adsabs/FOCAL huggingface dataset:
We now investigate relativistic reflection models with a primary Comptonisation continuum shape, which is more physical and has a sharper high-energy rollover compared to an exponential cutoff power law. In addition, such models have the advantage of having the hot corona temperature (kThot) as a physical parameter rather than a phenomenological exponential cutoff energy. We first apply the RELXILLCP model, which uses the NTHCOMP Comptonisation model (Zdziarski et al. 1996; Życki et al. 1999) as the incident spectrum. The other physical parameters are the same as those in the RELXILL model presented above. We find a good fit and infer similar hot corona temperatures of kThot ∼ 26 keV for all three epochs (Table 2). Then, we consider the REFLKERR where the hard X-ray Comptonisation spectrum is computed with the COMPPS model (Poutanen & Svensson 1996), which appears to be a better description of thermal Comptonisation when compared to Monte Carlo simulations (Zdziarski et al. 2020). Moreover, REFLKERR has, as physical parameter, either the Compton parameter (y) or the optical depth (τ). We choose to perform the fit with the optical depth as the direct inferred parameter. The temperature of the thermal seed photons (kTbb) Comptonised by the hot corona is an explicit physical parameter of this model. Here, we assume that the seed photons are provided by the cold disc, then kTbb was fixed at 10 eV corresponding to the expected maximum temperature of the accretion disc around a black hole mass of 1.4 × 108 M⊙ accreting at a ∼10% Eddington rate. In addition, REFLKERR allows us to choose either a slab or a spherical geometry for the hot corona. The latter corresponds to numerous active sphere regions above the disc surface. The hard X-ray shape of the reflected component is calculated using IREFLECT convolved with COMPPS rather than using XILLVERCP (see Niedźwiecki et al. 2019 for detailed explanations). Both models give good fits (Table 2, Fig. 3), though a larger χ2 value for the spherical geometry. Similar values of the hot corona temperature are measured for the three epochs: kThot ∼ 30−31 keV and kThot ∼ 21−22 keV for the slab and spherical geometries, respectively. From τ, we can infer the corresponding Compton-parameter of the hot corona (yhot) using the relation y = 4τ(kT/511 keV) (Beloborodov 1999). This correspond to yhot ∼ 0.5 and yhot ∼ 1.1 for the slab and spherical geometries, respectively.
Citation: Poutanen & Svensson 1996
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LPVs are known to exist on sets of sequences in period–luminosity space (e.g. Wood et al. 1999; Wood 2000) depending on their variability type and pulsation mode. The primary periods of Mira-like variables lie on the commonly called C and C′ sequences (Ita et al. 2004; Spano et al. 2011) with C′ lying at a lower period than C. Sequences A and B lie at lower periods still and are populated by the so-called OGLE small amplitude red giants or OSARGs (Soszynski et al. 2004a). Sequences D and E are populated by long secondary pulsators, whose nature is still unclear, as well as ellipsoidal and eclipsing binary systems (Soszynski et al. 2004b), with the latter two contained in a separate CRTS catalogue not utilized in this work. Soszyński & Wood (2013) also found that there is a low-amplitude population of SRVs whose primary period lies between the C and C′ sequences. Given that the selections of Fig. 1 produce a CRTS sample with a wide range of visual amplitudes, it is uncertain on which sequence our CRTS sources lie and to which class of LPVs they belong. The left-hand panel of Fig. 2 shows sources from Gaia’s DR2 LPV catalogue (Gaia Collaboration et al. 2016, 2018; Mowlavi et al. 2018) chosen to lie within a 15° aperture of the LMC. Cross-matching with 2MASS reveals three distinct sequences existing in the Ks band, with the middle sequence corresponding to the Mira C sequence. Belokurov et al. (2017) defined a variability parameter (the ‘Gaia amplitude’) based on Gaia’s flux information as:
(1)
\begin{eqnarray*}
\mbox{Amp} = \mbox{log}_{10} \left(\sqrt{N_{\mbox{obs}}} \: \frac{\sigma _{\overline{I_{\mathrm{ G}}}}}{\overline{I_{\mathrm{ G}}}} \right)
\end{eqnarray*}
where $\sqrt{N_{\mbox{obs}}}$ is the number of observations, $\sigma _{\overline{I_{\mathrm{ G}}}}$ is the mean flux error in the G band, and $\overline{I_{\mathrm{ G}}}$ is the mean flux in the G band. Requiring a Gaia amplitude greater than −0.55 isolates a sample of LPVs lying on a single period–luminosity sequence, as evident in Fig. 2. We produce a secondary cleaned CRTS sample located on this middle sequence by cross-matching with the Gaia DR2 source catalogue and enforcing that the Gaia amplitude > −0.55. Further, we restrict the sample to visual amplitudes >1 to eliminate small amplitude variables. Application of these two cuts gives a cleaned CRTS sample with 225 members. We utilize this cleaned CRTS sample in Section 3.1, where we exploit the period–age correlation of LPVs, the interpretation of which may be muddied if low visual amplitude variables were inadvertently retained. For example, AGB variables residing in a common Magellanic Cloud globular cluster have been observed by Kamath et al. (2010) both at high and low amplitudes. They found stars with low amplitudes also had lower stellar periods, implying differing period–age relations.
