rocky planets that form in our simulations as tracers of a larger number of planetary embryos present in the inner disk that will eventually merge into a handful of terrestrial planets. We note, however, that many short-period super Earths in extrasolar systems appear to contain significant amounts of gas, so these objects must have formed while the gas disk was still present. An additional potential problem with forming Earth or its precursors from pebbles is that most of these pebbles probably originated in the outer solar nebula and drifted inwards to Earth's location. Initially at least these pebbles would be volatile rich since they formed in a cold environment. Most of these volatile constituents would be lost as pebbles entered the hotter inner nebula. However, it is unclear whether the surviving components can be reconciled with Earth's current composition cosmochemically. The planet formation model used here was designed to be extremely efficient so that it could be run for the large number of times required by the optimization scheme. Some potentially important physical processes were simplified or neglected as a result. For example, we adopt a relatively simple model for type-I migration appropriate for tidal torques due to gas in a non-magnetized, isothermal disk \citep{tanaka:2002}. More detailed models for migration have been developed which could change the conclusions of this paper \citep{paardekooper:2011, baruteau:2011, benitez:2018}. We also have not considered planetary orbital evolution after the protoplanetary disk disperses due to planet-planet scattering, planetesimal driven migration, or stellar tides. In our model, the radial temperature profile varies as $1/\surd a$ where $a$ is orbital distance. This is roughly the thermal profile for a disk that is heated entirely by the central star \citep{chiang:1997}. However, temperatures in the inner part of the disk may be dominated by heat released by viscous accretion, which typically leads to a steeper radial temperature profile \citep{chambers:2009}. This would alter several parts of the planet formation model used here such as the pebble capture rate, the planetary gas accretion rate, and the migration of massive planets. Perhaps most importantly, changing the thermal profile would alter the aspect ratio of the inner disk and thus the pebble isolation mass. To test whether this effect is important, we carried out sets of 20 optimization runs for 1- and 4-planet systems using a modified temperature profile. Here, temperature varies as $1/\surd a$ outside 1 AU, and as $1/a$ inside 1 AU, normalized to 300 K at the transition point. For the 1-planet case, the new best-fit solution is poorer for the modified disk thermal proflie. However, the best-fit solution is somewhat improved for the 4-planet case with a score of 1.68 for the best optimization run compared to 1.79 previously. This improvement is almost entirely due to increased production of super-Earths (radii 1--4 Earth radii) with periods of less than 100 days---the region considered by \citet{fulton:2017}. The frequency of these planets increases by roughly 60\% for the best-fit case, due in part to a modest increase in the pebble isolation mass. This suggests future population synthesis studies may need to consider carefully the thermal structure of the disk, as done for example by \citet{ndugu:2018}. Pebbles in the simulations described here typically have Stokes numbers around 0.01 (sizes of order 1 cm). It is unclear whether these pebbles are large enough to form protoplanets via the streaming instability, although some recent simulations suggest this is possible \citep{yang:2017}. It is also uncertain whether protoplanets would be as large as the Pluto-sized objects favored by the optimization scheme. However, only a handful of such bodies need to form, so these may represent the high mass tail of the distribution seen in numerical simulations \citep{schafer:2017} . The best fit for any of the optimization runs is for a case with 8 planets with a score of 1.4. This solution lies 1.4 standard deviations from the estimated observational constraints on average. Other solutions are poorer, in particular when fewer planets are included. This mismatch could be due to some of the shortcomings of the model discussed above. It is also possible that the different sets of observational constraints used here are not mutually compatible within errors. Some of the published observational constraints that we use only provide statistical errors. However, \citet{burke:2015} found that there are probably systematic errors of at least comparable magnitude. If this is true of the other observational surveys it would improve our model fit, although the details of the fit could change as a result. In this study, we considered ensembles of 2000 stars with different protoplanetary disk properties and initial protoplanet orbits. Typically, a different ensemble of stars was used for each optimization run. Different runs sometimes found different best-fit model parameters, implying that an ensemble size of 2000 is small enough that the random, discrete nature of the initial conditions can influence the result. This is particularly true for the random initial orbits of the planets, since the outcome is quite sensitive to where planets are located. We could have considered larger ensembles. However, given that we currently have observations of only a few thousand extrasolar systems, using a larger ensemble of synthetic systems probably won't tell us much. If the planet formation model used here is a reasonable reflection of reality, the results suggest some features of the observed extrasolar planet population are affected by stochastic events and small-number statistics. Interpreting these features will require a better understanding of the earliest stages of planet formation, especially which disk locations are most likely to generate protoplanets. Despite this uncertainty, the optimization runs do yield some robust results. For example, the observations are best explained if protoplanets are initially large, have short atmospheric cooling timescales, and form in disks with low turbulence levels. This indicates that population synthesis studies can be a useful tool for constraining at least some aspects of planet formation. % % SUMMARY % | 18 | 8 | 1808.03293 |
1808 | 1808.07348_arXiv.txt | We constructed new sets of He-enhanced (Y=0.30, Y=0.40) nonlinear, time-dependent convective hydrodynamical models of RR Lyrae (RRL) stars covering a broad range in metal abundances (Z=0.0001--0.02). The increase in He content from the canonical value (Y=0.245) to Y=0.30--0.40 causes a simultaneous increase in stellar luminosity and in pulsation period. To investigate the dependence of the RR Lyrae distance scale on the He abundance we computed new optical ($RI$) and near-infrared ($JHK$) Period-Luminosity-Metallicity-Helium (PLZY) relations. Interestingly enough, the increase in He content causes a minimal change in the coefficients of both period and metallicity terms, since canonical and He--enhanced models obey similar PLZ relations. On the contrary, % the classical $B$- and $V$-band mean magnitude-metallicity relations and the R-band PLZ relation display a significant dependence on the He content. The He-enhanced models are, at fixed metal content, 0.2-0.5 mag brighter than canonical ones. This variation is only marginally affected by evolutionary effects. The quoted distance diagnostics once calibrated with trigonometric parallaxes ({\it Gaia}) will provide the opportunity to estimate the He content of field and cluster RRLs. Moreover, the use of either spectroscopic or photometric metal abundances will pave the way to new empirical constraints on the universality of the helium--to--metal enrichment ratio in old (t$\gtrsim$ 10 Gyr) stellar tracers. | During the last century RR Lyrae (RRL) stars have played a crucial role as standard candles and tracers of old stellar populations ~\citep{marconi15, madore17, neeley17}. They are old (t$\gtrsim$10 Gyr) low--mass radial variables in their central helium burning phase and are observed in the Milky Way ~\citep{vivaszinn2006,zinn14,drake13,pietrukowicz2015}, in Local Group ~\citep{soszynski10b,fiorentino12a,coppola15} and in Local Volume galaxies ~\citep[][]{dacosta10,sarajedini12}. RRLs are used as standard candles since they obey a relation between absolute visual magnitude and iron abundance~\citep[][]{caputo00,cacciari03a,dicriscienzo04}. This relation, whose linearity has also been questioned in the literature \citep{caputo00,catelan04,dicriscienzo04}, suffers from significant intrinsic errors and systematics. RRLs do not obey a Period-Luminosity (PL) relation in the optical bands but, thanks to the characteristic behaviour of near-infrared (NIR) bolometric corrections~\citep{bono01,bono03c}, they obey a PL relation in the NIR regime~\citep[][]{longmore1990,braga15,coppola15}. The advantages of these relations are the small dependence on reddening and evolutionary effects~\citep[][]{bono03c} and a milder dependence on metallicity when compared with $B$,$V$ magnitudes. Theory and observations indicate that more metal--rich RRLs are fainter than metal-poor ones, but we still lack firm constraints on the coefficient of the metallicity term in the NIR PL relations ~\citep[][] {bono03c,catelan04,dallora04,sollima06a,marconi15}. Optical and NIR Period-Wesenheit (PW) relations are solid diagnostics to determine individual RRL distances, but rely on the assumed reddening law~\citep[][]{dicriscienzo04,braga15,coppola15,marconi15}. These relations are reddening free by construction ~\citep[][]{madore82,ripepi12,riess2012,fiorentino13,inno13} and include a color term. This means that they mimic a Period-Luminosity-Color (PLC) relation, tracing the position of each variable inside the instability strip (IS). These are the reasons why Period-Wesenheit relations have been widely adopted to trace the 3D structure of highly reddened clusters in the Galactic Bulge ~\citep{soszynski14,pietrukowicz2015}. The main motivations for the current investigations are the following. a) The helium-to-metals enrichment ratio ($\Delta\,Y$/$\Delta\,Z$=1.4, with a primordial He abundance of 0.245) adopted in evolutionary \citep[][]{pietrinferni2006} and pulsation \citep[][]{marconi15} calculations is still affected by large uncertainties. RRLs are good laboratories for estimating the He content \citep{caputo1998}. To provide a new spin on the determination of this parameter we are investigating new pulsation observables together with spectroscopic measurements of the metal content for field and cluster RRLs.\par b) Using the $\Delta$S method \citet{walker1991a} found that Bulge RRLs approach solar metallicity. This finding was recently supported by \citet{chadid2017} using high-resolution spectra, since they found several RRLs at solar chemical compositions. This means a metallicity regime in which RRL pulsation properties are more prone to helium effects \citep{bono95a,marconi11}. \par c) The RRL distance scale is going to play a crucial role to constrain possible systematics affecting primary distance indicators \citep{beaton2016}. Sizable samples of RRLs have already been identified in Local Group galaxies \citep{monelli2017} and beyond \citep{dacosta10}. However, we still lack firm theoretical and empirical constraints on the $\Delta\,Y$/$\Delta\,Z$ ratio in extragalactic systems. \par To overcome the limitations of the current theoretical framework we computed new sets of pulsation models with the same metal abundances (Z=0.0001--0.02) adopted in \citet{marconi15} but helium enriched (Y=0.30 and Y=0.40\footnote{Metals (Z) and helium (Y) abundances by mass fraction.}, Marconi et al. 2018, in preparation). | We have presented new sets of He-enhanced (Y=0.30, Y=0.40) nonlinear, time-dependent convective hydrodynamical models of RRLs covering the same range of metal abundances investigated by \citet{marconi15}. The model mean magnitudes in the $RIJHK$ bands were used to obtain new Period-Luminosity-Metallicity-Helium (PLZY) relations in these filters (see Table \ref{plr}). The main effect of an increase in He is an increase in the luminosity level, and in turn, in the predicted pulsation period. Therefore, an increase in primordial He content from the canonical value (Y=0.245) to He-enhanced (Y=0.30,0.40) causes a minimal change in the coefficients of both period and metallicity terms, since the He--enhanced models obey similar PLZ relations. Owing to the sensitivity of the luminosity level to He variations, the classical relations connecting the $B$ and $V$ mean magnitudes to metallicity and the $R$-band PLZ relation display a significant He dependence. The He-enhanced models models are, at fixed metallicity, 0.2$\div$0.5 mag brighter than canonical ones. This is an interesting opportunity, because Gaia is going to provide accurate geometrical distances to calibrate both the zero-point and the slopes of the diagnostics adopted to estimate individual RRL distances. Spectroscopic RRL abundances based on ground-based measurements \citep{magurno2018} will pave the way for an empirical calibration of the PLZ relations. This means the opportunity to determine distance, reddening and chemical composition (metal, helium) for field RRLs for which are simultaneously available optical ($BVRI$) and NIR ($JHK$) mean magnitudes. Note that this approach applies to RRL in nearby stellar systems, and in turn, the opportunity to investigate the helium-to-metal enrichment ratio currently adopted in evolutionary and pulsation calculations is universal. | 18 | 8 | 1808.07348 |
1808 | 1808.02047_arXiv.txt | {We study the vertical stellar distribution of the Milky Way thin disk in detail with particular focus on the outer disk. We treat the galactic disk as a gravitationally coupled, three-component system consisting of stars, atomic hydrogen gas, and molecular hydrogen gas in the gravitational field of the dark matter halo. The self-consistent vertical distribution for stars and gas in such a realistic system is obtained for radii between 4-22 kpc. The inclusion of an additional gravitating component constrains the vertical stellar distribution toward the mid-plane, so that the mid-plane density is higher, the disk thickness is reduced, and the vertical density profile is steeper than in the one-component, isothermal, stars-alone case. We show that the stellar distribution is constrained mainly by the gravitational field of gas and dark matter halo in the inner and the outer Galaxy, respectively. We find that the thickness of the stellar disk (measured as the HWHM of the vertical density distribution) increases with radius, flaring steeply beyond R=17 kpc. The disk thickness is reduced by a factor of 3-4 in the outer Galaxy as a result of the gravitational field of the halo, which may help the disk resist distortion at large radii. The disk would flare even more if the effect of dark matter halo were not taken into account. Thus it is crucially important to include the effect of the dark matter halo when determining the vertical structure and dynamics of a galactic disk in the outer region.} | The stars in a spiral galaxy are distributed in a thin disk with a finite vertical extent. An important property that characterizes a gravitating disk is its vertical density profile, $\rho(z),$ where $\rho$ is the mass density and $z$ is along the direction normal to the disk plane. This is because the mass distribution and dynamics in a gravitating system are inter-related (e.g., Binney \& Tremaine 1987). For a one-component disk, the self-consistent vertical density distribution for an isothermal disk is given as a $sech^2 (z)$ distribution (Spitzer 1942). A real galaxy disk, however, consists of stars and interstellar gas. The gas contains 10-15 \% of the disk mass (Young \& Scoville 1991, Binney \& Merrifield 1998). However, because of its lower dispersion, the gas forms a thin layer and hence can affect the vertical distribution of stars significantly despite its low mass faction. The self-consistent vertical structure of such a multi-component disk (consisting of stars, atomic hydrogen gas, $\mathrm{HI,}$ and molecular hydrogen gas, $\mathrm{H_2}$) where the components are gravitationally coupled was studied by Narayan \& Jog (2002b) for the Galaxy for radii $< 12$ kpc, where it was shown that the vertical distribution of each component is affected by the others. This approach has subsequently been used by many papers in the literature, including those that derive the disk properties from observations (e.g., Comeron et al. 2011) as it gives a more accurate result for the disk thickness. In this paper, we adopt the method proposed by Narayan \& Jog (2002b) and extend their study to the outer disk of the Milky Way between radii of 12-22 kpc. The aim is to study the dynamical effect of gas as well as of the dark matter halo on the vertical stellar profile of the thin disk. The density of the disk components as well as the dark matter halo density decrease with radius, but in the outer parts of a galaxy, the halo is likely to dominate the dynamics. Further motivation for our study comes from recent observations that show that the stellar disk flares in the outer Galaxy (e.g., Momany et al. 2006, Lopez-Corredoira \& Molgo 2014). We show that the gravitational effect of any coupled component causes the stellar distribution to be constrained toward the mid-plane, thus reducing the disk thickness and making the density profile steeper than the single-component stars-alone case. We find that the constraining effect of dark matter halo starts to dominate beyond 14 kpc. The thickness of the stellar distribution increases with radius, flaring steeply beyond 17 kpc. If the effect of dark matter halo were not included, the stellar disk thickness would flare even more. Thus the dark matter halo plays a pivotal role in confining the stellar distribution closer to the mid-plane in the outer Galactic disk. Section 2 contains the formulation and numerical solution of equations, and the input parameters used. Section 3 contains results for the constraining effect of gas and dark matter halo. Section 4 contains a brief discussion to show that the gas distribution is also vertically confined by the effect of stars and the halo. The main conclusions are given in Section 5. | 18 | 8 | 1808.02047 |
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1808 | 1808.00568_arXiv.txt | \biceparray\ is the newest multi-frequency instrument in the \bicep/\keck\ program. It is comprised of four 550\,mm aperture refractive telescopes observing the polarization of the cosmic microwave background (CMB) at 30/40, 95, 150 and 220/270\,GHz with over 30,000 detectors. We present an overview of the receiver, detailing the optics, thermal, mechanical, and magnetic shielding design. \biceparray\ follows \bicepthree\,'s modular focal plane concept, and upgrades to 6" wafer to reduce fabrication with higher detector count per module. The first receiver at 30/40\,GHz is expected to start observing at the South Pole during the 2019-20 season. By the end of the planned \biceparray\ program, we project $ \sigma(r) \sim 0.003$, assuming current modeling of polarized Galactic foreground and depending on the level of delensing that can be achieved with higher resolution maps from the South Pole Telescope. | Measurements of the polarization of the Cosmic Microwave Background provide key information to further our understanding of the early universe. The $\Lambda$CDM model predicts an $E$-mode polarization pattern in the CMB at the level of a few $\mu$K as well as an arc-minute $B$-mode polarization arising from gravitational lensing of $E$-modes by the large-scale structure of the universe. Inflationary gravitational waves may be a source of degree-scale $B$-mode polarization and a detection of such signal can be used to constrain the tensor-scalar ratio $r$ and place limits on the energy scale and potential of Inflation~\cite{kamionkowski2016quest}. While classes of Inflation models could generate undetectably low levels of gravitational waves, a detection of $B$-mode polarization generated by primordial gravitational waves would be direct evidence for the theory of Inflation. However, in order to disentangle a potential CMB signal from polarized Galactic dust and synchrotron foregrounds, we need to probe the polarization of the CMB at multiple frequencies with high sensitivity. \begin{figure}[h] \centering \includegraphics[width=0.75\textwidth]{bk_15_status} \caption[Published $B$-mode polarization measurements]{ Published $B$-mode polarization measurement by different experiments. $B$-mode polarization from gravitational lensing has been detected by the \bicep/\keck, SPT, \polarbear, and ACT collaborations. } \label{fig:bk_15_status} \index{figure} \end{figure} The current constraint on tensor-to-scalar ratio is $r_{0.05}<0.06$ at 95\% confidence from \bicep/\keck\ data in conjunction with \planck\ temperature measurements~\cite{BK15} (Figure~\ref{fig:bk_15_status}). Over the past 10 years, our experimental strategy of utilizing small-aperture, cold, refracting telescopes has proven to be successful to probe the degree-scale polarization of the CMB. \bicep2\ observed the sky with 500 antenna-coupled transition-edge sensor (TES) bolometers at 150\,GHz from 2010 to 2012, and reported a 5\,$\sigma$ excess of $B$-mode power over the base lensed-$\Lambda$CDM model in the range $30<l<150$~\cite{ade2014detection}. The \keck\ consists of five 25\,cm aperture receivers, each similar to \bicep2, started observations at 150\,GHz in 2012. A joint analysis with \planck\ indicated the signal reported from \bicep2 is consistent with polarized emission from Galactic dust~\cite{BKP}. The interchangeable \keck\ receivers allowed us to diversify the frequency coverage by rapidly switching each receiver to 95, 220, and 270\,GHz. The latest instrument in our program, \bicepthree, replaced \bicep2 on its mount in 2015. \bicepthree\ uses a 0.52\,m telescope and 2500 detectors operating at 95\,GHz to realize an on-sky instantaneous sensitivity of 6.7\,\ukcmbrts \cite{kang2018bicep3}. Figure~\ref{fig:bk_program} shows the progression of the \bicep/\keck\ program, to larger apertures, larger focal planes, and wider frequency coverage. \begin{figure} \begin{center} \includegraphics[width=0.75\textwidth]{bk_program} \end{center} \caption{The progression of the \bicep/\keck\ program leading to the \biceparray. Bottom row: the beam patterns of the focal planes on the sky shown on a common scale. Each square represents a single receiver, and the colors indicate different observing frequencies in pink (30/40\,GHz), red (95\,GHz), green (150\,GHz), and blue (220/270\,GHz).} \label{fig:bk_program} \end{figure} \biceparray\ adopts the same interchangeable concept used in \keck\ and is comprised of four \bicepthree-class receivers, each optimized for a atmospheric window in the frequency range from 30 to 270\,GHz (Figure~\ref{fig:atmos}). The highest and lowest frequency receivers incorporate two bands within an atmospheric window, operating at 30/40~GHz and 220/270~GHz, by shifting bandpass in alternating focal plane modules over the focal plane (Figure~\ref{fig:bk_program}). The splitting of the band provides more information on polarized Galactic synchrotron and dust emission to test the parameters of the foreground model. The \keck\ telescope mount will be replaced by the new \biceparray\ mount\cite{crumrine2018biceparray} at the end of 2019. \biceparray\ receivers will be installed in the new mount in a staged approach over the next few years, with the first 30/40\,GHz receiver to be deployed at the end of 2019, followed by the 150, 95 and 220/270\,GHz receivers. We are currently planning on continuing observations with a subset of the \keck\ receivers installed in open slots until they are filled by available \biceparray\ receivers. The parameters of the \keck, \bicepthree\ and \biceparray\ receivers are given in Table~\ref{tab:rxs}. \begin{figure} \centering \includegraphics[width=1\textwidth]{atmos} \caption{ {\it Left:} Comparison of atmospheric transmission at the South Pole with the bandpasses of \bicep/\keck\ and \biceparray. Median atmospheric transmission during the observing season is shown in black, bracketed by the 10\textsuperscript{th} and 90\textsuperscript{th} percentiles. Transmission drops only slightly across 200--300\,GHz, making dust observations in the upper part of this window effective, with similar dust sensitivity to the 220 GHz band. {\it Right:} Minimally processed timestream pair-sum and pair-difference noise spectra from \keck. The stable Antarctic atmosphere enables observations at all of these frequencies that are low-noise across the indicated science band from 0.1--1\,Hz, corresponding to $25\lesssim\ell\lesssim250$. } \label{fig:atmos} \index{figure} \end{figure} \begin{table} \small \begin{center} \begin{tabular}[c]{|l|c|c|c|c|} \hline Receiver & Nominal & Nominal Single & Beam & Survey Weight \\ Observing Band & Number of& Detector NET & FWHM & Per Year \\ (GHz) & Detectors& (\ukcmbrts) & (arcmin) & (\ukcmb)$^{-2}$ yr$^{-1}$\\ \hline\hline \keck & & & & \\ \ \ \ 95 & \textbf{288} & 288 & \textbf{43} & \textbf{24,000} \\ \ \ \ 150 & \textbf{512} & 313 & \textbf{30} & \textbf{30,000} \\ \ \ \ 220 & \textbf{512} & 837 & \textbf{21} & \textbf{2,000} \\ \ \ \ 270 & \textbf{512} & 1310 & \textbf{17} & 800 \\ \hline\hline \bicepthree & & & & \\ \ \ \ 95 & \textbf{2560}& 288 & \textbf{24} & \textbf{213,000} \\ \hline\hline \biceparray & & & & \\ $\big \langle \hspace{-3pt} \begin{array}{l} 30 \\ 40 \end{array}$ & $\begin{array}{l} 192 \\ 300 \end{array}$ & $\begin{array}{l} 221 \\ 301 \end{array}$ & $\begin{array}{l} 76 \\ 57 \end{array}$ & $\begin{array}{l} 19,500 \\ 20,500 \end{array}$ \\ \ \ \ $95$ & $3456$ & $288$ & $24$ & $287,000$ \\ \ \ \ $150$ & $7776$ & $313$ & $15$ & $453,000$ \\ $\big \langle \hspace{-3pt} \begin{array}{c} 220 \\ 270 \end{array}$ & $\begin{array}{c} 8112 \\ 13068 \end{array}$ & $\begin{array}{c} 837 \\ 1310 \end{array}$ & $\begin{array}{c} 11 \\ 9 \end{array}$ & $\begin{array}{c} 37,000 \\ 15,000 \end{array}$ \\ \hline\hline \end{tabular} \end{center} \caption[Receiver parameters and sensitivity for \bicep\ program]{Receiver parameters as used in sensitivity projections. Boldface numbers are actual/achieved quantities for existing receivers. The remaining values in the survey weight column are scaled from the achieved survey weights using only the ratio of the number of detectors, plus, if necessary to change frequency, the ratio of nominal NET values squared. In 2017 the 270\,GHz \keck\ receiver realized single-detector NETs of 1310~\ukcmbrts\ from pair differences in the best channels, but with a large dispersion due to excess detector noise in many detectors. An improved 270\,GHz focal plane is fielded in 2018.} \label{tab:rxs} \end{table} | 18 | 8 | 1808.00568 |
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1808 | 1808.09449_arXiv.txt | We present here new insights about the merging galaxy cluster Abell 2034 ($\bar{z}=0.114$) based on a combined weak lensing and dynamical analysis. From our deep Subaru $BR_Cz'$ images plus Gemini-GMOS/N low-resolution spectra accompanied by available redshift data, we have obtained the individual masses of the colliding subclusters as well as estimated a timeline of the process. The collision event happened $0.56^{+0.15}_{-0.22}$ Gyr ago between the subclusters A2034S ($M_{200}^S=2.35_{-0.99}^{+0.84}\times 10^{14}$ M$_{\odot}$) and A2034N ($M_{200}^N=1.08_{-0.71}^{+0.51}\times 10^{14}$~M$_{\odot}$) with the gas content of both subclusters displaced in relation to their galaxy and dark matter distributions, in a scenario similar to that found in the Bullet Cluster. Following our data and modelling the collision event is, most likely, taking place not so far from the plane of the sky, with an angle of $27^\circ\pm14^\circ$ in relation to that. In spite of the intrinsic degeneracy inherent to the system (whether it has been observed incoming or outgoing), the comparison of our calculated time since the closest approximation with the estimated age of the observed X-ray shock front and the increment experienced by the velocity dispersion of the galaxy cluster members points toward an outgoing movement. Besides, we found a strong evidence for the presence of a third structure which we called A2034W. | Immediately after the equipartition epoch, matter became the dominant species relative to the radiation content of the Universe. Baryonic matter and radiation were strongly coupled preventing the growth of small baryonic matter overdensities. Dark matter (DM), on the other hand was not significantly coupled to this fluid, allowing their collapse and posterior growing to form more massive structures. After recombination, baryons started falling into the potential wells created by DM halos that were formed. From that period and on, structures have evolved due to the gravitational attraction. In the hierarchical (or bottom-up) scenario this is followed by a continuously ongoing process where smaller DM halos merge to form the large galaxy clusters we now observe, the largest gravitationally bounded structures in the Universe. Some of these mergers have reached energies not seen since the Big Bang \citep[$\approx 10^{64}$ ergs;][]{sarazin04}. Merging galaxy clusters have proved to be fruitful laboratories to study their three main components (DM, galaxies and the hot intra-cluster gas -- ICM) as well the interactions among them. Especially interesting are the dissociative mergers, where systems are observed with significant spatial detachment between the DM and ICM distribution \citep[e.g.][]{Monteiro-Oliveira17a,Monteiro-Oliveira17b}. The galaxy cluster Abell 2034 (A2034) was classified by \cite{abell58} as a richness 2 system. It has been exhaustively observed by ROSAT \citep{davidformanjones99}, ASCA \citep{white00}, XMM-{\it Newton} \citep{sevenmergers} and {\it Chandra} \citep{kempner03, owers14}, whose data revealed that A2034 is, indeed, a bimodal system comprised by two subclusters each one located at South and North of the field. In this work, we name them as A2034S and A2034N, respectively. The cluster field is presented in Fig.~\ref{fig:A2034.field}. \begin{figure*} \begin{center} \includegraphics[width=\textwidth, angle=0]{A2034_field.png}% \caption{Deep optical $R_C$ image of the A2034 field taken by the Suprime-Cam mounted on the Subaru telescope. The overlaid red contours correspond to the ICM X-ray emission mapped by {\it Chandra} telescope. The X-ray emission presents a single peak located between the two BCGs (green circles). The explanation for the south X-ray emission excess is still an open question. The boxes enclose the regions considered for the statistical subtraction aiming to identify the galaxies belonging to the cluster red sequence: the blue ones are sufficiently far from the cluster centre (box {\it magenta}) and are supposed to be dominated by field galaxies.} \label{fig:A2034.field} \end{center} \end{figure*} The subclusters were identified as galaxy overdensities near each brightest cluster galaxy (BGC), respectively BCG S and BCG N, which are $\sim$ 5 arcmin apart from each other. The X-ray morphology presents itself as a unimodal distribution, offcentred $91\pm1$ arcsec (this work) from the nearest BCG S, suggesting that A2034 as a whole is out of equilibrium. Moreover, \cite{kempner03} found an discontinuity in the X-ray surface brightness, distant $\sim$3 arcmin from the BCG S toward BCG N, that they classified as a cold front. In the same region, \cite{kempnersarazin01} found an elongated radio emission. All these evidences points towards a recent merger between A2034S and A2034N. An apparent contradiction to the post merger scenario was the detection by \cite{white00} of a high cooling flow rate (about $\sim 90-580$ M$\odot$ year$^{-1}$), a feature that is correlated with undisturbed systems. However, \cite{kempner03} recalculated this value to $23^{+21}_{-20}$ M$\odot$ year$^{-1}$ and argued that the previous high value was biased by the cold front and the cold gas found at the excess emission region in the Southern of the cluster. The cluster X-ray luminosity is also high, at $L_{[0.1-2.4~{\rm keV}]}=3.51\times10^{44}$ erg s$^{-1}$ \citep{Piffaretti11}. The excess emission region is made of a gas colder than that found elsewhere in the nearby field. \cite{kempner03} argued that its temperature is sufficiently low to be a post collisional gas. They also speculated that its origin could be a background galaxy cluster at $0.30 < z< 1.25$ according to the $L_X-T$ relation. For the whole cluster, \cite{davidformanjones99} found $L_{[0.5-2{~\rm keV}]}=5.2\times 10^{44}$ erg and an X-ray bolometric luminosity of $2.2\times 10^{45}$ erg s$^{-1}$. Based on $1.4 ~\mathrm{GHz}$ VLA data, \cite{giovannini09} confirmed the existence of an radio diffuse emission previously identified by \cite{davidformanjones99}. However, \cite{vanweeren11} revisited this cluster and questioned this classification since the source does not show a clear correlation with the X-ray distribution. They have suggested that the radio source could be a radio relic, but they could not agree to a definitive classification. More recently, \cite{Shimwell16} have presented a complete radio study of the field with Low-Frequency Array (LOFAR) data. They found that the ICM presents a complicated emission and puts A2034 as a more complex system than expected. The brightest X-ray region also coincides with a giant radio halo and a group of three radio relics candidates were also detected in the cluster neighbourhoods \cite{sevenmergers} performed mass reconstruction through weak lensing using Subaru $R_C$ (2800 s) and $g'$ (720 s) images. They found a more complex scenario than that revealed by the X-ray observations. They found mass counterparts for the subclusters A2034S and A2034N (placed in front of the cold front) and other ones to the west from the cluster. Based on SDSS spectroscopic data, they speculated that them could be background structures. In Tab.~\ref{tab:mass.comp} we show a compilation of the mass estimates for A2034 available in the literature. \begin{table} \begin{center} \caption[]{Compilation of the mass estimates for A2034 available in the literature.} \begin{tabular}{c c c} \hline \hline M$_{200}$ & Method & Reference \\ ($10^{14}$ M$_{\odot}$) $h_{70}^{-1}$ & & \\[5pt] \hline $5.00\pm0.03$ & Caustic & \cite{geller13}\\[5pt] $11\pm4$ & Caustic & \cite{owers14}\\[5pt] $4.30_{+4.00}^{-2.40}$ & WGL$^1$-NFW & \cite{delliou15}$^2$\\[5pt] $3.54_{-1.17}^{+0.95}$ & WGL - NFW & This work (Sec.~\ref{sec:wl})\\ \hline \hline \multicolumn{3}{l}{$^1$ Weak gravitational lensing.}\\ \multicolumn{3}{l}{$^2$ Based on M$_{\rm vir}=7.17\pm4.3\times 10^{14}$ M$_{\odot}$ $h_{70}^{-1}$ obtained by}\\ \multicolumn{3}{l}{\cite{sevenmergers}}\\ \end{tabular} \label{tab:mass.comp} \end{center} \end{table} Using deep {\it Chandra} images and radial velocities of the member galaxies, \cite{owers14} claimed that A2034 is, in fact, a post collisional system whose axis is nearly coincident with the N-S direction. They re-classified the X-ray feature behind A2034N as a shock front travelling with $v_{\rm shock}\simeq2057$ km s$^{-1}$ (corresponding to a Mach number $\mathcal{M}=1.59_{-0.07}^{+0.06}$). The gas content of A2034N then would have been ripped out due to the collision, which happened 0.3 Gyr ago and had a small impact parameter, as revealed by the X-ray morphology. Combining the radial velocities of the BCGs, they estimate that the collision axis is at an angle of $\sim 23^\circ$ respective to the plane of the sky. Moreover, they have found an absence of background galaxies in deep SDSS images which exclude the idea of an background cluster responsible for the X-ray excess emission as suggested by \cite{kempner03}. They have interpreted this feature as a gas lost by A2034N during the passage though A2034S. In this paper we intend to map the mass distribution in A2034 as well to measure the masses of the individual subclusters. We also plan to characterize the system through the dynamical point of view, aiming to recover the merger history of the system. To these purposes, we have used deep $B$, $R_c$ and $z'$ images obtained with the Subaru telescope and spectroscopic data taken with Gemini/N telescope complemented with data available in the literature. This paper is organized as follows. In Section~\ref{sec:wl} we present the weak lensing analysis, from the description of the data to the mass measurements. The dynamical overview based on the galaxy redshift analysis can be found in Section~\ref{sec:dynamical}. The proposed merger scenario for A2034 is described in Section~\ref{sec:two.body}. All results obtained are discussed in Section \ref{sec:discussion} and summarized in Section \ref{sec:summary}. Throughout this paper we adopt the following cosmology: $\Omega_m=0.27$, $\Omega_\Lambda=0.73$, $\Omega_k=0$, and $H_0=70$~km~s$^{-1}$ Mpc$^{-1}$. At the mean cluster redshift of $z=0.114$ we then have $1$ arcsec equals $2.06$~kpc, the age of the Universe $12.4$~Gyr, and an angular diameter distance of $424.6$~Mpc. | \label{sec:discussion} \subsection{The merger between A2034S and A2034N} \label{sec:merger} We present here the weak lensing analysis of A2034 based on three deep filters which have allowed us to select the galaxy populations (foreground, red members and background) more accurately than previous works \citep{sevenmergers,vanweeren11,owers14}. Regarding the cluster members, both numerical density maps (Fig.~\ref{fig:density}) point to the presence of three main concentrations, being two of them nearly aligned with the North-South direction and centred on bright giant galaxies (BCGs). The recovered total mass distribution map shows a somewhat crowded field. However, two of those mass clumps can be directly related to each optical subcluster due to their proximity in relation to each BCG. Actually, both BCG positions are coincident with the related mass peaks within 95\% c.l. in A2034S and 68\% c.l. in A2034N. The merging system total mass was evaluated as $3.54_{-1.17}^{+0.95}\times10^{14}$ M$_\odot$ giving a mass ratio of $M_S/M_N=2.2_{-1.7}^{+1.1}$, characterizing the system as a major merger or even a semi-major merger considering the error bars \citep{martel14}. According to our model, A2034S appears as the most massive, with $M_{200}^S=2.35_{-0.99}^{+0.84}\times 10^{14}$ M${\odot}$, whereas A2034N has $M_{200}^N=1.08_{-0.71}^{+0.51}\times 10^{14}$~M${\odot}$. Both subclusters are separated by a projected distance of $727_{-142}^{+131}$ kpc. Our proposed reconstruction for the mass distribution in the merging system A2034S\&N shows a reasonable agreement with the findings of \cite{sevenmergers} (see their Fig.~9). Their results show the presence of two mass concentrations with similar significance (regions ``C'' and ``S'') close to BCG S. The choice of the mass clump ``C'' as representative for A2034S was done based on its spatial concordance with the X-ray peak whereas the clump S was related to the Southern X-ray emission. However, this bimodal configuration in A2034S could not be reproduced in our data even when we changed the size of the smoothing filter used to make the mass map (Fig.