adsabs/FOCAL huggingface dataset:
Studies of the radial n-capture-to-iron abundance gradients are very scarce so far. We can only search for a broad agreement of our results with several studies of abundance gradients with galactocentric distances (Rgc). da Silva et al. (2016) studied n-capture elements across the Galactic thin disc based on Cepheid variables. Because the Cepheids are young stars, their Rgc may be rather close to their birthplaces and Rmean. da Silva et al. (2016) supplemented their sample of 111 Cepheids with 324 more stars from other studies and found that the [Y/Fe] distribution is flat throughout the entire disc. In our study, we confirm this finding not only based on the whole thin-disc sample of stars and on a subsample of younger ≤4 Gyr stars, but also add another light s-process dominated element strontium. Like in our study, da Silva et al. (2016) also obtained positive [El/Fe] radial gradients for La, Ce, Nd, and Eu. The slopes are rather similar. For [Eu/Fe], they differ just by 0.002 dex kpc−1. More recently, Luck (2018b) also investigated the gradients of n-capture element abundance-to-iron ratios with respect to Rgc for a sample of 435 Cepheids. It is interesting to note that the [Ba/Fe] versus Rgc slope according to this Cepheid sample is also negative, as in our study. [Ba/Fe] is the only n-capture element-to-iron ratio with a negative radial gradient in our sample of stars and in Luck (2018b). Overbeek et al. (2016) investigated trends of Pr, Nd, and Eu to Fe abundance ratios with respect to Rgc using 23 open clusters. As in our study, they found that these elements have positive linear trends with galactocentric radius (the linear regression slopes are of about +0.04 dex kpc−1). They also suggested that the [El/Fe] relation of Pr and Nd, but not Eu, with the galactocentric radius may not be linear because the [El/Fe] of these elements appears to be enhanced around 10 kpc and drop around 12 kpc. Because only a small number of stars lie at these large radial distances, we cannot address this question. For the thick-disc stars, the radial abundance-to-iron slopes are negligible, as was found for α-process elements by Li et al. (2018), even though the production sites of α-elements and s-processes dominated elements are quite different.
Citation: Luck (2018b)
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Citation: Luck (2018b)
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A more puzzling task was the resolution of the fine and hyperfine structure. First, the A–E split structure attributed to the methyl internal rotation motion can be generally used as a molecular “fingerprint” to search for a molecule in an astronomical line survey (Cernicharo et al. 2016; Belloche et al. 2017). Moreover, it is known that for interstellar searches using surveys that done at centimeter wavelengths (i.e., those conducted with telescopes such as GBT and Yebes 40m), the recognition of hyperfine patterns in the observed spectra significantly help in the proper identification of interstellar molecules (McCarthy & McGuire 2021). As an example, the analysis of the hyperfine structure of several N-bearing species, such as N-protonated isocyanic acid (H2NCO+, Rodríguez-Almeida et al. 2021b) which has been detected toward G+0693-0.027 along with ethyl isocyante (C2H5NCO, Kolesniková et al. 2018), was required for its conclusive line-by-line identification. In this context, sensitive data at centimeter wavelengths are crucial because in TMC-1 and in G+0.693-0.027 the excitation temperatures of the molecules are very low, and hence the brightest lines of this relatively complex species fall at those frequencies. These interstellar sources have recently been suggested to be among the best targets to search for large molecular systems (McGuire et al. 2021; Rivilla et al. 2021b), further supporting centimeter-wave studies of as-yet-undiscovered interstellar candidates. Hence, we employed our narrowband LA-MB-FTMW spectrometer (León et al. 2021) to completely resolve both fine and hyperfine structures. We initially used the SPFIT/SPCAT program package (Pickett 1991) to analyze hyperfine patterns for the previously measured A and E lines of Z-AHA (see Fig. 3). This scrutiny was followed by measuring additional hyperfine components of several a- and b-type R-branch lines. Once we completed the analysis, new R-branch transitions belonging to the E-conformer were identified.