Citation: Mowlavi et al. 2018
Citation function: Uses; Uses
The GALEX NUV observations span a wide wavelength range, from 1771 to 2831 Å. While these data provide no spectral information within that bandpass, we rely on solar and stellar flare studies to inform the likely contributors to the flare flux. The NUV spectral region has not had as many observational constraints as the far-UV region in flare studies. Few NUV stellar flare spectra exist at all, and the few that do were obtained either for solar flares or on nearby M dwarfs. Flares observed in the UV are often associated with the more impulsive phases of solar flares, starting with early observations showing a close temporal association between UV and hard X-ray emission (Cheng et al. 1981). Welsh et al. (2006) reported on high time-resolution NUV+FUV flares seen with GALEX on nearby M dwarfs, and Hawley et al. (2007) presented high spectral-resolution NUV flare measurements of an M-dwarf flare. From these two studies, the contribution of emission lines relative to continuum emission could be determined; the main emission lines in the flare NUV spectrum were Mg ii, Fe ii, Al iii, and C iii. However, the main emission component overall was a continuum component. Recent results from solar flares observed from space (Heinzel & Kleint 2014; Kleint et al. 2016) demonstrate an NUV spectrum originating largely from Hydrogen Balmer continuum emission. The formation of the NUV emission appears to originate from an impulsive thermal and nonthermal ionization caused by the precipitation of electron beams through the chromosphere. This explains the temporal correlation with solar flare hard X-ray emission observed previously. More recently, Kowalski et al. (2019) presented accurately calibrated NUV flare spectra at high time cadence on an M dwarf, again finding a large flux enhancement due to continuum radiation. They commented that the oft-used 9000 K blackbody used to describe blue-optical stellar flare emission (Hawley et al. 2003) under-predicts the NUV continuum flare flux by a factor of two. Based on general similarities in radiative properties between solar and stellar flares studied thus far (Osten 2016), it is likely that a combination of line and continuum emission enhancements are present in the NUV flare flux from the flares being considered, but we cannot speculate about the relative contribution of one versus the other. These sources originate from different layers of the stellar atmosphere: singly and doubly ionized emission lines likely originate in the chromosphere, and the Balmer continuum emission also originates from the chromosphere. Some lines, such as Mg ii, exhibit absorption components and self-reversals, indicating optical depth effects in the atmosphere, while other lines such as Fe ii appear to be optically thin. Any hot blackbody emission might originate further in the photosphere.