~\ref{fig:mass.map}). The same argument is valid for their clumps ``W1''and ``W2'' detected as a single structure in our mass map (\#3). Regarding A2034N, we found similar results showing that the mass concentration appears a little bit Eastern in relation to the BCG N. It is worth at this point to highlight that our source sample was comprised of 28 background galaxies by arcmin squared, whereas \cite{sevenmergers} had 52.4. Regarding the nature of the surrounding mass clumps (\#5, \#6, \#7 and \#8) we can only speculate. Those structures appear on the mass map but have no optical or X-ray counterparts. Actually, based on the weak-lensing mass reconstruction only we can not even say if those clumps correspond to individual structures or to a collection of small structures on the same line-of-sight \citep{Liu16}. In Tab.~\ref{A2034.tab:masses} we have estimated their masses assuming that they are at the same redshift as A2034, and those range from 0.6 to 1.8 $\times 10^{14}$M$_\odot$, excluding the peak \#8 whose location is too close to the border region to provide a trustful mass determination. However, given the non-detection of counterparts (Fig.~\ref{fig:optical.counter}), it is reasonable to assume that they could arise from the large scale structure. In that case, those masses can be overestimated by a factor up to 1.8 (at z$\sim$0.4; Fig.~\ref{fig:signal}). Given the depth of our data, one would imagine that a full investigation on the reality of the lowest S/N mass peaks or a search for the respective counterparts for those structures would probably require space-based imaging. Regarding the mass clumps \#3 and \#4 we are not mentioning them here because given their spatial positions they can be associated with A2034W, which we will discuss in Sec.~\ref{sec:A2034W}. \begin{figure} \begin{center} \includegraphics[width=1.0\columnwidth, angle=180]{A2034_signal.png} % \caption{Ratio between the surface critical density at A2034's redshift and at a given redshift ({\it black} continuous line). In this model, we have fixed all background sources at the mean redshift $\bar{z}=1.11$ as found in Sec.~\ref{A2034.tab:masses} after a colour cut in the COSMOS sample. The factor $\xi = \Sigma_{cr,{\rm A2034}}/\Sigma_{cr}$ allows us to estimate the expected mass at a given redshift by a direct comparison with the values presented in Tab.~\ref{A2034.tab:masses} ($M_{200}(z)=M_{200}(z=0.114)\times\xi^{-1}$). The horizontal {\it dotted} line corresponds to unitary ratio, for a comparison.} \label{fig:signal} \end{center} \end{figure} We can classify the system A2034S\&N as a dissociative merger \citep{dawson} since, within our mass map precision, the X-ray peak related to A2034S is found $169_{-42}^{+48}$ arcsec apart from the mass peak and $91\pm1$ arcsec in relation to the BCG S. In A2034N, as pointed by previous studies \citep{kempner03,owers14}, its gas content was stripped out during the collision in the sense that we cannot identify any gas concentration related to the mass clump. This configuration, therefore, resembles the famous Bullet Cluster \citep[e.g.][]{clowe04,clowe06} where the gas content of both subclusters was displaced in relation to the galaxies and the total mass distribution. The main difference is that, according to our dynamical analysis, the merger axis of A2034S\&N present comparatively a larger angle ($\alpha=27^\circ \pm14^\circ$) in relation to the plane of the sky. Although this estimate has been previously proposed by \cite{owers14}, they found this considering the subclusters' relative velocity to be equal to the shock propagation, which is known to be less precise \citep[e.g.][]{springel07,Machado+2015}. Moreover, the BCGs could not be at rest after the collision which may bias the value of $\delta v$ estimated by only their redshifts. The distribution of the members with spectroscopic data is well represented by a single Gaussian and this implies that the subclusters are no longer separated along the line-of-sight. In this case, the redshift alone cannot be used as a trustworthy tool to classify the member galaxies according to their host group. So, we turn to the galaxy projected spatial distribution. In fact, our results showed that A2034S\&N are separated in the line-of-sight by only $\delta v / (1+z)=403\pm228$ km s$^{-1}$. Departing from the mass posteriors we can apply the scaling relation proposed by \cite{biviano06} to evaluate the expected value of the subcluster velocity dispersion before the collision, under the assumption that it occurred with no mass loss. The comparison with the measured velocity dispersion can be used as an indicator of the dynamical effect of the merger on the galaxies. The pre-merger values are $675_{-96}^{+97}$ km s$^{-1}$ for A2034S and $548_{-117}^{+122}$ km s$^{-1}$ for A2034N which leads to boost factors $f\equiv\sigma_{\rm obs}/\sigma_{\rm pre}$ of $f_{\rm S}=1.42^{+0.29}_{-0.36}$ and $f_{\rm N}=1.59^{+0.42}_{-0.56}$. These results are another indication that the system is seen less than 1 Gyr after the pericentric passage \cite[see Fig.~29 in][]{pinkney} and suggest that our outgoing scenario is preferred in relation to the incoming configuration. In this case, we found that the system is seen $0.56_{-0.22}^{+0.15}$ Gyr after the pericentric passage. \subsection{A2034W} \label{sec:A2034W} After a complete description of the merging between A2034S and A2034N, which presents a very good match with the observational features, a question yet remains open: what is the nature of A2034W? The structure is characterized by $\bar{z}_W=0.1140\pm0.001$ and $\sigma_v/(1+\bar{z})=1061$ km s$^{-1}$ according to our dynamical analysis from 17 identified members with available redshift. In contrast with the red members spatial distribution, the recovered mass field does not show a clear counterpart for A2034W. In fact, it is located between two mass peaks (\#3 and \#4) as we can see in Figs.~\ref{fig:boot.peaks} and \ref{fig:Xray.galaxies}. Applying the same approach as used to determine the uncertainty on the A2034S\&N mass peaks centre, we found no spatial coincidence among A2034W and the mass peaks \#3 or \#4 (see Fig.~\ref{fig:boot.peaks}). Therefore, within the limitation of our analysis, the mass counterpart for A2034W remains unclear. The hypothesis that A2034W has no dark matter counterpart although possible is very unlikely. We found no remarkable galaxy concentration at the position claimed by \cite{owers14} and posteriorly by \cite{Shimwell16} (see their Fig.~3). We warn that this concentration was identified based only on the spectroscopic members. However, the choice of the targets for this kind of observation can induce a bias on their spatial distribution. On the other side, the use of photometric selected red galaxies avoid this effect since their selection is done based only on colour criteria. The red members distribution closely follow the radio emission map presented by \cite{Shimwell16}. In particular, we found that their region ``C'' (see their Fig.~1) is spatially coincident with the position of A2034W. We also have found a red members concentration around their region F. Another question we can propose is: what is the relationship between A2034W and the south emission excess in the X-ray distribution? As pointed out by \cite{kempner03} the properties of this emission are not consistent with an equilibrium state. However, it is cooler than the gas belonging to A2034S\&N to be considered part of the current merger. This is also corroborated by our general analysis of A2034S\&N. A possible hypothesis is that this cool gas could be a remaining part of a merger happened in the past and involving A2034W \citep[e.g.][]{kempner03}. \begin{figure} \begin{center} \includegraphics[width=\columnwidth,angle=-0]{A2034_map_compare.png} \caption{$R_C$ image overlaid with X-ray {\it Chandra} contours ({\it red}) and the spatial numeric distribution of the red cluster members ({\it green}). The numbers are positioned at the mass peak positions. The third galaxy concentration, A2034W, is located nearby the south emission excess present in the X-ray distribution. In spite of the proximity, the correspondence of the galaxy clump A2034W with the nearby mass clumps \#3 and \#4 is not obvious.} \label{fig:Xray.galaxies} \end{center} \end{figure} Finally, our analysis points towards the conclusion of \cite{Shimwell16} that the scenario found in the A2034 field is more complicated than proposed in previous analyses. This complexity can be a hint that previous events had taken place there. For example, a previous merger involving A2034W can explain the south excess in the X-ray distribution, as suggested by \cite{kempner03}. However, further hidrodynamical simulation is required to provide a comprehensive explanation for this complex scenario found in A2034. \begin{itemize} \item Despite the complexity of the field, the main system is well described by a bimodal merging between A2034N and A2034S closely aligned with the North-South direction. This fact is corroborated by the good match between the recovered age from our two-body analysis and those got by observational X-ray features; \item We found the individual masses $M_{200}^S=2.35_{-0.99}^{+0.84}\times 10^{14}$ M${\odot}$ and $M_{200}^N=1.08_{-0.71}^{+0.51}\times 10^{14}$~M${\odot}$, leading to a mass ratio $M_S/M_N=2.2_{-1.7}^{+1.1}$; \item The X-ray morphology presents only one peak, related to A2034S. It shows a detachment of $91\pm1$ arcsec and $169_{-42}^{+48}$ arcsec respectively from the BCG S and its related mass clump. On the other side, there is, within 95\% c.l., an agreement between both BCG S and its mass peak positions. Regarding A2034N, its gas content seems to have been also stripped out. \item The cluster member classification based on galaxy spatial distribution plus redshift has revealed that A2034S\&N are located not far from each other in relation to the line-of-sight. The relative velocity is only $\delta_v /(1+\bar{z}) = 497\pm255$ km s$^{-1}$; combined with an estimation of the perpendicular velocity based on the shock propagation, we found that the merger axis is located at $27^\circ \pm14^\circ$ from the plane of the sky; \item The two-body model predicts that the collision occurred $0.56_{-0.22}^{+0.15}$ Gyr ago with a 3D-velocity of $1767_{-334}^{+305}$ km s$^{-1}$. In spite of the model degeneracy, both shock presence and the measured boost in the velocity dispersion support the outgoing scenario; \item We confirm the existence of the Western structure A2034W although our analysis does not provide a way to identify its associated dark matter mass clump. \end{itemize} | 18 | 8 | 1808.09449 |
1808 | 1808.07489_arXiv.txt | Very strong \ion{Sc}{1} lines have been found recently in cool M~giants in the Nuclear Star Cluster in the Galactic Center. Interpreting these as anomalously high scandium abundances in the Galactic Center would imply a unique enhancement signature and chemical evolution history for nuclear star clusters, and a potential test for models of chemical enrichment in these objects. We present high resolution K-band spectra (NIRSPEC/Keck~II) of cool M~giants situated in the solar neighborhood and compare them with spectra of M~giants in the Nuclear Star Cluster. We clearly identify strong \ion{Sc}{1} lines in our solar neighborhood sample as well as in the Nuclear Star Cluster sample. The strong \ion{Sc}{1} lines in M~giants are therefore not unique to stars in the Nuclear Star Cluster and we argue that the strong lines are a property of the line formation process that currently escapes accurate theoretical modeling. We further conclude that for giant stars with effective temperatures below approximately $3800$\,K these \ion{Sc}{1} lines should not be used for deriving the scandium abundances in any astrophysical environment until we better understand how these lines are formed. We also discuss the lines of vanadium, titanium, and yttrium identified in the spectra, which demonstrate a similar striking increase in strength below $3500$\,K effective temperature. | With the advent of high resolution infrared spectroscopy, it has become possible to explore the spectra and composition of stars in the nuclear star cluster (NSC) just a few parsecs from the Galactic Center. Several chemical abundance studies have addressed the giants in the Galactic Center region and the NSC \citep[see e.g.,][]{ryde_schultheis:15,rich:17,do:18}. Spectroscopy in the Galactic Center poses a special challenge, as the high extinction generally at present restricts investigations to the infrared K band. Although there is considerable heritage in the 1.6 $\mu \rm m$ H band from e.g. APOGEE and earlier studies using NIRSPEC/Keck~II \citep[e.g.][]{origlia:11}, one is challenged by the paucity of weak lines suitable for abundance analysis, as well as the presence of molecular bands that cause blends. The cool, luminous, giants of the NSC are easiest to observe, but pose the greatest perils for analysis. Low resolution studies have advanced a scenario in which the NSC and nuclear disk have abundant metal rich stars, reaching to [Fe/H] $=+1$ \citep{do:15,Feldmeier-Krause2017}. \citet{rich:17} challenges this picture with new high resolution NIRSPEC spectra in the Galactic Center, and finding no stars above [Fe/H] $=+0.6$. \citet{do:18} reports high resolution spectra of NSC stars behind AO correction, arguing for extreme enhancements of scandium, vanadium, and yttrium; in some cases the analysis finds 10 times Solar abundance for these elements. Especially interesting are the strong \ion{Sc}{1} lines found in M~giants that are discussed in \citet{rich:17}. They suggest non-LTE effects as the cause for them, while \citet{do:18} argue for extreme over-abundance (as much as a factor of ten compared to iron) of scandium in the NSC. If confirmed this latter interpretation would be a chemical signature of the special environment in the Galactic Center, and potentially very important. There are good reasons to assume that the enrichment and star formation histories might be different in the Galactic Center, especially if the NSC has a unique formation history. Large magnetic fields, suppressed star formation, high turbulence, tidal forces, and the deep Galactic Center potential well might lead to a very different chemical history for the stellar populations in the Galactic Center. Furthermore, one might expect inhomogeneities in the trends with a larger scatter including outliers due to a possible mixture of populations that in principle might include the disk, inner, halo, NSC, nuclear disk, and bulge. A unique scandium abundance trend for the NSC would suggest that the Galactic Center and similar environments is a site for a new channel of nucleosynthesis of scandium and possibly other elements. Such a trend might also provide a powerful chemical tag for stars formed in nuclear environments. Scandium resides between the $\alpha$-elements calcium and titanium in the periodic table and is considered an iron-group element. Even titanium is sometimes considered an iron-group element, such as in the discussion of the metal-poor star HD84937 \citep{sneden:16}; therefore scandium can be seen as an intermediate element between the $\alpha$ elements and the iron-peak group. The precise origin of scandium and its only stable isotope, $^{45}$Sc, seems to be complex and is still a matter of debate. Scandium is produced in the innermost ejected layers of core-collapse SNe (type II) during neon burning and explosive silicon and oxygen burning via the radioactive progenitor $^{45}$Ti, as reviewed by \citet{woosley:95,romano:10}, while the contribution from type Ia SNe seems to be negligible \citep{iwamoto:1999,clayton:03}. However, the predicted trends of scandium with [Fe/H] disagree even when taking into account metallicity- and mass-dependent yields, which might be important \citep{woosley:95,nomoto:13,chieffi:02}. Chemical evolution models predict for too little scandium production. This could be due to the problems in the stellar yield calculations (see also \citealt{romano:10}). It should be noted, however, that although extreme enhancements of individual elements are known in stars of low metallicity, for stars with [Fe/H]$>$0, the total metal production is so high that no single supernova event can affect the abundance of a given species, save perhaps for an r-process production event that might result from a neutron star merger. These factors raise the bar significantly for any purported enhancements of metals in stars of high metallicity. From the observational point of view scandium seems to behaves like a typical $\alpha$-element, e.g.\ enhanced in the thick disk and the galactic Bulge \citep{Battistini15,lomaeva:18}. \citet{nissen:2000} and \citet{howes:2016} also find [Sc/Fe]$\approx +0.3$ abundances for halo stars and metal poor stars toward the bulge, however their studies have no stars more metal rich than [Fe/H]$\sim-1.5$. It is noteworthy that \citet{gratton:1991,prochaska:2000,ernandes:2018} report a constant $\rm [Sc/Fe] \sim 0$ (and as well for V) for galactic Bulge globular clusters spanning $-1.5 < \rm {[Fe/H]} < 0.0$. \myedit{\cite{smith:2002} report a solar mean scandium abundance for 12 red giants all having subsolar metallicity in the Large Magellanic Cloud---an investigation carried out in the K-band.} In the Galactic center, the picture has been more complicated. \citet{carr:00} measured scandium abundance for the cool supergiant star IRS7 located in the Galactic Center and found extremely strong \ion{Sc}{1} lines as well as \ion{V}{1} and \ion{Ti}{1} lines. An abundance of $\rm [Sc/Fe] \sim 0.9\,dex$ is required to fit the strength of the \ion{Sc}{1} lines. However, supergiant stars are affected by large velocity fields, depth-dependent turbulence, temperature inhomogeneities, etc., and it is known that fully realistic atmospheres for supergiants remain a challenge. We have employed high resolution K band spectroscopy to overcome the high and variable extinction toward the NSC \citep[see, e.g.,][]{ryde:16b,rich:17}. \myedit{High-resolution, K-band spectroscopy was pioneered already in the late 80s by Smith \& Lambert \citep{smith:85,smith:86,smith:90}.} We can now push fainter to observe cool M~giants with the largest telescopes, and thus avoid the supergiant stars. However, as we emphasized earlier, even the interpretation of the M~giant spectra is challenging. Our aim here is to test whether the strong scandium lines in cool M~giants in the NSC are due to either physical effects in the line-formation process \citep{rich:17} or due to intrinsically high scandium abundances \citep{do:18}. Toward this aim, we have acquired spectra for range of solar neighborhood stars, similar to those we have observed in the NSC in \citet{rich:17}, using the same instrument and telescope configuration (NIRSPEC on KECK II). These are used as a benchmark to compare with and we discuss different possible reasons for the strong \ion{Sc}{1} lines in the K-band spectra of M giants. | Our question in this paper is whether the abundance of scandium in stars observed in the Nuclear Star Cluster is really unusual. An unusual abundance would have a profound effect on our understanding of the formation of the Nuclear Star Cluster and its stellar population. There are several physical reasons for a spectral line to be strong, and by going through all these, we conclude that the strong \ion{Sc}{1} lines in the Nuclear Star Cluster most likely are due to line formation effects and certainly not due to an anomalous high scandium abundance. These \ion{Sc}{1} lines in the K-band are not good abundance diagnostics for the elements in cool M giants. Our conclusion is based on the fact that similar stars in the solar neighborhood, where the scandium abundances are known, show similarly strong \ion{Sc}{1} lines, much stronger than they can be modeled with a traditional synthetic spectroscopy. We conclude that the \ion{Sc}{1} lines in cool stars are strong everywhere in the Galaxy and is an inherent property of the line formation process, and should not be used to derive the scandium abundance of the stars. Lines of ionized scandium should be used instead or studies should use warmer stars. Non-LTE calculations for scandium, and perhaps other physical phenomena needing theoretical modeling, are needed before we can use the \ion{Sc}{1} lines at all. Our findings emphasize the perils of attempting to infer composition from a small number of lines that are far too strong for any conventional application of abundance analysis. Although high resolution infrared spectroscopy makes it possible to study the composition of stars in highly obscured regions, the relative paucity of weak lines requires that work in the infrared is approached with great caution. In the case where anomalies are suspected, it is important to turn to general established trends derived from optical measurements in the disk and bulge, as an additional essential check before claiming to discover unusual composition. We also note that at high metallicities, it would be quite surprising to see large enhancements of any species, that would imply extreme productions of individual elements in a metal rich environment. The cool giants we have analyzed are almost certainly on the asymptotic giant branch (AGB) and as such, they are likely to be the most luminous stars in their stellar population, contributing a substantial fraction of the K band integrated light. We know that the NSC is a complex stellar population with a wide range of age and we might expect many extragalactic nuclear star clusters to also have wide age ranges and possibly host very substantial populations of cool, luminous AGB stars. Individual supergiants can also contribute a substantial portion of the integrated light in the K band. There is a growing trend to attempt to infer detailed abundances (e.g. abundances of individual atomic species) from the integrated light of stellar populations. The temptation is strong for the extragalactic nuclear star clusters of low luminosity spirals, which have relatively low broadening due to the stellar velocity dispersion, making them very attractive targets for spectrum synthesis. We emphasize that if such a practice is carried out for the integrated light in the K band, that it be done using actual spectra of Solar neighborhood giants and a full population synthesis code. The numbers of cool AGB stars are small, and their contribution to the stellar population is stochastic \citep{frogel:90}. Even then, our understanding of line formation for cool giants is far from complete and therefore, we suggest that integrated light studies using only the K band be avoided. High resolution spectroscopy in the near-infrared is a relatively new subject area with great promise. The advent of very high resolution spectrographs on the current generation of 8-10m telescopes and future plans for ELTs give the subject great promise. It is important to support this endeavor with an equally serious investment in laboratory measurements, so that full advantage can be taken of the expected bounty of high quality data. | 18 | 8 | 1808.07489 |
1808 | 1808.05997_arXiv.txt | A recent study by \cite{Opher15} suggested the heliosphere has a ``croissant" shape, where the heliosheath plasma is confined by the toroidal solar magnetic field. The ``croissant" heliosphere is in contrast to the classically accepted view of a comet-like tail. We investigate the effect of the ``croissant" heliosphere model on energetic neutral atom (ENA) maps. Regardless of the existence of a split tail, the confinement of the heliosheath plasma should appear in ENA maps. ENA maps from the Interstellar Boundary Explorer (IBEX) have shown two high latitude lobes with excess ENA flux at higher energies in the tail of the heliosphere. These lobes could be a signature of the confinement of the heliosheath plasma, while some have argued they are caused by the fast/slow solar wind profile. Here we present ENA maps of the ``croissant" heliosphere, focusing on understanding the effect of the heliosheath plasma collimation by the solar magnetic field while using a uniform solar wind. We incorporate pick-up ions (PUIs) into our model based on \cite{Malama06} and \cite{Zank10}. We use the neutral solution from our MHD model to determine the angular variation of the PUIs, and include the extinction of PUIs in the heliosheath. In the presence of a uniform solar wind, we find that the collimation in the ``croissant" heliosphere does manifest itself into two high latitude lobes of increased ENA flux in the downwind direction. | The interaction between the solar wind and the partially ionized gas of the local interstellar medium (LISM) creates a bubble known as the heliosphere. Classically, the shape of the heliosphere has been regarded as comet-like, with a long tail pointed in the direction opposite the Sun's motion through the LISM \citep{Parker61, Baranov93}. In this view, the solar magnetic field was assumed to have a negligible effect on the global structure of the heliosphere. Computational models based on the magnetohydrodynamic (MHD) equations which include the LISM and a dipolar solar wind magnetic field have been able to reproduce this long tail structure \citep{Opher06, Pogorelov07, Opher09, Washimi11, Pogorelov13, OpherDrake13}. However, \cite{Opher15} and \cite{Drake15} used a unipolar solar magnetic field to limit artificial magnetic dissipation inherent in modeling the current sheet. They argued that the magnetic tension of the solar wind magnetic field instead alters the flows in the heliosheath and produces a ``croissant" heliosphere. While the inclusion of a dipolar solar magnetic field can weaken this effect, the use of this magnetic field configuration does not necessarily lead to a long tail structure \citep{Opher16, Michael18}. The \cite{Opher15} model is based on a 3D MHD simulation of a single ion fluid with four neutral fluids, and displays a shortened heliotail due to the presence of the lobes. The interstellar plasma is able to flow between these two jets, where a pressure balance between the thermal pressure of the interstellar plasma with the magnetic and thermal pressure of the lobes occurs. The jets were shown to exhibit turbulence, with turnover timescales on the order of years for the largest eddies. \cite{Drake15} confirmed the existence of the jets with an axisymmetric analytic model of the heliosheath. Using data from the Ion and Neutral Camera (INCA) on board the Cassini spacecraft and Voyager data, \cite{Dialynas17} argue that the heliosphere is tail-less, being ``bubble-like". While the ``bubble-like" shape is an idealization as the solar wind plasma needs to be able to escape, it does agree with the model of \cite{Opher15} in arguing for a tail-less heliosphere. It is important to emphasize that the two models do differ in the mechanism which structures the heliosphere. \cite{Dialynas17} suggest that the ``bubble-like" heliosphere is created by a strong interstellar magnetic field (as in \cite{Parker61}), while \cite{Opher15} show a structure and argue (see model by \cite{Drake15}) that the solar magnetic field plays a fundamental role in organizing the heliosphere. While both models suggest a shortened tail, these different structures would likely produce different observational signatures in ENAs. \cite{Pogorelov15} claimed that modeling the neutrals kinetically, as opposed to the four fluid approximation, would cause the disappearance of the two tails since the fluid-like treatment of the neutrals would suppress charge exchange across the region separating the lobes. \cite{Alexashov05} studied the effect of kinetic neutrals as compared to multi-fluid neutrals and found the plasma solutions to be similar within $\sim$5\% in the nose direction. Charge exchange strongly influences the heliosphere; however, the findings by \cite{Alexashov05} suggest that the inclusion of kinetic neutrals would not be sufficient to remove the jets in \cite{Opher15}. This is seen in \cite{Izmodenov15} where evidence of collimation of the heliosheath plasma by the solar magnetic field is seen even with neutrals modeled kinetically. Other possible effects that could suppress the two tails are the inclusion of solar cycle and the numerical erosion of the solar magnetic field between the two tails \citep{Pogorelov15}. The study of the effects of numerical erosion will be left to future work. The Interstellar Boundary Explorer (IBEX) was launched 19 October 2008 to study the heliosphere by observing energetic neutral atoms (ENAs) \citep{McComas09a}. ENAs provide an indirect method for studying the shape and thermodynamic properties of the heliosphere. IBEX is in a highly elliptical orbit around the Earth lasting roughly a week \citep{McComas11}, producing a global ENA map every six months for each energy band. There are two ENA cameras onboard IBEX: IBEX-Lo \citep{Fuselier09}, which measures ENAs from $\sim$10 eV to $\sim$2 keV, and IBEX-Hi \citep{Funsten09}, which measures ENAs from $\sim$300 eV to $\sim$6 keV. ENAs are created when an ion charge exchanges with a neutral atom, where the ion steals the electron from the neutral atom. Due to this charge exchange process, the energy of an ENA is dictated by its parent proton. Since different regions of the global heliosphere contain populations of ions with different energies, IBEX can probe these different regions by the energy signature of the ENAs. INCA was also able to observe ENAs while in orbit around Saturn. The primary objective of INCA was to observe ENAs originating from the plasma in Saturn's magnetosphere, but it was also able to image the heliosphere in directions away from Saturn \citep{Krimigis09, Dialynas13}. INCA imaged heliospheric ENAs in the $\sim$5 to $\sim$55 keV energy range, complementing observations of IBEX while also allowing a deeper look down the heliotail at energies $>30$ keV, where the charge exchange rate drops. Shortly after its launch, IBEX observed a ribbon-like structure in its ENA observations \citep{McComas09b}. This ribbon, known as the IBEX ribbon, is superposed on top of a globally distributed flux (GDF) of ENAs generated by processes within the inner heliosheath (IHS) and the LISM. \cite{Schwadron11} first attempted to separate the IBEX ribbon from the GDF by applying a mask over the region immediately surrounding the ribbon and interpolating with respect to the background flux. From this, ENA maps of both the GDF and the IBEX ribbon could be made separately. Due to this technique, ENAs originating in the IHS can be studied using the GDF maps. This technique was used again by \cite{Schwadron14} to study the GDF over the five years of IBEX data, which revealed how solar wind affected ENAs in the IHS at different energies. From a theoretical perspective, this separation of the GDF from the IBEX ribbon was done by \cite{Heerikhuisen14}, who investigated the ENA contributions from both the IHS and the interstellar plasma disturbed by the heliosphere to better understand the ribbon creation mechanism and to simulate the ribbon. \cite{Zirnstein15} also studied both the GDF and the IBEX ribbon from a theoretical perspective, relating it to the effects of solar cycle variations and focusing primarily on the effects in the nose direction. The first detailed analysis of IBEX tail measurements was published by \cite{McComas13}. The ENA maps showed two high latitude lobes in IBEX-Hi measurements, with an excess of flux from $\sim$2 keV to $\sim$6 keV. Additionally, at these energies a deficit of ENAs were also observed in the lower latitude heliotail \citep{McComas13, Schwadron14}. Unlike the ribbon, it is believed that the lobes seen in the heliotail originate from the IHS at these energies. \cite{McComas13} proposed that the observed lobe structures could be attributed to the fast and slow solar wind in the heliosphere. During solar minimum, the fast solar wind is deflected towards high latitudes while the slow solar wind exists at lower latitudes. It was suggested that the high latitude lobes exhibit an excess of ENA flux at higher energies since they originate from the fast solar wind. Likewise, the deficit of ENAs in the lower latitude heliotail was attributed to the presence of slow solar wind. This hypothesis of fast and slow solar wind being responsible for the lobes in the heliotail was investigated by \cite{Zirnstein16a} using the first five years of IBEX observations. \cite{Zirnstein16a} used a simple flow model of the heliosphere to simulate the deficit of ENAs in the lower latitude heliotail observed by IBEX at higher energies. While the presence of the fast and slow solar wind was shown to be a contributor to the lobes, it has not yet been shown that the fast/slow solar wind is entirely responsible for features in the tail. They found that the deficit was a result of the asymmetry in the solar wind. This study was followed by \cite{Zirnstein17} who created ENA maps of the heliotail after incorporating solar cycle dependence and extinction of pick-up ions (PUIs) in the IHS. Using a 3D MHD solution from \cite{Heerikhuisen13}, they were able to model the lobe structures seen in IBEX GDF measurements, but their model underpredicted the ENA flux. Here we intend to investigate how the ``croissant" heliosphere affects the GDF for a uniform solar wind. We first study the effect of ENAs from the confinement of the heliosheath flows by the solar wind magnetic field, a feature of the heliosphere that exists regardless of whether there is a comet-like tail or a split tail such as the ``croissant" heliosphere \citep{Michael18}. \cite{Izmodenov15} show similar confinement of the heliosheath plasma by the solar magnetic field, as evidenced by Fig. 6b in their work, which shows a peak in the mass flux around the azimuthal solar magnetic field; however, no synthetic ENA maps were created from this model. In the future, we will incorporate solar cycle time dependent effects as in \cite{Michael15}. We center maps on the nose as well as the tail to investigate both directions. \cite{Opher15} used a line of sight integration of the ion pressure multiplied by the neutral density to create a proxy for ENA integration. In this proxy map, it is shown that for high energy ENAs, two regions of strong emission should manifest themselves in the north and south. These regions of excess ENA emission show similar features to the ENA maps from \cite{McComas13}. Due to the cooling length at the energies of IBEX, we are only able to probe until certain distances. The cooling length is the distance out to which 1/e of the local ions have survived charge exchange along a streamline. At the energies between 1.7 keV - 6 keV, within the range of IBEX, the cooling length is around 100 AU (Fig. \ref{Stream}b) so we cannot observe the tail at distances much beyond this unless there is an additional mechanism that either scatters high energy ions into lower energies or drives low energy ions to higher energy. Turbulence within the heliospheric jets might act as an energy driver, but will not be included in the ENA analysis presented here. On the other hand, above 10 keV, the charge exchange cross-section drops exponentially \citep{Lindsay05}, increasing the cooling length. Therefore, at energies $>$50 keV such as those measured by INCA, the cooling length (Fig. \ref{Stream}b) is expected to increase to distances greater than 200 AU. While INCA may have been able to probe deeper down the tail, IBEX is able to make a more complete global ENA map. ENA maps from INCA required the removal of Saturn's magnetospheric contribution and at times the solar direction, whereas the maps of IBEX are corrected for Earth's magnetospheric contribution as IBEX enters the magnetosheath in addition to other sources of noise. The effect of Saturn and the Sun on INCA measurements was greater than that which IBEX experiences, making IBEX observations better for comparing the structure of the heliosphere, though the lower energies prevent a thorough exploration of ENAs from the tail due to the short cooling lengths. The purpose of this paper is to investigate the ENA maps produced by the ``croissant" heliosphere of \cite{Opher15}. We will explore the contribution of the heliospheric jets to ENA maps. In section 2, we present the model we use in calculating our simulated ENA maps. In section 3, we show our results. Finally in section 4 we discuss the implications of our results and how these results can drive future studies. | In the present work we show ENA maps of the ``croissant" heliosphere model. We included PUIs into our model, and allowed for both latitudinal and longitudinal variations of the PUIs. By exploring the effect of the ``croissant" heliosphere on ENA maps, we made the following conclusions: 1. \textit{Reproducing the IBEX lobes.} Using a uniform solar wind solution we were able to produce a high latitude lobe structure similar to those seen in IBEX ENA maps at energies greater than ${\sim}$2 keV. The presence of the lobes in our maps are caused by the collimation of the solar wind plasma via the solar wind magnetic field within the IHS, which affects the ENA flux signal even in the presence of extinction. 2. \textit{Required improvements.} We are unable to produce the strong ENA signal seen around the 1 keV energy band directly down the tail. Additionally, we are unable to quantitatively predict the flux values seen in IBEX observations. This will be explored in the future. | 18 | 8 | 1808.05997 |