Citation: Rodríguez-Almeida et al. 2021b
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The sudden brightening of NGC 2617 provides a good opportunity to investigate the properties of CL-AGNs. In general, the CL processes can be interpreted as changes in obscuration or accretion rate. The obscuration model can be well applied to some X-ray selected CL-AGNs (e.g., Bianchi et al. 2005). The crossing time for an intervening object orbiting outside a BLR can be estimated as (LaMassa et al. 2015)
2





where rorb is the orbital radius of the foreground object on a circular, Keplerian orbit around the central black hole, M8 is the black hole mass in units of 108M⊙, and rsrc is the true size of the BLR. The above equation is derived by evaluating the time needed for this object to travel the length of an arc that corresponds to the projected size of the BLR. The observational constraints of rdt and rdt/rBLR are discussed in estimating tcross for the double-peaked BEL CL-AGN NGC 3516, where rdt is a dust torus radius (Feng et al. 2021, and references therein). rdt/rBLR ∼ 4–5 is obtained for most reverberation-mapped AGNs with the Hβ BLR and dust torus lag measurements. NGC 3516 has rdt/rBLR = 7. Thus, it seems appropriate to use rdt/rBLR = 5 in the CL-AGN NGC 2617. NGC 2617 has averages of rBLR = 6.15 lt-day and M8 = 0.22 for Hα and Hβ, and we obtain tcross = 7 yr. NGC 2617 is Seyfert 2 in 1994 and on 2003 December 30, and Seyfert 1 on 2013 April 25 (see Figure 2 in Shappee et al. 2014). Our observations show a Seyfert 1 type of optical spectra from 2019 October to 2020 May. The poor sampling of long-term spectroscopic observations makes it impossible to give the reliable durations of the appearance and “disappearance” of BELs. A long-term monitoring with a good sampling, e.g., twice a year, may give a reliable duration of “disappearance” of BELs in the CL processes, especially, the continuous and complete process of changing of Seyfert 1 → Seyfert 2 → Seyfert 1. This will be important to test the obscuration model of CL-AGNs.
Citation: Feng et al. 2021
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The majority of observational evidence based on studies of SN Ia populations thus points towards a DD scenario for most, if not all, SNe Ia. However, finding self-consistent progenitor and explosion models that recreate the observed luminosity function as well as correlations between luminosity, light-curve parameters, and host galaxy properties has proven difficult. In particular, simulations based around explosions of MCh WDs (linked strongly with the SD scenario but also with many DD scenarios) find difficulty in reproducing the light curves of ‘normal’ SNe Ia as well as ‘peculiar’ objects (Ropke et al. 2007; Sim et al. 2013; see Maoz et al. 2014; Blondin et al. 2017; Jha, Maguire & Sullivan 2019 for overviews). In recent years, attention directed towards explosions of sub-MCh WDs triggered by double detonations (primarily related to a DD scenario) has led to promising results (e.g. Shen, Toonen & Graur 2017; Shen et al. 2018; Townsley et al. 2019; Gronow et al. 2020; Shen et al. 2021) although they still struggle to match observations at late times in the light-curve evolution (Gronow et al. 2021; Shen et al. 2021). An additional factor in support of the sub-MCh model is that the SN luminosity is related to the mass of the primary WD, which itself is likely to be related to its age (although this relation is probably complicated by other factors such as accretion rate, metallicity, and the composition of the companion), thereby providing an explanation for observed relation between light-curve stretch and stellar age (Rigault et al. 2013, 2020; Rose, Garnavich & Berg 2019; Nicolas et al. 2021). Other proposed scenarios include hybrid models in which standard CO WDs merge with hybrid helium-CO WDs (Zenati, Toonen & Perets 2019). With many models showing promising similarities to observations but each subject to its own drawbacks, it is becoming accepted that more than one progenitor scenario may contribute significantly to the overall population of ‘normal’ SNe Ia; detailed observations are thus required in order to place constraints on the relative fractions of each possible progenitor channel.
