adsabs/FOCAL huggingface dataset:
With this study, our group continues the effort to investigate the optical properties of cosmic dust analogues in the mid infrared (MIR) to the millimeter domain as a function of temperature, undertaken 20 years ago by different groups (we refer to Demyk et al. 2013, for the details of the studied samples). Briefly, Agladze et al. (1996) were the first to study relevant interstellar silicate dust analogues in the temperature range from 1.2 K to 30 K and in the wavelength range from 700 μm to 2.9 mm. Mennella et al. (1998) studied amorphous carbon samples and silicate samples in the 24 μm−2 mm spectral domain and 24−300 K temperature range. Boudet et al. (2005) investigated silica and silicate samples in the 10−300 K temperature range and in the spectral region 100−1000 μm. The work by Coupeaud et al. (2011) was focussed on pure sol-gel Mg-rich silicates, of composition close to enstatite (MgSiO3) and forsterite (Mg2SiO4), amorphous and crystalline, whose spectra were recorded in the 100−1000 μm spectral range and from 10 K to 300 K. These studies have brought important results about the spectroscopic characteristics and behavior of interstellar dust analogues in the FIR at varying temperature. They show that the MAC of the amorphous analogues (silicates and carbonaceous matter) increases with the grain temperature and that its spectral shape cannot be approximated with a single power law in the form λ− β. The dependence of the MAC on the temperature, which is not observed in crystalline samples, is related to the amorphous nature and to the amount of defects in the structure of the material (Coupeaud et al. 2011). A physical model was proposed by Meny et al. (2007) to explain these experimental results. This model, named the TLS model, is based on a description of the amorphous structure of the material in terms of a temperature independent disordered charge distribution and of a collection of atomic configurations modeled as two-level systems and sensitive to the temperature.
Citation: Mennella et al. (1998)
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Although most galaxies have an implied SFR that scatters within a factor two around the MS, some do show a significantly higher SFR. Those objects also exhibit a higher gas content, shorter gas depletion times (e.g., Genzel et al. 2015; Schinnerer et al. 2016; Tacconi et al. 2013, 2018), and higher dust temperatures (e.g., Magnelli et al. 2014). Likewise, the stellar-light radial distribution is different in these two galaxy populations; while MS galaxies are closely approximated by exponential disks (e.g., Bremer et al. 2018), those above (and below) it exhibit a higher central mass concentration (e.g., Wuyts et al. 2011). Based on this dichotomy and the parametrization of the MS over cosmic time, a scenario has been proposed to explain the evolutionary path of galaxies along the MS. Since the normalization of the MS, the gas fraction of galaxies, and cosmic molecular gas density decrease from z ∼ 2.5 to 0 at a similar pace (e.g., Speagle et al. 2014; Decarli et al. 2016; Tacconi et al. 2018), it is thought that MS galaxies evolved through a steady mode of star formation, possibly regulated by the accretion of cool gas from the intergalactic medium (e.g., Dekel et al. 2009; Kereš et al. 2009; Davé et al. 2010; Hodge et al. 2012; Romano-Díaz et al. 2014, 2017; Feng et al. 2015; Anglés-Alcázar et al. 2017). From theoretical predictions, the scatter of the MS could thus be explained as the result of a fluctuating gas inflow rate that is different in each galaxy (e.g., Tacchella et al. 2016; Mitra et al. 2017). In this context, a galaxy enhances its SFR and moves towards the upper envelope of the MS due to gas compaction. As the gas is depleted, the SFR decreases and the galaxy falls below the MS. As long as a SFG is replenished with fresh gas within a timescale shorter than its depletion time, it will be confined within the MS (Tacchella et al. 2016). On the other hand, the enhanced star formation efficiency of galaxies above the MS has been linked to mergers (e.g., Walter et al. 2009; Narayanan et al. 2010; Hayward et al. 2011; Alaghband-Zadeh et al. 2012; Riechers et al. 2013, 2014) and instability episodes in gas-rich disks (particularly at high redshift; e.g., Davé et al. 2010; Hodge et al. 2012; Wang et al. 2019).