1808 | 1808.02579_arXiv.txt | We consider photon signals arising from the annihilation or decay of low-mass (sub-GeV) dark matter which couples dominantly to quarks. In this scenario, the branching fractions to the various kinematically accessible hadronic final states can largely be determined from chiral perturbation theory. Several of these final states yield striking spectral features in the sub-GeV photon spectrum. New experiments, such as e-ASTROGAM and AMEGO, are in development to improve sensitivity in this energy range, and we discuss their potential sensitivity to this class of models. | There has been much recent interest in dark matter models in which the candidate particle has a mass $\lesssim {\cal O}(\gev)$. A variety of theoretical models have been developed in which such a candidate can have a weak coupling to Standard Model particles and can obtain a relic density consistent with observational limits (see, for example,~\cite{Dolgov:1980uu,Carlson:1992fn,Moroi:1993mb,Hall:2009bx,Hochberg:2014dra,Kuflik:2015isi}). Moreover, the sensitivity of direct detection experiments tends to be suppressed at such small masses, allowing these models to escape stringent experimental constraints. As such, a variety of experiments have been proposed to improve sensitivity to this class of models. In particular, a variety of astrophysical observatories, including e-ASTROGAM~\cite{DeAngelis:2017gra} and AMEGO~\cite{Caputo:2017sjw}, are being developed to fill in the current ``MeV-gap" in experimental sensitivity to photons, and these instruments will be well-positioned to probe the annihilation or decay of sub-GeV dark matter. It has recently been pointed out that, if sub-GeV dark matter couples predominantly to quarks, then the indirect detection signatures are particularly striking~\cite{Boddy:2015efa,Boddy:2015fsa,Boddy:2016fds,Bartels:2017dpb,Cata:2017jar,Dutra:2018gmv}. This is because there are few kinematically accessible particles when the center-of-mass energy of the annihilation or decay process is $\sqrt{s} \lesssim {\cal O}(\gev)$, and those particles tend to yield fairly striking photon signatures. In~\cite{Boddy:2015efa,Boddy:2015fsa}, the case $\sqrt{s} < 2m_{\pi^\pm}$ was considered. In this case, the dominant two-body final states are $\gamma \gamma$, $\gamma \pi^0$ and $\pi^0 \pi^0$; the particular final state is determined by the quantum numbers of the initial state, and the final photon spectra are particularly simple. If $\sqrt{s} > 2m_{\pi^\pm}$, then there are typically multiple final states for any choice of the initial state quantum numbers, and three-body final states are also important. But the branching fractions and kinematic distributions of the final state particles can be estimated using chiral perturbation theory (see, for example,~\cite{Cata:2017jar} and~\cite{CHPTReviews}). In this work, we will derive the photon spectra which arise in general for dark matter annihilation or decay to light mesons, assuming $\sqrt{s} \lesssim 1~\gev$. Our main assumption will be that primary electroweak interactions are negligible; the primary products of dark matter annihilation or decay will consist only of light mesons, with photons produced only by meson decay. For this purpose, we will find that the most important final states are those containing an $\eta$, which decays to $\gamma \gamma$ with $\sim 40\%$ branching fraction. We will limit ourselves to final states with at most three mesons. We will assume that dark matter couples to light quarks, and will find that the available final states can be classified by the quantum numbers of the initial state. We will determine the branching fractions and spectra of all relevant final states using chiral perturbation theory, and will assess the sensitivity of current and upcoming instruments. The plan of this paper is as follows. In section II, we will describe the application of chiral perturbation theory to sub-GeV dark matter. In section III, we will describe the photon spectra arising from meson decay. In section IV, we will present our results, and we conclude in section V. | We have considered the indirect detection of sub-GeV dark matter annihilation or decay. If dark matter couples to quarks, then the hadronic final states and branching fractions are largely determined by symmetry and kinematics, and can be derived in chiral perturbation theory. In particular, striking photon signals can be produced by the process $\eta \rightarrow \gamma \gamma$. Especially for the case of dark matter decay, the current lower bounds which can be obtained from EGRET data already exceed those obtained from Planck by orders of magnitude. Future data from an experiment such as e-ASTROGAM, looking at dwarf spheroidal galaxies, can provide an even greater improvement. In this work, we have utilized chiral perturbation theory at lowest order, and have focused on the photons arising directly from the decay process $\eta \rightarrow \gamma \gamma$, since the signal is well understood and the background is small. But dark matter annihilation or decay in this energy range generally produces a larger number of pions, especially after including $\eta$ decay, but the kinematics are more difficult. More generally, for slightly larger energies, a much wider range of final states is accessible and becomes relevant for indirect detection. But as increasing interest is shown in sub-GeV dark matter, it would be interesting to perform a more comprehensive study of the hadronic final states which can be produced. In a similar vein, it is worth noting that if sub-GeV dark matter couples primarily to light quarks, then it can potentially be produced at proton beam fixed-target or beam-dump experiments, such as NA62~\cite{NA62} or SeaQuest~\cite{Berlin:2018pwi}, or proposed experiments such as SHiP~\cite{Anelli:2015pba}, or related proposed experiments such as FASER~\cite{Feng:2017uoz}. In particular, one might hope that dark matter could be produced in the rare decays of heavier mesons. But to determine the available signals at such experiments, a more detailed study beyond lowest order in chiral perturbation theory would be necessary. | 18 | 8 | 1808.02579 |
1808 | 1808.05503_arXiv.txt | {Thanks to large dedicated surveys, large-scale magnetic fields have been detected for about 10\,\% of early-type stars. We aim to precisely characterize the large-scale magnetic field of the magnetic component of the wide binary \OmiLup, by using high-resolution ESPaDOnS and HARPSpol spectropolarimetry to analyse the variability of the measured longitudinal magnetic field. In addition, we investigate the periodic variability using space-based photometry collected with the BRITE-Constellation by means of iterative prewhitening. The rotational variability of the longitudinal magnetic field indicates a rotation period $P_{\mathrm{rot}}=2.95333(2)$\,d and that the large-scale magnetic field is dipolar, but with a significant quadrupolar contribution. Strong differences in the strength of the measured magnetic field occur for various chemical elements as well as rotational modulation for Fe and Si absorption lines, suggesting a inhomogeneous surface distribution of chemical elements. Estimates of the geometry of the large-scale magnetic field indicate $i=27\pm 10\,^{\circ}$, $\beta = 74^{+7}_{-9}\,^{\circ}$, and a polar field strength of at least 5.25\,kG. The BRITE photometry reveals the rotation frequency and several of its harmonics, as well as two gravity mode pulsation frequencies. The high-amplitude g-mode pulsation at $f=1.1057$\,\d dominates the line-profile variability of the majority of the spectroscopic absorption lines. We do not find direct observational evidence of the secondary in the spectroscopy. Therefore, we attribute the pulsations and the large-scale magnetic field to the B5IV primary of the \OmiLup system, but we discuss the implications should the secondary contribute to or cause the observed variability.} | \label{sec:Introduction} \subsection{Magnetic and pulsating early-type stars} \label{sec:intro_general} Large-scale magnetic fields are detected at the stellar surface of about 10\,\% of the studied early-type stars by measuring their Zeeman signature in high-resolution spectropolarimetry (e.g., MiMeS, \citet{2016MNRAS.456....2W}; the BOB campaign, \citet{2015IAUS..307..342M}; and the BRITE spectropolarimetric survey, \citet{2016arXiv161103285N}). These large-scale magnetic fields appear to be stable over a time scale of decades, have a rather simple geometry (most often a magnetic dipole), and have a polar strength ranging from about 100\,G to several tens of kG. Because the fields remain stable and their properties do not depend on any observed stellar parameters, we expect that these large-scale magnetic fields were produced during earlier stages of the star's life, relaxing into the observed configuration \citep[e.g.,][]{1999stma.book.....M, 2015IAUS..305...61N}. In addition, a dynamo magnetic field is likely to occur in the deep interior of early-type stars, produced by the convective motions of ionized matter in the convective core \citep{1989MNRAS.236..629M}. However, no direct evidence of such a magnetic dynamo has ever been observed at the stellar surface, nor is it expected, since the Ohmic diffusion time scale from the core to the surface is longer than the stellar lifetime. The large-scale magnetic fields detected at the surface of early-type stars have implications on the properties of the circumstellar environment, the stellar surface, and the interior, altering the star's evolution: \begin{itemize} \item The ionized wind material follows the magnetic field lines, and can create (quasi-)stable structures in the circumstellar environment or magnetosphere. The precise properties of the magnetospheric material depend on the stellar and magnetic properties \citep[e.g.,][]{2002ApJ...576..413U, 2005MNRAS.357..251T}. In general, magnetospheres are subdivided into centrifugal magnetospheres, where material remains trapped by the magnetic field and supported against gravity by rapid rotation, and dynamical magnetospheres. The latter has a region of enhanced density in the circumstellar environment that continuously accumulates new wind material and loses matter to accretion by the star. \item At the stellar surface, the large-scale magnetic field can affect the stratification and diffusion of certain chemical species at the surface, which can cause surface abundance inhomogeneities of certain chemical elements, and a peculiar global photospheric abundance composition. This would lead to rotational modulation of line profiles and photometric variability. Stars for which such peculiarities are observed are denoted by the Ap/ Bp spectral classification. \item The structure and evolution of the deep stellar interior is anticipated to be altered by the large-scale magnetic field, due to the competition of the Lorentz force with the pressure force and gravity. This leads to a uniformly rotating radiative envelope \citep[e.g.,][]{1937MNRAS..97..458F, 1992MNRAS.257..593M, 1999A+A...349..189S, 2005A+A...440..653M, 2011IAUS..272...14Z}, altering the depth over which material overshoots the convective core boundary into the radiative layer \citep[e.g.,][]{1981ApJ...245..286P, 2004ApJ...601..512B}. Currently, this effect has only been determined for two stars with a technique referred to as magneto-asteroseismology, which combines the analysis of the star's pulsations with that of its magnetic properties, and then performing forward seismic modelling. This was done for the magnetic $\beta$\,Cep pulsator V\,2052\,Oph \citep{2012A+A...537A.148N, 2012MNRAS.424.2380H, 2012MNRAS.427..483B} and the magnetic g-mode pulsator HD\,43317 \citep{2017A+A...605A.104B, 2018arXiv180500802B}. \end{itemize} Whenever present, the properties of the stellar pulsations depend on the strength, geometry and orientation of the large-scale magnetic field \citep[e.g.,][]{1982MNRAS.201..619B, 1984MmSAI..55..215G, 1985ApJ...296L..27D, 1990MNRAS.242...25G, 1992ApJ...395..307G, 1993PASJ...45..617S, 1995PASJ...47..219T, 1996ApJ...458..338D, 2000A+A...356..218B, 2005A+A...444L..29H, 2011A+A...526A..65M, 2017MNRAS.466.2181L}. For stars with a spectral type from O9 to B2, $\beta$\,Cep-type pulsations are expected. These are low-order pressure modes with periods of the order of several hours. For slightly less massive stars (spectral types B2 to B9) Slowly Pulsating B-type (SPB) oscillations are predicted. These are low-degree, high-order gravity modes with a period of the order of a few days. Moreover, their pulsation modes show a regular pattern in the period domain. Both the $\beta$\,Cep and SPB pulsations are driven by the $\kappa$-mechanism, related to the temperature dependent opacity of iron-like elements. In addition gravito-intertial modes, which are excited by the motions of the convective core and have the Coriolis force and buoyancy as restoring forces, are anticipated for early-type stars with periods longer than the stellar rotation period \citep[e.g.,][]{2014A+A...565A..47M}. Stellar pulsations remain the sole way to probe the interior of a single star, with the parameters of the pulsation modes dependent on the conditions inside the star. Variability due to gravity waves \citep[e.g.,][]{1981ApJ...245..286P, 2013ApJ...772...21R} was recently also found in photospheric and wind lines of the O9Iab star HD\,188209 \citep{2017A+A...602A..32A}, the B1Ia star HD\,2905 \citep{2018A+A...612A..40S}, and the B1Iab star $\rho$\,Leo \citep{2018MNRAS.476.1234A}, as well as in the close binary V380\,Cyg \citep{2014MNRAS.438.3093T}. The presence of gravity waves seems to be a common property of hot massive stars that have evolved beyond half of the core-hydrogen burning stage, irrespective of binarity or a magnetic field. \begin{figure*}[t] \centering \includegraphics[width=\textwidth, height = 0.33\textheight]{omilup_BRITE_lightcurves}% \caption{\textit{Top}: Satellite-orbit averaged and reduced BRITE light curves for \OmiLup, where the UBr photometry is indicated in red (\textit{left}) and the BAb photometry in blue (\textit{right}). Photometric variability is indicated in parts-per-thousand (ppt). \textit{Bottom}: Corresponding satellite-orbit averaged temporal variability of the on-board CCD temperature, showing different and discontinuous behaviour.} \label{fig:BRITE_lightcurves} \end{figure*} \subsection{omicron Lupi} \label{sec:intro_omilup} \OmiLup (HD\,130807, HR\,5528, HIP\,72683, B5IV, $V=4.3$\,mag) is an early-type star and a member of the Sco-Cen association, following its Hipparcos parallax \citep{2011MNRAS.416.3108R}. This region is a site of recent massive star formation at a distance of 118 -- 145\,pc, with the exact value depending on the sub-group of the association. Isochrone fitting to the Hertzsprung-Russell diagram indicates that the star formation occurred some 5 -- 20\,Myr ago. Using interferometry, \citet{1951CiUO..112...94F} detected a secondary component for \OmiLup, at an angular separation of 0.115\,arcsec. The most recent interferometric measurement indicated that the components have an angular separation of 0.043\,arcsec, with a contrast ratio of $0.28 \pm 0.06$\,mag \citep{2013MNRAS.436.1694R}. From the distance to the Sco-Cen association, the authors deduced that the components are 5.33\,au apart with a mass ratio of 0.91. Moreover, the distance to the Sco-Cen association implies that the largest measured angular separation is above 17\,au, such that the orbital period of the binary must be longer than 20\,years \citep{2011A+A...536L...6A}. Within the scope of the MiMeS survey, HARPSpol observations were collected for \OmiLup. \citet{2011A+A...536L...6A} concluded that \OmiLup hosts a large-scale magnetic field, with variability of the measured longitudinal magnetic field indicating a rotation period between one and six days. This agrees well with the small value of the projected rotation velocity, $v\sin i = 27 \pm 3$\,\kms \citep[determined by][]{2005ESASP.560..571G}. Moreover, \citet{2011A+A...536L...6A} determined $T_{\mathrm{eff}} = 18000$\,K and $\log g = 4.25$\,dex for \OmiLup from a comparison with synthetic spectra using \textsc{tlusty} non-local thermal equilibrium atmosphere models and the \textsc{synspec} code \citep{Lanz2007, Hubeny2011}. The authors also noted weaker \ion{He}{I} lines and stronger \ion{Si}{II} than expected from the solar abundances. Hence, the surface abundance of certain chemical elements seems to be peculiar. Lastly, Si, N, and Fe exhibited line-profile variations (LPVs) on a time scale of about one day. \citet{2011A+A...536L...6A} proposed surface abundance inhomogeneities as the cause of these LPVs. \OmiLup has recently been observed by the BRITE-Constellation of nano-satellites to monitor its photometric variability. This space-based photometry could aid in the determination of the rotation period of \OmiLup by observing the rotational modulation caused by surface abundance inhomogeneities due to the large-scale magnetic field. Moreover, it might permit us to determine the precise value and the physical process causing the variability with a period of about one day that was noted by \citet{2011A+A...536L...6A} as LPVs. Additional ground-based, high-resolution, optical spectropolarimetric data were collected to characterize the magnetic field of \OmiLup more precisely. We introduce the various observational data sets in Sect.\,\ref{sec:observations}, and indicate how these were prepared and corrected for instrumental effects when needed. In Sect.\,\ref{sec:st_params}, we estimate the stellar parameters of \OmiLup by fitting synthetic spectra to the observations and we search for evidence of the secondary component in the spectroscopy. The periodic photometric variability is investigated in Sect.\,\ref{sec:photometry}, while Sect.\,\ref{sec:magnetic} covers the analysis of the large-scale magnetic field. The sub-exposures of the spectropolarimetric sequences are employed to detect and characterize the LPVs in Sect.\,\ref{sec:LPV}. We end this work by discussing the obtained results in Sect.\,\ref{sec:discussion} and by drawing conclusions and providing a summary in Sect.\,\ref{sec:conclusions}. | \label{sec:conclusions} We combined HARPSpol and ESPaDOnS spectropolarimetry to study and characterize the large-scale magnetic field of \OmiLup. Using the variability of the measured longitudinal magnetic field, we determined the rotation period to be $P_{\mathrm{rot}} = 2.95333(2)$\,d, which agrees with earlier estimates that the rotation period would be of the order of a few days \citep{2011A+A...536L...6A}. We assumed that the primary component of the \OmiLup system hosts the large-scale magnetic field, given the lack of firm detection of a secondary component in the spectroscopy. Comparing the strength of the measured $B_l$ for various chemical elements, we noted large differences, indicative of chemical peculiarity and abundance structures at the stellar surface. The largest values were obtained for Fe, while the smallest values were derived from \ion{He}{I} lines. This suggests that Fe surface abundance inhomogeneities are located closer to the magnetic poles, while those for He are present near the magnetic equator. Yet, we cannot fully exclude a possible contamination by the secondary component of \OmiLup in the LSD Stokes\,I profiles. ZDI is needed to verify the locations of the suggested surface abundance inhomogeneities. Yet, this is not feasible with the current spectropolarimetric dataset, as we are lacking observations at several necessary rotational phases. Fitting models to the rotational variability of the measured $B_l$ values favors a description of a dipolar magnetic field with a quadrupolar contribution. This remains valid for the LSD profiles constructed with all metal lines, averaging out the effects of the surface abundance inhomogeneities, as well as for the LSD profiles from the Balmer lines. Typically, the strength of the quadrupolar contribution is about 10\,\% of that of the dipolar contribution. Using simple approximations, we estimated the inclination angle of the magnetic component of \OmiLup to be $i=27\pm10\,^{\circ}$, which then leads to an obliquity angle $\beta = 74^{+7}_{-9}\,^{\circ}$. A conservative lower limit on the polar strength of the large-scale magnetic field, measured from the LSD profiles of the Balmer lines, would be $5.25$\,kG. The BRITE photometry for \OmiLup shows up to six significant frequencies, indicating periodic photometric variability. Three of these frequencies ($f_1$, $f_2$, and $f_3$) correspond to the rotation frequency, and its second and third frequency harmonic. One frequency ($f_4$) is confirmed to be of instrumental origin, due to periodic variability of the satellite on-board temperature that was not perfectly accounted for during the correction process. The remaining two frequencies ($f_3$ and $f_6$) fall in the frequency domain of SPB pulsations. In case $f_3$ and $f_6$ originate from the magnetic component, \OmiLup\,A would be classified as a magnetic pulsating early-type star. However, the few detected pulsation mode frequencies are not sufficient for detailed magneto-asteroseismic modelling. Investigating selected absorption lines in the individual sub-exposures of the spectropolarimetric sequences indicates the presence of LPVs. The first moment of these absorption lines almost always indicate $f_3$ as the dominant frequency, except for the \ion{Fe}{II} line where $f_{\mathrm{rot}}$ was the dominant frequency. This is, again, suggestive of surface abundance inhomogeneities for Fe. Moreover, the equivalent width of the studied \ion{Fe}{II} and \ion{Si}{II} lines did change significantly with the rotation phase, demonstrating non-uniform surface abundances for these chemical species. The shape of the LPVs for the other selected absorption lines, where $f_3$ was dominant, agreed with a low-order pulsation mode, confirming that $f_3$ is a pulsation mode frequency. | 18 | 8 | 1808.05503 |
1808 | 1808.10700_arXiv.txt | We investigate the relationship between star formation activity and outflow properties on kiloparsec scales in a sample of 28 star forming galaxies at $z\sim$~2~--~2.6, using adaptive optics assisted integral field observations from SINFONI on the VLT. The narrow and broad components of the \Ha\ emission are used to simultaneously determine the local star formation rate surface density (\sfrsd), and the outflow velocity \vout\ and mass outflow rate $\dot{M}_{\rm out}$, respectively. We find clear evidence for faster outflows with larger mass loading factors at higher \sfrsd. The outflow velocities scale as \mbox{\vout~$\propto$ \sfrsd$^{0.34 \pm 0.10}$}, which suggests that the outflows may be driven by a combination of mechanical energy released by supernova explosions and stellar winds, as well as radiation pressure acting on dust grains. The majority of the outflowing material does not have sufficient velocity to escape from the galaxy halos, but will likely be re-accreted and contribute to the chemical enrichment of the galaxies. In the highest \sfrsd\ regions the outflow component contains an average of $\sim$45\% of the \Ha\ flux, while in the lower \sfrsd\ regions only $\sim$10\% of the \Ha\ flux is associated with outflows. The mass loading factor, $\eta$~=~$\dot{M}_{\rm out}$/SFR, is positively correlated with \sfrsd\ but is relatively low even at the highest \sfrsd: $\eta \lesssim$~0.5~$\times$~(380~cm$^{-3}$/n$_e$). This may be in tension with the \mbox{$\eta$ $\gtrsim$~1} required by cosmological simulations, unless a significant fraction of the outflowing mass is in other gas phases and has sufficient velocity to escape the galaxy halos. | \nocite{NMFS18a} \nocite{Newman12_406690} Galaxy scale outflows are expected to play a major role in regulating the star formation and chemical enrichment histories of galaxies \citep[e.g.][]{Dave12, Hopkins12, Vogelsberger13, Hirschmann13, Chisholm17}, mediating the co-evolution of galaxies and their central supermassive black holes \citep[e.g.][]{Silk98, Fabian12, King15}, and setting the sizes of galaxy disks \citep[e.g.][]{Okamoto05, Sales10}. Powerful winds driven by star formation and AGN activity transfer large amounts of mass and energy from galaxies to the surrounding circumgalactic medium \citep[e.g.][]{Peeples14, Tumlinson17}, depleting the supply of cold gas available for star formation within galaxies and preventing the circumgalactic gas from cooling and falling back onto galaxies \citep[e.g.][]{DiMatteo05, Springel05, Croton06, Hopkins06, Somerville08, Erb15, Beckmann17}. Outflows are therefore thought to play an important role in driving the low baryon fractions of galactic disks \citep{Dekel86, Efstathiou2000, Sales10}. The cosmic stellar mass density peaks at $\sim$20\% of the cosmic baryon density for a stellar mass of $\log(M_*/M_\odot$)~$\sim$~10.5, and drops to \mbox{$\sim$5-10\%} towards higher stellar masses (where black hole accretion feedback is most efficient) and lower stellar masses (where star formation feedback is most efficient) \citep[e.g.][]{Baldry08, Moster13, Moustakas13, Behroozi13}. Star formation driven outflows are expected to have the biggest impact on galaxies at $z\sim$~1~--~3, during the peak epoch of star formation \citep[see][and references therein]{Madau14}. Blueshifted absorption components are ubiquitous in the rest-frame UV spectra of $z~\sim$~2 star forming galaxies, revealing extended (tens of kiloparsec) reservoirs of diffuse outflowing material expelled from galaxies over long periods of time \citep[e.g.][]{Shapley03, Weiner09, Rubin10, Steidel10, Erb12, Kornei12, Bordoloi14}. Broad high velocity components in the rest-frame optical emission line spectra of star forming galaxies trace denser outflowing material within a few kiloparsecs of the launching points of the outflows. These emission components provide an instantaneous snapshot of the current outflow activity and are seen in $\sim$~10-30\% of star forming galaxies at $z~\sim$~2 \citep[when AGN host galaxies are explicitly excluded;][]{Genzel11, Newman12_global, Freeman17, NMFS18b}. Despite the prevalence of star formation driven outflows at high redshift, there are few quantitative constraints on their physical properties. Many studies have reported trends between the velocities of star formation driven outflows and the global $M_*$, SFR and/or \sfrsd\ of their host galaxies \citep[e.g.][]{Rupke05, Martin05, Weiner09, Rubin10, Steidel10, Erb12, Kornei12, Newman12_global, Talia12, Arribas14, Bordoloi14, Rubin14, Chisholm16, Heckman16, Sugahara17, NMFS18b}. Correlations between the outflow velocity \vout\ and star formation properties arise naturally because the level of star formation activity determines the amount of energy injected into the ISM by supernovae, stellar winds and radiation pressure from massive stars \citep[e.g.][]{Chevalier85, Strickland04, Murray11}. The stellar feedback combines with the turbulence driven by disk instabilities to counteract the disk gravity and launch outflows \citep[see e.g.][]{Ostriker11, Krumholz18}. The galaxy stellar mass drives the depth of the local potential which decelerates the outflowing material, but is also positively correlated with the SFR \citep[e.g.][]{Brinchmann04, Elbaz07, Noeske07, Peng10, Whitaker14} which determines the amount of energy available to accelerate the outflowing material. It is important to accurately characterise the SFR-\vout\ and \sfrsd-\vout\ relationships, because their scalings provide constraints on the primary outflow driving mechanism(s). If the outflows are driven by mechanical energy from supernovae and stellar winds, the outflow velocity is predicted to be weakly dependent on the level of star formation activity (\vout~$\propto$~\sfrsd$^{0.1}$; \citealt{Strickland04, Chen10}, \vout~$\propto$~SFR$^{0.2-0.25}$; \citealt{Ferrara06, Heckman00}). On the other hand, if the outflows are radiatively driven, the outflow velocity is predicted to scale strongly with the level of star formation (\vout~$\propto$~\sfrsd$^2$; \citealt{Murray11, Kornei12}, \vout~$\propto$~SFR; \citealt{Sharma12}). If the dominant outflow driving mechanism varies within individual galaxies, the power law scaling will be intermediate between the energy and momentum driven cases. The slopes of the \sfrsd-\vout\ and SFR-\vout\ relations remain a matter of debate. Some studies report relatively flat power law scalings with indices of 0.1-0.15 \citep[e.g.][]{Chen10, Arribas14, Chisholm16}, while other studies report somewhat steeper scalings with power law indices of 0.25-0.35 \citep[e.g.][]{Martin05, Weiner09, Heckman16, Sugahara17}. The discrepancies between different scalings reported in the literature are likely to originate from differences in the observed outflow tracers, adopted definitions of \vout, and range of probed outflow velocities (see discussions in e.g. \citealt{Kornei12, Heckman17}). The relationship between the level of star formation activity and the incidence and properties of outflows has, for the most part, only been investigated using galaxy integrated values. However, high spatial resolution observations of high SFR galaxies (both at $z~\sim$~2 and in the local universe) indicate that star formation driven outflows are launched from small ($\sim$~1~kpc) regions coincident with the most intense star formation events \citep{Shopbell98, Genzel11, Newman12_406690, Bolatto13}. Therefore, it may not be the global level of star formation which is most relevant for shaping the outflow properties, but the \textit{local} level of star formation. \citet{Bordoloi16} found that the properties of the outflowing material along different lines of sight to a lensed galaxy at $z~\sim$~1.7 are correlated with the properties of the nearest star forming region, suggesting that the outflows are indeed `locally sourced'. On the other hand, \citet{James18} found that the strongest outflow in a lensed galaxy at $z~\sim$~2.38 is associated with the most diffuse star forming region, suggesting that the outflow is `globally sourced'. Systematic studies of larger galaxy samples are required to conclusively determine whether the properties of star formation driven outflows are more strongly dependent on global or local galaxy properties. In this paper, we investigate the relationship between resolved $\sim$1-2~kpc scale (0.15-0.25'') star formation activity and the incidence and properties of outflows in a sample of 28 star forming galaxies at $z\sim$~2.3 from the SINS/zC-SINF AO Survey \citep{NMFS18a}. We build on the work of \citet{Newman12_global}, who explored the relationship between \textit{global} galaxy properties and the incidence and velocity of outflows in a similar sample of galaxies to the one used in this paper, and the work of \citet{Genzel11} and \citet{Newman12_406690}, who performed detailed analyses of the star formation and outflow properties of individual star forming clumps in 5 SINS/zC-SINF AO galaxies, of which 3 are included in our sample. Here we extend these analyses to study outflow properties as a function of resolved physical properties across 28 galaxies, considering not only the highly star forming clump regions but also the less active inter-clump regions. The outline of the paper is as follows. We probe the star formation and outflow properties on $\sim$0.17'' scales using adaptive optics assisted integral field observations of the \Ha\ emission line (described in Section \ref{sec:sample}). We stack the spectra of individual spaxels of the integral field datacubes to create high signal-to-noise (S/N) spectra in bins of resolved physical properties, and perform single and multi-component emission line fitting to analyse the properties of the outflow component in individual stacks (described in Section \ref{sec:sins_method}). We explore which physical properties are most closely linked to the presence of outflows in Section \ref{sec:broad_emission}, and present more detailed results on the relationship between the local \sfrsd\ and the incidence and properties of outflows in Section \ref{subsec:outflow_properties}. The implications of our findings are discussed in Section \ref{sec:discussion} and our conclusions are presented in Section \ref{sec:conc}. Throughout this work we assume a flat $\Lambda$CDM cosmology with \mbox{H$_{0}$ = 70 \kms\ Mpc$^{-1}$} and \mbox{$\Omega_0$ = 0.3}. | \label{sec:conc} We investigated the relationship between star formation activity and the incidence and properties of outflows on scales of 1-2 kpc in a sample of 28 star forming galaxies at $z\sim$~2~--~2.6 from the SINS/zC-SINF AO Survey. This work builds on previous studies of the relationship between \textit{global} galaxy properties and outflow properties in the SINS/zC-SINF AO sample \citep{Newman12_global}, and the relationship between the \textit{resolved} star formation and outflow properties of star forming clumps in 5 SINS/zC-SINF AO galaxies \citep{Genzel11, Newman12_406690}. With the aid of stacking we are able to probe not only the actively star forming clump regions which have been studied previously, but also the fainter inter-clump regions, spanning a factor of $\sim$50 in \sfrsd. We divided the spaxels from the 28 datacubes into bins of different physical properties (SFR, \sfrsd, stellar mass surface density \smsd, \sfrsd/\smsd, $A_V$ and galactocentric distance), and stacked the spectra of the spaxels in each bin to obtain high signal to noise \Ha\ line profiles. The \Ha\ profiles were used to simultaneously probe the star formation (from the narrow component of the line) and the outflows (which, when present, produce an additional broader line emission component). The width of the outflow component is a tracer of the outflow velocity, and the flux of the outflow component is a tracer of the mass in the outflow. Our main results are as follows: \begin{enumerate} \item The width of the \Ha\ line is most strongly dependent on the level of star formation (probed by the SFR and \sfrsd), supporting the notion that the observed broad emission is associated with star formation driven outflows. \smsd\ may also play a role in governing the incidence and properties of the outflows. \item The outflow component contains an average of $\sim$45\% of the \Ha\ flux emitted from the highest \sfrsd\ regions, but is less prominent at lower \sfrsd. \item The outflow velocity scales as \mbox{\vout~$\propto$~\sfrsd$^{0.34 \pm\ 0.10}$}. This scaling is shallower than the predicted \mbox{\sfrsd$^2$} dependence for outflows driven by momentum transport through radiation pressure, but steeper than the predicted \mbox{\sfrsd$^{0.1}$} dependence for outflows driven by kinetic energy from supernovae and stellar winds, suggesting that the observed outflows may be driven by a combination of these mechanisms. \item The outflow velocity is lower than the halo escape velocity in all but the highest \sfrsd\ regions, indicating that the majority of the outflowing material will not be expelled but will decelerate and fall back onto the galaxy disks. Simulations suggest that this material will likely be re-accreted after 1-2 Gyr, contributing to the chemical enrichment of the galaxies. \item The mass loading factor $\eta$ increases with \sfrsd. The normalisation of $\eta$ is uncertain due to the large uncertainties on the radial extent and electron density of the outflowing material, but we find \mbox{$\eta~\sim$ 0.4 $\times$ (380~cm$^{-3}$/n$_e$) $\times$} \mbox{(1.7~kpc/R$_{\rm out}$)}. This may be in tension with cosmological simulations (which typically require $\eta~\ga$~1 to explain the low efficiency of formation of low mass galaxies), unless a significant fraction of the outflowing mass is in other gas phases and is able to escape the galaxy halos. \item In 6/7 stacks the current star formation activity is powerful enough to drive the observed outflows. The energy and momentum outflow rates for the highest \sfrsd\ stack exceed the predicted rates for star formation driven outflows by a factor of 1.5. This may indicate that other energy sources (such as cosmic rays) contribute significantly to driving the outflows, that the SFR has changed since the outflows were launched, that the outflows have propagated away from their point of launch, that the adopted electron density or extent of the outflow is too low, and/or that the emission lines are partially broadened by mechanisms other than outflows (such as shocks or turbulent mixing layers). \end{enumerate} Our results confirm that \sfrsd\ is closely related to the incidence and properties of outflows on 1-2~kpc scales. In this paper we have only explored the average incidence and properties of outflows as a function of \sfrsd, which makes it difficult to draw strong conclusions on the relative importance of global and local galaxy properties in determining the properties of the outflows. In the future it will be important to investigate how much the outflow velocity and \Ha\ broad flux ratio vary at fixed local \sfrsd, and determine which local and/or global properties are responsible for driving these variations. | 18 | 8 | 1808.10700 |