Citation: Gronow et al. 2020
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Associations between ULXs and star clusters have also been studied in interacting galaxies. The latter are known to host a higher average number (>5) of ULXs. Therefore, they are good candidates in which to examine the properties of the population of ULXs. Poutanen et al. (2013) extensively examined the significant associations between ULXs and stellar clusters in the Antennae galaxies. Using data from HST and the Very Large Telescope supplemented with theoretical stellar isochrones, they estimated the ages of these clusters as 6 Myr. It was discussed that these ULXs were probably ejected from the cluster in the evolutionary process; thus these sources might be high-mass X-ray binaries instead of intermediate-mass black holes. Another well-known interacting galaxy is NGC 4490/NGC 4485 at a distance of 7.8 Mpc (Tully 1988). NGC 4490 is a late-type spiral galaxy and NGC 4485 is an irregular galaxy. Their linear sizes are 15 kpc for NGC 4490 and 5.6 kpc for NGC 4485. Radio observations show that star formation in NGC 4490 has been ongoing at a constant rate of ∼4.7 M yr−1 (Clemens et al. 1999). Also, NGC 4490 has a giant H i envelope that probably originated from the star formation (Clemens & Alexander 2002). Our aim in this study is to identify the possible optical counterparts of ULXs in this galaxy pair and to investigate their associations with star groups or clusters. Previously, three ULXs in this pair were detected by ROSAT HRI observations (Roberts & Warwick 2000). Later, using a Chandra ACIS-S observation, three more ULXs were identified by Roberts et al. (2002). The calculated unabsorbed luminosities of six ULXs (in the 0.5–8 keV band) fall into the range (2.6–4.9) × 1039 erg s−1. In addition, Roberts et al. noted that these ULXs appear to be spatially coincident with the star formation regions in the pair. Further, Fridriksson et al. (2008) searched for long-term variability of 38 X-ray sources in this galaxy pair using three Chandra observations. Eight of these sources were classified as ULXs in the luminosity range (0.6–3) × 1039 erg s−1. One of them is a transient ULX detected in a single observation (ID 4726). Gladstone & Roberts (2009) investigated spectral and temporal features of seven ULXs (except for ULX X-5—this source was ignored because of its low luminosity of LX ∼ 6 × 1038 erg s−1, which was given in Table 5 of Fridriksson et al. 2008) using the same Chandra and XMM-Newton data sets. The LX values of these sources are given in the range (0.9–4) × 1039 erg s−1 within the 0.5−8 keV energy band. Six of these seven sources (except the transient one, which is X-7 in this paper) were classified as ULXs by Swartz et al. (2011).
Citation: Clemens & Alexander 2002
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Chandrasekhar & Fermi (1953) assumed that “turbulent motions are isotropic”, and they adopted 

$f\,{=}\,1/\sqrt{3}$f = 1/3
. If the field strength is weak, turbulent motions will drag the field lines in random directions and turbulence will be isotropic (super-Alfvénic turbulence). However, there is overwhelming observational evidence that magnetic fields in the ISM have well-defined directions, indicating that turbulence is sub, trans-Alfvénic and hence turbulent properties are highly anisotropic (see, for example, Montgomery & Turner 1981; Shebalin et al. 1983; Higdon 1984; Sridhar & Goldreich 1994; Goldreich & Sridhar 1995, 1997). Heyer et al. (2008), using CO data, found that velocity structures in Taurus are highly anisotropic. In the same region, Goldsmith et al. (2008) reported the existence of highly anisotropic density structures, which are aligned parallel to the mean field, known as striations. Striations have also been observed in the Polaris Flare (Panopoulou et al. 2015) and Musca (Cox et al. 2016; Tritsis & Tassis 2018), and they are formed due to magnetosonic waves (Tritsis & Tassis 2016) in sub-Alfvénic turbulence (Beattie & Federrath 2020). More evidence for ordered magnetic fields in molecular clouds can be found in Franco et al. (2010), Franco & Alves (2015), Pillai et al. (2015), Hoq et al. (2017), and Tang et al. (2019). Stephens et al. (2011) explored the magnetic field properties of 52 star forming regions in our Galaxy and concluded that more than 80% of their targets exhibit ordered magnetic fields. The diffuse atomic clouds in our Galaxy are preferentially aligned with the magnetic field (Clark et al. 2014), implying the importance of the magnetic field in their formation. Planck Collaboration Int. XXXV (2016) studied a larger sample of molecular clouds in the Goult Belt and concluded that density structures align parallel or perpendicular to the local mean field direction. This is also consistent with sub, trans-Alfvénic turbulence (e.g., Soler et al. 2013). In addition, Mouschovias et al. (2006), using Zeeman data, concluded that turbulence in molecular clouds is slightly sub-Alfvénic as well. All these lines of evidence indicate that ISM turbulence is sub, trans-Alfvénic, and hence anisotropic.