Citation: Tacconi et al. 2018
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Other important components of the X-ray sky are X-ray binaries, the hot gas present in our own Galaxy and the extragalactic hot gas. Amongst these X-ray sources, the large hot gas reservoir (Tvir ∼ 107 K) filling the space between the galaxies in the clusters, known as the ICM, has been observed in its X-ray emission for a long time (Reichert et al. 1981; Jones & Forman 1984; Branduardi-Raymont et al. 1985; Oukbir, Bartlett & Blanchard 1997; Diego et al. 2003a; Diego, Silk & Sliwa 2003b; Cavagnolo et al. 2009; Hurier et al. 2015). However, the hot gas (Tvir > 106 K) present in the form of circumgalactic medium (CGM) in massive galaxies (Mh ∼ 1012–1013 h−1 M⊙; Birnboim & Dekel 2003; Kereš et al. 2005; Singh et al. 2015) is less explored in X-rays due to its fainter X-ray emission. Some of the recent observations and studies (Grcevich & Putman 2009; Anderson & Bregman 2011; Dai et al. 2012; Putman, Peek & Joung 2012; Anderson, Bregman & Dai 2013; Bogdán et al. 2013a,b; Gatto et al. 2013) indicate that the CGM can account for a good fraction of the baryons in these galaxies. The X-ray emission from the CGM, therefore, is a promising tool to put strong constraints on the distribution and energetics of the gas (Singh et al. 2016) with eROSITA. At energies above 2 keV, the extragalactic point sources like AGN completely dominate the X-ray sky (Lehmann et al. 2001; Kim et al. 2007). Even below 2 keV, where the X-ray emission from the hot gas in the ICM and CGM is significant, the major contribution to the observed X-ray sky comes from AGNs (Sołtan 2007). Therefore, studying the X-ray emission from the AGN is crucial to understand the origin and evolution of the AGN as well as to extract the X-ray signal from the subdominant components. We thus also compute the angular power spectrum of the unresolved AGNs that are expected to contribute to the diffuse X-ray background of eROSITA and contaminate the angular power spectrum due the ICM/CGM in the 0.5–2 keV X-ray band.
Citation: Bogdán et al. 2013a
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For the first time, we detected [C ii] 158-μm emission from a GRB host galaxy at z > 2. This is the second detection of [C ii] 158-μm emission among known GRB host galaxies, following GRB 980425 (Michałowski et al. 2016). The [C ii] 158-μm fine structure line is the dominant cooling line of the cool interstellar medium, arising from photodissociation regions (PDR) on molecular cloud surfaces. It is one of the brightest emission lines from star-forming galaxies from FIR to metre wavelengths, almost unaffected by dust extinction. [C ii] 158-μm luminosity, L[C II], has been discussed as an indicator of SFR (e.g. Stacey et al. 2010). If L[C II] scales linearly with SFR, the ratio to FIR luminosity, L[C II]/LFIR, is expected to be constant, since LFIR is a linear function of SFR (e.g. Kennicutt 1998a). However, LC II/LFIR is not constant, but declines with increasing LFIR, known as the ‘[C ii] deficit’ (e.g. Luhman et al. 1998, 2003; Malhotra et al. 2001; Sargsyan et al. 2012; Díaz-Santos et al. 2013, 2017; Spilker et al. 2016). The [C ii] deficit persists when including high-z galaxies (e.g. Stacey et al. 2010; Wang et al. 2013; Rawle et al. 2014). In Fig. 5, we compare the [C ii] deficit in the GRB 080207 host and other star-forming galaxies. Two GRB hosts are shown by stars: GRB 080207 (orange star) and 980425 (blue star). The comparison sample is compiled from the literature up to z ∼ 3 (Malhotra et al. 2001; Cormier et al. 2010, 2014; Ivison et al. 2010; Stacey et al. 2010; Sargsyan et al. 2012; Farrah et al. 2013; Magdis et al. 2014; Brisbin et al. 2015; Gullberg et al. 2015; Schaerer et al. 2015). Active galactic nuclei are separated from star-forming galaxies based on either (i) the explicit description in the literature or (ii) EWPAH 6.2μm  0.1 (Sargsyan et al. 2012). As reported by previous studies (e.g. Maiolino et al. 2009; Stacey et al. 2010), high-z galaxies are located at a different place from local galaxies in the L[C II]/LFIR–LFIR plane.
Citation: Stacey et al. 2010
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The detailed kinematic reconstruction of the Sgr tidal debris by Law & Majewski (2010) used an initial total mass of 