1808 | 1808.01116_arXiv.txt | We measure the scale dependence and redshift dependence of 21 cm line emitted from the neutral hydrogen gas at redshift $1<z<5$ using full cosmological hydrodynamic simulations by taking the ratios between the power spectra of \HI--dark matter cross correlation and dark matter auto-correlation. The neutral hydrogen distribution is computed in full cosmological hydrodynamic simulations including star formation and supernova feedback under a uniform ultra-violet background radiation. We find a significant scale dependence of \HI\ bias at $z>3$ on scales of $k\gtrsim 1 \hmpci$, but it is roughly constant at lower redshift $z<3$. The redshift evolution of \HI\ bias is relatively slow compared to that of QSOs at similar redshift range. We also measure a redshift space distortion (RSD) of \HI\ gas to explore the properties of \HI\ clustering. Fitting to a widely applied theoretical prediction, we find that the constant bias is consistent with that measured directly from the real-space power spectra, and the velocity dispersion is marginally consistent with the linear perturbation prediction. Finally we compare the results obtained from our simulation and the Illustris simulation, and conclude that the detailed astrophysical effects do not affect the scale dependence of \HI\ bias very much, which implies that the cosmological analysis using 21 cm line of \HI\ will be robust against the uncertainties arising from small-scale astrophysical processes such as star formation and supernova feedback. | \label{sec:introduction} The acceleration of the Universe has been one of the greatest mysteries since it was first discovered by the observations of type Ia supernovae \citep{Perlmutter+1999}. One of the most natural explanations of the accelerated expansion is the dark energy in the regime of general relativity or modified theory of gravity \citep[e.g.][for review]{Clifton+2012}. Because the acceleration only becomes effective at the late epoch of $z\lesssim 1$, the most promising probe of dark energy or modified gravity is the large-scale structure of the Universe. Baryon acoustic oscillation (BAO) is recognized as a useful technique which is least affected by the systematics to constrain the dark energy models \citep[e.g.][]{Albrecht+:2006}. After the first detection of BAO by the clustering of luminous red galaxies (LRG) in the Sloan Digital Sky Survey (SDSS) \citep{Eisenstein+:2005}, \rev{\RA{significant attention has been paid to constrain the dark energy using BAO in the power spectrum and correlation function}} \citep[e.g.][]{Ross+2015,Beutler+2011}. As the BAO is a measurement of the oscillation peak scales, an accurate prediction of the peak scales is required. It is well known that the oscillation peak scale is readily changed by the non-linear clustering of matter \citep{Nishimichi+2007} or the non-trivial couplings among different fluctuation modes due to galaxy bias \citep[e.g.][]{Cole+:2005, Dalal+2008b, McDonald+:2009}. Another important aspect of the BAO is the combination of parallel and perpendicular components to the line of sight \citep{AP+:1979}. Although the AP-test makes the BAO a more powerful tool to constrain cosmological parameters, the systematic effect due to redshift space distortion (RSD) has to be taken into account as it is degenerate with the AP effect. As we can observe galaxies only in redshift space, the distance to the galaxies are contaminated by the peculiar velocities of galaxies; on large scales, galaxies are coherently attracted toward the overdensity regions which makes the anisotropic two dimensional correlation function squashed, while on small scales, non-linear random motion makes correlation function elongated along the line of sight \citep[e.g.][]{Matsubara:2004}. The RSD is important not only for correctly understanding the distortion of the correlation function to utilize the AP effects in the BAO, but also to gain an independent cosmological information from the BAO. Since the RSD is a direct measure of the velocity field, it is sensitive to the potential fluctuation $\Phi$ and thus to the theory of modified gravity \citep[e.g.][for review]{Hamilton:1998}. The current measurements of the BAO and RSD have been mainly focused on the galaxy or QSO distribution as they are considered to be good tracers of the large-scale structure \citep{Kirshner+1981, Alam+2017, Padmanabhan+2007, Torre+2017, Zhao+2016, Busca+2013, Slosar+2011}. However, due to the difficulty of taking the spectrum to accurately measure the redshift of the sources, we have studied the BAO only at $z<3$. The first detection of the reionisation absorption signature by the EDGES observation of 21 cm line absorption associated with neutral hydrogen (\HI) gas \citep{Bowman+:2018} has also opened a new window to probe the large-scale structure. Even after the epoch of reionisation, some fraction of neutral hydrogen is confined within the high-density regions such as inside the galaxies preventing the ultra-violet photons to penetrate. Several surveys to map the 21 cm distribution by intensity mapping are proposed where individual objects are not resolved but a continuous smoothed sky distribution is mapped out. For example, the Square Kilometre Array (SKA) will cover the 25,000 (SKA1) square degree of sky with 50kHz frequency resolution, which is adequate for accurate redshifts at $0.35<z<3.06$ (for SKA1-MID), but with a moderately coarse angular resolution for the single-dish observation \citep{Santos+:2015, Bull+2015}. There are interferometer mode in SKA which has significantly better angular resolution (depends on the configuration), however, its small field of view is not suited for a wide sky coverage and thus for cosmological analyses. Another example is the Baryon acoustic oscillations In Neutral Gas Hydrogen (BINGO) which will target much lower redshifts at $0.13< z<0.48$ \citep{BINGO:2012}. BINGO will cover $15\times 200$ square degree of sky with 1MHz frequency resolution with a resolution of 40 arcmin for single-dish observation. Much attention has been paid to the cosmological application of \HI ~observations \citep[e.g.][]{Camera+:2015, Bull+2015, Raccanelli+2015, Olivari+:2018, Villaescusa-Navarro+:2017, Obuljen+:2018, Dinda+:2018}. It has been shown that the Square Kilometre Array (SKA) has a capability to constrain the dark energy parameters comparable to that from the galaxy redshift surveys such as Euclid \cite{Bull+2015} which is a Stage-IV survey according to the Dark Energy Task Force \citep{Albrecht+:2006}. However, it is also shown that the difficulty of the \HI\ observation lies in the foreground removal \citep[e.g.][]{Wyithe+:2005} and in the modeling of \HI\ bias. As in the case for galaxies or QSOs, it is important to understand the connection between \HI\ gas and dark matter distribution, as theoretical predictions are often only for dark matter, the most dominant component of matter in terms of gravitational interaction. While \cite{Bull+2015} assumed the simplest constant bias, \cite{Umeh:2017} found that the non-linear coupling between different fluctuation modes or shot-noise can modulate the amplitude of power spectrum even on large scales. Therefore, considering more realistic models or measurements from numerical simulations is of great importance for robust cosmological analyses. To predict the \HI\ bias, halo model approach has been considered \citep{Padmanabhan+:2017, Penin+:2018}, and several work have been done using pure N-body simulations \citep{Sarkar+2016,Sarkar+2018a}. They do not fully solve the radiative transfer but populate the \HI\ according to the mass of the host dark matter halo. In this paper, we measure the scale- and redshift-dependent \HI\ bias using the full cosmological hydrodynamic simulation developed by the Osaka group, which takes gas dynamics into account with an appropriate UV background radiation and star formation with supernova feedback. We also perform the same analyses using the publicly released data of the Illustris simulation, and compare with our results to examine the impact of AGN feedback on the cosmological signals. % \rev{{It is well known that $\Omega_{\rm HI}$ depends on the mass resolution of the simulation \citep[e.g.][]{Nagamine2004a,Dave2013}, therefore we use the Illustris-3 simulation which has similar resolution to our fiducial run. This allows us to perform a fair comparison and to minimize the effect of mass resolution. }} This paper is organized as follows. In Section~\ref{sec:simulation}, we describe the details of cosmological hydrodynamic simulations used in this work, and explain how we generate the mock \HI\ data based on these simulations. In Section~\ref{sec:pspec}, we show the results for measuring the \HI\ bias in real space to explore the redshift and scale dependence of the bias. In Section~\ref{sec:modeling}, we show the anisotropic \HI\ power spectra in redshift space, and compared them with the linear and non-linear clustering models. Section~\ref{sec:results} is devoted to the interpretations of our results, and then we give a summary in Section~\ref{sec:summary}. Throughout this paper, we assume the cosmological parameters consistent with the WMAP-9 year result \citep{WMAP9}. | \label{sec:summary} Future 21 cm surveys will reveal the three dimensional distribution of neutral hydrogen gas over cosmological scales, which will potentially be a new probe of the large-scale structure of the Universe. In this paper, we explore the properties of \HI\ clustering using two different set of cosmological hydrodynamic simulations, the Illustris and the Osaka simulations that include nonlinear baryonic effects of star formation and feedback. We first measure the scale and redshift dependences of \HI\ bias in real space by taking the ratio of power spectra of {\HI}--dark matter cross correlation and dark matter auto correlation. Fitting with the constant plus linearly-scaled bias with $k$, \rev{{we find that the \HI\ bias monotonically increases with redshift for both simulations. This result is consistent with the Illustris-TNG simulation \citep{TNG+2018}, but the redshift evolution is stronger in the Illustris-1\,\&\,3 than in TNG100-1.}} We also find that the \HI\ bias shows a significant scale dependence at $z>4$ up to the scales where the perturbation theory holds, but it is consistent with being constant at $z\leq 3$. If we limit our analysis to the large scales of $k<0.25\,\hmpci$, we find no evidence of scale dependence at $1<z<5$. In both cases, the best-fitting bias parameters are fairly consistent between Illustris and Osaka simulations, which implies that the scale dependence of \HI\ bias on large scales is not sensitive to the details of the small-scale astrophysics. This means that, as far as we use the large scale modes, the cosmological analysis such as the determination of BAO scale is unlikely to be affected by the astrophysical uncertainties of feedback on small scales. However, at the same time, if one use the data more aggressively up to higher $k$, we certainly need accurate knowledge on the astrophysical effects such as supernova or AGN feedback. We leave more detailed and thorough investigation of the astrophysical impact of feedback on the \HI\ power spectrum as a future work. \rev{\RA{Further discussion on the evolution of $\Omega_{\rm HI}$ is given in the appendix.}} We then measure the redshift space distortion using the anisotropic two dimensional power spectrum. We jointly fit the monopole and quadrupole of the Legendre expanded power spectra including the peculiar velocity effect to the models widely applied for galaxy redshift surveys, with the free parameters of bias $b_{\HI}$ and velocity dispersion $\sigma_v$. We note that, since we only have two simulations, the cosmic variance largely affects the amplitude of large scale fluctuations. Therefore, we fit the data only in the range of $0.18<k<k_{\rm max}$ where $k_{\rm max}$ is given by Eq.~(\ref{eq:k_max}). We find that the measured bias parameter in redshift space is consistent with the one directly measured in real space from the ratio of power spectra. We also find that the velocity dispersion of \HI\ gas is systematically below the prediction from linear perturbation theory but marginally consistent with the prediction. Compared with the previous work by \citet{Sarkar+2018a}, we find significant disagreement on the values of \HI\ bias for the entire redshift range, which may mainly arise from the prescription of the \HI\ gas assignment to the dark matter halos in an $N$-body simulation by \citet{Sarkar+2018a}. On the other hand, the best-fitting values of our velocity dispersion are fairly consistent with the previous work within the statistical error of the single box simulation. Although the simulations used in this paper solve baryonic distribution and hydrogen ionization process in a more realistic manner, more detailed analysis will be required to fully understand the discrepancy. In this paper, we also introduced a new empirical model for RSD. The model is a simple replacement of linear power spectrum $P^{\rm lin}(k) \rightarrow P^{\rm NL}(k)$ in the Kaiser formula with the FoG prefactor. This model is consistent with the full TNS model on scales $k<k_{\rm max}$, and it gives a better fit of monopole at $k>k_{\rm max}$ but slightly off from the data for quadrupole on those scales. We construct a mock simulated data assuming that the 21 cm line is observed by future SKA-like survey for both interferometer and single-dish modes. We find that, for single-dish observation (i.e. low angular resolution observation), the models systematically underestimate the bias parameter and velocity dispersion. It will require a model in which the coarse angular resolution has been taken into account. In this paper we have limited ourselves to the discussion of the clustering properties of \HI\ gas, however it is straightforward to extend our analysis to the cosmological parameter recovery, such as the growth rate $f$ or dark energy parameters $w$ or $\Omega_{\rm DE}$. We leave these analysis to our future work, in which we will employ a larger number of hydrodynamic simulations with larger box-sizes. \appendix | 18 | 8 | 1808.01116 |
1808 | 1808.03922_arXiv.txt | One peculiar feature of the solar cycle which is yet to be understood properly is the frequent occurrence of double peaks (also known as the Gnevyshev peaks). Not only the double peaks but also multiple peaks and spikes are often observed in any phase of the cycle. We propose that these peaks and spikes are generated due to fluctuations in the \bl\ process (the poloidal field generation from tilted bipolar magnetic regions). When the polar field develops, large negative fluctuations in the \bl\ process can reduce the net polar field abruptly. As these fluctuations in the polar field are propagated to the new toroidal field, these can promote double peaks in the next solar cycle. When fluctuations in the polar field occur outside the solar maximum, we observe their effects as spikes or dips in the following sunspot cycle. Using an axisymmetric \bl\ dynamo model we first demonstrate this idea. Later, we perform a long simulation by including random scatter in the poloidal field generation process and successfully reproduce the double-peaked solar cycles. These results are robust under reasonable changes in the model parameters as long as the diffusivity is not too larger than $10^{12}$~cm$^2$~s$^{-1}$. Finally, we analyze the observed polar field data to show a close connection between the short-term fluctuations in the polar field and the double peaks/spikes in the next cycle. Thereby, this supports our theoretical idea that the fluctuations in the \bl\ process can be responsible for the double peaks/spikes in the observed solar cycle. | \label{sec:int} Sun's magnetic activity, commonly measured using the sunspot number or sunspot area, oscillates with a period of about 11 years. This is popularly known as the solar cycle or sunspot cycle. Interestingly, every solar cycle is different from the previous ones in terms of the cycle duration and amplitude. Apart from this variation, there exist several short-term variations in the observed solar data \citep{LB89,Baz14,Scott15,Maetal}. One distinct and puzzling observable among these short-term variations is the occurrences of double peaks. It has been observed that during the solar maximum, when sunspot number reaches its maximum value, solar cycle occasionally shows two peaks \citep{FS97,NG10,Geor,Baz14}. These are also known as {\it Gnevyshev peaks} and the gap between these two peaks at the solar maximum is known as {\it Gnevyshev gap} \citep{Gnev67,Gnev77}. Although observed in many earlier cycles, this double-peak feature has received special attention in recent years mainly because of the last three solar cycles being double-peaked; see NASA science report\footnote{\href{https://science.nasa.gov/science-news/science-at-nasa/2013/01mar_twinpeaks}{https://science.nasa.gov/science-news/science-at-nasa/2013/01mar$\_$twinpeaks}}. We note that these double peaks are not the artefacts of insufficient observations but are real features \citep{NG10}. We also note that this feature is not only limited to the sunspot number or area data, but also observed in other proxies of the solar activity e.g., coronal activity \citep{Gnev63,Kane09, Kane10} One could argue that the double peak is a result of the fact that when two hemispheres reach their maxima at two different times, the combined solar activity can have two peaks. By making a careful analysis of the solar data, we shall show that a time difference between the maxima of two hemispheric solar activity may lead to a double peak, however, this happens rarely. In fact, most of the times, the double-peak occurs only in one hemisphere \citep{NG10}. Importantly, the double-peak type spikes are not only observed during solar maximum, they are also seen at any phase of the solar cycle. When the spike appears near a solar maximum, we see it as a double-peak. The double peaks and spikes are possibly the manifestrations of the recently discovered quasiperiodic ``burst" or oscillations with periods of 6 -- 18 months in the solar activity \citep{Scott15}. Using magnetohydrodynamics shallow-water model, \citet{Dik17,Dik18} have shown that the energy exchange among magnetic fields, Rossby waves and differential rotation in the solar tachocline can lead to quasi-peridic nonlinear oscillations, which possibly correspond to the observed burst of solar activity. Also see \citet{Za10,Za18} for studies connecting the Rossby waves in the tachocline with the short-term oscillations. However, there could be a different mechanism of producing double peaks and spikes in the solar cycle. Irregular fluctuations are inherent in the solar dynamo which can appear in any phases of the solar cycle. When strong fluctuations appear near the solar maximum, we may see them as double peaks. In this study, using a dynamo theory we shall identify the source of these fluctuations and explore how these fluctuations can promote double peaks in the solar cycle. Over last two decades, the solar magnetic cycle has been modelled with great details using the \bl\ dynamo models, also named as the flux transport dynamo models \citep{CSD95,Dur95,DG09,Cha10,Kar14a}. In this model, the poloidal field is generated from the decay and dispersal of tilted bipolar magnetic regions (BMRs) near the solar surface. This field is largely transported to the poles through meridional flow. From the surface, the poloidal field is then transported down to the deep convection zone (CZ) through meridional circulation, turbulent diffusion and pumping, where differential rotation stretches this field to produce a toroidal field. The toroidal field then rises up the surface due to magnetic buoyancy and gives tilted BMR. It is believed that the tilt is introduced due to the Coriolis force during the rise of the toroidal flux in the CZ \citep{DC93}. The observed correlation between the surface polar flux and the next cycle strength supports this part of the dynamo model \citep{Das10,KO11,Muno13,Priy14}. The new BMRs again decay and produce poloidal flux which forms the seed for the next cycle. The tilt angle of a BMR is crucial in generating a net poloidal flux as has been realized in the surface observations \citep{Das10}, as well as surface flux transport models \citep{JCS14}, and a 3D (or 2$\times$2D coupled) dynamo model with explicit BMR depositions \citep{HCM17, KM17, LC17}. In observations, we find a considerable scatter of the mean BMR tilt around its systematic variation with the latitude---Joy's law \citep{How91,SK12,MNL14,pavai15,MN16}. This scatter is the primary cause of the variation in the polar field \citep[e.g.,][]{JCS14,HCM17,KM17, Nagy17}. The effect of scatter is very profound when BMRs appear near the equator \citep{Ca13,KM18}. Other effects such as the fluctuations in the net BMR flux, BMR emergence rates, time delay of BMR emergence, meridional circulation speed etc can also introduce additional variation in the polar flux \citep{LC17,KM17,Nagy17}. Ultimately, it is the fluctuations in the \bl\ process which is the primary cause of the variation in the polar field and consequently in the sunspot cycle as has been pointed out earlier by \cite{CD00, CCJ07, CK09}. In this study, we shall show that the fluctuations in the \bl\ process can also occasionally produce short-term fluctuations in the polar field. These fluctuations can be propagated to the toroidal field and therefore can cause double peaks in the next solar cycle. As these fluctuations can occur at any phase of the polar field build up, the fluctuations can appear at any phase of the solar cycle. When they occur outside the solar maxima, we observe them as spikes and dips. We shall explicitly identify the fluctuations in the polar field from the observed data and show that this can be responsible for the double peaks in the solar cycle. \begin{figure*} \centering \includegraphics[scale=0.70]{sunspot_area_double_peak_2.eps} \caption{Temporal variation of the monthly sunspot area (in millionths of a solar hemisphere) for individual hemispheres as well as the combined data. Cycle numbers are also printed on the figure. } \label{fig:obsssn} \end{figure*} | \label{sec:conc} In this study, we have theoretically modelled the double peaks and spikes observed in the solar cycle. We have shown that due to large negative fluctuations in the \bl\ process can abruptly decrease or even reverse the polar field for a short time. This is observed in the form of frequent polar surges of wrong polarity field in surface polar field data available for last four solar cycles (\Sec{sec:supp}). The proxy of polar field data for the previous cycles (for which the polar field measurement is not available; \citet{muno12}) also shows occasional fluctuations. As the polar field is the seed for the next cycle, the fluctuations in the polar field can be propagated in the subsequent solar cycle and they can cause short-term fluctuations in the solar cycle. When the abrupt decrease in the polar field happens in the growing phase of the polar field, we observe a clear double peak in the subsequent sunspot cycle. We have presented this idea by making a clean experiment in which we have artificially flipped the source of the poloidal field ($\alpha$) for six months and as a result a momentary reversed polar field promotes a clear double peak in the next sunspot cycle. Next, we performed a set of long simulations by including random scatter in the $\alpha$ and reproduce many double-peaked solar cycles. To the best of our knowledge, this is the first systematic effort of modeling the double-peaked solar cycle, although three previous attempts exist. First, \citet{Gnev67} argued that the double-peaked solar cycle is caused by two different processes occurring at two different latitude bands. When the time interval between maxima of these two processes is large, the double-peak is seen in the latitude-averaged solar activity. If this is the correct explanation of double peaks, then the question remains what are these two physical process and what determines the time lag between them. Second, \citet{Geor} based on the mechanism of flux transport dynamo showed that the double peaks are the manifestation of two surges of the toroidal field. One surge is generated from the poloidal field that is advected due to meridional circulation all the way on the poles, down to the base of CZ and finally to low latitudes, and the other surge is generated from the poloidal field that is diffused to the base of the CZ directly from the surface. She suggested that when the timescales involved in two surges of toroidal field do not coincide, double peak in the sunspot cycle is observed. However, no modelled double-peak solar cycle was presented. To our knowledge, if this idea applies to the flux transport dynamo model, then we would have observed double peak even without including fluctuations in the \bl\ process. Without introducing fluctuations in $\alpha$, we however have not observed any double peak in any simulations in the parameter ranges we have explored. Therefore, this idea does not work, at least, in our model. Finally, quasi-periodic nonlinear oscillations in the tachocline \citep{Dik17,Dik18}, as discussed in the Introduction, could be a possible cause of the doubled peaked solar cycle, although a detailed model is needed. One strong support of our idea is that the fluctuations in the \bl\ process are identified in the observed polar field as well as in the proxy of the polar field data (as discussed in \Sec{sec:supp}). Recent independent studies \citep{Ca13, JCS15, Mord16, Kit18} also support the polar field fluctuations and they have proposed that these fluctuations are the cause of the variation in the subsequent solar cycle. Another strong support is the fact that the double peaks are observed independently in two hemispheres and in any phase of the solar cycle. If the double peaks are caused by the fluctuations in the dynamo process, then they are expected to appear in any hemisphere and in fact, in any phase of the solar cycle. This is exactly observed in the Sun. Occasional spikes or dips observed in the rising or declining phases of the solar cycle are also caused by the same origin. \begin{appendix} \begin{figure} \centering \includegraphics[scale=0.75]{diff_alflc.eps} \caption{ (a) Same as Figure 3 but obtained from simulations with $250\%$ (top) and $100\%$ fluctuations in $\alpha$. } \label{fig:diffpara} \end{figure} \end{appendix} | 18 | 8 | 1808.03922 |
1808 | 1808.08492_arXiv.txt | The bright transient AT2018cow has been unlike any other known type of transient. Its high brightness, rapid rise and decay and initially nearly featureless spectrum are unprecedented and difficult to explain using models for similar burst sources. We present evidence for faint $\gamma$-ray emission continuing for at least 8 days, and featureless spectra in the ultraviolet bands -- both unusual for eruptive sources. The X-ray variability of the source has a burst-like character. The UV-optical spectrum does not show any CNO line but is well described by a blackbody. We demonstrate that a model invoking the tidal disruption of a $0.1 - 0.4\;M_\odot$ Helium White Dwarf\,(WD) by a $10^5-10^6\;M_\odot$ Black Hole\,(BH) {\revv located in the outskirts of galaxy {\tt Z~137-068}} could provide an explanation for most of the characteristics shown in the multi-wavelength observations. A blackbody-like emission is emitted from an opaque photosphere, formed by the debris of the WD disruption. Broad features showing up in the optical/infrared spectra in the early stage are probably velocity broadened lines produced in a transient high-velocity outward moving cocoon. The asymmetric optical/infrared lines that appeared at a later stage are emission from an atmospheric layer when it detached from thermal equilibrium with the photosphere, which undergoes more rapid cooling. The photosphere shrinks when its temperature drops, and the subsequent infall of the atmosphere produced asymmetric line profiles. Additionally, a non-thermal jet might be present, emitting X-rays in the $10-150$\,keV band. | The transient AT2018cow/ATLAS18qqn/SN2018cow was discovered at an offset of 6\arcsec (1.7\,kpc) from galaxy {\tt Z~137-068} \citep{2018ATel11742....1S} by the ATLAS wide-field survey \citep{2018TNSTR.838....1T} on 2018-06-16 10:35:38 UT (MJD 58285.44141, referred to in this paper as the discovery date \Td) at an AB magnitude $o$ = ${14.74\pm0.10}$\,mag (the $o$-band covers $560-820\;{\rm nm}$,\footnote{http://www.fallingstar.com/specifications}). A previous observation by \cite{2018ATel11738....1F} on MJD~58282.172 (3.3\,d before the discovery date) with the Palomar 48-inch in the $i$-band did not detect a source down to a limiting magnitude of $ i > 19.5$\,mag, while on MJD~58286 (\Td+0.75\,d) $i = 14.32\pm0.01$\,mag, nearly 5 magnitudes brighter - a rapid rise. Maximum light occurred at MJD~58286.9 \citep[\Td+1.46\,d,][]{2018arXiv180705965P}). Spectroscopic follow-up by \citet{2018ATel11732....1P}, and \citet{2018ATel11776....1P} using the SPRAT on the Liverpool Telescope (402-800\,nm, with 2\,nm resolution) on MJD 58287.951 (\Td+1.56\,d) found a smooth spectrum. \citet{2018ATel11736....1J} reported the Ca\,II~H and K absorption lines close to the redshift of the co-located galaxy, proving that the transient was near that galaxy. Spectra taken on the Xinglong 2.16-m Telescope using the BFOSC showed weak broad bumps or dips in the spectrum \cite[][]{2018ATel11740....1X, 2018ATel11753....1I} {\revv which may be interpreted as highly velocity-broadened lines though Perley et at (2018a) considered the features as an absorption trough.} The velocity derived from the broadening of the presumably He emission was $\approx 1.6\times 10^4\,{\rm km\,s}^{-1}$ on \citep[\Td+4.1\,d,][]{2018arXiv180705965P}. Intrinsic optical polarization was measured on days \Td+4.9 and 5.9\,d by \citet{2018ATel11789....1S}. At high energies the transient was detected by the {\em Neil Gehrels Swift Observatory} \citep[hereafter {\em Swift},][]{2004ApJ...611.1005G} XRT \citep{2005SSRv..120..165B} in the $0.3-10$\,keV band \citep{2018arXiv180706369R}, NICER \citep[$0.5-10$\,keV]{2018ATel11773....1M}, NuSTAR \citep[$3-60$\,keV]{2018ATel11775....1M}, and INTEGRAL IBIS/SGRI \citep[$30-100$\,keV]{2018ATel11788....1F}. A search for impulsive emission by Fermi/GBM \citep[$10-1000$\,keV]{2018ATel11793....1D}, Fermi/LAT \citep[$\geq 100$\,MeV]{2018ATel11808....1K}, the INSIGHT HXMT/HE \citep[$80-800$\,keV]{2018ATel11799....1H} and Astrosat CZTI \citep[$20-200$\,keV]{2018ATel11809....1S} was unsuccessful. In the radio a search of pre-outburst data by \cite{2018ATel11744....1D} found $3\sigma$ upper limits of $370\,\mu$Jy at 3\,GHz and $410\,\mu$Jy at 1.4\,GHz. The transient was detected at 90 and 150\,GHz on \Td+4.5\,d with a flux density of $\approx$ 6\,mJy at 90 GHz \citep{2018ATel11749....1D}, at 350\,GHz on day 5.8 with flux density of $30.2\pm1.8$ mJy/beam \citep{2018ATel11781....1S} and with a $5\sigma$ detection at 15.5\,GHz of 0.5 mJy on \Td+6.3\,d \citep{2018ATel11774....1B}. Further detections were reported on days \Td+10 and 11\,d at 9\,GHz and 34\,GHz, and on \Td+12 also at 5.5\,GHZ \citep{2018ATel11795....1D,2018ATel11818....1D}. We will adopt a distance to the transient consistent with it being associated with the nearby galaxy {\tt Z\,137-068}, which has a red-shift $z = 0.01414\pm0.00013$\footnote{NED refcode 2007SDSS6.C...0000} The reddening towards the galaxy is low, $E(B-V) = 0.077$ \citep{2011ApJ...737..103S} and \NH$_{\rm galactic} = 6.57\times 10^{20}\,{\rm cm}^{-2}$ \cite[][]{2013MNRAS.431..394Willingdale}. Adopting cosmological parameters $H_0 = 71.0\;{\rm km~s}^{-1}{\rm Mpc}^{-1}$, $\Omega_{\rm m}=0.27$, $\Omega_v=0.73$, \cite[][]{2011ApJS..192...14J} % the distance is $60\pm 4$\,Mpc\footnote{using NED/IPAC}. During our studies and preparation of this paper three other studies were published in preprint form \citep{2018arXiv180705965P,2018arXiv180706369R,2018arXiv180800969Perley}, and we discuss and use their results in our discussion of the nature of the transient whilst extending their analysis. {\revv As in this paper, \citet{2018arXiv180800969Perley} proposed that the transient could be a TDE and they discussed constraints on the TDE properties using recent models. In two further papers a more general analysis was made in terms of a central engine \citep[][]{2018arXiv181010880Ho,2018arXiv181010720M_Raf}, leaving open the nature of the source. } We discuss our observations, and present a model derived from the observations in terms of the tidal disruption of a He white dwarf by a non-stellar mass black hole, i.e. a TDE-WD event, where the debris forms a photosphere which produces blackbody-like emission in the UV-optical bands and with emission lines formed above the photosphere. Moreover, a rapidly expanding cocoon has become detached from the photosphere and envelops the system initially. It produces very broad emission features attributed to velocity broadened lines, i.e., the bumps seen by \cite{2018arXiv180705965P} {\revv and \cite{2018arXiv180800969Perley}}. Finally, a jet is associated with the event, responsible for the high-energy $\gamma$-ray and X-ray emission. | The {\em Swift} data of AT2018cow, supported by other studies and reports suggests that possibly this was the tidal disruption of a He WD on a relatively small non-stellar mass black hole, resulting in a large, hot, ionised debris cloud emitting a thermal spectrum with weak, broad line emission. A jet may also be associated with the event and emit a non-thermal X-ray spectrum; it is not as luminous as those seen in long GRBs. The X-ray component due to optically thick Compton scattering on the hot debris cloud is negligible, and X-ray emission from a shock caused when a high velocity cocoon encounters the circum-system medium is estimated to be non-detectable. We explain the multi-wavelength temporal behaviours of the source with a WD-TDE model, and the observations give a constraint for the WD mass to be $\approx 0.1 - 0.4\,M_\odot$ and the BH mass to be $\sim 1.3 \times 10^5 - 1.3 \times 10^6\,M_\odot$. The model also predicts the total accreted rate of mass accretion is $\approx 8.5\times 10^{24}{\rm g\,s}^{-1}$, see section~\ref{photosphere}, while we can use $\zeta \sim 3-6$ to also derive the ejected mass loss rate is $\approx 6.3\times 10^{24}\,{\rm g\,s}^{-1}$. The model is consistent with the observed bursty peaks in the X-ray luminosity which are also seen in other TDE as would be expected from an unsteady accretion process. | 18 | 8 | 1808.08492 |
1808 | 1808.00992_arXiv.txt | We present multi-epoch photometry and spectroscopy of a light echo from $\eta$ Carinae's 19th century Great Eruption. This echo shows a steady decline over a decade, sampling the 1850s plateau of the eruption. Spectra show the bulk outflow speed increasing from $\sim$150 km s$^{-1}$ at early times, up to $\sim$600 km s$^{-1}$ in the plateau. Later phases also develop remarkably broad emission wings indicating mass accelerated to more than 10,000 km s$^{-1}$. Together with other clues, this provides direct evidence for an explosive ejection. This is accompanied by a transition from a narrow absorption line spectrum to emission lines, often with broad or asymmetric P Cygni profiles. These changes imply that the pre-1845 luminosity spikes are distinct from the 1850s plateau. The key reason for this change may be that shock interaction with circumstellar material (CSM) dominates the plateau. The spectral evolution of $\eta$ Car closely resembles that of the decade-long eruption of UGC~2773-OT, which had clear signatures of shock interaction. We propose a 2-stage scenario for $\eta$ Car's eruption: (1) a slow outflow in the decades before the eruption, probably driven by binary interaction that produced a dense equatorial outflow, followed by (2) explosive energy injection that drove CSM interaction, powering the plateau and sweeping slower CSM into a fast shell that became the Homunculus. We discuss how this sequence could arise from a stellar merger in a triple system, leaving behind the eccentric binary seen today. This gives a self-consistent scenario that may explain interacting transients across a wide range of initial mass. | The underlying physical mechanism for $\eta$ Car's astounding brightness variation and prodigious mass ejection has been the central mystery associated with this object since John Herschel first drew attention to its erratic flashes and relapses in the mid-19th century \citep{herschel1847}. Because an extremely luminous and massive star appears to have survived this event, it has been discussed as a prototype for a growing and diverse class of non-terminal eruptive transients seen in external galaxies that have luminosities between traditional novae and supernovae (SNe), often referred to as giant eruptions of luminous blue variables (LBVs) or ``SN impostors'' (see \citealt{smith+11,vdm12}). Unlike these extragalactic transients, though, $\eta$ Car is nearby enough that it affords us the opportunity to dissect the properties of its spatially resolved bipolar ``Homunculus'' nebula \citep{gaviola} that was ejected in the event \citep{currie96,morse01,sg98,smith17}. There is a vast literature concerning multiwavelength observational details of the Homunculus (see a recent review by \citealt{smith12}), but the main ingredients to note here are its high ejected mass of about 15 $M_{\odot}$ \citep{smith03,sf07}, its high expansion speeds that also imply a large kinetic energy \citep{smith06}, and that the majority of the mass is concentrated in very thin walls of the mostly hollow bipolar shell \citep{smith06}. Such extreme mass loss suggests that brief eruptions may be important in the evolution of massive stars \citep{so06}. There are also complex ejecta outside the Homunculus, called the Outer Ejecta \citep{thackeray50,walborn76}, which have elevated N abundances \citep{davidson82,davidson86,sm04}. Some of these Outer Ejecta have very fast expansion speeds of 3000-5000 km s$^{-1}$ indicating an origin in the 19th century Great Eruption \citep{smith08}, while the majority are slower and implicate at least two major mass-loss eruptions 300-600 years before the 19th century event \citep{kiminki16}. Were it not for this last point of recurring major eruptions, the one-time merger of a binary system (discussed several times; \citealt{jsg89,iben99,pz16,smith16b}) might seem like a natural explanation for the energy and mass ejection of the Great Eruption. The system is still a close and highly eccentric binary system today \citep{damineli,dcl97}, which requires a triple system initially in any merger model. The orbital parameters of the surviving binary are constrained surpisingly well \citep{madura12}, considering that we have not yet detected the secondary star. The high eccentricity dictates that the two stars come very close to one another at periastron and may even collide or exchange mass during an eruption \citep{soker01,soker04,ks09,smith11}, adding complexity to any binary model. Indeed, brief luminosity spikes in the historical light curve of $\eta$ Car are seen to coincide with times of periastron passage \citep{sf11}. Binary interaction is likely to be very important in the physics of $\eta$~Car's eruption and in other SN impostors, but the details of a working scenario are still a matter of much debate; our main interest in this paper is to characterize the observed properties of the mass loss during the eruption to help guide our understanding of how so much mass left the system in such a short time. Two qualitatively different models have emerged for the driving physics of $\eta$ Car's mass-loss that can be summarized plainly as either a wind or an explosion, although perhaps neither is quite so simple. The eruptive mass loss exhibited by $\eta$ Car occupies a grey area between winds and explosions --- it is either a heavily mass-loaded and energy starved wind, or a relatively weak explosion that only unbinds the outer envelope. This is between the opposite extremes of either a line-driven wind or a core-collapse SN explosion. The more traditional interpretation (traditional in the sense that it has been around longer and is more developed) involves a strong radiative luminosity that pushes the star above the classical Eddington limit and initiates a strong outflow of matter. This is interpreted in the context of the theory for continuum-driven super-Eddington winds \citep{owocki04,owocki17,os16,quataert16,shaviv00,so06,vanmarle08}. In this picture, the outflow is a result of the increased radiative luminosity, and the emitting surface is expected to be a relatively cool pseudo-photosphere in the outflowing wind \citep{davidson87,hd94,dh97,os16}. Consequently, the roughly 20-year duration of the eruption indicates that the star was exceeding its classical electron-scattering Eddington limit by about a factor of 5 the entire time. The other type of scenario for the Great Eruption mass loss is primarily as a hydrodynamic explosion \citep{smith13}. This picture is different from the previous one in the sense that in the former, it is the momentum of escaping photons that accelerates the outflowing material. In an explosion model, the radiation we observe is largely a byproduct of heating by the shock interaction between fast explosively ejected matter that overtakes slower circumstellar material (CSM). This scenario is generally referred to as ``CSM interaction'', and is similar to the standard model for CSM interaction in SNe~IIn, but with a non-terminal and lower-energy explosion. A key point is that in the CSM interaction model, we avoid the puzzle of how a star's envelope can persist in a strongly super-Eddington state for 20 years, because here the emitting material is not bound. Instead, the primary source of luminosity during the plateau of the eruption resides in the shock itself as it plows through the dense CSM. A simple 1-D model shows that one can account for the observed decade-long plateau of $\eta$ Car's eruption while also matching the present-day observed properties of the massive shell nebula \citep{smith13}. Both types of models have the shortcoming that they lack a clear explanation for the ultimate source of the energy. In the super-Eddington wind model, the star's radiative luminosity is assumed to increase temporarily and then decrease as dictated by the observed light curve, and is the primary agent driving the mass loss. The source of this extra luminosity is unknown. In the CSM interaction scenario, the power source is the relatively sudden deposition of energy deep inside the star for some unknown reason, and the radiation is a byproduct. That energy source may be from binary orbital energy, as mentioned above, or from nuclear burning instabilities akin to those that have been suggested in eruptive SN progenitors \citep{qs12,sq14,sa14,woosley17}. The wind model may have problems accounting for several aspects of the observed nebula (see below), whereas the CSM interaction requires us to invoke some slow pre-existing CSM for the fast ejecta to collide with. Both models can potentially account for the bipolar shape if we invoke either rapid rotation \citep{do02,og97,st07} or equatorial CSM \citep{frank95,langer99}. A key ongoing challenge is to understand how a merger or some other physical model might provide the required energy on the appropriate timescale, in a way that yields the observed results. Primary sources of empirical information that have guided these two mass-loss scenarios involve the historical visible-wavelength light curve \citep{sf11} and the physical parameters of the remnant of the explosion - the ``Homunculus Nebula'' and its surrounding debris, which can be studied in exhaustive detail. The historical record provides the observed fact that the object's luminosity did exceed the Eddington limit for the star's presumed mass for more than a decade, motivating the wind model. On the other hand, continued study of the present-day nebula gives several clues that together point strongly toward a hydrodynamic explosion. These are: (1) the large mass of 12-20 $M_{\odot}$ in the Homunculus combined with its fast expansion speeds gives a large kinetic energy of order 10$^{50}$ ergs \citep{smith03,smith06}, which exceeds the radiative energy budget of $\sim$10$^{49}$ ergs. This low ratio of luminious to kinetic energy is more characteristic of radiation from expanding and cooling SN envelopes than of stellar winds (although winds with extreme photon tiring might also achieve this; \citealt{owocki04}). (2) Observations of material outside the Homunculus indicate very high expansion speeds reaching 5000 km s$^{-1}$, which is easier to explain with shock acceleration \citep{smith08}. This outer fast material also raises the total kinetic energy budget of the event even more. (3) Most of the mass in the Homunculus resides in the extremely thin walls of the bipolar lobes, which points to compression in a radiative shock \citep{smith06,smith13}. Other details of the structure in the nebula also point toward a shock rather than a steady wind (see discussion in \citealt{smith13}). \begin{figure*} \includegraphics[width=6.3in]{fig1.eps} \caption{The environment around the EC2 light echo. (a) Large field of view 3-color composite image at visible wavelengths, with [O~{\sc iii}] $\lambda$5007 in blue, H$\alpha$ in green, and [S~{\sc ii}] $\lambda\lambda$6717,6731 in red. The images were obtained in 2003 with the MOSAIC camera on the CTIO 4m telescope \citep{smith03a}. An arrow points toward the location of $\eta$ Car itself, off the top of the image. (b) Same field of view as (a), but showing IR images obtained with the IRAC camera on {\it Spitzer}, in Bands 1 (blue), 2 (green), and 3 (red) \citep{smith10spitz}. (c) Same image and color scheme as (a) but zoomed in on a smaller field around EC2. (d) same field as (c) but showing only the $i$-band CTIO4m/MOSAIC image in grayscale, also obtained in 2003. The red box shows the location of our most commonly used IMACS slit aperture. (e) Same field as (c) and (d) but in the IR, with the same {\it Spitzer} images and color scheme as (b). (f) A sketch showing the global geometry involved. An Earth-based observer is to the left, looking through the cold clouds on the near side of the nebula, which appear dark in optical images and glow in PAH emission in the IR. They are seen in silhouette against the bright screen of H~{\sc ii} region emission that fills the interior of the nebula. Cold clouds on the far side are also seen in PAH emission in the IR, but cannot be seen in silhouette at visible wavelengths, because they are behind the line emission. This is the case for the EC2 cloud, as well as the EC1 group of echoes discussed in our previous papers \citep{rest12,prieto14}. Dashed curves denote the rough locations of the light echo parabolas corresponding to the light curve peaks in the 1830s-1840s, as well as the 1850s plateau.