Citation: Shebalin et al. 1983
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An alternative possibility to coronal heating is that the release and dissipation of magnetic energy actually takes place in the chromosphere. In this scenario, the chromospheric plasma is directly heated to coronal temperatures, rather than by thermal conduction fronts that are driven into the lower atmosphere as a consequence of heating localized in the corona (Hansteen et al. 2010). One manifestation of this process might be the type II spicules (De Pontieu et al. 2007) that have been suggested to play a role in supplying mass and energy to the corona (De Pontieu et al. 2009, 2011; Moore et al. 2011). N. E. Raouafi et al. (2015, in preparation) have argued that type II spicules have much more in common with so-called “classical” spicules (Beckers 1968, 1972; Pasachoff et al. 2009) than do type I spicules. In particular, type II and classical spicules are both commonly observed in quiet Sun and coronal hole regions, whereas type I spicules appear to be exclusively confined to active regions. Furthermore, Pereira et al. (2013) artificially coarsened Hinode/SOT observations of type II spicules to demonstrate that their lifetimes and ejection speeds are consistent with the earlier, ground-based observations of classical spicules (∼5 minutes and 25 km s−1). (N. E. Raouafi et al. 2015, in preparation) suggest that the term “classical spicules” be used for the earlier objects (see Sterling 2000, for a review) and the terms type I and II spicules be reserved for spicular phenomena observed during the era of Hinode observations. The key properties of type II spicules are their faster velocities (30–110 km s−1) and shorter lifetimes (50–150 s) compared with type I spicules. Type II spicules are typically observed in Ca ii emission by SOT before fading out of that passband and appearing in the SDO/AIA 304 Å channel, indicating that some degree of heating takes place as they rise. There is some observational evidence for a transition region or coronal component of their emission, visible as a bright, moving front in the AIA 171 Å channel (e.g., De Pontieu et al. 2011). The two possibilities for producing this warmer emission are: (1) pre-existing coronal material is shock-heated as the upflowing spicular material rams into it (Klimchuk 2012; Petralia et al. 2014a); or (2) the tip of the spicule is heated by some in situ process, which must be impulsive because the cooler emission quickly disappears from view (De Pontieu et al. 2007).
Citation: Pasachoff et al. 2009
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The mass estimate for each cluster was obtained by fitting a reduced tangential shear profile predicted by a projected Navarro–Frenk–White (NFW) profile (e.g. Bartelmann 1996) to the observed ellipticities. We derive the best-fitting profile parameters R200 and c200 by minimizing the merit function 
(4)
\begin{eqnarray*}
\chi ^{2}\!=\!\sum _{i=1}^{N}{\frac{\left|g_{i}(\theta _{i},\beta _i;R_{200},c_{\mathrm{NFW}})\!-\! \tilde{\varepsilon }_{\mathrm{t},i}(\theta _{i})\right|^{2}}{\tilde{\sigma }_{\!i}^{2}\left(1\!-\!\left| g_{i}(\theta _{i},\beta _i;R_{200},c_{\mathrm{NFW}})\right|^{2}\right)^{2}}}. 
\end{eqnarray*}
Here gi(θi, Σcrit, i;; R200, cNFW) is the model prediction for galaxy i and $\tilde{\varepsilon }_{\mathrm{t},i}$ the observed ellipticity times 1.08 for the same galaxy. The factor 1.08 is the multiplicative shear calibration bias of the used KSB+ pipeline (Kaiser et al. 1995; Erben et al. 2001) to convert from measured to true ellipticity. This calibration bias has an uncertainty of ∼5 per cent. This uncertainty is a dominant source of systematic uncertainty in the mass measurements. Each shear profile was centred on the BCG, using distances in the range of 0.2–4.2 Mpc for the fitting procedure. We minimized the χ2 on a grid of R200 and c200. Finally, we used the mass–concentration relation described by Bhattacharya et al. (2013) to put priors on the concentration parameter to break the degeneracies in the profile models. The initial mass estimates from equation (4) are biased. In evaluating the NFW shear profile we make use of the ratio in equation (2) when averaging the value of β over the reference catalogue sources. However, $\frac{\gamma (\langle \beta \rangle)}{1-\kappa (\langle \beta \rangle)}\ne \left\langle \frac{\gamma (\beta)}{1-\kappa (\beta)}\right\rangle$. Given the finite width of the β distribution that are averaged over when calculating βi from a reference catalogue (equation 3), we find a biased point estimate for βi. Especially in the inner regions of the cluster, this would model the shear profile incorrectly. We estimate the final masses by correcting for the averaging over β in two subsequent iterations. We utilize the best-fitting mass estimate from the zeroth iteration to predict the reduced shear, g, at the projected distance θ from the cluster centre, and βk. We then introduce $\beta _i^{\prime }$, which satisfies the equation: 
(5)
\begin{eqnarray*}
g(\beta _{i}^{\prime })=\!\frac{\sum _{k=1}^{N}w_kg(\theta _{i},\beta _k)}{\sum _{k=1}^{N}w_k}\frac{1}{v_{b}(c_{1},c_{2})}. 