 M⊙ for the Sgr satellite. However, several studies point to a Sgr remnant mass significantly exceeding that value. Ibata et al. (1997) and Ibata & Lewis (1998) estimate lower bounds of 109 M⊙ for the mass of the dwarf today. With the discovery of previously unseen branches of the stream, the total luminosity budget of the progenitor galaxy is now believed to be on the order of 108 L⊙ (Niederste-Ostholt et al. 2010). As a result, recent studies have shifted to using dark-matter halo masses as large as 1011 M⊙ (Purcell et al. 2011; Gòmez et al. 2015) based on halo abundance matching arguments. Such high values are comparable to the mass of the LMC progenitor (see e.g., Jethwa et al. 2016; Peñarrubia et al. 2016) and imply a mass ratio relative to the MW on the order of 1:10. However, unlike the Magellanic Clouds, which may be on their first passage near the MW (Besla et al. 2007), Sgr is known to have experienced multiple close passages in the past. If true, such high Sgr progenitor masses would have important implications for the formation and evolution of the MW disk (e.g., Purcell et al. 2011; Gòmez et al. 2013; D’Onghia et al. 2016). Because the dynamical friction force is proportional to the square of the satellite mass, we expect drag to play a much more important role in slowing down the Sgr satellite and bringing it to closer galactocentric distances. Here, we perform an exploration across orbital angular momentum parameter space analogous to Section 3.1, this time using a Sgr progenitor mass of 6×1010 M⊙ according to the recent estimates of Gibbons et al. (2017). We consider two possibilities: a “slow sinking” scenario in which, as in Section 3, Sgr crosses the MW virial radius at 





