adsabs/FOCAL huggingface dataset:
One of the major endeavors of modern observational cosmology is to paint a coherent picture of the history of the universe. To this end, the final frontier remains the identification and characterization of the first sources that appeared in the universe, those which played a significant role in reionizing the intergalactic medium from a neutral state to a fully ionized one over the first billion years (corresponding to redshifts of 6 ≲ z ≲ 12). Extragalactic surveys (of deep fields as well as lensing clusters; Grogin et al. 2011; Koekemoer et al. 2011; Bradley et al. 2012; Ellis et al. 2013; Bradley et al. 2014; Schmidt et al. 2014; Treu et al. 2015; Lotz et al. 2017; Salmon et al. 2018; Coe et al. 2019) with the Hubble Space Telescope (HST) have yielded impressive gains in the number of galaxy candidates at redshifts z = 7–10, with samples reaching over 1000 objects, and revolutionized our understanding of galaxy evolution therein. Complementing these observations, the spectroscopic confirmation (e.g., Finkelstein et al. 2013; Oesch et al. 2015; Zitrin et al. 2015; Roberts-Borsani et al. 2016; Hoag et al. 2017; Stark et al. 2017; Hashimoto et al. 2018) and characterization (e.g., Laporte et al. 2017; Mainali et al. 2018; Endsley et al. 2021) of over a dozen sources has seen impressive advances with ground-based spectroscopy (e.g., probing the rest-frame UV and FIR with Keck/MOSFIRE, VLT/X-Shooter, and ALMA), particularly for the brightest and rarest objects. For the rest-frame optical, however, the Spitzer Space Telescope has, until now, afforded the only realistic means for statistical analyses. However, the Infrared Array Camera’s (IRAC) coarse spatial resolution and the limited depth probed by many surveys make robust and uncontaminated constraints on galaxy properties a challenging feat. Further advances with current facilities are challenging owing to the limited wavelength coverage of HST and the observed faintness of star-forming galaxies as one approaches redshifts of z > 10. The imminent arrival of the James Webb Space Telescope (JWST) has the potential to detect galaxies well beyond the current frontier of z ∼ 12 (e.g., Behroozi et al. 2020) thanks to the unprecedented resolution and sensitivity of its near-IR (NIR) imaging and spectroscopic capabilities, and revolutionize our current understanding of galaxy evolution.
Citation: Schmidt et al. 2014
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Regardless of whether the NGC 6334 filament will form massive stars, our ArTéMiS result that the filament inner width is within a factor of 2 of 0.1 pc has interesting implications. Our NGC 6334 study is clearly insufficient to prove that interstellar filaments have a truly universal inner width, but it shows that the finding obtained with Herschel in nearby clouds is not limited to filaments in low-mass star-forming regions. It is quite remarkable that the NGC 6334 filament has almost the same inner width as the faint subcritical filaments in Polaris (cf. Men’shchikov et al. 2010; Arzoumanian et al. 2011), the marginally supercritical filaments in Musca and Taurus (Cox et al. 2016; Palmeirim et al. 2013), or the lower-mass supercritical filaments in Serpens South and Vela C (Hill et al. 2012), despite being three orders of magnitude, two orders of magnitude, and at least a factor of ~3 denser and more massive than these filaments, respectively (see Table 1). While not all of these filaments may have necessarily formed in the same way, this suggests that a common physical mechanism is responsible for setting the filament width at the formation stage and that the subsequent evolution of dense filaments, through accretion of background cloud material (cf. Heitsch 2013; Hennebelle & André 2013) for example, is such that the inner width remains at least approximately conserved with time. A promising mechanism for creating dense filaments, which may be quite generic especially in massive star-forming complexes, is based on multiple episodes of large-scale supersonic compression due to interaction of expanding bubbles (Inutsuka et al. 2015). With about seven bubble-like HII regions per square degree (Russeil et al. 2013, see also Fig. 2), there is ample opportunity for this mechanism to operate in NGC 6334. More specifically, at least in projection, the NGC 6334 filament appears to be part of an arc-like structure centered on the HII region Gum 63 (see Fig. 2a), suggesting the filament may partly result from the expansion of the associated bubble. Interestingly, the background column density is one order of magnitude higher for the NGC 6334 filament than for the other filaments of Table 1, which is suggestive of a significantly stronger compression. More observational studies are needed to investigate the structure and environment of a larger number of filaments in massive star-forming regions to determine whether the characteristics of the NGC 6334 filament are generic. More theoretical work is also needed to better understand the physics controlling the width of interstellar filaments.
Citation: Hill et al. 2012
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The purpose of this paper is to demonstrate a complementary approach to single-dish IM experiments, capable of the statistical detection of H i field and therefore the fluctuations in the underlying matter field by measuring the H i power spectrum with interferometric observations. Interferometers have inherent advantages over single-dish measurements. Besides providing high angular resolutions, they are less sensitive to systematics which poses a major problem to the autocorrelation power. However, the smallest k-modes accessible to an interferometer is determined by the shortest baselines which may hinder probing the BAO scales. Interferometers such as CHIME (Bandura et al. 2014), TIANLAI (Xu, Wang & Chen 2014), and HIRAX (Newburgh et al. 2016) are custom designed to probe the BAO scales using the 21 cm signal in the redshift range z ∼ 0.5–2. In this paper, we study the feasibility of detecting the cosmological 21 cm signal with MeerKAT and present forecasts on the statistical measurement of H i with MeerKAT L-band (856 MHz  ν  1712 MHz) observations on quasi-linear scales. We present the sensitivity estimates at z ∼ 0.27 from our newly developed simulation pipeline, which is our first attempt towards the measurement of H i power spectrum with MeerKAT. Our pipeline is based on the methods being developed for similar statistical measurement of H i from the Epoch of Reionization from a series of experiments at lower radio frequencies such as LOFAR (Van Haarlem et al. 2013), GMRT (Paciga et al. 2013), PAPER (Parsons et al. 2014), HERA (DeBoer et al. 2017), and MWA (Tingay et al. 2013). We show that with MIGHTEE (Jarvis et al. 2016; Maddox et al. 2021), one of MeerKAT’s large survey projects, we can achieve constraints on the H i power spectrum at z = 0.27. There are implicit advantages with such survey projects with a specific emission line. First, it provides a one to one correspondence between observed frequency and redshift, thereby delivering a very high redshift resolution. Secondly, these are generally less time consuming compared to an optical spectroscopic galaxy survey which requires very high sensitivity to detect individual galaxies.
Citation: Xu, Wang & Chen 2014
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In Figure 8, we also find small differences in the [Na, Al, O/Fe] abundance ratios between the stars in the bright and faint RC groups, although it is not as clear as in the case of [Na, Al, O/H] abundances. In particular, unlike Figure 7, stars in the bRC group are more enhanced in [Na/Fe] but appear to be more depleted in [Al/Fe] and [O/Fe] than those in the fRC group. The mean differences are 0.053 ± 0.021 dex, 0.032 ± 0.018 dex, and 0.071 ± 0.045 dex in [Na/Fe], [Al/Fe], and [O/Fe], respectively, which are marginally significant at p-values of 0.22, 0.18, and 0.23. When the relative fraction of RC stars is taken into account (27%; see Section 4), the difference in [Na/Fe] between the genuine RC stars would correspond to Δ[Na/Fe] ∼ 0.20 dex, which is comparable to that expected from our chemical evolution model for the bulge stars (Δ[Na/Fe] = 0.2 ∼ 0.3 dex; Kim & Lee 2018; Lee et al. 2019).10