} \label{fig:img} \end{figure*} Recent studies have added significantly to the already tremendous repository of observational information about $\eta$~Car. Namely, the discovery\footnote{Historical aside: Light echoes from $\eta$ Car have been reported previously. \citet{walborn+liller} discovered that clouds in the Keyhole nebula were reflecting the peculiar spectrum of $\eta$ Car. \citet{elliott79} obtained spectra of these features as well, but interpreted the relatively broad line wings as evidence that the Keyhole was a supernova remnant. Additional spectra of these reflected echoes were also obtained and interpreted as echoes with minor spectral variability over time \citep{lm84,lm86,boumis98}. However, these were not strongly variable echoes of the Great Eruption, but rather, reflected light from the star in its modern post-eruption state. Interestingly, though, \citet{walborn+liller} pointed out that if these nearby clouds are scattering light from the star today, then this may explain why drawings of the Keyhole by \citet{herschel1847} look different from its appearance today \citep{gratton63}. If so, then John Herschel was arguably the first to record light echoes from the Great Eruption.} of light echoes from $\eta$~Car's Great Eruption \citep{rest12} and the evolution of light echo brightness and spectra over time \citep{prieto14} allow us to probe deeper, providing a unique and crucial link between the historical brightness record, the kinematics and structure of the nebula, and potential similarity to modern extragalactic analogs. \citet{rest12} showed that light echo spectra near the peak of the eruption showed a characteristic temperature that was significantly cooler (G-type) than published expectations for pseudo-photospheres of LBV eruptions \citep{davidson87} and observed spectra of LBV eruptions \citep{hd94}. This sparked a debate. \citet{dh12} argued that if one were to extrapolate the published pseudo-photosphere models of \citet{davidson87} in the appropriate way, the wind photosphere might be consistent with temperatures as cool as observed. \citet{os16} noted inconsistencies in the analysis by \cite{davidson87}, but also showed that by properly accounting for opacities in radiative equilibrium, wind mass-loss rates of the order of that inferred for $\eta$ Car's Great Eruption are compatible with temperatures around 5000~K after all. In any case, spectroscopy of the subsequent fading of that same echo \citep{prieto14} showed that the temperature became cooler still, dropping to 4000-4500~K and forming molecular bands commonly seen in extremely cool carbon stars. This behavior with time contradicts simple expectations for a pseudo-photosphere model, where the apparent temperature should increase as the photosphere recedes to deeper wind layers \citep{davidson87}. The development of such cool temperatures and molecular features in the spectra presented by \citet{prieto14} correspond to one of the brief luminosity spikes (e.g., 1843, 1838, etc.) observed in the early stages of the eruption \citep{sf11}. As described in this paper, the temporal evolution of echo spectra shows clear disagreement with a wind pseudo-photosphere interpretation of the eruption, but gives unambiguous evidence of an explosive component to the mass loss. There are a number of important implications for the nature of the Great Eruption and the evolutionary history of the $\eta$ Car system. | \subsection{Overview} A key result from the observed spectral evolution of the EC2 echo combined with $\eta$~Carinae's other light echoes is that the basic character of the spectrum changes considerably during the decade-long 19th century eruption. These changes are indicative of a line of sight that views $\eta$ Carinae's Great Eruption from a vantage point at low latitudes near the equator, and might not be indicative of all viewing angles. While there are a number of complicated changes that occur (including rapid changes during fading from bright peaks), the underlying evolution is gradual and can be summarized as basically a 2 stage event: 1) A preperatory wind phase in the lead up to the Great Eruption during the early 1840s (and perhaps for decades before that), which may be strongly influenced by binary interaction, and 2) an explosive event with fast outflow speeds and sustained high luminosity. These two stages are annotated in Figure~\ref{fig:etaLC}. {\it Stage 1 (1845 and preceding decade or two):} This initial phase has slower outflow velocity (150$-$200 km s$^{-1}$) in the equator, and the overall appearance of the spectrum is dominated by narrow absorption lines with little or no emission at luminosity peaks, except perhaps narrow H$\alpha$ wind emission that may be contaminted by unresolved nebular emission \citep{rest12}. It has somewhat cooler apparent temperatures of 5000$-$5500~K near peaks \citep{rest12}, and cooler temperatures of $\sim$4,000 K and evidence of molecule formation in the fading after peaks \citep{prieto14}, with increasing emission-line strength as it fades. From this viewing direction, the bulk of outflowing material (traced by the absorption trough) is moving at about 200 km s$^{-1}$ at times that likely correspond to the 1830s to the early 1840s (although there may be some faster material at lower density indicated by the absorption line wings extending out to several hundred km s$^{-1}$). There may be faster material at other latitudes as well. {\it Stage 2 (late 1840s - 1850s plateau):} While the overall continuum shape is similar to earlier phases (a slight increase in the apparent temperature to 6000~K), there are distinct changes in line properties that trace the outflowing material. The 1850s plateau as viewed in EC2 spectra developes broader line widths and increasing emission strength in the narrow components indicating an increase in the bulk outflow speed to 600 km s$^{-1}$, which is much faster than the 50$-$200 km s$^{-1}$ speed of the Homunculus at low latitudes near the equator \citep{smith06}. This time period also shows much stronger emission lines in general, a decrease of line blanketing absorption strength at shorter wavelengths, P Cygni profiles instead of pure emission in many lines, and signs of higher excitation (including possible He~{\sc i} emission). Most remarkably, epochs corresponding to the mid-1850s show the appearance and strengthening of very broad emission wings from $-$10,000 to $+$20,000 km s$^{-1}$. Narrow absorption components at $-$150 km s$^{-1}$ persist from Stage 1, suggesting that this is slower, previously ejected material along the line of sight at a somewhat larger radius. {\it A key point is that at least three very different expansion speeds are seen simultaneously in Stage 2.} A major implication for the nature of $\eta$ Car's eruption is that the observed changes show that a steady state wind is clearly {\it not} a good approximation. The bulk outflow speed increases dramatically over a time period of a few years, and the fastest material in Stage 2 is two orders of magnitude faster than in Stage 1 (and far exceeds the escape speed). Line strengths increase while the apparent continuum temperature and luminosity stay relatively constant; this probably signifies an increase in density and strong departures from LTE that may be indicative of shock excitation. The increasing velocity along the same line of sight also has important physical implications. It requires that fast ejecta follows after much slower material in the same direction, making it inevitable that fast material will catch slow material and will collide in a strong shock. The fact that the outflow speed changes with time (increasing from slow to fast) is therefore direct evidence supporting earlier claims, based on multiple lines of circumstantial evidence, of a strong CSM interaction component that helps power the visible luminosity of the event \citep{smith13,smith08,smith03}. Another interesting outcome, noted in Figure~\ref{fig:etaLC}, is that the ejection date of the Homunculus from its measured kinematic age is solidly in between Stage 1 and Stage 2. From a recent study of available archival {\it HST} imaging over more than a decade, the proper motion expansion of the Homunculus gives a fairly precise date of origin for the Homunculus (extrapolating from linear motion observed today) of 1847.1 $\pm$0.8 yr \citep{smith17}. Of course, if there was strong CSM interaction that accelerated slow pre-shock material and decelerated the fast ejecta, as in a Type~IIn supernova, the true ejection date might be slightly different (most likely a short time after this). Alternatively, if the material was ejected over a longer period (for example, a more gradually increasing ouflow speed over many years as opposed to an instantaneous pulse), then this is a mass-weighted average ejection date. In any case, it is remarkable that the Homunculus date of origin lies in between the slow material and fast material seen in our echo spectra. This strongly supports a picture wherein fast ejecta swept up and shocked slower material, making a radiative shock that cooled rapidly to form the thin walls of the Homunculus \citep{smith13}. We consider the physical interpretation of the 2 Stage event in more detail below. \subsection{Comparison with UGC~2773-OT} The relatively nearby LBV-like transient in the dwarf galaxy UGC~2773, named UGC~2773-OT \citep{smith10,foley11,smith16}, has been compared with $\eta$ Car's Great Eruption before. \citet{rest12} showed that early spectra of UGC~2773-OT at peak luminosity (soon after discovery) were quite similar to light echo spectra of $\eta$ Car that correspond to early peaks in the Great Eruption light curve (EC1). \citet{smith16} showed that after several years had passed, UGC~2773-OT faded very slowly, sustaining a high luminosity for a decade, very similar to the slow light curve evolution of $\eta$~Car during its 1850s plateau. \citet{smith16} presented a series of spectra of UGC~2773-OT that document its spectral evolution over several years. Light echo spectroscopy of EC2 spanning several years now allows us to extend the comparison. The evolution of EC2 spectra shows remarkable similarity to the spectral evolution of UGC~2773-OT, further supporting the case that they are close analogs. As seen in Figures~\ref{fig:spec} and \ref{fig:halpha} (a more densely sampled series of UGC~2773-OT spectra can be seen in \citealt{smith16}), the evolution of the overall low-resolution spectrum and of velocities and excitation is very similar between the two objects. They have similar continuum temperatures, both increasing a small amount as they evolve. They show many of the same lines, which show mostly narrow absorption at early times, transitioning into stronger emission at later times. Both show narrow emission from [Ca~{\sc ii}] $\lambda\lambda$7291,7324, which is seen in a subset of SN impostors. Similarities in H$\alpha$ are particularly remarkable: In both $\eta$ Car and UGC~2773-OT, the H$\alpha$ emission line profile starts out as a weak and narrow P Cygni profile but then gets stronger and broader (from around 100-200 km s$^{-1}$ initially up to 600-1000 km s$^{-1}$ at later times), with a very similar asymmetric emission line profile (Figure~\ref{fig:halpha}). The spectral similarity is interesting because these two also have similar light curves. If they are close analogs, perhaps UGC~2773-OT can help us fill-in gaps in our knowledge due to limitations of light echo spectra. The most significant limitations of the echo spectra are relatively low signal to noise because they are faint, a lack of access to other wavelengths, and contamination from narrow nebular line emission from the Carina Nebula. Data for UGC~2773-OT do not present the same limitations. For example, from light echo spectoscopy alone, it is ambiguous if the residual narrow He~{\sc i} emission (Figure~\ref{fig:extractHa}) comes from $\eta$ Car itself, or if it is narrow nebular He~{\sc i} emission arising on the globule's surface that is exposed to radiation from O-type stars in the Carina Nebula. The He~{\sc i} $\lambda$6678 line shows an asymmetric and possibly P Cygni profile, but it is narrow and weak, so this profile could potentially arise from background subtraction. It is therefore extremely interesting that UGC~2773-OT shows very similar narrow He~{\sc i} emission, which is absent at first and then grows with time while the eruption has almost constant luminosity and temperature (it is not H~{\sc ii} region contamination). In UGC~2773-OT, the He~{\sc i} emission is also much narrower than the H$\alpha$ line, qualitatively very similar to $\eta$ Car. In the case of UGC~2773-OT, the narrow He~{\sc i} must be intrinsic to the object, arising from slow pre-shock gas that is photoionized by a shock (it must be a shock, since radiative excitation from a 6000 K photosphere would not produce strong He~{\sc i} emission). This gives a possible indication that the narrow He~{\sc i} emission in $\eta$ Car arises in a similar fashion. The very strong contamination from H$\alpha$ in the Carina Nebula (intrinsically narrow and unresolved in our spectra) makes it impossible to say anything conclusive about narrow H$\alpha$ emission in echoes, and hence about narrow H$\alpha$ emission from pre-shock gas in the eruption. UGC~2773-OT does not have the same ambiguity, and narrow emission is present. Some of this arises from a surrounding H~{\sc ii} region, but high-resolution echelle spectra show resolved widths of $\sim$50 km s$^{-1}$, indicating that expanding CSM is partly responsible for this narrow emission \citep{smith10}. Perhaps UGC~2773-OT hosts extended expanding nebulosity, similar to the Outer Ejecta of $\eta$ Car \citep{kiminki16,mehner16}. The light echoes of $\eta$~Car are faint and require a significant allocation of time on large telescopes in the southern hemisphere; for practical reasons, this has limited our cadence to roughly 1--2 observations per year. Moreover, different echoes trace different epochs in the Great Erupton, adding uncertainty to the exact time evolution. The time sampling of spectra for UGC~2773-OT is better and more clearly understood. UGC~2773-OT exhibits a quite gradual transition over a few years from Stage 1 to Stage 2. This is an important clue; the changes in the spectrum (notably the presence of broader and stronger emission lines) was not a sudden change on a dynamical timescale, but rather, the transition happened gradually. Either the star is changing slowly, or optical depth effects govern the slow emergence of radiation from faster material deeper in the expanding envelope. A limitation of the light echo spectra is that they become very noisy in the blue part of the spectrum, due to a combination of ISM reddening and detector/grating efficiency. This makes it difficult to study the spectrum at blue wavelengths, while the UV range of the spectrum is not available to us. The extreme faintness of the echoes combined with bright background emission from the surrounding star-forming region make it difficult to study the echo spectra in the IR. For all these similarites, though, UGC~2773-OT is not an identical twin of $\eta$ Car's eruption. A few interesting differences between the spectral evolution of these two objects are: 1) UGC~2773-OT has no brief spikes in its light curve, which in $\eta$ Car have been attributed to periaston interactions as noted above. Evidently UGC~2773-OT did not experience these sorts of interactions with a wide companion, but suffered a similar decade-long outburst anyway. This provides another indirect suggestion that interactions with this wide companion were not critical in powering $\eta$~Car's event. 2) UGC~2773-OT has no 6500~\AA \ (presumably Fe~{\sc ii}) line exhibiting a P~Cygni profile that develops at late times (Figure~\ref{fig:fe2}). The significance of this difference is unclear, since UGC~2773-OT does show other Fe~{\sc ii} lines with similar P Cyg profiles. Since this line is absent at some epochs for $\eta$ Car, but then appears at epochs when the broad emission is seen, this line deserves a closer look. 3) The Ca~{\sc ii} IR triplet lines grow in strength at late times much more than seen in $\eta$~Car's echo spectra. This is probably an optical depth effect at late times, and it will be interesting to see how the echo spectra continue to develop. 4) UGC~2773-OT does not show the absurdly broad emission wings of H$\alpha$ that are seen in $\eta$ Car. We do not yet know if this is a viewing angle effect, a timing issue (the broad lines appear late in $\eta$ Car's spectra, and begin to fade after 1-2 yr), an optical depth effect, or a fundamental difference in shock ejection in the two events. Examining echoes that view $\eta$ Car from other latitudes will help clarify any angle dependence of the fast ejecta. \subsection{Transitioning between two stages of the eruption: Winds vs. Explosions} Continued monitoring of the spectral evolution of $\eta$~Car's light echoes has demonstrated clearly that observed spectra from later in the eruption show fundamental differences compared to earlier spectra. Among the major differences are a faster bulk outflow speed, the appearance of extremely fast ejecta, stronger emission lines, and weaker absorption. This transition occurs well after the eruption was already underway, and coincides roughly with the ejection date imprinted on the expanding nebula as measured from its proper motion expansion (Fig.~\ref{fig:etaLC}). This points to a dramatic change in the physical state of the star's envelope after the peak of the eruption in 1845. A critical question for interpreting $\eta$ Car's eruption is what caused this transition from Stage 1 to Stage 2. Before examining that, we first consider the basic transition of outflow properties in the two stages as deduced from available light echo spectra. The simple fact that a major transition in physical state occured mid-way through the eruption is an important physical clue. This change clearly indicates that a single physical mechanism does not govern the mass loss throughout the whole eruption. For example, the traditional picture of the star increasing its luminosity above the Eddington limit and driving a strong wind cannot adequately explain both the precursor luminosity spikes (in 1838 and 1843) and also the long 1850s plateau, since these two phases clearly have fundamentally different outflow properties at roughly the same luminosity. Similarly, the observed transition occurring after a time when the eruption was already underway indicates that the origin of the eruption was not as simple as a single, instantaneous deposition of energy that blasted off the star's envelope on a dynamical timescale in 1847. Instead, there was a long preparation phase (years to decades) leading up to the peak of the eruption, perhaps due to a building instability inside the envelope, or perhaps due to increasing intensity of binary interaction as the orbital parameters changed before a merger (see below). We must therefore seek a physical explanation for the eruption that naturally accounts for both of these observed phases and the transition from one physical regime to the next in the correct order. Regardless of the underlying physical trigger of the eruption, the plain fact that slow outflow velocities observed at earlier epochs were followed by faster outflow velocities at later epochs along the same direction (i.e., echoes probing essentially the same line of sight) necessarily {\it requires that CSM interaction play an important role in the event}. Fast material must overtake the slower previously ejected material and shock. In doing so, some kinetic energy of the fast material is thermalized and converted to luminosity. \subsubsection{Stage 2 as an explosion} The observed spectra in Stage 2 show at least three different outflow speeds simultaneously (slow $\sim$200 km s$^{-1}$; intermediate 500-1000 km s$^{-1}$; and very fast 10,000-20,000 km s$^{-1}$). This fact is not easily explained by a steady wind. It is, however, a commonly observed trait of standard SNe~IIn powered by an explosion crashing into dense CSM. Thus, the two-stage empirical description of the eruption from light echoes outlined above is remarkably compatible with the CSM interaction scenario proposed earlier by \citet{smith13}. This model envisioned Stage 1 as a relatively slow (200 km s$^{-1}$) super Eddington wind, which is quite similar to the value observed in light echoes (note that the absorption speed in light echoes is seen from the equator; outflow speeds are probably higher at other latitudes, and some of this could cause the more extended absorption wings). This was followed in Stage 2 by an explosive energy injection of roughly 10$^{50}$ erg, and \citet{smith13} showed that the ensuing CSM interaction luminosity could in principle account for the 1850s plateau in the historical light curve of the Great Eruption. The changes seen in light echo spectra therefore provide direct confirmation of a CSM interaction model like that of \citet{smith13}. As noted by \citet{smith13}, these numbers are somewhat malleable, though, and can be adjusted depending on the desired level of complexity. The Homunculus nebula of course dictates that there must be a range of speeds and densities at various latitudes \citep{smith06}, while the \citet{smith13} model was a simple 1-D estimate. Even in 1-D, one can adjust the time dependence of physical parameters to achieve a similar end result that fits the light curve. In particular, different choices for the relative amount of mass and speeds of the outflows in Stage 1 and Stage 2 can lead to similar CSM interaction luminosities. Light echoes combined with observations of the present-day nebulosity help constrain possible values. For example, light echoes provide strong evidence that the outflow speed in Stage 1 was indeed roughly 150 km s$^{-1}$. A signifficant difference compared to the values adopted by \citet{smith13}, however, is that the speed of the fast material appears to be even higher than assumed. In a CSM interaction scenario, the dominant observed outflow speed of $\sim$600 km s$^{-1}$ arises from the cold dense shell, where fast ejecta and shocked CSM pile up in a thin cooled layer. This also corresponds well to the final coasting velocity of the Homunculus nebula \citep{smith06}. The very high speeds seen in light echoes give more freedom in this type of model, since we don't know {\it a priori} what fraction of the Stage 2 mass loss is contained in the fastest outflowing material. If Stage 2 is characterized by outflow speeds of 10,000 - 20,000 km s$^{-1}$, then it is easy to have a situation where most of the mass is ejected at slow speeds in Stage 1, while most of the kinetic energy and momentum is supplied by a fast wind with much lower mass-loss rate. The maximum mass in the fastest ejecta can be derived by assuming that most of the mass is supplied in the slow Stage 1 wind, whereas the fast ejecta in Stage 2 provide essentially all the kinetic energy that powered the event. Under this limiting assumption, the total mass contained in the fastest ejecta can be expressed as \begin{displaymath} M_{fast} \le 0.1 M_{\odot} \times \frac{E_{50}}{V_4^2} \end{displaymath} \noindent where $E_{50}$ is the total energy of the event in units of 10$^{50}$ erg, and $V_4$ is the speed of the fast ejecta or wind in units of 10$^4$ km s$^{-1}$. This is similar to but slightly smaller than the mass of the extremely fast Outer Ejecta estimated by \citet{smith08}. Thus, the fast material seen in the broad wings in light echoes must represent a small fraction of the total mass budget of the Great Eruption, unless the total energy of the event is much higher than generally believed. The radiated energy inferred from the historical light curve (with zero bolometric correction) is only about 2$\times$10$^{49}$ erg \citep{smith+11}, the kinetic energy of the Homunculus is almost 10$^{50}$ erg \citep{smith03}, and the fast Outer Ejecta \citep{smith08} are thought to contain a similar amount of kinetic energy as the Homunculus. Significantly increasing the total energy would require either hotter temperatures (and hence, a larger bolometric correction) during the event, which seems incompatible with the relatively cool apparent temperatures seen in light echoes \citep{rest12,prieto14}, or instead, a much larger amount of invisible mass in the fast ejecta outside the Homunculus. There must be some additional mass ejected in Stage 2 at lower speeds, however, in order to account for the final momentum of the Homunculus. For example, with 10 $M_{\odot}$ ejected in the slow Stage 1 wind at 150 km s$^{-1}$ and only 0.1 M$_{\odot}$ ejected in Stage 2 at 10$^4$ km s$^{-1}$, the final coasting speed of the Homunculus would only be $\sim$250 km s$^{-1}$. In any case, the significant changes seen in light echoes are highly constraining for any model of the event. While Stage 1 can be explained quite well with existing models of a quasi-steady, continuum-driven super-Eddington wind \citep{owocki04,vanmarle08,vanmarle09,os16}, the transition to Stage 2 requires time-dependent energy input well beyond this. \subsubsection{Stage 2 as a time-dependent wind} Rather than a slow wind followed by a single hydrodynamic explosion, as in a Type IIn supernova, the observed properties of the Great Eruption might be accounted for with a more complicated, time-dependent wind with slow outflow transitioning to a faster wind. This has been predicted in models that include super-Eddington energy deposition below the surface of a massive star \citep{quataert16}. In this sort of model, the observed differences between an explosion and a wind become less obvious. Much of the radiated energy arises from internal shocks in the wind, and the photosphere can reside in the compressed post-shock zone itself \citep{quataert16,owocki17}, which is qualitatively similar to the prediction of the explosion plus CSM interaction model. If one envisions a fast wind rather than a hydrodynamic explosion, then light echoes require that the fast wind must be able to achieve extremely high speeds of 10$^{4}$ km s$^{-1}$ and a mass-loss rate (spread over the 3-4 year duration of the broad wings seen in light echo spectra) of $\dot{M} \simeq 3 \times 10^{-2}$ M$_{\odot}$ yr$^{-1}$. The corresponding mechanical luminosity of such a wind is at least 2$\times$10$^8$ L$_{\odot}$, or about $\Gamma$=40. \subsubsection{Grey area} The central question of whether Stage 2 is better described as a wind or an explosion ventures into muddy waters. Limiting cases of these two are well defined. An explosion will result if energy deposition occurs faster than the dynamical timescale with a total energy well exceeding the gravitational binding energy of layers above. A strong wind will result when energy is carried efficiently through the star's envelope by convection, and photon diffusion at the surface can power a steady radiatively driven wind. Observational estimates, however, clearly place the Great Eruption of $\eta$~Car precariously between these two extremes. This is in an interesting regime similar to that discussed by \citet{rm17}, where the energy that eventually emerges from the star's surface as kinetic energy or radiation must be transported through the envelope by acoustic waves that steepen to shocks, and may dissipate their energy in the outer envelope \citep[see also][]{piro11}. Light echoes paint a picture where $\eta$~Car was relatively stable at first, but then underwent a transition past some critical point where the outflow changed dramatically. One can imagine a physical scenario where the rate of energy deposition grows with time to exceed a critical limit, or where there is a sudden change in the deposition rate or depth in the envelope. Regardless of specific mechanism, it is tempting to ascribe the two stages to (1) an early phase that is trans-Eddington, where radiative damping or weak shock dissipation is sufficient to inhibit strong shock formation, depositing energy in the outer envelope at a rate that can can be carried away by radiation, thus driving a strong super-Eddington wind, and (2) later phases that exceed the critical wave luminosity, where radiative damping and shock dissipation are no longer able to suppress strong shock formation, and shocks grow in strength, removing mass from the surface of the star hydrodynamically. Thus, $\eta$~Car is probably an object where {\it we have directly witnessed the transition from a quasi-steady wind to explosive mass loss}. This motivates continued theoretical investigation of time-dependent energy deposition in massive star envelopes, in order to ultimately reconcile the observed physical parameters with the central engine that caused the outburst. The energy deposition required in this picture could in principle arise from one of multiple possible physical causes, including: inspiral of a companion during a stellar merger, runaway instability in shell burning, wave driving, or the pulsational pair instability \citep{qs12,sq14,sa14,smith+11,woosley17}. A number of observational facts (but perhaps most importantly the axisymmetry of the Homunculus nebula and the decade-long duration of the ramping-up of luminosity in Stage 1) suggest that a binary merger event is a plausible explanation for the Great Eruption, although not all the others are necessarily implausible. A general model for such a merger with CSM interaction is discussed in Section 4.6, followed in Section 4.7 by a discussion of details pertaining specifically to $\eta$ Car. \subsection{Origin of the Fast outflow?} The biggest surprise in our study of $\eta$~Car's light echoes has been the discovery of extremely fast ejecta indicated by the broad H$\alpha$ line wings extending from $-$10,000 to $+$20,000 km s$^{-1}$. This is discussed more in a companion paper \citep{smith+18}, so the reader is referred to that paper for additional observational details. So far, no model proposed for $\eta$~Car or stellar mergers in general predicts such extreme outflow velocities that produce a SN-like blast wave. A speed of 20,000 km s$^{-1}$ is much faster than any escape speed envisioned in the $\eta$ Car system; it is even 10 times faster than the wind of a WR star. Super-Eddington winds that drive strong mass loss are generally expected to have relatively slow outflow speeds comparable to the escape speed at large radii where the wind originates \citep{owocki04,quataert16,vanmarle08,vanmarle09}. Binary mass loss from L2 predicts relatively slow outflow speeds no more than several hundred km s$^{-1}$, even at very high luminosities \citep{pejcha16a}. Bipolar jets driven by accretion onto a companion \citep{ks09} would be expected to be no more than a few times the surface escape speed of the accreting star, and one would not expect to see a fast outflow from bipolar jets in the equator (especially if those same bipolar jets are invoked to shape the Homunculus). Clearly, the fast speeds place important fundamental constraints on the nature of the event. If the outflow traces a steady flow, the extremely high speed would imply an outflow from a massive compact object, such as a jet from an accreting neutron star or black hole. These are the only objects with such high escape speeds. The presence of such a companion in the 1850s might be reconciled with a lack of any such companion seen in data at the present epoch if the compact object is the thing that merged with a companion star in the Great Eruption, making the present-day primary star a Thorne-$\dot{\rm Z}$ytkow object (T$\dot{\rm Z}$O). This would open a direction of inquiry far beyond the scope of this paper, and would be a departure from current ideas about $\eta$~Car --- but it is interesting to note that a T$\dot{\rm Z}$O might be consistent with reports of unusual lines seen in one particular equatorial region in the nebula that shows very strong emission from species such as Sr, Y, and Zr, plus Sc, Ti, V, Cr, Mn, Fe, Co, Ni, etc \citep{hartman04,bautista06,bautista09}. This is the so-called ``Strontium Filament'' in the equatorial ejecta. A bipolar jet from a compact object is not, however, a very satisfying explanation for the origin of the fast material because of geometrical reasons. Namely, such a jet is expected to be highly collimated and bipolar, as in the case of SS~433 \citep{paragia99}. Yet, the light echo in which the very fast material is seen views $\eta$ Car {\it from near the equator} of the Homunculus. Such a jet could therefore have had little impact on shaping the bipolar Homunculus nebula (which has an orthogonal orientation), and we would need to invoke some other explanation for the very fast ejecta with speeds of $\sim$5,000 km s$^{-1}$ in the polar regions of the Outer Ejecta seen at the present epoch \citep{smith08}. Instead of relatively steady mechanisms including jets, the extremely fast ejecta seen in light echoes of $\eta$~Car more naturally point to a wide-angle explosive outflow driven by strong shock acceleration, which is not necessarily expected to be close to any escape speed in a system (providing that it exceeds the escape speed) because it is determined mainly by the energy in the shock and the density gradient where the shock acceleration occurs. What could be the origin of such a shock? Energy deposition deep in the stellar envelope, by whatever mechanism, will be transported outward by waves and will steepen to a shock if the energy deposition rate exceeds the steady stellar luminosity \citep{rm17}. When such strong shocks exit the star, they will accelerate a small amount of mass to very high speeds. So then the question is shifted to what the source of this energy deposition would be. One mechanism that can most likely be ruled out here is energy deposition by wave driving from core convection in late evolutionary phases \citep{qs12,sq14,quataert16,fuller17}. The reason it doesn't work in the particular case of $\eta$~Car is because of timescales. While this is an efficient way to suddenly dump energy into the stellar envelope, it is only expected to be significant in the latest Ne and O core burning phases \citep{qs12}, which last just a couple years. It has been 170 years since the Homunculus was ejected, and $\eta$ Car has apparently not yet undergone core collapse. A mechanism that is harder to rule out, and which may indeed be a plausible explanation, is sudden energy deposition via the pulsational pair instability (PPI). This is explored more in Section 4.5. Another possibility is that the Great Eruption of $\eta$ Car was a stellar merger event, which is an attractive hypothesis for several reasons, as noted earlier and discussed further in Section 4.6. If the lead-up to the eruption corresponds to the inspiral and L2 mass loss phase, and the decades-long eruption is the common envelope ejection phase with CSM interaction, then what specifically launches a small fraction of the total mass to extremely high speeds while most of the mass is ejected at only 600 km s$^{-1}$? How does a stellar merger eject material at speeds much faster than the escape velocity of either star? This is a central question for models of massive star mergers that remain unanswered. One speculative possibility is that the energy deposition arises from unsteady nuclear burning as fresh fuel is mixed to deeper shell burning layers, leading to an outburst \citep{sa14}. Explosive common envelope ejection was discussed in a very different scenario (merger of a post-He core burning star with a low-mass companion) by \citet{ppod10}, but perhaps something similar might happen when the cores of the two stars merge, mixing unburned fuel into the core. Perhaps such a mechanism could lead to explosive energy deposition that travels outward through the star and steepens to a shock, mimicking a scenario like a PPI eruption. While it is still uncertain how this would work, observations seem to require that some violent process like this must be an ingredient of any merger model for $\eta$~Car in order to explain the extremely fast ejecta seen in light echoes. An even more speculative origin for the fast ejecta involves an external mechanism. Namely, any binary merger model for the Great Eruption of $\eta$ Car must involve a hierarchical triple system, since a wide and eccentric binary remains today. The current companion on its eccentric orbit would have plunged into the bloated star or common envelope during the event \citep{smith11}, and perhaps that violent collision led to the ejection of a small amount of material to high speeds in certain directions. Whether this could provide enough energy to power the fast ejecta seen in light echoes is uncertain. This is discussed more in our companion paper on the fast ejecta \citep{smith+18}. A related point has to do with the influence of the shock breaking out of the star. Broad emission lines in light echoes reveal that a small amount of material appears to have been accelerated to very high speeds by a shock. If a strong shock with $\sim$10$^{50}$ ergs breaks out of the surface of the star or the outer boundary of the common envelope, it coud be accompanied by a UV flash from the shock breakout itself. It would be interesting to search for observational evidence of this in future data, either by high-ionization nebular lines in echo spectra, or in blue light curves of echoes --- especially in echoes that view the eruption from polar directions that do not need to peer through dense equatorial CSM. One could also search for signals of a UV flash from shock breakout in extragralactic SN impostors. Note, however, that the broad H$\alpha$ wings that are seen in echo spectra probably do not arise from direct emission by the fast ejecta immediately after ejection, because in that case the high velocities would be seen for only an extremely brief window of time. It is more likely that the broad emission wings arise as the fastest freely expanding ejecta approach the reverse shock in CSM interaction, consistent with their more persistent appearance in Stage 2 of the eruption. This fast material will either be excited radiatively by inward propagating X-rays from the shock front, or collisionally as it crosses the reverse shock. This is commonly observed in late phases of CSM interaction in SNe IIn, and is the explanation for the very broad H$\alpha$ wings still remaining in the spectrum of SN~1987A, more than a decade after explosion \citep{smith05rs87a,heng06}. \begin{figure*} \includegraphics[width=3.3in]{fig16.eps} \caption{A sketch of the possible geometry in a hypothetical stellar merger model for $\eta$ Car's eruption, showing the two phases discussed in Sections 4.3 and 4.5. {\it Top:} Phase I corresponds to the decades leading up to the Great Eruption, which in a merger model is the inspiral phase when the orbit decays and there is prodigious mass loss from the system through L2. This is adapted from the scenario for lower-mass mergers like V1309~Sco discussed by \citet{pejcha16a,pejcha16b,pejcha17}. In the case of $\eta$ Car, light echoes from this period indicate a relatively slow outflow of 150-200 km s$^{-1}$. The luminosity in this phase is a combination of shock heating as the L2 outflow collides with itself in a ``death spiral'' \citep{pejcha17}, as well as stellar photospheric luminosity. {\it Bottom:} Phase II corresponds to the 1850s plateau phase of the Great Eruption, when higher velocities are seen. The very broad wings in H$\alpha$ correspond to a 10$^4$ km s$^{-1}$ explosive ejection or very fast wind (light blue), which is excited as the fast ejecta approach the reverse shock. The intermediate-width $\sim$600 km s$^{-1}$ lines (broader than the 150 km s$^{-1}$ seen in Phase I) correspond to the thin swept up post-shock shell (red). The post shock gas expanding at 600 km s$^{-1}$ will cool and will eventually form the dense walls of the Homunculus. This is adapted from the scenario discussed previously by \citet{smith13}, where strong CSM interaction dominates the luminosity in this phase. During Phase II, the slow 150 km s$^{-1}$ outflow can still be seen in absorption from favorable directions, although it is eventually swept up and becomes part of the walls and pinched waist of the Homunculus.