adsabs/FOCAL huggingface dataset:
We detected radio emission for six sources, but while the morphology of radio emission in four over six cases is compact on $\le \,$arcseconds scales, predominantly unresolved, for NGC 3185 and NGC 3941 we find a more complex morphology. The estimated radio luminosities at this frequency are L$\, \sim \, 5\, \times \, 10^{20}$ and ${\sim }\, 6\, \times \, 10^{18}$ W Hz−1 for NGC 3185 and NGC 3941, respectively, well below the radio-loud/radio-quiet threshold of L$\, \sim \, 10^{23}$ W Hz−1 defined by Condon (1992). Considering NGC 3185, the structure observed in L band has a size of $5\, \times \,$3.8 arcsec, which translates into a linear scale of $0.5\, \times \,$0.4 kpc.6 Analogous considerations can be made for the C and X Bands, leading to linear scales of $0.38\, \times \,$0.36 kpc. There is evidence for radio emission spread over the host galaxy scale for a number of Seyferts when observed with adequate angular resolution and sensitivity (Orienti et al. 2015). In particular, circumnuclear starburst rings with knots of star formation are observed. Orienti & Prieto (2010) observed diffuse radio emission with knots of star formation for a number of Seyferts, on scales smaller than 1 kpc (NGC 5506, NGC 7469, and NGC 7582). Analogously, the morphology of the radio emission in the case of NGC 3185 as observed in our radio maps may be consistent with emission from circumnuclear rings of star formation. In order to check this hypothesis, we overlapped L band and C band radio contours to an archivial HST image7 for NGC 3185, as shown in Fig. 4 . The radio contours overlap with optical emission, from which a nearly circular, ring-like structure emerges. However, further work is needed to understand the link between radio emission and star formation in the form of circumnuclear rings in this source. In particular, intermediate angular resolution scales radio observations (such as those of e-MERLIN) combined with H α maps would allow us to confirm the star formation hypothesis. If we consider the parent sample with our new data, then we find that the average spectral slope is compatible with flat (0.30 ± 0.10), with a nearly equal number of flat and steep slopes. If we consider the two sub-populations of type-1 and type-2 Seyferts, then no significant differences are found with respect to the average spectral index and the average radio-loudness parameter, the only difference being that type-1 sources are slightly more luminous. Considering the X-ray radio loudness parameter, the black hole mass and the Eddington ratio, we do not find significant differences between the type-1 and type-2 sub-classes. We note that, even though the sources in the Reference sample have an Eddington ratio nearly an order of magnitude smaller than the average for the Parent sample, the X-ray radio-loudness parameter does not exhibit anomalous values.
Citation: Orienti et al. 2015
Citation function: Background
Concerning case (1), the first explanation that has been proposed is an increasing relative velocity among the building blocks with decreasing distance from the planet leading to substantial water loss in the case of the most energetic impacts (Estrada & Mosqueira 2006), which occurred closer to Jupiter. Nonetheless, this scenario has been discarded by a detailed study showing that Io and Europa analogs exhibit an overabundance of water when they are formed via an N-body code simulating imperfect accretion and water loss during collisions (Dwyer et al. 2013). A second explanation is that the observed water gradient among the satellites results from an outwardly decreasing temperature of the CPD, leading to the existence of a snowline at a given radial distance from Jupiter (see, e.g., Lunine & Stevenson 1982). In this case, bodies that formed inward of the snowline (Io) accreted from essentially water-poor building blocks, whereas bodies that formed outward of the snowline (Ganymede, Callisto) formed from a primordial mixture of water ice and silicates (e.g., Canup & Ward 2002; Mosqueira & Estrada 2003a, 2003b; Mousis & Gautier 2004). Within this scenario, the low water content of Europa is puzzling. So far, Europa’s water content has been mostly attributed to its formation both outward and inward of the snowline due to either (i) its migration inward of the snowline during formation (i.e., growth), (ii) the progressive cooling of the disk and thus inward migration of the snowline during its formation, or (iii) an interplay between the two mechanisms (Alibert et al. 2005; Canup & Ward 2009). However, the evolution of the CPD has been systematically modeled using an ad hoc parametrization of the turbulent viscous disk (the so-called α-viscosity, Shakura & Sunyaev 1973) which governs the temperature evolution and lifetime of the disk. While providing a good starting point for evolutionary disk models, this kind of parametrization has been highly questioned in recent years (Bai & Stone 2013; Simon et al. 2013; Gressel et al. 2015). Hence, using a predefined α-viscosity prescription to describe the CPD’s evolution and provide hints on Europa’s formation remains questionable. The same remark holds for planet (or satellite) migration, which has also been extensively studied within recent years (see, e.g., Paardekooper et al. 2010; Bitsch et al. 2014). These studies have shown that in realistic disk conditions, migrating planets can behave significantly differently from what was previously thought, i.e., a persistent inward motion (e.g., Tanaka et al. 2002), due to the existence of regions where the migration is halted and even reversed. Because the studies of satellite formation have been based so far on the migration formulation of Tanaka et al. (2002; e.g., Canup & Ward 2002, 2006; Alibert et al. 2005; Sasaki et al. 2010) their proposed growth/migration scenario is questionable.