10
The previous study by Lee et al. (2019) noted a clear separation of the two groups according to Na abundance among bright RGB stars in the outer bulge. The apparent lack of such a distinct difference between the two groups in this study may be due to a larger uncertainty on abundances of relatively faint sample stars.
 The overall chemical patterns, however, are not identical to those observed in typical GCs, where the later-generation stars are more enhanced in [Na, Al/Fe] and more depleted in [O, Mg/Fe] than the first-generation stars at a given metallicity, although the trend of [Na, Al O/Fe] between the two RCs is less clear. Figure 9 shows the comparison of stars in this study with stars in metal-rich GCs ([Fe/H] > −1.0) on the Na–O diagram. The stars used in this study have a different distribution from stars in GCs. Although the bRC group is slightly more enhanced in [Na/Fe] and more depleted in [O/Fe] than the fRC group, the [Na/Fe] variation of RC stars is smaller than that of GC stars. This discrepancy might imply the different chemical evolution between stars in the bulge and typical GCs. We note, however, that even though we employ only metal-rich GC stars for the comparison, the majority of stars are still far more metal-poor ([Fe/H]  −0.5) than stars in the bulge. Because the relatively small [Na/Fe] variation is expected from the chemical evolution model for metal-rich bulge stars and the O-depletion is indistinct in some metal-rich GCs, such as NGC 6121 and 47 Tuc (see Kim & Lee 2018; Lee et al. 2019), the direct comparison of bulge stars with similarly metal-rich GCs on the Na–O plane would require further spectroscopic observations for such GCs in the bulge.
Citation: Lee et al. (2019)
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Lower limit on μ: at the highest accretion rates, to observe pulsations, the magnetic field must be at least high enough to truncate the accretion disc at or above the neutron star surface. Thus by setting rt = Rs we obtain the lower limit on the dipole moment as