} \label{fig:merger} \end{figure*} \subsection{Pulsational Pair Instability Eruption} As noted in the previous section, one potential explanation for the energy deposition required to power the fast ejecta (and for the global energetics of the event) is a pulsational pair instability (PPI) eruption. Recently this has been discussed in detail -- including applying it to the specific case of $\eta$ Carinae -- by \citet{woosley17}. Very massive stars that approach the ends of their lives with He core masses above about 30 $M_{\odot}$ will encounter the pair instability during the latest nuclear burning phases \citep{fh64,barkat67,rs67}. The core then implodes and triggers explosive burning, which may lead to a diverse range of outcomes. When the explosive burning exceeds the core binding energy, a terminal pair instability supernova (PISN) destroys the star \citep{bond84,hw02}. In the lower range of final He core masses (about 30-60 $M_{\odot}$), the energy depositon from explosive burning may not be energetic enough to completely unbind the star. Instead, a repeating cycle occurs where explosive burning expands the core, which cools, and eventually contracts again to reignite explosive burning. As a result, a series of repeated pulsations occur that cause severe eruptive mass loss before the star finally collapses to a black hole \citep{hw02,woosley17}. The mass ejected, energy, and recurrence timescales for these eruptions are diverse. A comprehensive overview of PPI eruptions has recently been discussed by \citet{woosley17}. We will not repeat that discussion here, except to note some pros and cons of PPI eruptions as an explanation for $\eta$~Car, specifically informed by clues in our new light echo spectroscopy. {\it Pro:} A key argument in favor of the PPI eruption mechanism to explain $\eta$ Car's Great Eruption is that this is a well-established fundamental mechanism that is predicted to occur in very massive stars -- even single massive stars. Indeed, $\eta$ Car is a very massive star whose likely initial mass, based on its current luminosity, is in the right ballpark to experience PPI eruptions \citep{woosley17}. Moreover, the total mass lost \citep{smith03}, total kinetic plus radiated energy of the Great Eruption \citep{smith03,smith06,smith08}, shock driven mass loss and the presence of CSM interaction \citep{smith13}, and repeated eruptions over centuries \citep{kiminki16}, fall nicely within the range of expectations for PPI eruptions \citep{woosley17}. The usual assumption that the PPI and PISN are resticted to low metallicity (because line-driven winds reduce the star's mass too much at $Z_{\odot}$) may not be such a concern, since mass-loss rates adopted in most stellar evolution models are too high anyway \citep{smith14}, and since other effects such as rotational mixing or binary interaction can influence the mapping of initial mass to final He core mass \citep{cw12}. Particularly well-matched to evidence from light echo spectrosopy and the nebulosity around $\eta$ Car is that the PPI eruptions may produce strong shocks with fast explosive outflows, giving a natural explanation for the broad wings we observe in light echo spectra (this work) and fast Outer Ejecta \citep{smith08}, but there may also be a wide range of outflow speeds. Moreover, the recurring nature of the PPI eruptions allows faster material to overtake previous eruptions and power a transient with strong CSM interaction \citep{woosley07}, as inferred for $\eta$ Car \citep{smith13}, and seems consistent with $\eta$ Car's eruptive history \citep{kiminki16}. {\it Con:} One potential objection arises because the PPI is generally expected, over most of the applicable initial mass range, to occur only within a few years before the star's final collapse to a black hole \citep{woosley07}. Yet, it has been $\sim$170 yr since the Great Eruption and $\eta$~Car doesn't seem to have collapsed to a black hole yet; moreover, $\eta$ Car appears to have suffered previous major eruptions 300-600 yr before the Great Eruption \citep{kiminki16}, so its delay between pulsations would need to be quite long. As noted by \citet{woosley17}, however, the expected outcomes of the PPI are diverse, and the delay between pulsations can in some cases be much longer -- up to centuries or even millennia. These longer delays occur at the highest part of the initial mass range just below the threshold for true PISNe, where the PPI flashes are not energetic enough to fully unbind the massive core -- but they {\it almost} do it, leading to very long recovery times as the expanded and cooled core undergoes Kelvin-Helmholtz contraction. Here we run into a potential snag, though, because applying these long PPI delays to $\eta$ Car may be problematic in two ways. (1) The longest delays occur for a fairly narrow range of mass at the highest masses (He core masses of roughly 60-64 $M_{\odot}$ for low metallicity models; \citealt{woosley17}). This would make $\eta$ Car's eruption an extremely rare event, which is in itself perhaps not a debilitating objection. However, the requirement that this long delay occurs for the highest part of the mass range exacerbates the difficulty mentioned above with reaching this end at roughly $Z_{\odot}$. Getting massive stars at $Z_{\odot}$ into the lower end of the PPI range despite mass loss is difficult enough, but getting them to reach their end with the most massive He cores seems unlikely. (2) By necessity, the PPI flashes that are {\it almost} energetic enough to unbind the star, thereby achieving a long delay before the next pulse, also produce explosive events with high energy exceeding 10$^{51}$ erg \citep{woosley17}. The kinetic plus radiative energy budget of $\eta$ Car's 19th century Great Eruption can accomodate about 10$^{50}$ erg \citep{smith03}, but a total energy exceeding 10$^{51}$ erg seems difficult to reconcile with observational estimates. A different counterargument to the PPI for $\eta$ Carinae has to do with the properties of its nebula and companion star. First, the Homunculus nebula is highly bipolar in shape, with a tightly pinched waist. The PPI doesn't give a clear explanation for this geometry. Instead, this suggests either a strong influence of shaping by interaction with a close companion star, or perhaps very rapid rotation (rapid rotation late in life despite already having suffered extreme mass loss may also require interaction with a companion). Moreover, previous major mass-loss events had a different geometry \citep{kiminki16}, which seems hard to reconcile with a single rapid rotator. Second, the PPI model gives no explanation for why $\eta$ Car's current companion star is so unusual, with a highly eccentric wide orbit and having an extremely strong and fast wind. As discussed in the final section of this paper, these properties would seem to require violent binary interaction of some sort. (Although it adds complexity, there is no clear reason to rule out a PPI event occuring in a binary system.) However, it is worthwhile to ask if a binary or multiple system interaction could account for the Great Eruption on its own, even without a PPI event. \subsection{A generic model for eruptive transients: Binary merger with CSM interaction} The $\eta$ Car system, its surrounding nebula, its historical record, and light echoes represent one of the most observationally rich objects in massive star research. Any model for $\eta$~Car must face a daunting gauntlet of observational constraints. This may act to repel conservative theorists, or to prematurely quash potentially interesting models. To any proposed simple theory applied to $\eta$~Car, one can usually respond with ``Yes, but what about... [insert obscure observational detail here]?'' For now, we momentarily stow such objections in order to discuss a general scenario, and we will return to specific complexities of $\eta$~Car later in Section 4.7. The two stages of the event described above (in Section 4.3) arose from an empirical description of the spectral evolution and historical light curve, but we can also attempt to ascribe a physical cause to these observed changes. Here we discuss a promising physical model to account for the properties of $\eta$ Car's Great Eruption: a massive star merger event. A merger of two massive stars provides an attractive model for the trigger and the energy supply of the Great Eruption, and has been discussed before (see the Introduction). However, the merger model has never been reconciled with detailed observational constraints for $\eta$~Car, and there have been significant unanswered problems with a simple merger as proposed in previous models. Here we describe how a merger model may indeed be reconciled with many of the observed properties of the Great Eruption. The model described below differs from previously proposed merger models for $\eta$~Car in that it adopts the hypothesis that CSM interaction is a main engine for producing the observed radiative luminosity. This provides a self-consistent explanation for the two observed phases of the event and their properties. In a simplified scenario involving a binary merger, these two phases may be understood basically as: \textit{\textbf{Phase 1 (1840s and before):}} This is the inspiral phase with mass transfer and/or common envelope, when mass and angular momentum are shed from the L2 point, allowing the orbit to degrade and for the binary to eventually merge. This creates a relatively slow outflowing disk or torus (Fig.~\ref{fig:merger}; top). In this phase, the mass shed in an equatorial spiral from L2 quickly catches and shock heats previous L2 mass loss as the spiral winds up. The shock heating of this torus may make a considerable contribution to the total luminosity, and gradually rises as the stars move closer together. A such, the ``photosphere'' in this phase may be a composite of the two stellar photospheres plus optically thick shock-heated material in the circumbinary disk. This inspiral phase and the associated luminosity has been discussed in detail for lower-mass stellar mergers \citep{pejcha14,pejcha16a,pejcha16b}. \textit{\textbf{Phase 2 (1850s plateau):}} Relatively sudden energy deposition from a final merger event (i.e. the merging of the two stellar cores inside the common envelope) steepens to a strong shock in the bloated envelope, and triggers an explosive outflow or very strong fast wind. As noted in previous sections, a large deposition of energy deep inside the star's extended envelope is likely to steepen to a strong shock \citep{rm17}, and we suppose that this is the driving mechanism of the fastest ejecta in $\eta$~Car. The fast ejecta from this explosive outflow or fast wind then overtake and shock the slower outflow from Phase 1 (Fig.~\ref{fig:merger}; bottom). This sudden ejection will naturally lead to a sustained phase of high luminosity that could last for years, because the fast ejecta take time to expand through the previous mass loss as in a scaled-down version of a Type~IIn supernova \citep{smith13}. This is the plateau phase of the Great Eruption, when $\eta$~Car was thought to have exceeded its own classical Eddington limit for about a decade. A key point, however, is that in the model proposed here, the observed ``super-Eddington'' light comes largely from CSM interaction \citep{smith13}, and not from stellar luminosity diffusing out through a bound stellar atmosphere as in super-Eddington wind models \citep{owocki04,owocki17,os16,quataert16}.\footnote{Note, however, that the super-Eddington model discussed by \citet{quataert16} is somewhere in between, where much of the emergent radiation has been processed by internal shocks as the fast super-Eddington wind overtakes slower outflowing material upstream, similar to that described here. An interesting direction for future work on this model would be to determine if a super-Eddington wind with such internal shocks could achieve the extremely fast speeds observed in our light echo spectra.} A few key attributes make a simple model like this plausible: First, it is self consistent, in the sense that the initial inspiral during Phase I must shed mass and angular momentum in order for the orbit to degrade and for a merger to proceed, but in doing so, it also naturally provides the CSM required for the interaction that will occur in Phase II. We do not need to invoke a different mechanism or a previous outburst to provide the CSM into which the shock expands. Lost from the L2 point, this Phase I mass loss should be relatively slow and concentrated in the equator. CSM interaction where a fast outflow overtakes a slow and dense equatorially concetrated outflow (i.e. a torus) may naturally lead to a bipolar shape in the resulting nebula \citep{frank95,langer99}. Second, CSM interaction is an extremely efficient engine for converting outflow kinetic energy into radiation. Whereas a recombination plateau may typically tap only $\sim$1\% or less of the available kinetic energy, CSM interaction is much more efficient, typically converting $\sim$50\% or more of the available kinetic energy into light. The conversion is most efficient when a fast and relatively low-mass explosion collides with slow and massive CSM. Thus, for a given explosion or eruption energy, transients with CSM interaction will be much more luminous than those where the radiation is powered by a recombining H envelope. This is a key way to get super-luminous SNe II from normal core-collapse energy \citep{sm07}, and it is what converts kinetic energy to radiation in PPI eruptions \citep{woosley17}. Similarly, it will allow events with modest kinetic energy to achieve substantial outburst luminosity. In any magnitude-limited sample of transients found in nearby galaxies, then, the ones powered by CSM interaction are likely to make a dominant contribution simply because they tend to be brighter than those without interaction for the same explosion energy. Third, this two-phase merger scenario is qualitatively consistent with the time sequence of outflow velocities seen in light echo spectra of $\eta$~Car, which initially show a slow velocity of 150-200 km s$^{-1}$. As time proceeds, extremely fast material appears, while simultaneously the ``narrow'' emission component broadens from 200 to about 600 km/s, as if it is being accelerated by a shock. Some slower velocities (150 km/s) are still seen in absorption along this line of sight, even at later times when the broad emission appears, requiring the simultaneous presence of slow moving CSM outside much faster ejecta. It would be difficult to avoid strong CSM interaction, based on this observational evidence alone. As noted by \citet{smith13}, the radiative shock that forms will lead to a very thin post-shock layer as in models of SNe~IIn \citep{vanmarle10,chugai04}, which explains many attributes of the structure of the Homunculus around $\eta$ Car that are harder to explain with wind mass loss. While the $\eta$~Car system has additional complexities that will be discussed below, this basic sort of merger plus CSM interaction model may be widely applicable to other non-terminal transients. \citet{smith13} proposed that CSM interaction may be responsible for a wide range of non-terminal transients besides $\eta$ Car, including other LBV giant eruptions, SN impostors, SN~2008S-like events, pre-SN eruptions, infrared transients, so-called luminous red novae, or other events. All these objects have quite similar spectra at peak resembling SNe~IIn, although they exhibit a wide range of outflow speed and luminosity \citep{adams15,kochanek11,prieto09,smith+11}. A similar model with CSM interaction was applied recently to luminous red novae like V1309~Sco as well \citep{mp17}. In this physical scenario, $\eta$~Car's Great Eruption can be seen as one of the most energetic examples of stellar merger events that span a wide range of initial masses for massive stars, extending to similar transients that have been associated with mergers from low-mass stars. These include clear mergers like the spectacular example of V1309~Sco \citep{tylenda11,pejcha14}, and may include other suspected mergers like V838 Mon \citep{bond03}. Dusty transients like SN~2008S, NGC 300-OT, and their kin have been discussed variously as terminal events like electron capture SNe, eruptive transients, or possibly mergers as well \citep{prieto08,prieto+08,berger09,bond09,botticella09,thompson09,smith09,smith+11}. \citet{kochanek14} estimated rates and argued that merger events from low-mass binaries are common, with the rate falling toward higher initial mass, and they noted that the peak luminosity is a steep function of the initial mass. \citet{smith16b} discussed ways that $\eta$ Car and the recent transient in NGC~4490 seem to be an extension of stellar merger transients to higher initial mass, and noted that the duration of the bright transient may also depend on initial mass. \citet{smith16b} noted that these objects can show quite similar spectra at various times in their evolution. \citet{blagorodnova17} suggested that the recent transient in M101 may be a massive star merger that fits with this scenario as well. Although the common envelope phase and stellar mergers are still far from being well understood, theoretical models of low-mass events are able to reproduce some aspects of the observed transients, in most cases with the mass ejection and luminosity powered by recombination \citep{ivanova13,nandez14}. For the case of $\eta$ Car, the high speed outflow suggests strong shock excitation, the decade-long plateau seems too long to be powered by recombination, and the thin walls of the resulting nebula appear to have been compressed in a radiative shock \citep{smith13}, so a merger model with CSM interaction would seem more appropriate for $\eta$~Car, closer to the models for V1309 Sco discussed by Pejcha and collaborators (see above). In fact, one might argue that a merger model where CSM interaction does not constitute a significant fraction of the luminosity could be problematic; the inspiral must eject mass and angular momentum in order for a merger to occur, so an ejected envelope will likely collide with it unless the final ejecta are very slow. Even in some low-mass meger events like V1309 Sco, some evidence of relatively high speed outflow is seen \citep{mason10}, making CSM interaction seem likely (although the speeds are not as extreme as seen in echoes of $\eta$ Car). In the specific case of $\eta$~Car, a merger scenario as described above is quantitatively plausible in terms of the energy budget of the event and the resulting nebula. \citet{smith13} already showed that the luminous plateau phase of the Great Eruption (1845-1860) can be powered by CSM interaction with a $\sim$10$^{50}$ erg explosion running into previous mass loss. This energy can easily be supplied by a massive-star merger event. The gravitational potential energy of two massive stars (say $\sim$60 $M_{\odot}$ each) that are about to merge (separated by roughly 100 $R_{\odot}$, which is roughly the present-day radius of the primary star; \citealt{hillier01}) would be about 1.4$\times$10$^{50}$ erg. This is independent of the source of energy for that explosive outflow (i.e. a PPI event works as well, as noted above). However, a merger scenario provides a natural explanation for why the $\eta$~Car primary seen today, which would be a merger product, appears to still be a very rapid rotator, rotating at a substantial fraction of its expected critial rotation rate \citep{do02,smith02,smith06,st07,smith03a}. An explanation for the high current rotation and for the bipolar nebula is harder to find in any single-star model (including the PPI model), because the primary has expanded significantly from its ZAMS radius and it has already shed huge amounts of mass and angular momentum. One difference indicated by our light echo spectroscopy, as compared to the simple CSM interaction model in \citet{smith13}, is that the post-1843 mass-loss rate may be lower and its bulk outflow speed is higher. Radiative transfer models of the light echo spectra are needed for quantitative estimates of how much mass is contained in the outflow responsible for the broad H$\alpha$ wings we observe, but with more kinetic energy per unit mass in the fast outflow, this may suggest that more of the mass in the Homunculus was provided by the slow L2 mass loss in Phase 1, while most of its kinetic energy came from the explosive event in Phase 2. It will be interesting to see if future numerical simulations of an explosive outflow expanding into a compact torus similar to the one formed by L2 mass loss in lower-mass models \citep{pejcha16a,pejcha16b} can explain observed structures in the nebula of $\eta$ Car, such as the apparent ``holes'' in the polar caps \citep{smith06,smith03}. \citet{gonzalez18} recently simulated the shaping of the Homunculus via an explosion with shock interaction following the scenario proposed by \citet{smith13}, finding that it could indeed explain structural features of the Homunculus. With the adopted parameters, the Homunculus simulated by \citet{gonzalez18} was expanding too quickly, and so that author favored a super-Eddington wind scenario. The adapted pre-explosion mass loss and the explosion parameters are not tightly constrained, however, and a different ratio of mass-loss rates in Stage 1 and Stage 2 can change the resulting final speed of the Homunculus. The fast ejecta speeds reported here, indeed, alter the expected explosion parameters one might adopt and point directly to an explosive outflow. Overall, this simplified two-phase merger scenario therefore gives a somewhat satisfactory (although incomplete) explanation for the evolution of observed light echo spectra of the Great Eruption, basic energetics of the event, and basic structural properties and kinematics of the Homunculus seen today. It does not, however, resolve questions of greater complexities in the presently observed system, including the unusual surviving companion star, a history of previous outbursts, and some of the more detailed structures in the nebula. These are discussed below. \subsection{Can we reconcile this generic binary merger picture with the complexity of Eta Car and its wide companion?} In a standard picture of a stellar merger event, guided by spectacular recent examples such as V1309~Sco, one envisions a relatively long (many orbits) inspiral phase with L2 mass loss shedding both mass and angular mementum, followed by a relatively brief plunge-in (common envelope) phase that may give rise to a sudden but one-off transient event. The end product should be a (potentially dust-obscured) rapid rotator surrounded by a toroidal nebula. As such, this type of scenario provides a plausible origin for the unusual properties of LBVs and B[e] supergiants \citep{justham14,ppod10}, both of which have blue-straggler environments that are more isolated from O-type stars than expected in single-star evolution \citep{st15}. Mergers may also help account for LBVs as immediate SN progenitors \citep{justham14}. In some basic respects, $\eta$~Car seems to fit this general two-phase merger picture, as noted above, and there are attractive aspects of a merger model for $\eta$ Car as noted previously by several authors \citep{jsg89,iben99,ppod10,pz16,justham14,smith16}. It is the most luminous star in its birth cluster Trumper 16 \citep{smith06b}, consistent with being a blue straggler (even though the lifetimes are all similarly short for such massive stars), and there are clues that it is currently a rapid rotator \citep{smith02,smith03wind,do02}. However, $\eta$~Car also has a lot of observational ``baggage''. There are several apparent contradictions between simplified merger models and the vast amount of observational data for $\eta$ Car, which requires that the situation is more complicated if indeed a stellar merger is the correct physical description. To reconcile a merger model with the specific observed case of $\eta$ Car, the most critical problems posed by observations are that: (1) the system is still a binary (and the surviving companion star is weird), and (2) the Great Eruption was not a one-off event, because the star suffered major mass-loss eruptions at least twice before. There are a number of other complexities as well. These challenges are discussed below. We warn the reader that the following sections are speculative; it is not our goal here to provide the final word, but to contemplate how many pieces of the complex puzzle fit together. \begin{figure}\begin{center} \includegraphics[width=2.3in]{fig17.eps} \end{center} \caption{A sketch of the proposed orbital interaction in a hierarchical triple. {\it Panel A:} As the primary star nears the end of the main sequence, its expansion initiates mass transfer. {\it Panel B:} After RLOF ends, the primary is a stripped-envelope WR star, and the mass gainer is spun up, overluminous, and enriched (a blue straggler). The orbit of the inner binary widens at the end of RLOF, and apastron may widen further due to mass loss or Kozai cycles. {\it Panel C:} Eventually the orbit widens enough that the WR interacts with the wide tertiary companion and they exchange places (TEDI; \citealt{pk12}), sending the WR star out on an eccentric orbit, and sending the tertiary inward. Various processes influence the orbit until a series of colisions and a final merger occur.} \label{fig:triple} \end{figure} \subsubsection{But what about the weird surviving companion?} In its presently observed state $\sim$170 yr after the Great Eruption, $\eta$ Car is known to be a binary system with a roughly 5.5 yr period and a highly eccentric ($e \simeq 0.9$) orbit \citep{damineli,dcl97,madura12}. Thus, if the Great Eruption was powered by a stellar merger event, then its persisting binarity requires that the eruption must have occurred in a hierarchical triple system. Explaining $\eta$ Car's eruption as binary merger in a triple system was proposed about two decades ago \citep{iben99}, although it is now amusing to note that Iben described such a model as ``preposterous''. In that model, the wide tertiary (now binary) companion was essentially an innocent bystander that is a relatively inert main sequence O-type star. However the surviving companion is probably not a normal main sequence O-type star (see below). More recently, \citet{pz16} proposed a conceptually similar model of a binary merger in a triple system, although this time with the surviving companion playing a more active role by helping to initiate the merger of the inner binary via the Kozai-Lidov effect \citep{kozai62,lidov62}. This model, however, has discrepancies with observational parameters. For instance, because of the precarious nature of the original triple system, a merger or collision is triggered by the Kozai-Lidov mechanism after only 0.1-1 Myr \citep{pz16}, which is in tension with the 3-4 Myr age of the Tr16 cluster, and also at odds with the fact that the strong enrichment of nitrogen in $\eta$ Car's ejecta requires that it is a $>$3 Myr old (post-main-sequence) object \citep{davidson82,davidson86,sm04}. This model doesn't explain why the orbit of the wide tertiary that survives today as the eccentric binary companion would be aligned with the equatorial plane of the Homunculus (presumably the plane of the binary merger), since the Kozai-Lidov torque is most effective when the two orbits are misaligned (indeed, \citealt{pz16} adopted a relative inclination between the inner and outer orbital planes of about 90$^{\circ}$). Moreover, in this model \citet{pz16} explained the ejection of the Homunculus via an enhanced wind caused by tidal heating several decades before the 1840s with the final merger occuring in 1838; this is ruled out by the 1847.1 ($\pm$0.8 yr) ejection date from proper motions \citep{smith17}, and the radiatively driven wind of 500 km s$^{-1}$ in their model provides no explanation for the fast Outer Ejecta \citep{smith08}, the very fast ejecta seen in light echoes, the extremely thin walls of the bipolar Homunculus, multiple major eruptions, or the very luminous 1850s plateau in the light curve. In this model as well, the surviving companion should be a normal main sequence O-type star, which it is probably not (see below). Nevertheless, it is interesting to pursue the model of a merger in a triple system, to find a satisfying scenario that is in agreement with available observational constraints. \citet{smith16b} showed that, at least observationally, the light curve and spectra of $\eta$ Car's eruption have similarities with other transients that have been proposed as mergers. Any model that aims to explain the Great Eruption as a merger event needs to account not only for the highly eccentric orbit of $\eta$~Car's surviving companion seen today, which is in the same plane as the equator of the Homunculus, but also its unusual nature and likely evolutionary state. Previous models assumed that the companion was a main sequence star, but this seems inconsistent with its observed wind properties. The inferred physical properties of the surviving companion star in the $\eta$~Car system indicate that it was probably a more active player than previously suggested. In particular, estimates of the companion's wind indicate extreme physical parameters. Typical values for the mass-loss rate and wind speed of $\eta$ Car's companion derived by comparing models of the colliding winds with observed X-ray emission are $\dot{M}$=(1-2)$\times$10$^{-5}$ $M_{\odot}$ yr$^{-1}$ and $v_{\infty}$$\approx$3000 km s$^{-1}$, respectively \citep{parkin11,parkin09,corcoran10,corcoran05,pc02,okazaki08,russell16,hamaguchi16}. This is a much denser and faster wind than any normal main-sequence O-type star, especially when considering constraints on the companion's luminosity and ionizing flux that would point to an initially $\sim$30 $M_{\odot}$ star or less \citep{mehner10,teodoro08,verner05}. A typical 30 $M_{\odot}$ main sequence star, by contrast, would have $\dot{M}$ \, = \, 10$^{-7}$ $M_{\odot}$ yr$^{-1}$ and $v_{\infty}$=1000 km s$^{-1}$ \citep{smith14}. Instead of a main-sequence companion, then, the extreme wind properties of $\eta$ Car's companion are much more consistent with it being a fairly typical hydrogen-poor early-type Wolf-Rayet (WR) star. The surprising strength and speed of the companion's wind, as well as its possible similarity to those of WR stars, has been noted several times before, but to our knowledge {\it the rather profound evolutionary implications have so far not been discussed in the literature}. Namely, finding that $\eta$ Car's surviving companion is likely a WR star forces us to rethink the interaction history of this system. If the primary star were an extremely massive single star of $M_{\rm ZAMS} \simeq 200 M_{\odot}$ or so, or even if it is a merger product from two 60-80 $M_{\odot}$ stars, then its short main-sequence lifetime should prohibit its lower-mass companion (currently the secondary) from being a WR star already as a consequence of its own unaltered evolution with mass loss. More massive stars have shorter lifetimes, and there is not enough time for this to have happened if the companion has evolved effectively as a single star. Indeed, a 30 $M_{\odot}$ star that has made it to the WR phase already would be older than the 3 Myr stellar population in the Carina Nebula \citep{smith06b}, so one would need to invoke the unlikely scenario that a wandering star ejected from another cluster was captured to make the $\eta$ Car multiple system. A companion star of $M_{\rm ZAMS} \simeq 30 M_{\odot}$ should still be midway through its main-sequence H-burning phase, and it certainly would not yet have shed its H envelope --- especially when its much more massive companion is still an LBV that has retained its H envelope. At this point we must resign ourselves to the fact that something fairly complicated has happened to $\eta$~Car. We have what appears to be an initially $\sim$30 $M_{\odot}$ star that has reached its He core burning phase as a compact fast-winded WR star, while orbiting around one of the most luminous blue supergiant stars in the Milky Way that still retains its H envelope and has presumably just finished its core H-burning evolution. How can this be? Without the assistance of time travel, {\it we conclude that the surviving wide secondary must have participated in close binary evolution in order to alter the masses and lifetimes of the system components that have been inferred}. This requires previous close interaction and mass exchange, followed by dramatic orbital evolution. Here is a scenario that appears plausible, given observational constraints, although it is admittedly speculative and probably not a unique explanation. \begin{itemize} \item Suppose that the $\eta$ Carinae system was originally a high-mass hierarchical triple system (Figure~\ref{fig:triple}a). Toward the end of its main sequence, the initially most massive star in the inner close binary initiates stable mass transfer, being stripped of its H envelope in the process, and donating that envelope to its mass-gainer companion. \item This mass transfer produces a stripped-envelope WR star (initially the primary, now the wide secondary), and makes the mass gainer into a rapidly rotating, overluminous, N-enriched, blue supergiant or LBV that is now the most massive star in the system (Figure~\ref{fig:triple}b). \item Up until this point in the evolution of the system, the wide tertiary has been inert. However, after mass transfer and reversal of the mass ratio, the inner binary system's orbit widens \citep{paczynski71}. With the center of mass closer to the more massive LBV-like mass gainer, it is the stripped WR star who moves outward (Figure~\ref{fig:triple}b). \item The speculative part about this scenario is that this widening of the orbit may then trigger a violent 3-body interaction, wherein the stripped WR star and the outer tertiary companion (still a main sequence O-type star) exchange places. This interaction kicks the WR star out on a wide and highly eccentric orbit (as observed today), while kicking the previously inert tertiary star inward to interact and eventually merge with the mass gainer (Figure~\ref{fig:triple}c). \item The orbital evolution of the inner binary may then be pushed to a merger by Kozai cycles, by grazing collisions with the bloated mass-gainer, and/or by interaction with a disk that still surrounds the mass gainer. \item This final merger, potentially preceded by a few violent grazing collisions, eventually culminates in the Great Eruption. \end{itemize} Is this type of model really so far-fetched? For earthling astronomers accustomed to orbiting a single star, stellar binary and triple systems often seem exotic. However, observed statistics indicate that for massive stars, binary interaction is the norm, and hierarchical triple systems are common rather than a rare exception \citep{moe17,chini12,kiminki12,kk12,sana12,eggleton07,kf07}. Mass loss and mass transfer in the inner binary in triple systems can have dramatic effects in tandem with Kozai cycles and tides \citep{naoz16,st13,mp14,kem98}, and companions can exchange places \citep{kea94,keo94,pk12}. \citet{pk12} have discussed in detail how mass loss and interaction from the inner binary in a hierachical triple system can lead to chaotic orbital evolution, including collisions, exchanging places, and the formation of highly eccentric systems, in an instability they refer to as the triple evolution dynamical instability (TEDI). The present configuration of the $\eta$ Car system seems to naturally fit expectations from TEDI. Triple systems have been invoked to help explain the observed fraction of close binaries \citep{tokovinin06,moe17}, and mergers in triples akin to the scenario above may be necessary to explain the origin of blue stragglers \citep{pf09,iben99b,pk12}. This may be particularly relevant to LBVs, since the observed environments of LBVs suggest that they are indeed massive blue stragglers \citep{st15}. Interestingly, a triple-system encounter that led to an exchange of the original wide tertiary and an inner companion was already suggested for $\eta$ Car \citep{lp98}, although not in the context of triggering a merger. While the scenario outlined above is admittedly somewhat ad hoc and in need of further quantitative exploration to assess its probability, the basic picture is self consistent and provides a single plausible explanation for a large number of pecularities of the $\eta$ Car system: The exchange of partners explains the highly eccentric and aligned orbit of the secondary, it accounts for the apparent age discrepancy where a lower-luminosity WR star is at a more advanced evolutionary stage than the more luminous H-rich primary, and the exchange of partners is the event that sent one of the stars inward to trigger the merger that powered the Great Eruption. It may also shed some light on the previous eruptions (discussed in the next subsection). We have a plausible explanation for why this event occured at the 3-4 Myr age of the Tr16 cluster and not sooner, which is that the process began as a result of mass transfer when the initial primary finished its main sequence evolution; the widening orbital evolution following mass transfer is what initiated the 3-body interaction in a hierachical system that had been relatively stable during the main sequence. Without having mass transfer and an exchange of partners that kicked out the stripped envelope star (originally the primary), it is almost unfathomable that we could have a WR star that has a significantly lower presumed initial mass while being in a more advanced stage of evolution than its more luminous primary star, and also in a highly eccentric orbit. Capturing an unrelated star from the host cluster might explain the high eccentricity, but would still conflict with the apparently advanced evolutionary state of a star whose H-burning lifetime is longer than the Tr16 cluster. We encourage quantitative studies of the 3-body parameter space that might lead to the current observed properties of the $\eta$ Car system in the qualitative scenario that we described, but such an exploration is well beyond the scope of our current paper. \subsubsection{But what about the previous eruptions?