Citation: Mosqueira & Estrada 2003a
Citation function: Background; Differences
(1a) Work is underway for an updated ExoMol line list for SiO which will extend into the ultraviolet. The current line list only considers vibration-rotation transitions and so the current maximum wavenumber was set at 6049 cm−1. (1b) The HITEMP line list for NO includes data from the ExoMol NOname line list (Wong et al. 2017). (1c) The Ames line list (Huang et al. 2017) and the CDSD-4000 databank (Tashkun & Perevalov 2011) are also available for CO2, as well as the HITEMP compilation (Rothman et al. 2010). (1d) The previous ExoMol line list for H2O, BT2 (Barber et al. 2006), is only complete up to temperatures of 3000 K, whereas the more accurate ExoMol POKAZATEL line list Polyansky et al. (2018) is complete up to 4000 K. (1e) There is also a line list for MgH from ExoMol Yadin (Yadin et al. 2012). However, since it only covers the ground electronic X2Σ+ state, and so is less complete than the more recent MoLLIST line list of Gharib-Nezhad et al. (2013), we use the latter. (1f) Previous to the ExoMol aCeTY line list of Chubb et al. (2020b), the main sources of data for acetylene were from HITRAN (Gordon et al. 2017) and ASD-1000 (Lyulin & Perevalov 2017). The data from HITRAN are only applicable for studies performed at room temperature, and were shown in Chubb et al. (2020b) to be inadequate for high-temperature applications. ASD-1000 was a vast improvement, although there does seem to be opacity missing from some of the hot bands when compared to ExoMol aCeTY in Chubb et al. (2020b). (1g) The previous ExoMol line list for CH4, called 10–10 (Yurchenko & Tennyson 2014), is only complete up to 1500 K. The updated 34–10 line list is therefore recommended instead. Future updates of the database will investigate using data for methane based on recent line lists from either TheoReTs (Rey et al. 2017) or HITEMP (Hargreaves et al. 2020); these are currently expected to be more accurate when considering high-resolution applications. For low-resolution applications, we expect the quality of the ExoMol line list used here to be sufficient, particularly because completeness is more important than accuracy at lower resolutions (Yurchenko et al. 2014). (1h) The energy states from Coppola et al. (2011) are used in the Amaral et al. (2019) line list for HD+. (1i) The energy states from Engel et al. (2005) are used in the Amaral et al. (2019) line list for HeH+. (1j) The pressure and temperature broadened profiles for the resonance doublets of Na and K are computed using Allard et al. (2016, 2019). See Sect. 3.2.3 for a discussion on the broadening profiles of these atoms.
Citation: Polyansky et al. (2018)
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Using the high-level ab initio calculations implemented in the MOLPRO 2015 software package (Werner et al. 2015, 2020), the electronic structures of CO have been determined. In our calculations, the molecular orbitals (MOs) and ground-state energies are computed by the HF method. Then, the CASSCF method (Knowles & Werner 1985; Werner & Knowles 1985) is applied to perform the state-averaged calculation to generate multi-configuration wave functions by utilizing the HF MOs as the starting orbitals. Finally, based on the CASSCF wave functions, the internally contracted multi-reference configuration interaction method with the Davidson correction (icMRCI+Q) (Knowles & Werner 1988; Werner et al. 1988; Knowles & Werner 1992; Shamasundar et al. 2011) is performed to consider the dynamic correlation and size-consistency error. The augmented correlation consistent polarized weighted core valence quintuplet aug-cc-pwCV5Z-DK Gaussian basis set is selected to describe the carbon and oxygen atoms, which is found to give an excellent production of the potential energy and dipole moments for electronic states as the CO molecule dissociates, as mentioned in previous publications (Dunning 1989; De Jong et al. 2001; Peterson & Dunning 2002). The scalar relativistic effect is considered by the third-order Douglas–Kroll Hamiltonian approximation (Reiher & Wolf 2004, 2004). The PECs, PDMs, and TDMs for the singlet and triplet states were computed at the internuclear distances from 0.8–7.5 Å with step sizes of 0.05 Å for 0.8–0.9 Å, 0.02 Å for 0.9–1.5 Å, 0.05 Å for 1.5–2.6 Å, 0.1 Å for 2.6–6 Å, and 0.5 Å for 6–7.5 Å, and those for the quintet states were computed at the internuclear distances from 0.8–15 Å with step sizes of 0.05 Å for 0.8–0.9 Å, 0.02 Å for 0.9–1.5 Å, 0.05 Å for 1.5–2.6 Å, 0.1 Å for 2.6–6 Å, and 0.5 Å for 6–15 Å. The obtained PECs can be also used to determine the dissociation energy D