(5)

\begin{equation}
\mu _{\rm min} = \gamma _{\rm B}^{-1/2}(GM)^{1/4}\dot{M}_{\rm max}^{1/2}R_{\rm s}^{7/4}.
\end{equation}

We assume the mass accretion rate can be estimated from the bolometric luminosity as $L = GM\dot{M}/R_{\rm s}$, and estimate L from the observed luminosity in the X-ray band by applying a bolometric correction factor (L = ϵbolLX). The typical reported values of the bolometric correction factor have a range of ϵbol ∼ 1–2 (Gilfanov et al. 1998; Galloway et al. 2002, 2008; Campana et al. 2003; Migliari & Fender 2006; Casella et al. 2008). The mass accretion rate then follows as

(6)

\begin{eqnarray}
\dot{M} &amp;=&amp; 10^{16} \mbox{ g s}^{-1} \left( \frac{ \epsilon _{\rm bol} L_{\rm X}}{1.87 \times 10^{36}\mbox{ erg s}^{-1}} \right)\nonumber\\
&amp;&amp;\times\,\left( \frac{M}{1.4 \,\mathrm{M}_{\odot }} \right)^{-1} \left(\frac{R_{\rm s}}{10 \mbox{ km}}\right).
\end{eqnarray}

Expressing the X-ray luminosity in terms of the observed X-ray flux and the assumed distance (LX = 4πd2F), we obtain the lower limit on the magnetic moment

(7)

\begin{eqnarray}
\mu _{\rm min} &amp;=&amp; 9.36 \times 10^{24} \mbox{ G cm}^{-3} \left( \frac{ \gamma _{\rm B} }{ 1 } \right)^{-1/2} \left( \frac{ \epsilon _{\rm bol} }{ 1 } \right)^{1/2} \nonumber \\
&amp;&amp;\times \left( \frac{ F_{\rm max}}{10^{-10}\, \mbox{erg}\mbox{ cm}^{-2}\mbox{ s}^{-1}} \right)^{1/2} \left( \frac{d}{10 \mbox{ kpc}}\right) \nonumber \\
&amp;&amp;\times \left(\frac{M}{1.4 \,\mathrm{M}_{\odot }}\right)^{-1/4} \left( \frac{R_{\rm s}}{10 \mbox{ km}}\right)^{9/4},
\end{eqnarray}