} Another key challenge for any merger model of the Great Eruption is that a binary system merging into one star should be a singular event, whereas observational evidence indicates that $\eta$~Car has erupted repeatly. Proper motions of the Outer Ejecta around $\eta$~Carinae reveal that it suffered at least two major eruptive mass-loss episodes prior to the 19th century Great Eruption; these occurred aproximately 300 and 600 yr before \citep{kiminki16}. These precursor eruptions also had a different geometry than the 19th century eruption, with the first being almost entirely one-sided (to the northeast and blueshifted) and the second sending material in a few different directions, but with neither sharing the axisymmetry or orientation of the bipolar Homunculus \citep{kiminki16}. In the scenario discussed above, where a triple-system exchange leads to a stellar merger, one might envision the multiple eruptions as a consequence of several ``near misses'' or grazing collisions before the final stellar merger event. This is an expected outcome of the TEDI instability described above. After the exchange of partners that kicks out the WR star and kicks in the original tertiary, the inner binary will likely be eccentric and its orbital evolution can be influenced by Kozai cycles as noted above. This may drive the binary to closer periastron distances such that they eventually have a near miss or their outer envelopes collide, leading to significantly asymmetric mass ejection from the primary star's loosely bound envelope. Additionally, the companion may interact with a remnant disk or bloated envelope around the rapidly rotating mass gainer \citep{munoz15}. The tidal or dynamical friction of these grazing collisions may degrade the orbit enough that that the binary finally merges together after a few violent encounters \citep{pk12}. (Collisions resulting from the TEDI instability were originally discussed in the context of enlarged red giant envelopes in lower mass stars, but the bloated envelope of a massive LBV also has a very large radius.) Since the first pass is likely to be the most eccentric, this might be an interesting explanation for why the first of the three historical eruptions sent ejecta in largely one direction \citep{kiminki16}. The final event (the actual merger of the two stars) produced strong mass loss that was, by contrast, highly axisymmetric (the Homunculus). If a significant fraction of the primary star's outer envelope is ejected in these repeating eruptions, there may be a physical mechanism for the delay between eruptions. Namely, the star's envelope may recover and re-establish equilibrium on a thermal timescale (centuries for the mass of the outer envelope). When the envelope reaches a radius comparable to the periastron separation, this may trigger the next grazing collision. The outer companion star seen today shares the same orbital plane as the equator of the Homunculus, whereas the previous eruptions did not share the same geometry. One might also then envision this series of grazing collisions or failed mergers as a process by which the merging stars exchanged angular momentum, so that the inner orbits became aligned with the primary star's rotation and the orientation of the outer binary. This is admittedly speculative and unrefined, but at least it gives some plaubile explanation for the differing geometry in subsequent historical eruption events. This is lacking from any other model for $\eta$ Car's eruption discussed so far, and is a fruitful topic for numerical simulations. \subsubsection{But what about the structure of the equatorial skirt and the Strontium Filament?} The structure of $\eta$ Car's highly non-spherical ejecta provide important clues about its recent violent mass-loss history, as noted previously (see, e.g., \citealt{smith12,smith13}). The detailed structure of the material in the equatorial plane has particular relevance to any merger model, since the inspiral phase leading to a merger should shed a large amount of mass through L2 \citep{pejcha16a}. The ragged spray of debris referred to as the Equatorial Skirt is a particularly recognizable feature of the Homunculus \citep{morse98}, and the so-called ``Strontium filament'' is one particular location in the equatorial skirt with unusually strong emission lines from [Sr~{\sc ii}] and a number of other low-ionization metals \citep{hartman04}. Most of the equatorial skirt has the same age as the rest of the Homunculus \citep{morse01}, but some features appear to have been ejected decades before or afterward \citep{smith17}. One expects these features to have some satisfactory explanation in a model of the Great Eruption. At the pinched waist of the Homunculus Nebula, multiwavelength observations (especially in the infrared) have revealed the existence of a complicated toroidal structure. This torus has been discussed by numerous authors, the history of which is summarized by \citet{smith12}. More recently, high-resolution observations of $^{12}$CO 2--1 with the Atacama Large Millimeter Array (ALMA) have revealed new clues to the structure of the equatorial ejecta \citep{smithALMA}. The CO emission shows a toroidal structure with a size similar to previous IR data, but the density structure of the CO torus departs strongly from azimuthal symmetry. In particular, the torus is a series of clumps, with higher density on the far southest side and an opening on the northwest side. It basically shows a ``C'' shape, with the gap in the ``C'' pointing toward us (i.e. the near side of the equator to the northwest direction). The connection to any model for the Great Eruption is evident when we consider the relative orientation of the torus compared to the currently observed eccentric binary system. Namely, the direction of apastron in the eccentric binary seen today is also to the northwest \citep{madura12}, toward the middle of the gap in the torus \citep{smithALMA}. This suggests that the wide eccentric companion seen today had roughly the same orbital orientation during the lead up to the Great Eruption (i.e. the inspiral or Phase 1 as discussed above), which it maintained as the equatorial ejecta expanded past it. This is consistent with the notion that the orbital period did not change by more than a few percent before and after the Great Eruption \citep{smith11}. In other words, the wide secondary star we study today was not a major source of energy to power the Great Eruption. It may be that periastron encounters were able to enhance the mass loss toward the far side of the nebula, whereas equatorial material ejected toward us (i.e. toward apastron) by the central binary may have been diverted or disrupted by the eccentric companion after being shed by the inner binary. This is an area where hydrodynamic simulations could be very useful to understand the interaction. It is interesting to note, however, that the polar lobes of the Homunculus do not depart so strongly from azimuthal symmetry. In any case, the resulting density structure of the torus helps explain why the currently observed equatorial skirt appears as a ragged spray of streamers in optical images. Namely, the equatorial skirt is illuminated by scattered light that escapes preferentially through holes and gaps between clumps in the inner torus. Because of the azimuthal asymmetry, more light escapes in the direction of apastron, which is why the equatorial skirt is seen mostly on our side of the Homunculus. The fast post-eruption wind from the secondary can also escape more easily through the gaps in the torus, influencing structures on the near side of the equatorial ejecta. See \citet{smithALMA} for a more detailed discussion. The peculiar nature of the Sr filament may have less to do with the Great Eruption. As discussed recently by \citet{smithALMA}, the Sr filament is downstream from the dense inner clumps known as the ``Weigelt knots'', and its pecular low ionization may arise because it is shadowed by them. The Weigelt knots appear to have been ejected later, probably in the 1890 eruption or afterward \citep{dorland04,smith04acs}, but they are clearly in the direction of apastron in the equatorial plane. The asymmetry of the mass ejection during 1890 is beyond the scope of our paper. \subsubsection{But what about the S Condensation and NN jet?} Further outside the Homunculus, emission line images reveal a spectacular array of clumpy nebular structures known as the Outer Ejecta \citep{walborn76}. Most of these are from older eruptions centuries before the Great Eruption \citep{kiminki16}. Prominent among the Outer Ejecta are the so-called ``NN Jet'' and ``S Condensation'', which resemble collimated outflows in the equatorial plane \citep{walborn76,walborn95,kiminki16,mehner16}. The S Condensation and NN Jet have proper motions that indicate they were either ejected in the Great Eruption or in the decades leading up to it, but not in the eruptive events 300 and 600 years before \citep{kiminki16}. While detailed study of the kinematics has revealed that they are ballistic ejections rather than true hydrodynamic jets \citep{meaburn93,glover97,morse01,kiminki16,mehner16}, one may still wonder how $\eta$ Car managed to simultaneously send two large bullets out in opposing directions in its equatorial plane in an event that was otherwise highly axisymmetric. Here, again, it is likely that the wide eccentric secondary plays an important role in this non-axisymmetric structure. As noted by \citet{smith11}, the apparent magnitude in the historical light curve and apparent color in the years leading up to the Great Eruption peak dictate that the emitting photosphere was larger than the periastron separation of the current binary, and that as such, some sort of violent interaction like a collision must have occurred at times of periastron. This remains true whether this photosphere was a true hydrodynamic surface of the star, or (more likely) the photosphere of a common envelope or the emitting radius of the shock heated circumbinary disk during the inspiral phase of a merger. These periastron collisions have been interpreted as causing the brief luminosity spikes seen in the historical light curve in 1838 and 1843 \citep{smith11}. The wide companion must have plunged into one side of the dense envelope around the star and popped out the other side, possibly multiple times. This is discussed in more detail in our companion paper on the fast ejecta seen in $\eta$ Car's light echoes \citep{smith+18}, where we point out that with the known orbital geometry from \citet{madura12}, the S Condensation roughly matches the direction of ingress and the NN Jet roughly matches the trajectory of egress in this collision. The wide companion may have left a tunnel in its wake as it drilled its way through the extended toroidal envelope, through which fast ejecta from the Great Eruption may have been able to squirt. We noted earlier that the known echo geometry of EC2 indicates that it views $\eta$ Car roughly in the equatorial plane, and its position angle is within 10-20 deg of the S Condensation's trajectory (also in the equatorial plane). In other words, if the S Condensation is a bullet of fast ejecta, then the EC2 echo seems to be looking nearly down the barrel of the gun. This is probably related to the very broad wings of H$\alpha$ seen in the EC2 echo reported here. \subsubsection{But what about the brief 1838 and 1843 luminosity spikes?} Following the scenario discussed in the previous two subsections, the star that we now see as the wide secondary was orbiting around a close binary in the process of merging, as it ejected substantial amounts of material in the equator. In addition to profoundly influencing the structure of the outflowing ejecta seen today around $\eta$~Car, the close periaston encounters may have played an important secondary role in the light curve. These periaston passes were not enough to power the total kinetic and radiated energy of the whole 10$^{50}$ erg event. However, shock heating from the companion ripping through the L2 mass loss disk or plunging into the bloated common envelope of the merger may have powered the brief luminosity spikes in the light curve \citep{smith11,sf11}. In this view, the main 1850s plateau and the brief 1838 and 1843 precursor spikes have two different specific physical causes, even though they both stem from the same merger event. This may help explain why similar brief luminosity spikes are absent in UGC~2773-OT, even though the light curve and spectral evolution of its plateau phase are almost identical to $\eta$ Car's (i.e. the outer tertiary may have been on a wider orbit that didn't cause periastron collisions in the case of UGC~2773-OT, or it may have simply been a merger not in a triple system). Even if we are restricted to merger events in triple systems, the influence of periastron encounters by the outer companion could cause significant diversity from one merger event to the next, depending on the configuration of the outer orbit and the mass of that companion. This may be related to the tremendous diversity in extragalactic SN impostors \citep{vdm12,smith+11,pastorello10}. \subsubsection{But what about the age of the Homunculus?} In the two-stage merger scenario proposed above, the ejected mass that constitutes the Homunculus Nebula actually leaves the star over a time period of several decades or more. The single, well-constrained apparent ejection date for the Homunculus of 1847.1 $\pm$0.8 yr from proper motions \citep{smith17,morse01} arises because a strong shock from an explosive event sweeps up this previously ejected material into a thin cooled shell with a single dynamical age. The shock cooling, which provides the luminosity of the plateau, also allows this swept up material to collapse to a very thin layer as seen today in the walls of the Homunculus \citep{smith06,smith13}. This shock essentially erases the previous mass-loss history and creates the illusion, when we measure its proper motion expansion, that the whole nebula was ejected instantaneously \citep{smith17}. Doppler velocities seen in light echoes, on the other hand, show that strong mass loss was occurring over several years, but that the fastest material appeared after 1847. We have associated the fast, explosive outflow with a shock that results from energy deposition due to the final stellar merger event, whereas the slower 150-200 km s$^{-1}$ outflow seen prior to 1847 was due primarily to mass loss associated with the inspiral phase of the merger. \subsubsection{But what about...?} Admittedly, there are remaining open questions associated with $\eta$ Car and its nebula that are not close to being settled by the speculation in this final part of the paper. The general scenario outlined above certainly is a challenge for quantitative models to match, but light echoes combined with the properties of $\eta$ Car's ejecta narrow the range of possible configurations and free parameters considerably. While it may be fun to entertain even more complicated scenarios (multiple mergers, precessing jets, exotic compact object interactions, a T$\dot{\rm Z}$O, etc.), the triple-interaction scenario described above with an exchange of partners leading to a merger seems like the minimum level of complexity needed to simultaneously account for a WR-like companion in a wide eccentric orbit around a much more massive and more luminous rapidly rotating primary that still retains its hydrogen envelope, which can, moreover, achieve disk plus bipolar geometry, heavily nuclear-processed ejecta, and multiple eruptions with one of those being a $\ge$10$^{50}$ erg explosive ejection event that accounts for the spectroscopic evolution seen in light echoes. | 18 | 8 | 1808.00992 |
1808 | 1808.02445_arXiv.txt | Nearby pulsars have been suggested as sources of $\sim$TeV $e^+/e^-$ Cosmic Ray (CR) excess on Earth. The High-Altitude Water Cherenkov Observatory (HAWC) detected extended TeV emission regions in the direction of two nearby middle-aged pulsars, Geminga and PSR B0656+14. By modeling the TeV emission as inverse Compton emission from electron-positron pairs diffusing in the interstellar medium (ISM), the HAWC collaboration derives a diffusion coefficient much smaller than the standard value in the vicinity of the two pulsars, which make them unlikely the origin of the positron excess. We propose that the observed $\gamma$-ray emission originate from the relic pulsar wind nebula. A two zone diffusion model with a slow diffusion in the nebula and a fast diffusion in the ISM can explain the HAWC surface brightness profile and the positron excess simultaneously. Inefficient diffusion in the $\gamma$-ray emission region surrounding a middle-aged pulsar maybe a common phenomenon that can be tested by future observation. The implied diffusion coefficient in the ISM is smaller than the one suggested by the standard CR propagation model, but it is fully consistent with the predictions of the spiral arm model. | The origin of cosmic ray (CR) is a long standing problem. CRs are composed of primarily protons and atomic nuclei with a small fraction of positrons and electrons. CR particles with energy below the so-called CR knee ($10^{15}$eV) are confined within our Galaxy by the Galactic magnetic fileds, thus they must have a Galactic origin. CR particles with energy above the CR knee are instead considered to be extra-galactic. The standard picture for the origin of Galactic CR assumes that the CR particles are accelerated at Galactic supernova remnant shocks and then diffuse to Earth. Recently, {\it PAMELA} \citep{Adriani10} and {\it AMS-02} \citep{Aguilar13} detected an excess of positron flux at energies from tens of GeV to hundreds of GeV compared with the prediction of the standard picture \citep{Strong07}. Both dark matter particle interaction \citep[eg.,][]{Ibarra14} and inhomogeneity of astronomical sources \citep[eg.,][]{Hooper09,Shaviv09} have been proposed to explain the observed positron excess. \cite{Hooper09} attribute the positron excess purely to the pulsars in our Galaxy as they are good sources of positron. \cite{Shaviv09} show that the positron excess below $\sim 300$GeV can also be explained by a spiral arm model, in which the supernova rate is higher in the spiral arm region than in the inter-arm region due to the higher concentration of supernova remnants. In the spiral arm model, the pulsar contribution is important to the positron excess at $\gtrsim 300$GeV, which requires nearby pulsars located within a few hundreds of pc from the Earth. It worth noting that the spiral arm model can not only explain the positron excess but also other CR anomaly like the rising spectrum of CR below 1GeV, the boron to carbon ratio \citep{Ben14} and sub-Fe/Fe ratio \citep{Ben16}. Therefore, unless specifically noted, in the following discussion we focus on the positron excess $\gtrsim 300$GeV. Two promising candidate pulsars for the positron excess above $\sim 300$GeV are Geminga and PSR B0656+14 with distances of 250pc and 288pc respectively \citep{Brisken03,Verbiest12}. Energetic electron-positron pairs that escape from the pulsar produce inverse Compton (IC) and synchrotron $\gamma$-ray through interaction with radiation and magnetic field in the vicinity of the pulsar. Extended emission regions in the direction of Geminga and PSR B0656+14 pulsar \citep{Abe17} have been revealed recently by {\it HAWC} between 1 and 50TeV. By modeling the TeV emission as IC emission from electron-positron pairs diffusing in the ISM, the {\it HAWC} collaboration derives a diffusion coefficient of $D\approx (4.5\pm 1.2) \times 10^{27}\rm cm^2/s$ at $100$TeV using a joint fit of the surface brightness profiles of Geminga and PSR B0656+14. This fitted $D$ value in the $\gamma$-ray emisison region is much smaller than the standard value of CR diffusion coefficient in the ISM, which is $D_{ISM}\sim 10^{28}\rm cm^2/s$ at 1GeV with a rigidity dependence of 0.3-0.6 \citep{Strong07}. Assuming such a small $D$ in the ISM between the pulsar and the Earth, \cite{Abe17} argue that the positron flux from Geminga pulsar contributes only a few percents of the observed positron excess, while the positron flux from PSR B0656+14 is negligible. In this work, we propose that the $\gamma$-ray emission detected by {\it HAWC} is originated from the relic pulsar wind nebula (PWN). Then the $HAWC$ surface brightness profile and the positron excess can be explained simultaneously by a two zone diffusion model with a slow diffusion in the relic PWN and a fast diffusion in the ISM. In this paper, we focus on Geminga pulsar and its PWN, but note that PSR B0656+14 is also potentially important for the positron excess. As according to the X-ray morphology PSR B0656+14 is possibly receding from us at a velocity $\gtrsim 400 \rm km/s$ \citep{Birzan16}. Since the pairs flux reaching Earth around 1TeV have been injected at an early phase when the pulsar was closer to us, the positron flux from PSR B0656+14 is expected to be enhanced after we take the recession into account. In section \ref{sec:observation}, we describe the observations. Our model for the Geminga pulsar and nebula is presented in section \ref{sec:model}. In section \ref{sec:IC_emission}, we compare the calculated IC emission of the Geminga nebula with the $\gamma$-ray data. In section \ref{sec:spatial_distribution}, we compare the two zone diffusion model results with the $HAWC$ surface brightness profile and the positron excess data. We discuss the implication of our results in section \ref{sec:discussion}. | {\label{sec:discussion}} We propose that the $\gamma$-ray emission detected by $HAWC$ and $MILAGRO$ in the direction of Geminga originated from the relic PWN surrounding the pulsar. We developed a two zone diffusion model with a slow diffusion $D_1$ in the PWN and a fast diffusion $D_2$ in the ISM. This model would explain both the surface brightness profile and the positron excess $\gtrsim 300$GeV. With $D_1/D_2 \sim 0.05$, the Geminga pulsar can supply a significant fraction of the positron excess above $\sim 300$GeV. The diffusion coefficient in the ISM is $D_{ISM}\sim 2\times 10^{27}\rm cm^2/s (E/1GeV)^{1/3}$. This value is one order of magnitude smaller than the standard value but more consistent with the low value required in the spiral arm model for the CR propagation \citep[e.g.,][]{Shaviv09}. There are several factors that can affect the $\gamma$-ray emission and the surface brightness profile of the Geminga PWN and the positron flux at Earth. The {\it MAGIC} upper limit implies either a broken power law injection spectrum or a single power law but with a larger initial period $P_0$. This can be tested by future multi-wavelength observation. We assume that the particle acceleration efficiency $\eta $ is a constant. If instead $\eta$ gradually decreases with time, then our calculation underestimates the positron flux from the Geminga pulsar. This helps to relieve the discrepancy between the HAWC detection \citep{Abe17} and the idea that the Geminga pulsar is a main source of the positron excess. The proper motion of the Geminga pulsar was studied with a single zone model. We found that a bow shock nebula morphology likely has appeared in the GeV emission. This can be tested by future $MAGIC$ observation. The main caveat of our model is that we neglect the dynamical evolution of the PWN and focus only on the time evolution of the pulsar spin-down power. In other words, we investigate a model with a dynamical pulsar and a static nebula with time independent diffusion coefficient and magnetic field. The peak of the positron flux on Earth corresponds to the pairs injected at an early phase of the Geminga pulsar, when the pulsar spin down luminosity was still large. Therefore, the understanding of the diffusion of TeV pairs in young PWNe, like the Crab, is crucial for explaining the observed positron excess at $\gtrsim 300$GeV. The slow diffusion revealed by \cite{Abe17} in the $\gamma$-ray emission region instead characterizes the properties of the relic nebula at the late phase and is likely irrelevant. In young PWN like the Crab with an age of a few thousand years, the diffusion coefficient of TeV positrons is found to be $\sim 5\times 10^{26}\rm cm^{2}/s$ through a spectral index fitting of the synchrotron emission\footnote{Note we extend the diffusion coefficient provided in Table 3 of \cite{TC12} to 1TeV with a Komogorov type energy dependence.} \citep{TC12}. The corresponding escape time is estimated to be \begin{equation} t_{esc}\sim \frac{R^2}{4D_{PWN}} \sim 150 \rm yr\,\left(\frac{R}{1\rm pc}\right)^2 \left(\frac{5\times 10^{26} \rm cm^2/s}{D_{PWN}}\right). \end{equation} In the early phase of the pulsar evolution, the nebula is small and $t_{esc}$ is much smaller than the age of the Geminga pulsar $t_{age}$. The diffusion time of TeV positrons in the ISM is approximately the age of the pulsar, i.e., $t_{age}=d^2/4D_{ISM}$. This implies \begin{equation} D_{ISM} \sim \frac{d^2}{4t_{age}}\sim 1.5\times 10^{28}{\rm cm^2/s}\,\left(\frac{d}{250 \rm pc}\right)^2 \left(\frac{320 \rm kyr}{t_{age}}\right) \end{equation} at 1TeV, where $d$ is the distance of the pulsar. If we assume a Komogorov type turbulence, then the diffusion coefficient in the ISM becomes $1.5\times 10^{27}\rm cm^2/s (E/1GeV)^{1/3}$ which again is more consistent with the value required by the spiral arm model. Based on the above discussion, the time evolution of PWNe can affect the positron flux on Earth. This will be addressed in future work. In summary, by considering a physically motivated PWN model for the $\gamma$-ray observation of $HAWC$ and $MILAGRO$ we have shown that the Geminga pulsar is a good candidate to be source of a significant fraction of the observed positron excess at $\gtrsim 300$GeV . | 18 | 8 | 1808.02445 |
1808 | 1808.09082_arXiv.txt | We carry out a comprehensive study of HI 21 cm line observations and $^{13}$CO line observations of 21 supernova remnants (SNRs). The aim of the study is to search for HI absorption features to obtain kinematic distances in a consistent manner. The 21 SNRs are in the region of sky covered by the Very Large Array Galactic Plane Survey (HI 21 cm observations) and Galactic Ring Survey ($^{13}$CO line observations). We obtain revised distances for 10 SNRs based on new evidence in the HI and $^{13}$CO observations. We revise distances for the other 11 SNRs based on an updated rotation curve and new error analysis. The mean change in distance for the 21 SNRs is $\simeq25\%$, i.e. change of 1.5 kpc compared to a mean distance for the sample of 6.4 kpc. This has a significant impact on interpretation of the physical state of these SNRs. For example, using a Sedov model, age and explosion energy scale as the square of distance, and inferred ISM density scales as distance. | \label{sec:intro} \indent Supernova remnants (SNRs) play an important role in determining the state of the interstellar medium of the Galaxy, as described in, e.g., the reviews by \cite{2005Cox} and \cite{2001Ferriere}. However, in order to understand SNRs, reliable % distances are necessary. Other objects, including HII regions and molecular clouds, are often associated with SNRs. To determine whether any given association is real, the SNR distance is needed. \\ \indent One common method to find the distance to SNRs is to obtain kinematic distance from analysis of HI absorption spectra. This method has yielded reliable distance estimations, but it has been mainly applied to brighter SNRs. The HI absorption spectra of fainter SNRs are usually noisy. This is caused by real fluctuations, i.e. HI emission occurring at many random positions and velocities in both source and background regions used to create the absorption spectrum. For bright SNRs, these fluctuations are small compared to the absorption signal but for faint SNRs they are comparable. For the fluctuations in HI emission, the HI and continuum images are poorly correlated, whereas for real absorption there is a clear match between the absorption signal (decreased HI intensity) and the continuum emission. Thus careful investigation of the HI channel images is essential to assess the reality of any features which are seen in the HI absorption spectrum. Here we follow the analysis methods described by \cite{Leahy2010} and \cite{2017Ranasinghe}. \\ \indent The sample of SNRs that we consider are those that cover the region of the VLA (Very Large Array) Galactic Plane Survey (VGPS) \citep{Stil}. There are 59 SNRs in the sky area covered by the VGPS survey. Of these, we have previously studied 10 SNRs without previous distance determinations, including SNRs G31.9+0.0 and G54.4‑0.3 \citep{2017Ranasinghe}, 4 SNRs with new molecular cloud associations \citep{2017RanasingheLeahy}, and 4 SNRs without molecular cloud associations \citep{2017RLT}. In this work we analyze data for SNRs with previously published distances and for which the HI line data quality is high enough for analysis. This results in 21 SNRs in our sample presented here. For 10 of these SNRs, we revise previous distances based on new evidence from the HI and $^{13}$CO data. For the remaining 11 SNRs, we confirm published kinematic velocities and revise distances using an updated error analysis and a more recent Galactic rotation curve. In Section \ref{sec:DA}, we present a brief description of the data, construction of HI absorption spectra and kinematic distances. The details for each of the 10 SNRs with new evidence are given in Section \ref{sec:results}. The discussion and the final table summarizing our work for all 21 SNRs is described in Section \ref{sec:dissummary}. | \label{sec:dissummary} We have analyzed HI 21 cm line observations and $^{13}$CO line observations of 21 supernova remnants (SNRs) which are located in the sky area of the VGPS survey. The Galactic rotation curve we use for that area of the Galaxy is the URC rotation curve with \cite{Reid2014} parameters. We examined 1420MHz continuum emission, HI emission and absorption spectra, $^{13}$CO emission spectra and the HI line $^{13}$CO line channel maps. We obtain new observational evidence for 10 SNRs and thus revise their distances to be different than previously published values. For the other 11 SNRs, we confirm the kinematic velocities but use an updated error analysis and the updated rotation curve to revise distances. The uncertainties in $V_r$ and the resulting distance are discussed in \cite{2017Ranasinghe}. Table \ref{tab:table1} presents our results for the 21 SNRs. The literature distances are presented with a note on the method, whether by HI absorption, association with molecular cloud or by the $\Sigma-D$ relation. The various assumed values of radial velocity for the SNR and parameters of the rotation curve ($R_0$ and $V_0$) used in previous distances are listed in columns 4 to 6. Our current most likely radial velocity, $V_r$ is listed in column 8. Column 9 notes whether the SNR is at the near side or far side of the tangent point, or at the tangent point. The new distance is listed in column 10. Column 11 notes whether there is an association of the SNR with a molecular cloud. We note that previous work has used a wide range of Galactic rotation curves, usually assuming $V(R)$ is constant. In some cases the assumed rotation curve is not quoted. The mean and standard deviation for the quoted $R_{0}$ values is 8.25 kpc and 0.56 kpc, and for the $V-{0}$ values is 223 km s$ ^{-1}$ and 12 km s$ ^{-1}$. This range in $R_{0}$ and $V_{0}$ can by itself can lead to a range of distances. For many cases our radial velocities are similar to those in the literature. They agree within 5 km s$ ^{-1}$ for 9 cases, are very different for 3 cases, and for the remaining cases there is no well defined value in the literature. Next we discuss the improvement to the distances to the SNRs, in addition to the fact that we use a consistent rotation curve here, rather than the different rotation curves used for each SNR in the literature. The distance has significantly changed (more than 2$\sigma$), based on our consistently derived errors, for 9 SNRs G$18.1-0.1$, G$18.8 +0.3$, G$20.0 -0.2$, G$23.3 -0.3$, G$23.6 +0.3$, G$27.4 +0.0$, G$33.6 +0.1$, G$39.2 -0.3$ and G$54.1 +0.3$. For 8 SNRs the old and new distances agree within 2$\sigma$: SNRs G$18.6 -0.2$, G$21.5 -0.9$, G$22.7 -0.2$, G$32.8 -0.1$, G$34.7 -0.4$, G$35.6 -0.4$, G$41.1 -0.3$ and G$46.8 -0.3$. For 4 SNRs the previous distance was very poorly known G$21.8 -0.6$, G$29.7 -0.3$, G$43.3 -0.2$ and G$49.2 -0.7$. Including the first two groups of 21 SNRs, the mean change in distance is 1.5 kpc. This is quite significant, considering the distance range is 3 to 13.8 kpc with mean 6.4 kpc. I.e. on average the distance improvement is $23.4\%$. The largest change in distance is for G$20.2-0.2$ which went from an old estimate of 4.5 kpc to a new value of 11.2 kpc. The improved distances have a significant effect on the physical interpretation of supernova remnants. For example, we use the Sedov model equations as given in \citet{1972Cox}. The inferred shock radius $R_s$ of an SNR depends linearly on distance, $d$. As described in \citet{2017LeahyWilliams}, the ISM density $n_{0}$ depends on the emission measure, $EM$, as $\sqrt{EM}$, and $EM$ depends on distance as $d^2$. Explosion energy $E_{0}$ depends on $EM$, thus on $d^2$. For the Sedov model, the important quantity is $E_{0}/n_{0}$, which depends linearly on distance. SNR age in the Sedov model depends on $R_s^{5/2} (n_{0}/E_{0})^{1/2}$, thus depends on $d^2$. For more sophisticated models, the scaling with distances are not as simple as this, but the change in results is similar. In summary, we have obtained distances to a significant number of SNRs using a consistent method and rotation curve. Follow-up work will investigate the implications for the ages and evolutionary states of these SNRs by incorporating other data. For example, with X-ray spectra the shock temperature and the emission measure can be determined. These can be used to estimate physical properties of SNRs using explosion models. % Easy-to-use SNR explosion models can be based on analytical fits to numerical explosion calculations. Such models have been presented by \cite{2017LeahyWilliams}, based on the \cite{1999TM} models. These models have been applied to the set of LMC SNRS by \cite{2017Leahy} to obtain important properties of the SNR population, including finding a log-normal distribution of explosion energies, the LMC SNR birthrate, and the distribution of ISM densities around SNRs in the LMC. Similar methods will be applied to learn about the properties of Galactic SNRs. \begin{deluxetable*}{cccccccccccc} \label{tab:table1} \tabletypesize{\scriptsize } \tablecaption{Distances to supernova remnants} \tablenum{1} \tablehead{ \colhead{\#} & \colhead{Source} & \colhead{} & \colhead{} & \colhead{Literature} & \colhead{} & \colhead{} & { } & \colhead{} & \colhead{New results} & \colhead{} & \colhead{} \\ \cline{3-7} \cline{9-11} \colhead{} & \colhead{} & \colhead{Dist\tablenotemark{a}} & \colhead{V$_{r}$} & \colhead{R$_{0}$} & \colhead{V$_{0}$}& \colhead{Refs} & &\colhead{V$_{r} $} & \colhead{KDAR\tablenotemark{b}} & \colhead{Dist} & \colhead{$^{13}$CO\tablenotemark{c}} \\ \colhead{} & \colhead{} & \colhead{(kpc)} & \colhead{(km s$^{-1}$)} & \colhead{(kpc)} & \colhead{(km s$^{-1}$)}& \colhead{ } & & \colhead{(km s$^{-1}$)} & \colhead{} & \colhead{(kpc)} & \colhead{} } \startdata 01 & G18.1 -0.1 & $5.6$\tablenotemark{H} & 100 & 8.5 & 210 & 1 & & 103.74 & N & $6.4 \pm 0.2$ & Possible\\ 02 & G18.6 -0.2 & $4.6 \pm 0.6$\tablenotemark{H} & 62 & 8.5 & 220 & 2 & & 62.84 & N & $4.4 \pm 0.2$& No\\ 03 & G18.8 +0.3 & 12.1\tablenotemark{H}\tablenotemark{M} & 20 & 7.6 & 214 & 3& & 21.35 & F & $13.8 \pm 0.4 $ & Yes\\ 04 & G20.0 -0.2 & 4.5\tablenotemark{M} & 66 & 8.5 & 220 & 4 & & 66.40 & F & $11.2 \pm 0.3 $ & Yes\\ 05 & G21.5 -0.9 & $4.7 \pm 0.4$\tablenotemark{H} & 68 & 8 & 220 & 5 & & 67.79 & N & $4.4 \pm 0.2 $ & No\\ 06 & G21.8 -0.6 & 5.5\tablenotemark{H}\tablenotemark{M} & 86 & 8.0 & 220 & 6A & & 93.35 & N &$ 5.6 \pm 0.2 $ & Yes \\ & & 5.2\tablenotemark{H}\tablenotemark{M} & 85 & 8.0 & 220 & 6B & & & & \\ 07 & G22.7 -0.2 & $4.4 \pm 0.4$\tablenotemark{M} & 77 & 8.31 & 241 & 7 & & 76.63 & N & $4.7 \pm 0.2 $ & Yes\\ 08 & G23.3 -0.3 & $3.9 - 4.5$\tablenotemark{H}\tablenotemark{M} & 66 - 80 & 7.6 & 214 & 8 & & 78.51 & N & $4.8 \pm 0.2 $ & Yes\\ 09 & G23.6 +0.3 & 6.9\tablenotemark{M} \tablenotemark{S} & \nodata & \nodata & \nodata & 9 & & $ 99.95$ & N & $5.9 \pm 0.2$ & Possible\\ 10 & G27.4 +0.0 & $7.5 - 9.8$ \tablenotemark{H} & $ V_{TP} - 84 $ & 8.5 & 220 & 10 & & 99.95 & N & $ 5.8 \pm 0.3$ & No\\ 11 & G29.7 -0.3 & $5.1 - 7.5$\tablenotemark{H}\tablenotemark{M} & 95 - 102 & 7.6 & 220 & 11A & & 95.00 & N & $ 5.6 \pm 0.3 $ & Yes\\ & & 10.6\tablenotemark{M} & 54 & 8.0 & 220 & 11B & & & & &\\ 12 & G32.8 -0.1 & 4.8\tablenotemark{H}\tablenotemark{M} & $\sim81$ & 8.0 & 220 & 12 & & 81.81 & N &$ 4.8 \pm 0.3 $ & Yes\\ 13 & G33.6 +0.1 & 7.1\tablenotemark{M}\tablenotemark{S} & \nodata & \nodata & \nodata & 9 & & 57.90 & N & $ 3.5 \pm 0.3 $ & No \\ 14 & G34.7 -0.4 & $2.5 - 2.6$\tablenotemark{H} & 42 & 8.5 & 220 & 14 & & 50.48 & N & $3.0 \pm 0.3 $ & No\\ 15 & G35.6 -0.4 & $3.6 \pm 0.4 $\tablenotemark{H}\tablenotemark{M} & $\sim61$ & 8.5 & 220 & 15 & & 63.67 & N & $ 3.8 \pm 0.3 $ & Possible\\ 16 & G39.2 -0.3 & 6.2\tablenotemark{H}\tablenotemark{M} & 84 & 8.0 & 220 & 16 & & 69.39 & F & $8.5 \pm 0.5 $ & Yes\\ 17 & G41.1 -0.3 & $ 8 - 9.7$\tablenotemark{H} & \nodata & 8.33 & 218 & 17 & & $ V_{TP} - 63.01 $ & F & $ 8.5 \pm 0.5$ & No\\ 18 & G43.3 -0.2 & $8 - 11$\tablenotemark{H} & \nodata & \nodata & \nodata & 18 & & 12.55 & F & $11.3 \pm 0.4 $ & No\\ 19 & G46.8 -0.3 & $6.8 - 8.6$\tablenotemark{H} & $ V_{TP} - 59$ & 10.0 & 250 & 19 & & $V_{TP} - 0$ & TP $-$ F & $5.7 \pm 0.9 - 11.4 \pm 0.5$ & No\\ 20 & G49.2 -0.7 & 4.3\tablenotemark{H} & 70.7 & 8.4 & 254 & 20A & & $V_{TP}$ & TP & $5.4 \pm 0.6 $ & No\\ & & 6\tablenotemark{M} & $ V_{TP}$ & 8.5 & \nodata & 20B & & & & & \\ 21 & G54.1 +0.3 & $5.6 -7.2$\tablenotemark{H}\tablenotemark{M} & $53\pm12$ & 7.6 & 220 & 21 & & 53.66 & TP & $4.9 \pm 0.8 $ & Yes\\ \enddata \tablecomments{ \tablenotetext{a}{Literature distance estimation method - Superscript H: HI absorption, M: Molecular cloud association/interaction and S: $\Sigma$-D relation.} \tablenotetext{b}{KDAR- Kinematic Distance Ambiguity Resolution, indicating whether the SNR is at the near (N), far (F) or tangent point (TP) distance.} \tablenotetext{c} {Associated with $^{13}$CO.} \tablerefs { (1) \cite{Leahy2014}, (2) \cite{Johanson}, (3) \cite{Tian2007}, (4) \cite{Petriella2013}, (5) \cite{Tian2008}, (6A) \cite{Tian2008} , (6B) \cite{Zhou2009}, (7) \cite{Su2014}, (8) \cite{Leahy2008}, (9) \cite{Kilpatrick2016}, (10) \cite{TianLeahy2008}, (11A) \cite{LeahyTian2008}, (11B) \cite{Su2009}, (12) \cite{ZhouChen2011}, (14) \cite{Cox1999}, (15) \cite{Zhu2013}, (16) \cite{Su2011}, (17) \cite{LeahyRanasinghe2016}, (18) \cite{Brogan2001}, (19) \cite{Sato1979}, (20A) \cite{TianLeahy2013}, (20B) \cite{Koo1995}, (21) \cite{LeahyTianWang2008}.}} \label{tab:table1} \end{deluxetable*} | 18 | 8 | 1808.09082 |