Citation: Kleint et al. 2016
Citation function: Background; Background
From the X-ray point of view, old open clusters are interesting for a number of reasons. First, X-ray observations efficiently detect different classes of close, interacting binaries, enabling the study of processes such as tidal coupling and the link between X-rays and rotation. The X-ray luminosity of late-type stars strongly depends on stellar rotation. As single stars age, they spin down due to magnetic braking (Pallavicini 1989). As a result, their X-ray emission decreases accordingly. An old star like our Sun (∼4.5 Gyr) has an X-ray luminosity of about 1026 to 1027 erg s−1 (0.1–2.4 keV; Peres et al. 2000). Even with the deepest exposures of a sensitive X-ray telescope like the Chandra X-Ray Observatory, this is nearly impossible to detect except for the nearest stars. Nevertheless, an early ROSAT observation of the old open cluster M 67, which lies at ∼840 pc (Pasquini et al. 2008) and is about as old as the Sun (





; Dinescu et al. 1995), revealed a large number of X-ray sources among the cluster members (Belloni et al. 1993). Many of these turned out to be close, tidally interacting binaries, where the stellar rotation is locked to the orbital period and therefore kept at a level that can sustain magnetically active coronae. Subsequent XMM-Newton (Gondoin 2005; Giardino et al. 2008; Gosnell et al. 2012) and Chandra (van den Berg et al. 2004, 2013; Giardino et al. 2008) observations of old open clusters have detected many such active binaries (ABs). ABs can be binaries of two detached stars, or they can have a contact or semi-detached configuration such as in W UMa and Algol binaries, respectively. In terms of the number of sources, ABs are the most prominent X-ray source class in old open clusters, but other classes of interacting binaries are represented as well. In cataclysmic variables (CVs), the X-rays are the result of accretion from a late-type main-sequence donor onto a white dwarf. In fact, the first ROSAT observation of M 67 was aimed at studying the X-rays from a CV that was discovered in the optical (Gilliland et al. 1991). The origin of the X-ray emission from more exotic open-cluster binaries, like blue stragglers, is less understood, but in X-rays they are more similar to the ABs than to the mass-transfer sources (van den Berg 2013).
Citation: Pasquini et al. 2008
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Comparison of apparent effective diameters of these sources to direct size measurements supports a similar conclusion. Simpson et al. (2015) present ALMA observations of 23 SCUBA-2-selected SMGs with a median physical half-light diameter of 2.4 ± 0.2 kpc, while Ikarashi et al. (2015) show ALMA observations of 13 AzTEC-selected SMGs with a median physical half-light diameter of $1.34^{+0.26}_{-0.28}$ kpc. ALMA observations of four SPT-selected lensed SMGs give a mean physical half-light diameter of 2.14 kpc (Hezaveh et al. 2013b). This measurement is consistent with a recent lensing analysis of a significantly expanded SPT-selected DSFG sample (Spilker et al. 2016). These high-resolution ALMA observations constrain the FIR sizes of the sources to be 1.0–2.5 kpc. Earlier observations of the physical sizes of SMGs by CO detection and 1.4 GHz imaging suggest larger sizes (e.g. Tacconi et al. 2006; Biggs & Ivison 2008; Younger et al. 2008). However, Simpson et al. (2015) point out that the submillimetre sizes are consistent with resolved 12CO detections, while the sizes derived from 1.4 GHz imaging are about two times larger because of the cosmic ray diffusion, which can explain the results before higher frequency observations at ALMA were possible (Chapman et al. 2004; Tacconi et al. 2006; Biggs & Ivison 2008; Younger et al. 2008). Similarly, Ikarashi et al. (2015) reveal that the 12CO detected sizes and the 1.4 GHz imaging sizes of similar sources are greater than their submillimetre sizes as well. Furthermore, observations of local galaxies also show the submillimetre sizes are smaller than the CO detected sizes (e.g. Sakamoto et al. 2006, 2008; Wilson et al. 2008) and the 1.4 GHz continuum sizes (e.g. Elbaz et al. 2011). Our photometrically derived $\sqrt{\mu }d$ is best compared to the submillimetre continuum sizes. With a median apparent effective diameter of $4.2^{+1.7}_{-1.0}$ kpc, the $\sqrt{\mu }d$ of our sample is one to six times the observed intrinsic diameters (1.0–2.5 kpc). Lensing (or multiplicity) increases the apparent effective size of a source, so this comparison favours a lensing (or multiplicity) interpretation for the ACT-selected sources.