where Fmax is the highest observed X-ray flux with pulsation and we adopted the boundary values of γB and ϵbol that give the lowest magnetic moment. For a detailed discussion of these assumptions we again refer to Section 4.
Citation: Galloway et al.; 2008
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Cassiopeia A (Cas A) is a Galactic SNR located $3.4 \, {\rm kpc}$ away (Reed et al. 1995), with an age of approximately $330 \, {\rm yr}$ (Fesen et al. 2006) and a radius of $1.7 \, {\rm pc}$ (Reed et al. 1995). It provides a unique laboratory to test the efficiency of dust condensation in SN ejecta, and the subsequent destruction of dust by the reverse shock, and as such has been studied extensively in the past. IR/sub-mm observations have led to derived dust masses ranging from ${\sim } 10^{-4} \, {\rm M}_\odot$ of hot ($T \sim 100 \, {\rm K}$) dust (Arendt, Dwek & Moseley 1999; Douvion, Lagage & Pantin 2001) to $2{\text{-}}4 \, {\rm M}_\odot$ of cold dust emitting at sub-mm wavelengths (Dunne et al. 2003), although this higher mass has been attributed to foreground dust emission in the ISM (Krause et al. 2004). Analyses of integrated fluxes from Spitzer and Herschel observations (Rho et al. 2008; Barlow et al. 2010; Arendt et al. 2014) found ${\sim } 0.01 \, {\rm M}_\odot$ of hot dust, with ${\sim } 0.1 \, {\rm M}_\odot$ of cold, unshocked dust present in the central regions, in agreement with simulations of the dust formation and evolution in Cas A by Nozawa et al. (2010). Dunne et al. (2009) suggested the observed polarization of the sub-mm emission could be explained by ${\sim } 1 \, {\rm M}_\odot$ of cold dust, similar to the value of $1.1 \, {\rm M}_\odot$ given by Bevan et al. (2017) as the most likely mass based on the shape of emission-line profiles affected by extinction. De Looze et al. (2017) utilized spatially resolved Herschel and Spitzer observations of Cas A to fit the dust continuum emission, following the removal of line and synchrotron contamination, using a four-component model including ISM dust emission and three SNR dust temperature components. They found a large mass of unshocked cold dust in the centre of the SNR (up to $0.6 \, {\rm M}_\odot$), significantly above previous estimates based on the IR/sub-mm emission.
Citation: Dunne et al. 2003
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The traditional way of studying instabilities of non-linear dynamical systems is via linear stability analysis. In many cases, it leads to much simpler system of linearized equations, which can be solved analytically. However, in many other cases even the linearized system does not allow general solutions in terms of analytic functions for arbitrary equilibrium configuration. One way to overcome this problem is to restrict the analysis to special equilibrium configurations which allow us to simplify the linearized system of equations even further. Accordingly, most early studies of jet stability assumed simplified jet structure, including the magnetic field topology (e.g. Hardee 1979, 1982; Cohn 1983; Payne & Cohn 1985; Istomin & Pariev 1996; Begelman 1998; Lyubarskii 1999; Tomimatsu, Matsuoka & Takahashi 2001; Narayan et al. 2009). This allowed to obtain the solution in terms of Bessel and hypergeometric functions. Although very useful in many respects, this approach still cannot address the stability of jets with more complex and more realistic structure. In particular, all these equilibrium jets included surface currents, whereas Gourgouliatos et al. (2012) argued that such current sheets are likely to promote resistive jet instabilities. A jet that is free of current sheets would be free of resistive instabilities in computer simulations. As an alternative, they constructed equilibrium solutions which are current-sheet-free structures. There is an intuitive reason for expecting a magnetized jet without a current sheet to be more stable – such jets carry no volumetric net current. If the surface of a magnetized jet with current sheet is perturbed, two neighbouring perturbations could behave analogously to two parallel current carrying wires that are carrying current in the same direction. Since such wires attract, one might analogously expect the surface of such a jet to become more corrugated, i.e. the perturbations can grow. By avoiding current sheets on the surface, a current-sheet-free jet avoids this mode of destabilization. Another interesting feature of these solutions is that the jets carry zero net current and magnetic flux. In some part of the jet cross-section the poloidal electric current flows outwards and in the rest of the cross-section exactly the same amount of the current flows in the opposite direction. Thus, one does not have to worry of having a return current outside of the jet on large scales. The same applies to the poloidal magnetic field. Unfortunately, the magnetic structure of these solutions is too complex for the linearized equations to allow analytical solutions.