 (approximately 8 Gyr ago), and a “rapid sinking” scenario, in which we examine a first infall around 





, about 4 Gyr ago.
Citation: Purcell et al. 2011
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For many hot channel eruption events, their exact origin and initiation process still remain under debate. Some observations supported that before the hot channel eruption, a corresponding MFR may already exist (e.g., Liu et al. 2016; Wang et al. 2016; Yang et al. 2019) and slow photospheric flows might be important for its gradual formation (e.g., Hou et al. 2018; Yan et al. 2018; Vasantharaju et al. 2019); while others argued that they can also be newly built up via rapid coronal reconnection during the eruption (Cheng et al. 2011; Song et al. 2014; Palacios et al. 2015). For their loss-of-equilibrium, some researchers found that the loss-of-equilibrium of the hot channel is initially facilitated by the breakout type (e.g., Shen et al. 2012; Chen et al. 2016; Mitra et al. 2018) or tether-cutting type (e.g., Chen et al. 2014, 2018; Joshi et al. 2015; Dudík et al. 2016) preflare reconnection, but others believed the loss-of-equilibrium of the hot channel may be directly triggered by the ideal MHD instabilities (e.g., Liu et al. 2007; Bi et al. 2015; Vemareddy et al. 2017). Note that Duan et al. (2019) recently carried out a survey on the initiation mechanism of all the major solar flares, either eruptive or confined, from 2011 to 2017. Statistically, their study results indicate that magnetic reconnection and ideal MHD instabilities seemingly play an equal role in the triggering of major flares and eruptions. In the present study, our NLFFF extrapolation and topology analysis shows that a seed MFR existed within an HFT configuration prior to the later solar eruption. Near the HFT, two sets of sheared arcades serve as an enveloping field holding the seed MFR, and have rooted their opposite footpoints at canceling fluxes. Combining these results with imaging observations, we suggest that the preflare 3D reconnection is most likely first triggered at the HFT in the precursor phase due to the preceding flux cancellation (Savcheva et al. 2012a), and the newborn hot channel should originate from the preexisting seed MFR. Through preflare 3D reconnection, the seed MFR was rapidly heated up to 10 Mk manifesting as a newborn hot channel, and its spatial size also underwent an enlarging process at the same time (as revealed by Figures 2 and 5). As a result, the enlarged hot channel may readily lift into the abovementioned torus-unstable domain (around 12 Mm above the photosphere), and soon enter its exponential accelerated eruption.
Citation: Dudík et al. 2016
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As in Flowers et al. (2019), to compare our model transmission spectra directly against the Ehrenreich et al. (2020) results, we must calculate the transmission spectra as a function of orbital phase throughout the duration of transit. To account for orbital phase dependencies, we apply the following procedure:1.Account for phase-dependent backlighting of the planet (i.e., stellar limb-darkening effects). At different points of its transit, the planet will occult regions of its host star of varying brightness. Furthermore, at a fixed orbital phase, different regions of the planet’s limb will be backlit by varying intensities of stellar light. Similar to Flowers et al. (2019), we calculate the normalized stellar intensity at the center of each cell of the 2D projected planetary grid produced by our GCM at each modeled orbital phase of the planet. We use the quadratic limb-darkening coefficients reported by Ehrenreich et al. (2020) to establish the stellar center-to-limb intensity profile, and we take into account the 89.°623 orbital inclination of WASP-76b (Ehrenreich et al. 2020) to determine where the planet resides on the stellar disk as a function of its orbital phase. We make the assumption of constant impact parameter b over the course of transit.
9