1808 | 1808.10857_arXiv.txt | Motivated by the recent discovery of rare Enormous Lyman-Alpha Nebulae (ELAN) around $z\sim2$ quasars, we have initiated a long-term observational campaign with the MUSE/VLT instrument to directly uncover the astrophysics of the gas around quasars. We present here the first 61 targets of our effort under the acronym QSO MUSEUM ({\it Q}uasar {\it S}napshot {\it O}bservations with {\it MU}se: {\it S}earch for {\it E}xtended {\it U}ltraviolet e{\it M}ission). These quasars are characterized by a median redshfit of $z=3.17$ ($3.03<z<3.46$), absolute $i$ magnitude in the range $-29.67\leq M_i(z=2)\leq-27.03$, and different levels of radio-loudness. This sample unveils diverse specimens of \lya\ nebulosities extending for tens of kiloparsecs around these quasars (on average out to a maximum projected distance of 80~kpc) above a surface brightness SB$>8.8\times10^{-19}$~\unitcgssb ($2\sigma$). Irrespective of the radio-loudness of the targets, the bulk of the extended \lya\ emission is within $R< 50$~kpc, and is characterized by relatively quiescent kinematics, with average velocity dispersions of $\langle \sigma_{\rm Ly\alpha}\rangle < 400$~km~s$^{-1}$. Therefore, the motions within all these \lya\ nebulosities have amplitudes consistent with gravitational motions expected in dark matter halos hosting quasars at these redshifts, possibly reflecting the complexity in propagating a fast wind on large scales. Our current data suggest a combination of photoionization and resonant scattering as powering mechanisms of the \lya\ emission. We discover the first $z\sim3$ ELAN, which confirms a very low probability ($\sim1\%$) of occurrence of such extreme systems at these cosmic epochs. Finally, we discuss the redshift evolution currently seen in extended \lya\ emission around radio-quiet quasars from $z\sim3$ to $z\sim2$, concluding that it is possibly linked to a decrease of cool gas mass within the quasars' CGM from $z\sim3$ to $z\sim2$, and thus to the balance of cool vs hot media. Overall, QSO MUSEUM opens the path to statistical and homogeneous surveys targeting the gas phases in quasars' halos along cosmic times. | \label{sec:intro} In the current paradigm of structure formation, most of the baryons at high redshift ($z\gtrsim1.5$; \citealt{Meiksin2009} and references therein) are distributed in a web of diffuse filamentary structures in which galaxies form and evolve. The complex interplay between this rich reservoir of gas and the galaxies themselves is still a matter of investigation, e.g. importance and strength of feedback from active galactic nuclei (AGN), existence of a cold mode of accretion onto galaxies, astrophysics of galactic outflows, build up of super-massive black holes in short timescales, angular momentum evolution (e.g., \citealt{Dekel2009,Brooks2009,DiMatteo2012,Shen2013,Dubois2013,Woods2014,Feng2014,AnglesAlcazar2014,Nelson2016, Stewart2016, Obreja2018}). Intergalactic and circumgalactic large-scale structures thus encode fundamental information to test our current galaxy evolution theories. So far, whether one focuses on large scales, i.e. on the intergalactic medium (IGM), or on smaller scales, ie. on the hundreds of kiloparsecs close to galaxies often referred to as the circumgalactic medium (CGM), the strongest constraints on the physical properties of these diffuse gas phases are obtained by analyzing absorption features along background sightlines (e.g., \citealt{Croft2002, Bergeron2002, Hennawi2006, QPQ2, Tumlinson2011, Farina2013b, Rudie2013, Turner2014, Farina2014, QPQ5, Lee2014}). Direct imaging of the same gas phases greatly complement these absorption studies, allowing a spatial, morphological, physical, and kinematical characterisation, which is simply not possible with sparse one-dimensional information inherent to the absorption technique. Yet, predicted to be diffuse, the IGM and CGM is expected to be hard to detect in emission (SB$_{\rm Ly\alpha}\sim10^{-20}$~erg~s$^{-1}$~cm$^{-2}$~arcsec$^{-2}$; e.g., \citealt{Lowenthal1990, GW96, Bunker1998, Rauch2008}). Nevertheless, recently, the deployment of new advanced integral field spectrographs on 10-m class telescopes, i.e. the Multi-Unit Spectroscopic Explorer (MUSE; \citealt{Bacon2010}) on the ESO/VLT and the Keck Cosmic Web Imager (KCWI; \citealt{Morrissey2012}), opened up the possibility of routinely performing such an experiment by targeting very low levels of surface brightness (SB). Indeed, the latest observations are able to directly studying in emission at least the CGM of galaxies with reasonable observational times ($\sim$tens of hours). In particular, these observations usually show extended \lya\ emission on scales of tens of kiloparsecs around the targeted $z\gtrsim1.7$ galaxies (\citealt{Wisotzki2016, Leclercq2017}), opening up a new parameter space for the study of the CGM gas phases, and ultimately of galaxy evolution. Even before the advent of the current new instrumentations, it has been shown that such extended \lya\ emission on CGM scales is more easily detected around high-redshift quasars or active galactic nuclei (AGN). In this case, the \lya\ emission was detected around most of the objects ($50-70$\%) on $R<50$~kpc (e.g., \citealt{HuCowie1987,heckman91a,heckman91b,Moller2000b,Weidinger05,Christensen2006,North2012,qpq4,fab+16}), extending up to $R>100$~kpc in few rare bright cases known as Enormous Lyman-Alpha Nebulae (ELAN; \citealt{cantalupo14, hennawi+15,Cai2016, fab+2018,Cai2018}). The ELAN are indeed structures characterized by high observed surface brightnesses (SB$_{\rm Ly\alpha}> 10^{-17}$~erg~s$^{-1}$~cm$^{-2}$~arcsec$^{-2}$) spanning continuously out to hundreds of kiloparsecs, and thus resulting in luminosities $L_{\rm Ly\alpha}> 10^{44}$~erg~s$^{-1}$ (\citealt{Cai2016}). The comparison between all these pioneering observations however is hampered by (i) the heterogeneity of the technique used (longslit spectroscopy, narrow-band imaging, integral-field spectroscopy), (ii) by the lower and different sensitivities inherent to the previous observational instruments (SB$_{\rm limit}> {\rm few}\times 10^{-18}$~\unitcgssb for some works and SB$_{\rm limit}\gtrsim 10^{-17}$~\unitcgssb for others), (iii) by uncertainties in the redshift of the targeted quasars, and (iv) by the difficulties in achieving a clean removal of the unresolved emission from the quasar, which can easily outshine the faint large-scale emission (e.g., \citealt{Moller2000a}). The advent of the aforementioned new sensitive IFU intruments together with the discoveries of the ELAN have motivated intense research on quasar halos, and have resulted in frequent new clear detections of the cool CGM gas. Overall, mainly due to the new sensitivities achieved ($<10^{-18}$~\unitcgssb), recent studies now routinely show detections on $R\sim50$~kpc around $z\gtrsim3$ quasars (e.g., \citealt{Husband2015, Borisova2016, Fumagalli2016, Ginolfi2018}), i.e. the redshift range for which \lya\ is visible with MUSE. In very rare cases these newly discovered \lya\ nebulosities at $z\sim3$ match the observed SB and extent of the ELAN previously unveiled at $z\sim2$ (\citealt{fab+2018}). In the presence of a quasar, the detected \lya\ emission on halo scales has been usually explained as (i) recombination radiation produced after quasar photoionization ({\it a.k.a.} fluorescence; e.g.,\citealt{Weidinger04, Weidinger05, qpq4, cantalupo14, fab+15b, Borisova2016}), and/or (ii) \lya\ photons resonantly scattered within the gas distribution surrounding the quasar (e.g., \citealt{qpq4, cantalupo14, Borisova2016}), and/or (iii) recombination radiation produced after photoionization by several sources, e.g. quasar and active companions (e.g., \citealt{Husband2015, fab+2018, Husemann2018}). A clear determination of the contribution from the different mechanisms cannot be easily achieved by using only the information enclosed in the \lya\ emission. Indeed \lya\ photons within the CGM gas are likely affected by resonant scattering (e.g., \citealt{Dijkstra2017,Gronke2017} and references therein), which could lead to strong modification of the \lya\ spectral shape (e.g., \citealt{Dijkstra2017} and references therein), and the relative strength of the \lya\ line with respect to other diagnostics (e.g., \citealt{Neufeld_1990}). Notwithstanding these challenges, observations of the extended \lya\ emission -- together with the aforementioned absorption studies (e.g., \citealt{qpq3,QPQ5,QPQ7,qpq9}) -- currently paint a scenario in which quasar's halos are hosting a large reservoir of cool ($T\sim10^4$~K) gas. This gas is possibly tracing a complex set of astrophysical processes: gas/substructures infalling onto the central quasar (e.g., \citealt{Hu1991,Weidinger04,fab+2018}); strong turbulences or outflows (e.g., \citealt{Ginolfi2018}); interactions between substructures (e.g., \citealt{Hu1991,Husband2015,fab+2018,Husemann2018}); large-scale filaments (\citealt{cantalupo14}). Further, in the case of ELANe there are evidences that the \lya\ emitting gas on hundreds of kpc is composed by a population of cool and dense (volume density $n_{\rm H}> 1$~cm$^{-3}$) clumps (\citealt{cantalupo14}). Indeed, if one assumes an ELAN to be powered by the radiation from the associated brightest quasar, the high levels of \lya\ emission together with the current stringent limits on the \ion{He}{ii}/\lya\ ratio can be matched by photoionization models only if very high densities ($n_{\rm H}\gtrsim3$~cm$^{-3}$; thus low ionization parameters log$U\lesssim-2$), and low column densities ($N_{\rm H}\lesssim10^{20}$~cm$^{-2}$) are used (\citealt{fab+15b,fab+2018}). The same framework thus requires the clumps to have compact sizes $R \equiv N_{\rm H}/n_{\rm H}\lesssim20$~pc (\citealt{fab+15b,hennawi+15,fab+2018}). Current simulations are not able to achieve the resolutions needed to resolve such clumps in the CGM of galaxies, predicting low densities ($n_{\rm H}\sim10^{-2}-10^{-3}$~cm$^{-3}$) for such a medium (see discussions in \citealt{cantalupo14,hennawi+15}). This tension between observations and simulations motivated new research on hydrodinamical instabilities (e.g., \citealt{Mandelker2016}). In particular, very high resolution idealised hydrodinamical (\citealt{McCourt2018}) and magneto-hydrodinamical (\citealt{Ji2018}) simulations have shown that a mist of cool gas clouds could form and survive in the halo of galaxies, possibly explaining the high densities required by the observed levels of \lya\ emission around quasars. Also, the effects of the quasar activity (e.g. radiation, outflows) on the diffuse gas phases and on the production of \lya\ photons is not fully understood, and could lead to effects on the morphology and physical properties of the surrounding gas distribution. Extreme examples of what a powerful AGN can do are the high-redshift radio galaxies (HzRGs), whose powerfull radio jet and UV radiation clearly alter the surrounding gas (e.g., \citealt{Nesvadba2017,Silva2018}). These objects show extended \lya\ emission with active kinematics (FWHM$>1000$~km~s$^{-1}$) associated with the radio emission, but a more extended (100 kpc scales) quiescent (FWHM$<700$~km~s$^{-1}$) gas phase consistent with gravitational motions (e.g., \citealt{vanOjik1997, VillarMartin2002, Humphrey2007, VillarM2007}), and/or large-scale structures (\citealt{Vernet2017}). This quiescent large-scale gas phase is similar to what is seen around quasars (\citealt{Husband2015, Borisova2016, fab+2018}). Understanding the effects of AGN disruption mechanisms is fundamental given that these are invoked in cosmological simulations to modify the gas properties around and within massive galaxies to match observational constraints, e.g. star formation, halo mass versus stellar mass relation (e.g., \citealt{Silk1998, Sijacki2007, Booth2009, Richardson2016}). In this context, we have started to survey the $z\sim3$ quasar population with the main aim to characterise (i) the physical properties of the CGM/IGM in emission associated with such expected massive dark-matter halos (M$_{\rm DM}\sim10^{12.5}$~M$_{\odot}$, \citealt{white12,Trainor2012}), and (ii) the frequency of detection of ELAN [${\rm SB}_{\rm Ly\alpha}\gtrsim 10^{-17}$\unitcgssb\ out to 100~kpc]. In this paper we focus on presenting the observations at the \lya\ transition for the first 61 targeted quasars under the acronym QSO MUSEUM ({\it Q}uasar {\it S}napshot {\it O}bservations with {\it MU}se: {\it S}earch for {\it E}xtended {\it U}ltraviolet e{\it M}ission). This work is part of an on-going multi-technique and multi-wavelength effort to unveil any dependence on the nature of each system (e.g., geometry, environment, radio activity, luminosity) in the detection and properties of extended gaseous structures. This work is structured as follows. In Section~\ref{sec:obs}, we describe our observations and data reduction. In Section~\ref{sec:analysis} we explain our analysis procedures, and the current uncertainties on the systemic redshifts of the quasars in our sample. In Section~\ref{sec:results}, we present the observational results. In particular, we reveal the discovery of extended \lya\ emisssion around our targeted quasars, and show their diverse morphologies, individual radial profiles, stacked profiles, average covering factor for the \lya\ emission, compact line emitters associated with the targeted systems, kinematics, and spectral shape of the \lya\ emission. In Section~\ref{sec:disc} we discuss our results in light of the current statistics for extended \lya\ emission around quasars, and the usually invoked powering mechanisms. Finally, Section~\ref{sec:summ} summarises our findings. Throughout this paper, we adopt the cosmological parameters $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\Omega_M =0.3$ and $\Omega_{\Lambda}=0.7$. In this cosmology, 1\arcsec\ corresponds to about 7.6 physical kpc at $z=3.17$ (median redshift for our sample). All magnitudes are in the AB system (\citealt{Oke1974}), and all distances are proper, unless otherwise specified. | 18 | 8 | 1808.10857 |
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1808 | 1808.05265_arXiv.txt | We estimate the gravitational wave amplitude as a function of frequency produced during the creation of pulsars from the gravitational collapse of a massive star. The three main quantities needed are the magnitude of the magnetic field producing pulsar kicks, the temperature which determines the velocity of the pulsar and the duration time for the gravitational radiation. | Shortly after the collapse of a massive star often a neutron star is created. Because of the strong magnetic field and high temperature the neutron star has a large velocity and emits electromagnetic radiation, which is why it is called a pulsar. The origin of a pulsars velocity is called a pulsar kick. In research on pulsar kicks\cite{hjk07,kj12} the magnitude of B, the magnetic field, and T, the temperature of the neutron star at the time of the pulsar kick, were estimated. We use these parameters in the present estimate of gravitational wave creation from pulsar creation, for which there are no published previous estimates. In a study of gravitaional waves (GW) generated by magnetic fields\cite{kks10} it was found that the GW created during the EWPT could be detected by Lisa Interfer0meter Space Antena (LISA) while the GW created during the QCDPT cannnot be detected by LISA. There is also an article\cite{kk15} on the circular polarization of GW created during the EWPT and QCDPT. Previous research on gravitational radiation from neutron stars\cite{mp10, lbm13} was based in part on gravitational radiation from cosmological turbulance\cite{kmk02}. There have been studies of gravity waves generated by cosmological phase transitions, the Electroweak Phase Transition (EWPT) and the Quantum Chromodynamic Phase Transition (QCDPT). In a study of gravity waves generated by magnetic fields\cite{kks10} it was found that the gravitational waves generated during the EWPT might be detected by the Laser Interoferometer Space Antenna (LISA) while gravitational waves generated during the QCDPI canot be detected by LISA. With current theory the EWPT and QCDPT are first order phase thansitions, so bubbles of the new universe form within the older universe during the phase transition. A study of gravitational waves created by bubble collisions during the EWPT and QCDPT estimated the degree of circular polarization of the gravitational waves\cite{kk15}. The formalism used Refs\cite{kks10,kk15} is not suitable for estimates of the creation of gravitational radiation from pulsar creation. However, the research in Refs\cite{mp10,lbm13,kmk02} was based in part on gravitational radiation from cosmological turbulance which we use in our present research. \newpage | Our conclusion, based on our results shown in Figure 2, is that the gravitational wave amplitude produced via pulsar creation is smaller than that produced via the Cosmological Electroweak Phase Transition\cite{kks10} and is too small to be measured by the Lisa Interferometer Space Antenna (LISA)\cite{cornish01} or any other gravitational wave detector in the near future. \newpage \Large {\bf Acknowledgements:} \vspace{2mm} \normalsize {\bf Authors thank Prof. Tina Kahniashvili for helpful suggestions. Authors Leonard S. Kisslinger and Bijit Singha thank the CMU Physics Depart for support. Author Zhou Li-juan acknowledges the support in part by the National Natural Science Foundation of China (11365002),\\ the National Natural Science Foundation of China (11865005),\\ Guangxi Natural Science Foundation (2015GXNSFAA139012).} \vspace{5mm} | 18 | 8 | 1808.05265 |
1808 | 1808.05115_arXiv.txt | We analyze differences in positions of active galactic nuclei between \Gaia data release~2 and VLBI and compare the significant VLBI-to-\Gaia offsets in more than 1000 objects with their jet directions. \aplc{Remarkably} at least 3/4 of the significant offsets are confirmed to occur downstream or upstream the jet representing a genuine astrophysical effect. \aplc{Introducing redshift and \Gaia color into analysis can help distinguish} between the \aplc{contribution} of the host galaxy, jet, and accretion disk emission. \yykb{We find that strong optical jet emission at least 20-50pc long is required to explain the \Gaia positions located downstream from VLBI ones}. Offsets \aplc{in the upstream direction} of up to 2~mas are at least partly due to the dominant impact of the accretion disk on the \Gaia coordinates and by the effects of parsec-scale radio jet. The host galaxy was found not to play an important role in the detected offsets. BL~Lacertae object and Seyfert~2 galaxies are observationally confirmed to have a relatively weak disk and consequently downstream offsets. The disk \aplc{emission drives upstream offsets in a significant fraction of quasars and Seyfert~1 galaxies when it dominates over the jet in the optical band}. The observed behaviour of the different AGN classes is consistent with the unified scheme assuming varying contribution of the obscuring dusty torus and jet beaming. | Both VLBI and Gaia \citep{r:Gaia} provide positions for various objects, including active galactic nuclei (AGNs) with sub-milliarcsecond accuracy. Comparison of VLBI positions with those from \Gaia data release~1 (GDR1) revealed that 7\,\% of AGNs have statistically significant differences that cannot be explained by random measurement errors \citep{r:gaia1}. Subsequently, we discovered that the directions of VLBI-to-\Gaia position offsets are not isotropic and have a strong concentration downstream and upstream the jet \citep{r:gaia2}. The discovery presented a strong evidence that these position differences are not artifacts of radio or optical data analysis, but are due to the structure of the sources. We suggested in \cite{r:gaia2,r:gaia3} that the presence of strong parsec-scale optical jet \aplc{emission} causes these offsets. We predicted that further improvements in the VLBI and \Gaia accuracy will not reconcile position differences, and even more sources will have accurately measured significant offsets. We note that astrophysical properties of active galactic nuclei at parsec scales are still sometimes not assumed as playing the major role in significant VLBI-\Gaia offsets \cite[e.g.][]{r:Gaia-CRF2}. \cite{2018AJ....155..229F} discussed significant VLBI-\Gaia offsets for 20 sources and attributes the majority of them to binary or extended objects. The \Gaia data release~2 \citep[GDR2,][]{r:GDR2cat} brought a substantial improvement with respect to the GDR1: the number of sources increased by almost 50\,\%, the position accuracy of VLBI-\Gaia matched AGNs improved by a factor of 5, and magnitudes at red and blue passbands became available. This prompted us to update and expand the previous analysis, check our predictions, and study the origin of these offsets in more detail using additional information. \begin{figure*} \centering \includegraphics[width=0.95\textwidth,trim=0cm 0cm 0cm -0.3cm]{drawing.pdf} \caption{A cartoon explaining two opposite VLBI-to-\Gaia offset directions \aplc{with respect to the parsec-scale jet: downstream $\Psi = 0\degr$ and upstream $\Psi = 180\degr$}.} \label{f:diagram} \end{figure*} \begin{figure*} \centering \includegraphics[width=\textwidth,trim=0cm 0.3cm 0cm -0.2cm]{ltr_hists.pdf} \caption{Distribution of the offset direction relative to the jet, $\Psi$, for the whole set of sources and for two subsets filtered by restricting to small angle error estimates.} \label{f:ltr_hists} \end{figure*} Throughout this paper we define and use the offset $\vec{VG}$ as the vector from VLBI to \Gaia coordinates of a source. Its direction is described by the angle $\Psi$ relative to the parsec-scale jet direction, see the cartoon in \autoref{f:diagram}. The \Gaia position is shifted downstream the jet relative to VLBI when $\Psi\approx0\degree$, i.e.\ the optical centroid is farther away from the nucleus. The \Gaia position is shifted upstream the jet relative to VLBI position when $\Psi\approx180\degree$, i.e.\ the optical centroid is closer to the nucleus than the VLBI position. We consider the optical emission of an AGN as coming from three main components: the jet, the accretion disk, and the host galaxy. \Gaia coordinates derived from CCD measurements correspond to the optical centroid of these components. VLBI coordinates derived from correlated visibility data correspond not to the radio-band centroid position, but to the most compact feature. \apla{The difference in the measurement principles implies that even if radio- and optical structure of AGNs are similar and colocated, the coordinates measured by VLBI and \Gaia will still be different \citep{r:gaia3}.} The most compact feature typically coincides with the so-called radio core which is offset by less than 1~mas downstream the jet from its origin due to synchrotron opacity \citep{Kovalev_cs2008,r:Porcas09,r:MOJAVE_cs,r:cs_var}. \apla{Different aspects} of this three-component interpretation are further tested using the difference of \aplc{optical} spectra of the jet, accretion disk, and host galaxy. | \label{s:summary} We confirm using \Gaia Data Release 2 that the \Gaia offsets with respect to VLBI position have predominant directions downstream or upstream the jet. This means the results of \cite{r:gaia2} represent a genuine astrophysical effect and were not caused by some unaccounted errors in DR1. We estimate that for at least $73$\,\% of sources with large offsets which have $\sigma_\mathrm{P.A.} < 15\degr$ the VLBI-\Gaia offsets are dominated by the effect of jets. This is a significant increase from our previous bound of 53\,\% based on \Gaia DR1, and we expect this fraction to increase further as the accuracy of VLBI and \Gaia measurements is continously improving. This means that the fraction of offsets caused by other effects in AGNs may be much lower than our current bound of 27\,\%. We find that the host galaxy emission which may partially be obscured by the dusty torus does not have a dominant effect on the VLBI-\Gaia offsets. We explain VLBI-\Gaia offsets in the downstream direction by the presence of bright and extended optical jets which shift the \Gaia centroid. From the magnitudes of such offsets we infer that optical jets of projected length at least 20-50~pc are quite common among AGNs. VLBI-\Gaia offsets in the upstream direction indicate that the VLBI positions do not coincide with the nuclei confirming predictions by e.g.~\cite{Kovalev_cs2008,r:gaia3}. The VLBI positions are shifted by up to 2~mas or 20~pc projected downstream the jet. Unaccounted structure of radio jets and frequency-dependent synchrotron opacity may contribute to these offsets. However, their contribution is expected typically at a level of 0.2~mas or less. We surmise that either the observed offsets represent the tail of the distribution, or the contribution of source structure and core shift is underestimated, or another, yet unknown effect, causes such large shifts. Note that the typical expected shifts are shorter than 1~mas and can not yet be detected within the errors of VLBI and \Gaia positions. For objects with upstream VLBI-\Gaia offsets the optical position should be very close to the nucleus. Such AGNs are found to be typically located at higher redshifts and are bluer than those with offsets downstream the jet. As the accretion disk mostly emits at UV wavelengths, both higher redshifts and bluer colors indicate that more of the accretion disk flux falls into the optical passband of \textit{Gaia} affecting the measured position. We conclude that in the cases of upstream offsets \Gaia coordinates are mostly determined by the dominant emission of the accretion disk. Examining quasars, BL Lacs and Seyfert galaxies separately we find more support for our claims and get new insights on the classification itself. The unified scheme of active galaxies explains different AGN classes by orientation effects which involve Doppler boosting and obscuration by the dusty torus. It can also satisfactorily explain the found color index -- offset direction diagrams for different AGN classes. Quasars and Seyfert~1 galaxies may shift in both directions depending on relative contributions of their disks and jets. BL Lacs and Seyfert~2 galaxies have relatively weak BLR and accretion disk emission or the central region is obscured, and we find that almost none of them have upstream VLBI-\Gaia offsets. This is a strong evidence that \Gaia coordinates correspond to position of the nucleus due to the strong accretion disk and not just the jet origin being brighter in the optical band. } \vspace*{0pt} | 18 | 8 | 1808.05115 |
1808 | 1808.00730_arXiv.txt | The {\sl AKARI\/} Infrared Astronomical Satellite produced the all-sky survey (AFASS) maps in the far-IR at roughly arc-minute spatial resolution, enabling us to investigate the whole sky in the far-IR for objects having surface brightnesses greater than a few to a couple of dozen MJy\,sr$^{-1}$. While the AFASS maps are absolutely calibrated against large-scale diffuse emission, it was uncertain whether or not an additional flux correction for point sources was necessary. Here, we verify that calibration for point-source photometry in the AFASS maps is proper. With the aperture correction method based on the empirical point-spread-function templates derived directly from the AFASS maps, fluxes in the {\sl AKARI\/} bright source catalogue (BSC) are reproduced. The {\sl AKARI\/} BSC fluxes are also satisfactorily recovered with the 1-$\sigma$ aperture, which is the empirical equivalent of an infinite aperture. These results confirm that in the AFASS maps far-IR photometry can be properly performed by using the aperture correction method for point sources and by summing all pixel values within an appropriately defined aperture of the intended target (i.e., the aperture photometry method) for extended sources. | The {\sl AKARI\/} Infrared Astronomical Satellite ({\sl AKARI\/}; \cite{akari}) is the Japanese infrared space mission launched on 2006 February 21 (UT). The primary mission of {\sl AKARI\/} was to conduct an all-sky survey in the mid- and far-infrared at higher spatial resolutions and over a wider spectral range than its predecessor, the Infrared Astronomy Satellite ({\sl IRAS}; \cite{Neugebauer_1984}). For this purpose, {\sl AKARI\/} was outfitted with a cryogenically-cooled telescope with a diameter of 68.5\,cm and two instruments, the Far-Infrared Surveyor (FIS; \cite{Kawada_2007}) and the Infrared Camera (IRC; \cite{Onaka_2007}), covering a wavelength range of 2 -- 180\,$\micron$. {\sl AKARI\/} carried out its 550-day cryogen mission until it exhausted liquid Helium on 2007 August 26, and continued its post-cryogen mission in the near-infrared until the satellite was finally switched off on 2011 November 24. The FIS instrument covers the wavelength range of 50--180$\,\mu$m with two sets of Ge:Ga arrays, the Short-Wavelength (SW) and Long-Wavelength (LW) detectors in the N60 (50 -- 80\,$\micron$) and WIDE-S (60 -- 110$ \mu$m) bands and the WIDE-L (110 -- 180$\,\mu$m) and N160 (140 -- 180$ \mu$m) bands, respectively \citep{Doi_2002,Fujiwara_2003}. Pixel scales in the detectors were designed to be similar to the diffraction limit of the telescope, at around $30\arcsec$. For the all-sky survey, the natural 100-minute sun-synchronous orbit of {\sl AKARI\/} was used to achieve a scan speed of 3\farcm6\,s$^{-1}$ during the survey period from 2006 April to 2007 August, allowing to cover over 96 per cent of the sky with two or more scans, and in fact resulted in 99 per cent sky coverage \citep{Doi_2015,Takita_2015}. The primary absolute surface-brightness calibration of the FIS instrument was done through (1) pre-launch laboratory measurements of a blackbody source which indicated a 5\,$\%$ accuracy, and (2) on-orbit measurements of infrared cirrus regions without significant small-scale structures and comparisons of the results between FIS and the DIRBE instrument on-board the {\sl COBE} satellite \citep{Matsuura_2011}. The additional absolute calibration of the {\sl AKARI} far-infrared all-sky survey (AFASS) images was done, especially for very bright regions along the Galactic plane, via iterative comparison between the surface brightnesses in the AFASS images and expected surface brightnesses based on the DIRBE zodi-subtracted mission average (ZSMA) data set \citep{Hauser_1998} augmented with the Gorjian zodiacal emission model \citep{Gorjian_2000}. The resulting uncertainties were determined to be less than 10\,\% for surface brightnesses greater than $10, 3, 25$, and 26\,MJy\,sr$^{-1}$ at the N60, WIDE-S, WIDE-L,and N160 bands, respectively, at the spatial resolution of DIRBE at approximately half a degree \citep{Boggess_1992,Takita_2015}. Hence, the presently archived AFASS mapping data (the Public Release Version 1, AFASSv1 hereafter)\footnote{In the Data ARchives and Transmission System (DARTS) maintained by ISAS/JAXA (http://www.darts.isas.jaxa.jp/astro/akari/).} should give correct surface brightnesses of diffuse background emission. However, when aperture photometry was performed for a set of infrared flux calibrators detected in the AFASS maps in the N60, WIDE-S, and WIDE-L bands, the resulting fluxes came out to be roughly 30 -- 60\,\% less than expected \citep{Arimatsu_2014}. In addition, \citet{Takita_2015} compared fluxes listed in the {\sl AKARI\/}/FIS All-Sky Bright Source Catalogue Ver.\,1 (BSCv1; \cite{Yamamura_2009}) and those measured from the corresponding sources in the AFASS images via aperture photometry (with the 90$\arcsec$ radius aperture with the 120--300$\arcsec$ radius sky annulus) and noted that the ratio of the AFASS-to-BSC fluxes as a function of the number of sources within 5$\arcmin$ of the detected source (NDENS) was found to become large for $\mathrm{NDENS} \ge 1$ because of contamination by nearby sources. Note that the absolute flux calibration for the BSCv1 fluxes was done directly from the time-series detector signal readouts of photometry reference objects without making surface brightness maps \citep{Yamamura_2009}. As for the origin of the observed discrepancy between the observed and expected fluxes of point sources, \citet{Arimatsu_2014} offered two possible explanations: (1) the observed point source fluxes are always underestimated because of the extended point-spread-function (PSF) component beyond the aperture radius (e.g., \cite{Arimatsu_2011}), and (2) the sensitivity of the FIS detector may be lower for point sources than for diffuse sources because the signal from the point sources can be missed due to the slow transient of the FIS detector \citep{Shirahata_2009,Ueta_2017}. However, \citet{Doi_2015} discussed in detail about the slow transient effects in the AFASS data and how they were mitigated and removed. Thus, the observed flux underestimates is more likely caused by the unobserved extended point-spread-function (PSF) component beyond the 90$\arcsec$ aperture radius. Thus, in the present work, we examine if the missing extended surface brightness component is truly the cause of the observed flux underestimates for point sources. \citet{Arimatsu_2014} and \citet{Takita_2015} previously measured fluxes of stars in the Cohen catalogue \citep{Cohen_1999}, which contains IR photometric standard stars with a flux range of 0.02 --10\,Jy in the WIDE-S band. Because of the relative faintness of these Cohen standard stars, many of them fell outside of the {\sl AKARI\/} detection limit and their subsequent analyses were conducted using a stacking method to improve on the signal-to-noise ratio (S/N). Here, we take a different approach to verify the point-source flux calibration of the AFASS images. First, we establish the empirical PSF of the AFASS images in all four bands from the AFASS maps themselves (\S\,\ref{S:psf}). Then, we perform point-source photometry using the aperture correction factors derived from the empirical PSF templates as well as contour photometry with the 1-$\sigma$ aperture that simulates an infinite aperture (\S\,\ref{S:photom}), before summarizing the entire work (\S\,\ref{S:summary}). | \label{S:summary} We present an independent photometric analysis of point sources detected in the {\sl AKARI\/} far-IR all-sky survey maps (AFASSv1; \cite{Doi_2015}). By comparing with their corresponding flux entries in the the {\sl AKARI\/} bright-source catalogue ver.\,2 (BSCv2; \cite{Yamamura_2016}), we establish that the surface brightness of the AFASSv1 maps has been calibrated properly for point sources and extended sources alike, in addition to large-scale diffuse emission against which the AFASS maps were originally absolutely calibrated. Especially, we find that the suspected flux underestimates for point sources in the AFASS maps reported earlier by \citet{Arimatsu_2014} and \citet{Takita_2015} can be mitigated by applying a proper aperture correction. In doing so, we construct the empirical super-PSF templates ourselves (Fig.\,\ref{F:psfmad}) and derive our own point-source aperture correction factors (Table\,\ref{T:encircle}). The BSCv2 point-source fluxes are reproduced reasonably well even when direct photometry is done with a 1-$\sigma$ photometry aperture (i.e., all the surface brightness pixel values within the aperture are summed up to compute the source flux). This means that direct photometry for any source (point, extended, and diffuse) detected in the AFASS maps yields a correct flux measurement when a sufficiently large aperture is used. On the whole, the AFASS maps are useful tools to directly measure far-IR fluxes. This is true especially for sources that are not exactly point sources (i.e., extended), because the BSC catalogue flux entries may not be appropriate for such extended sources. \begin{ack} This research is based on observations with AKARI, a JAXA project with the participation of ESA. TU acknowledges support and encouragement by the {\sl AKARI\/} team at the Institute of Space and Aeronautical Science of JAXA, especially from Dr.\ Issei Yamamura. \end{ack} | 18 | 8 | 1808.00730 |
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