Citation: Hardee 1979
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Type II supernovae (SNe) are defined by the prominent hydrogen lines in their spectra. They are believed to originate from the collapse of an iron core of massive stars (≳8 M⊙) that retain their hydrogen envelope. The most common sub-type, comprising ∼70 per cent of all Type II SNe, is characterized by a phase of a roughly constant magnitude in the optical bands, hence their name type II Plateau (II-P). This plateau phase typically starts 1–2 weeks after the explosion and lasts for ∼100 d. Pre-explosion images have revealed that the progenitors of this class are red supergiants, in the mass range of (7–16 M⊙) (Smartt 2015; for individual progenitor detections see e.g. Van Dyk, Li & Filippenko 2003a; Van Dyk, Li & Filippenko 2003b; Van Dyk et al. 2012). Type II Linear (II-L) SNe constitutes another subclass of type II SNe (e.g. Patat et al. 1994; Arcavi et al. 2012; Faran et al. 2014a,b). They are spectroscopically very similar to type II-P events (Faran et al. 2014b, see), but their light curves are declining in all bands. In both types (II-P and II-L), there is typically a sharp drop in the luminosity after ∼100 d and the luminosity starts to follow roughly the exponential decay expected from emission powered by the decay of 56Ni. The distinction between these two classes is not well defined, and studies have used different definitions for II-L SNe. However, several recent works have shown that there exists a continuum of decline rates between slow declining and fast declining SNe (Anderson et al. 2014; Faran et al. 2014a), which suggests that a separation into two different classes may be artificial. Other type II sub-classes will not be discussed here. In this paper, we focus on the light curves of type II-P and type II-L SNe, without making the distinction between the two types, and refer to them in short as Type II SNe. Our goal here is to perform a uniform analysis of the bolometric luminosity and temperature evolution of a large sample of Type II SNe and to compare our findings to theoretical models, focusing on the transition to the plateau, which takes place during the first 2 weeks.
Citation: Patat et al. 1994
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The strong radio flare observed on MJD 58523 is likely to be associated with the ejection of RK1, but the mechanism capable of launching and accelerating these relativistic plasma knots is still not understood. The current picture links the launch of the jets with the transition from the HS to the SS, i.e. during the IMS, where usually ejections take place prior to strong radio flares (e.g. Fender 2006; Miller-Jones et al. 2012) and are marked by changes in the X-ray timing and spectral properties (Belloni et al. 2005). RK1 fits very well in this picture, as the inferred ejection date tej, lin + Sedov = MJD 58518.9 ± 2.4 (see Section 3.3) is only 4 d before the strong radio flare observed with MeerKAT, and when MAXI J1348–630 was in the IMS. The presence of Type-B QPOs (tentatively linked to discrete ejections, e.g. Soleri et al. 2008; Fender et al. 2009; Miller-Jones et al. 2012; Homan et al. 2020) has been detected with NICER close to the first hard-to-soft state transition during the main outburst of MAXI J1348–630 (Belloni et al. 2020; Zhang et al. 2020a). In particular, Type-B QPOs started to be detected on MJD 58522.6 (Zhang et al. 2020a). This is ≳3 d after our inferred ejection date, similar to what was already observed for H1743–322 (Miller-Jones et al. 2012) and MAXI J1535–571 (Russell et al. 2019a). However, a detailed discussion of the NICER timing results in relation to the inferred ejection dates will be discussed in a forthcoming paper. The ∼5 arcsec MeerKAT resolution does not allow us to resolve the two components soon after the ejection, while the ATCA observations were short (providing a non-optimal uv-coverage). As the synchrotron-emitting component is launched and expands, it is predicted to have an optically thick rising phase due to an increasing surface area, followed by an optically thin decay produced by adiabatic expansion losses (van der Laan 1966). Flares produced by discrete ejecta are observed to rise on very different time-scales (from minutes or hours to days, e.g. Brocksopp et al. 2007; Tetarenko et al. 2017; Bright et al. 2020), depending on the size and energy of the ejected component. It is possible that the flare had a fast rise, and thus we missed the optically thick phase with our coverage. Another possibility is that the flare was entirely optically thin, as already observed for some sources (e.g. Fender et al. 1997). The ATCA observation taken on MJD 58524 and the decreasing flux observed on MJD 58523 (see Section 3.2) seem to confirm that we are observing the optically thin decay of the transient jet (see Fig. 10).
Citation: Belloni et al. 2020
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