9
In reality, a planet on an inclined orbit will not have a constant b over the entire duration of transit; rather, the planet’s distance from the stellar equator will be decreased at ingress and egress, reaching its maximum at center of transit. Our tests reveal that, for WASP-76b, the relative error induced by the constant b assumption is on the order of 4% in distance, which results in a change on the order of 1 m s−1 at the blueshift level (see Section 3.1). Hence, our b treatment is justified. This procedure allows us to calculate a backlighting factor f, which ranges from 0 to 1, effectively replacing the constant I

λ,0 from Equation (3) with a variable 



Iλ,0×f(θ′,z,φ)

, for a given orbital phase φ and 2D projected polar angle 



θ′

.2.Account for the decreasing of the continuum by interpolating a light curve produced by the batman code (Kreidberg 2015). Step 1 ensures that less light is transmitted through the planet’s atmosphere than a uniform stellar disk would emit. Step 2 further enforces that the inner, optically thick core of the planet is simulated crossing a limb-darkened star, as opposed to a star of uniform brightness.3.Account for the planet’s rotation over the course of transit. Because the planet is continually rotating as it travels across the face of its host star, we must transform the GCM coordinate system so that the correct observer-facing hemisphere is modeled at each instance during transit. For simplicity, we assume zero obliquity, which allows us to calculate the coordinate transform simply by assigning a linear offset to each planetary longitude; i.e., ϕ
rotated = ϕ + φ.
Citation: Ehrenreich et al. 2020
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Citation: Ehrenreich et al. (2020)
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The physical parameters governing the MS are still being investigated, although some basic inferences have been made. The main parameters that describe quasars as accreting black holes are the black hole mass (MBH, ranging from 106$\rm M_{\rm \odot }$ to 109.5$\rm M_{\rm \odot }$; Kormendy & Richstone 1995), the accretion luminosity, Eddington ratio (LBol/LEdd, Boroson & Green 1992; Sulentic et al. 2000b; Marziani et al. 2001, 2003b), and the black hole spin (e.g. Wang et al. 2014). The MBH can be estimated by assuming that the gas motions are predominantly Keplerian around the black hole, and by applying the virial theorem for a system whose mass is entirely concentrated in the center of gravity. Reverberation mapping provides a measurement of the radial distance rBLR of the line emitting gas from the central black hole (Peterson 1993, 2004). The so-called ‘virial mass’ have been computed for large samples of quasars employing several different emission lines, various measures of line width, and exploiting correlation between the emitting region radius and luminosity (see Marziani & Sulentic 2012; Shen 2013; Popović 2020, for reviews). An independent method relies on the scaling law between MBH and the stellar velocity dispersion of the galaxy bulge (σ⋆; Kormendy & Ho 2013). Recently, spectropolarimetric observations of the Hydrogen Balmer lines have allowed to measure the intrinsic line FWHM due to a Keplerian velocity field and to compute the black hole mass independently from orientation (Afanasiev & Popović 2015; Afanasiev, Popović & Shapovalova 2019, and references therein). All methods to compute black hole mass in active galactic nucleus (AGN) are subject to caveats and suffer considerable uncertainties (e.g. Dalla Bontà et al. 2020, and references therein). None the less, the LBol/LEdd, which is proportional to the luminosity-to-black hole mass ratio (L/MBH), has been revealed to be a fundamental driver of the E1 MS, closely related to several observational parameters (Marziani et al. 2001, 2003b; Kuraszkiewicz et al. 2004; Shen & Ho 2014; Panda et al. 2018; Panda, Marziani & Czerny 2019). With decreasing Eddington ratio, source properties change from the ones of Pop. A to the ones of Pop. B, which tend to show broader Balmer emission lines, weaker RFeII,1 and more asymmetric Balmer line profiles (Sulentic et al. 2000b; Marziani et al. 2003b; Shen & Ho 2014).
Citation: Peterson; 2004
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SN 2011fe, discovered a mere ≈11 h after explosion (Nugent et al. 2011) by the Palomar Transient Facility (PTF; Law et al. 2009), is the brightest SN Ia since the advent of modern astronomical detectors. Located at just $d_L \approx 6.5~\rm {Mpc}$ (e.g. Shappee & Stanek 2011; Beaton et al. 2019), SN 2011fe exploded in a region of M101 uncontaminated by intervening dust (Patat et al. 2013) providing an ideal location for testing SN Ia progenitor and explosion models. The early detection allowed extensive follow-up observations across the electromagnetic spectrum (e.g. Matheson et al. 2012; Parrent et al. 2012; Pereira et al. 2013; Hsiao et al. 2013; Johansson, Amanullah & Goobar 2013; Tsvetkov et al. 2013; Munari et al. 2013; Mazzali et al. 2014; Zhang et al. 2016) and provided direct constraints on the radius of the exploding star (Nugent et al. 2011; Bloom et al. 2012). Stringent non-detections in radio (Chomiuk et al. 2012; Horesh et al. 2012; Kundu et al. 2017) and X-ray (Horesh et al. 2012; Margutti et al. 2012) observations exclude nearby CSM at high significance. Early ultraviolet (UV) photometry did not show any evidence for the ejecta encountering a nearby companion star (Brown et al. 2012) and nebular spectra lacked the Balmer emission lines from material ablated off the donor star by the ejecta impact (Shappee et al. 2013b; Lundqvist et al. 2015; Tucker et al. 2022). Pre-explosion imaging excludes the presence of a RG or He donor star (Li et al. 2011) and disfavor an accreting WD in the ∼105 yr prior to explosion (Graur, Maoz & Shara 2014). Multi-epoch spectropolarimetry reveal consistently-low continuum polarization suggestive of a symmetric ejecta distribution with evidence for minor oblateness (Milne et al. 2017). Finally, nebular-phase observations at ≳ 1 year after maximum light allows a direct view to the inner ejecta and provides unique constraints on the explosion conditions (McClelland et al. 2013; Kerzendorf et al. 2014, 2017; Mazzali et al. 2015; Graham et al. 2015; Taubenberger et al. 2015; Dimitriadis et al. 2017; Friesen et al. 2017; Shappee et al. 2017; Tucker et al. 2022). SN 2011fe is one of the best-studied astronomical objects in the past decade and remains a key benchmark for any SN Ia theory or model.
Citation: Matheson et al. 2012
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Various studies have investigated the color-magnitude relation at high redshifts (z ≳ 2), and the results are inconclusive. Reddy et al. (2008) found weak-to-no correlation between dust or UV color (related as shown in Fig. 4) and R mag for UV-selected galaxies at 1.5 ≲ z ≲ 2.6. Similarly, Heinis et al. (2013) through their investigation of far-infrared/dust properties of UV selected galaxies at z ~ 1.5 found that the average UV slope is mostly independent of the UV luminosity. On the other hand, Bouwens et al. (2009) studied the relation between the UV slope and MUV for a sample of U-band dropout galaxies at z ≃ 2.5 and found a positive color-magnitude correlation. At higher redshifts (z > 3), similar disagreements have been observed for samples based on the Lyman break color selection, in the sense that, some studies find no (or very weak) trends between β and UV magnitude (e.g., Ono et al. 2010; Finkelstein et al. 2012; Dunlop et al. 2012; Castellano et al. 2012), while other studies find a strong color-magnitude relation (e.g., Wilkins et al. 2011; Bouwens et al. 2012; Rogers et al. 2014; Bouwens et al. 2014). Bouwens et al. (2014) also argue that this relation could be described by a double power-law with a different slope for fainter magnitudes. The differences in these various studies could be because of many systematic and/or physical reasons, such as selection of galaxy samples (especially at lower redshifts), UV magnitude and/or β measurements, intrinsic scatter in β, dynamic range in UV luminosity. Our study relies on a UV selected sample with spectroscopic redshifts but it lacks a significant dynamic range in UV luminosity as seen in Fig. 5. A detailed study of a uniformly selected spectroscopic sample covering a wider range in UV luminosity is needed to understand the true nature of the color-magnitude relation at high redshifts. At least for ~L∗ (~0.5L∗ – 3.0L∗) galaxies in a uniformly selected VUDS sample at 2  z  2.5 there appears to be no correlation between β and M1500.
Citation: Bouwens et al. (2009)
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