Citation: Tacconi et al. 2006
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Citation: Tacconi et al. 2006
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The EW(λ8542)/EW(λ8498) values are obtained around ∼1.4 in the first observing night, which are somewhat smaller than the values around ∼1.6 derived in the second night, when a strong optical flare decay was detected (also see the Section 4.2). These low ratios are consistent with the values found for several other stars with strong chromospheric activity (e.g. Montes et al. 2000; Gu et al. 2002; López-Santiago et al. 2003; Zhang & Gu 2008; Gálvez et al. 2009; Cao & Gu 2014, 2015, 2017; Zhang et al. 2016), which suggests that the $\rm{Ca\,{\small II}}$ IRT line emission arises from plage-like regions. The EHα/EHβ ratios have also been usually used as a diagnostic for discriminating the presence of different structures on the stellar surface. As Huenemoerder & Ramsey (1987) discussed, the low ratios in RS CVn-type stars are caused by plage-like regions, while prominence-like structures have high values. Similar results have also been reported by Hall & Ramsey (1992) who found that low ratios (∼ 1–2) can be achieved both in plages and prominences viewed against the disc, but high values (∼ 3–15) can only be obtained in extended prominence-like structures viewed off the stellar limb. We obtain the ratios on SZ Psc during the flare decay phase change from 2.40 to 3.40, which are not especially high and anticorrelated with the variation of activity emission shown in Fig. 3, in which the EWs of Hα and Hβ line subtraction and the EHα/EHβ values are plotted as a function of orbital phase. Especially when the EWs have an increasing oscillation during the gradually decrease, which is resulted from flare ejection (see the discussion in Section 4.2), the ratios have a similar anticorrelation feature. These facts suggest that the low EHα/EHβ ratios and its variation might be dominantly associated with the flare decrease, and accompanied cool post-flare loops (see the Section 4.3) might play a part of role for the low ratios, which may have a state like absorbing prominence projected against the disc.
Citation: Zhang et al. 2016
Citation function: Similarities
The young starburst inferred by the detections of high ionization emission line of He II λ4686 and the blue WR bump in this and previous works (Guseva et al. 2000; Brinchmann, Kunth & Durret 2008) is confirmed by the age estimates made here for the bright and faint regions in Mrk 22 as ∼4 and ∼10 Myr, respectively. Unlike previous works, we carried out abundance analysis for both the regions separately. We found an appreciable metallicity difference of ∼0.5 dex between the bright region [12 + log (O/H) ∼ 8] and the faint region [12 + log(O/H) ∼ 7.5]. The separation between two regions is ∼0.6 kpc. Typical metallicity gradients in normal spiral galaxies have been found between −0.009 and −0.231 dex kpc−1, with an average gradient of −0.06 dex kpc−1 (Zaritsky, Kennicutt & Huchra 1994). The observed metallicity difference between the two regions in Mrk 22 is too large to be explained as a normal galactic metallicity gradient. The chemical composition as measured from the gas-phase metallicity [12 + log(O/H)] shows various degree of spatial variations in different types of dwarf galaxies. For instance, shallow gradient in metallicity is seen in SBS 0335−052 (Papaderos et al. 2006) while no significant variations were seen in Mrk 35 (Cairós et al. 2007). A study on a large sample indicates that normal BCD galaxies are chemically homogeneous (Kobulnicky & Skillman 1996; Papaderos et al. 2006; Kehrig et al. 2008; Cairós et al. 2009; Pérez-Montero & Contini 2009; Pérez-Montero et al. 2011; Hägele et al. 2011; García-Benito & Pérez-Montero 2012; Lagos & Papaderos 2013). On the other hand, the metallicity of extremely metal-poor galaxies is usually not homogeneous within the galaxy, with the low metallicity seen in regions of intense star formation (Papaderos et al. 2006; Izotov & Thuan 2009; Levesque et al. 2011; Sánchez Almeida et al. 2013, 2014, 2015). However, large metallicity gradients are not common in dwarf galaxies. The simplest explanation for large metallicity difference in a single system is a recent merger of two galaxies with different metallicity. In a few cases, significantly large metallicity differences between star-forming regions in dwarf galaxies were seen and understood in terms of recent tidal interactions or mergers (López-Sánchez, Esteban & Rodríguez 2004a,b; López-Sánchez, Esteban & García-Rojas 2006; López-Sánchez & Esteban 2009, 2010). The evolution in terms of metallicities in interacting dwarf galaxies is fairly complex as it can depend on various factors such as mixing of metals with the interstellar medium (ISM), possible outflows of metals, and inflow of metal-poor gas in tidally interacting systems.