\end{eqnarray*}
Citation: Kaiser et al. 1995
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One way to explain the observations is to postulate that the cores of planets are formed in the very first Myr of the protoplanetary disk evolution, or even in the embedded phase while the disk is still forming. The disks for which masses have been measured have ages >1 Myr, and a general trend of declining disk mass with ages older than 1 Myr has been observed (e.g., Barenfeld et al. 2016). Thus, it is possible that disks were massive enough to form the cores of planets at younger ages. This idea has been suggested by Greaves & Rice (2010), Williams (2012), and Najita & Kenyon (2014), among others. In this scenario, the vast majority of the material composing planets must already be in the form of planetesimals, for rocky planet formation, and of already-formed planetary cores. The latter condition is necessary as gas giant planets need to accrete gas from the gas-rich disk. Assuming a gas-to-dust ratio of 100, disks at ∼1–3 Myr have just the right amount of gas mass to explain the population of gas giants (see Fig. 1). Thus, cores must already be in place at this age. Scenarios have been proposed to explain that pebble accretion can form planetesimals very early (0.1 Myr) in disks, when these are massive and possibly gravitationally unstable (Booth & Clarke 2016). However, it is expected that the formation of planetary cores is highly inefficient, with ∼350 M⊕ of pebbles needed to grow the core of Jupiter from half a lunar mass to 20 M⊕ (e.g., Morbidelli et al. 2016). Thus, such an inefficiency would imply that disks were initially ∼10–100 times more massive than is observed at ages >1 Myr. This would also imply that the vast majority of disks were initially gravitationally unstable. One issue of this scenario is that, if early disks are ∼10–100 times more massive than observed for disks older than 1 Myr, an extremely efficient gas-removal mechanism has to be found, consistent with the observations of Class 0 outflows (Frank et al. 2014).
Citation: Greaves & Rice (2010)
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It is now largely understood that the dust-to-gas mass ratio ϵ in protoplanetary discs is unlikely to match the canonical value in the interstellar medium of 0.01 (Mathis, Rumpl & Nordsieck 1977). For example, recent observations of CO isotopologues and dust continuum emission have found disc-averaged ϵ values to be much higher (∼0.2; see Ansdell et al. 2016), and simulations have demonstrated that discs can be formed with significant dust enrichment (Lebreuilly, Commerçon & Laibe 2020). There is also an observed discrepancy between the radial extents of the gas and dust components of protoplanetary discs, with dust discs being significantly more compact (Andrews et al. 2012; Pérez et al. 2012, 2015b). In some cases, such as IM Lup, the gas disc may be 10 times larger in radial extent than the mm dust disc (e.g. Cleeves et al. 2016). This discrepancy can explained by a combination of grain growth and radial drift. Essentially, the radial velocity, vr, of dust particles has two components, namely drag, vdrag, and drift, vdrift, and is given by
(1)$$\begin{eqnarray*}
v_{\mathrm{r}}=v_{\mathrm{drag}}+v_{\mathrm{drift}}=\frac{v_{\mathrm{g}, \mathrm{r}}}{1+{St}^{2}}+\frac{1}{{St}+{St}^{-1}} \frac{1}{\rho _{\mathrm{g}} \Omega } \frac{\partial P}{\partial R},
\end{eqnarray*}$$where vg, r is the radial velocity of the gas, ρg is the gas mass surface density, Ω is the Keplerian angular velocity, and P is the pressure (Weidenschilling 1977). The Stokes number, St, determines the degree of coupling to the gas. For a vertically isothermal disc (as used here), this is expressed simply as
(2)$$\begin{eqnarray*}
St = \frac{\pi a\rho _\mathrm{s}}{2\Sigma _\mathrm{g}},
\end{eqnarray*}$$where a is grain size, ρs is intrinsic bulk grain density, and Σg is the gas surface density. Examining equations (1) and (2 ), we can see that the drag term decreases as grain size increases. However, the drift term peaks at St = 1, which means that the grain size for which radial drift is fastest depends on the surface density, and therefore location within the disc.
Citation: Andrews et al. 2012
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