Citation: Papaderos et al. 2006
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However, this outward migration of the inner edge should stop at the corotation radius rco, where the Keplerian frequency of the disk equals the spin frequency of Jupiter. When rcav > rco, there will be two possibilities: the angular momentum will be transferred from Jupiter to the disk and then the gas accretion will stop, or otherwise the corotation radius and the disk edge will move outward together and then the accretion will continue (Takata & Stevenson 1996; Liu et al. 2017). Although the current corotation radius is rco ≈ 2.25 RJ, Jupiter at its time of formation was much larger than it is today (Burrows et al. 1997; Fortney et al. 2011), and this means that the corotation radius was also larger than the current one if the conservation of the angular momentum of Jupiter is assumed. Considering the transport of the angular momentum from Jupiter to the disk, the angular momentum should have been conserved since the disk disappeared. According to a formation model of Jupiter, the radius of the planet was ≈1.75 RJ after its rapid gas accretion and it decreased little by little (Lissauer et al. 2009). When the radius of Jupiter is 1.75 RJ, the corotation radius should be rco ≈ 4.7 RJ. We can then consider two scenarios of Io formation. In the first one, Io formed around r ≈ 4.7 RJ, slightly interior to the r = 5.89 RJ of our fiducial model, and then moved outward after the disk dissipated. The satellites, especially the inner ones, could move outward by the tidal force from Jupiter (Yoder & Peale 1981). The outer ones would be pushed by the inner ones and move outward with them because of the resonance. In this case, the position of the snow line should have been more inside than the fiducial case in this work, but this thermal condition could easily be reproduced by another parameter set. The second possibility is that Io was not the innermost satellite. If a body was present at r = 3.7 RJ, Io would have been situated at r = 5.7 RJ if they were trapped in a 2:1 resonance. This orbit is consistent with the corotation radius when the radius of Jupiter is ≈1.5 RJ, and this radius can be achieved during the contraction of Jupiter. The innermost body may have been broken by the tidal force of Jupiter when it has entered inside the Roche limit. Current Io, trapped in the Laplace resonance, actually moves inward little by little because of the tidal dissipation, and the innermost body may have also experienced such inward migration (Lainey et al. 2009).
Citation: Burrows et al. 1997
Citation function: Uses; Uses
In this paper, we are interested not only in the mathematical properties of the ambipolar diffusion as a nonlinear diffusion process, but also in the inclusion of ambipolar diffusion terms in MHD codes. In astrophysics, over the past few decades, multidimensional MHD computer codes have been developed that model a variety of physical processes including ambipolar diffusion. Representative examples of such codes and simulations outside solar physics can be found in Basu & Mouschovias (1994), Mac Low et al. (1995), Padoan et al. (2000), Basu & Ciolek (2004), Kudoh & Basu (2008), Choi et al. (2009), Gressel et al. (2015), Tomida et al. (2015), O’Sullivan & Downes (2007), Masson et al. (2012), Viganò et al. (2019), and Grassi et al. (2019). The consideration of ambipolar diffusion processes in solar physics started many decades ago (e.g. Parker 1963), but it has undergone a true explosion in terms of its use in large numerical models (e.g. Leake et al. 2005; Leake & Arber 2006; Arber et al. 2007; Cheung & Cameron 2012; Leake & Linton 2013; Martínez-Sykora et al. 2012, 2017a,b, 2020a,b; Ni et al. 2015, 2016, 2021; Khomenko et al. 2017, 2018, 2021; González-Morales et al. 2018, 2020; Nóbrega-Siverio et al. 2020a,b; Popescu Braileanu & Keppens 2021). Such numerical calculations often encounter a problem: given the comparatively high values of χa in different cosmic environments, the advance in time may grind to a halt in magnetised regions when a standard Courant-Friedrichs-Lewy condition is adopted for the timestep based on ηa. The recent papers by González-Morales et al. (2018) and Nóbrega-Siverio et al. (2020b) describe the construction of so-called super-time-stepping (STS) modules for the Mancha code and for the Bifrost code, respectively, designed with the aim to overcome that stiffness problem. In at least three of the papers cited above (Masson et al. 2012; Viganò et al. 2019; Nóbrega-Siverio et al. 2020b), the basic ZKBP solution in cylindrical coordinates was used to test the ambipolar diffusion module, given its simplicity and the analytical expression available for it. However, as already mentioned, the ambipolar diffusion problem tends to give rise to sharp current sheets and singularities. The ZKBP solution is comparatively smooth in that sense and it would be good to have some other canonical solutions on hand that include current sheets having higher degrees of singularity, which naturally occur in the ambipolar diffusion problem.
Citation: Nóbrega-Siverio et al. 2020a
Citation function: Motivation; Motivation; Motivation