e
, the electronic excitation energy relative to the ground state T

e
, the internuclear separation R

e
, the harmonic frequency ω

e
, the first-order anharmonic constant ω

e

χ

e
, the rotational constant B

e
, and the vibrational coupling constant α

e
.
Citation: Shamasundar et al. 2011
Citation function: Uses; Uses
The photoevaporation model successfully explains the ‘two-time-scale’ nature of protoplanetary disc evolution, where the inner regions of protoplanetary discs appear to evolve slowly on Myr time-scales, before dispersing on a much more rapid time-scale (e.g. Kenyon & Hartmann 1995; Ercolano, Clarke & Hall 2011; Koepferl et al. 2013; Ercolano et al. 2014). Furthermore, slow-moving (∼5–10 km s−1) ionized winds are observed to be occurring in many nearby discs hosting young stars (e.g. Hartigan, Edwards & Ghandour 1995; Pascucci & Sterzik 2009; Rigliaco et al. 2013) and are consistent with the photoevaporation model (Alexander 2008; Ercolano & Owen 2010; Pascucci et al. 2011; Owen, Scaife & Ercolano 2013a; Ercolano & Owen 2016). The photoevaporation model can also explain a large fraction of observed ‘transition discs’ (e.g. Owen & Clarke 2012; Espaillat et al. 2014), specifically those with holes ≲10 au and accretion rates ≲10−9 M⊙ yr−1 (Owen et al. 2011b) and even those with larger holes and higher accretion rates in more recent models that incorporate CO depletion in the outer disc (Ercolano, Weber & Owen 2018). Transition discs are protoplanetary discs with evidence for a large hole or cavity in their discs (e.g. Espaillat et al. 2014), but they are known to be a heterogeneous class of objects (e.g. Owen & Clarke 2012; van der Marel et al. 2016) and their origins are not always clear. However, a specific prediction of the standard photoevaporation scenario is that there should be a large number of transition discs with hole sizes ≳10 au but that are no longer accreting. This final long-lived stage of disc dispersal gives rise to transition discs which have lifetimes between 105 and 106 yr, but remain optically thick – ‘relic discs’ (Owen et al. 2011b). The long disc lifetimes emerge from the simple fact that discs store most of their mass at large radii, but photoevaporative clearing proceeds from the inside out, so it will always take longer to remove the larger disc mass that resides at larger distance. While several discs satisfy this criterion (Dong et al. 2017), the number of observed non-accreting transition discs with large holes falls far below the theoretical expectations (Owen et al. 2011b; Owen & Clarke 2012). Studies by Cieza et al. (2013) and Hardy et al. (2015) showed that optically thick relic discs are rare and many non-accreting stars that show evidence for a circumstellar disc are more consistent with young, radially optically thin, debris discs.
Citation: Owen et al. 2011b
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In the line cores of Q/I and U/I profiles, we see depolarization and rotation for fields in the range 





 G. These are due to the Hanle effect. For 





 G, we see the signatures of level-crossings in the line cores of (Q/I, U/I) profiles, namely they tend toward the non-magnetic value (see Figures 1(b) and 2(b)). We recall that, traditionally the loops in the polarization diagram (namely, a plot of Q/I versus U/I for a given wavelength and for a range of magnetic field strength or orientation values) are identified to be due to the level-crossings in the incomplete PBE regime (see, e.g., Bommier 1980; Landi Degl’Innocenti & Landolfi 2004, see also Sowmya et al. 2014). When a given curve in the polarization diagram forms a loop the Q/I and U/I values tend toward the non-magnetic value. Based on this we identify the above noted behavior in the line cores of (Q/I, U/I) profiles for the mentioned field-strength regime as to be the signatures of level-crossings in the incomplete PBE regime. Polarization diagrams require the use of very fine grids of magnetic field strength or orientation. With the radiative transfer calculations presented in this paper, it is computationally difficult to produce such diagrams. For 





 G, transverse Zeeman effect like signatures are seen in the line core of (Q/I, U/I) profiles (see Figures 1(c) and 2(c)). The Faraday rotation (del Pino Alemán et al. 2016; Alsina Ballester et al. 2017; Sampoorna et al. 2017), which results in depolarization in the wings of Q/I and generation of U/I in the wings, strongly influences the wings of U/I profiles for the entire field-strength regime considered here, while it shows up in Q/I for 





 G. For the cases of theoretical model line and the isothermal model atmosphere considered in this section, the Voigt effect starts to show up in U/I for 





 G and in Q/I for 





 G, and its signatures are similar to those discussed in Sampoorna et al. (2019a). Also we see the signatures of incomplete PBE in the V/I profiles, which are now asymmetric about the line center for fields up to 30 G.
Citation: Sowmya et al. 2014
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Achondritic meteorites are fragments of differentiated asteroids or planetary bodies of the solar system. Ureilites are coarse-grained ultramafic (olivine–pyroxene) achondrites. They include accessory minerals of metal and sulfide associated with high abundances of carbon phases (on average 3 vol% and up to ∼7 vol%; Goodrich et al. 2015), including graphite and high-pressure diamond (Goodrich 1992; Mittlefehldt et al. 1998; Nabiei et al. 2018). According to their petrology, ureilites are further divided into main group ureilites (formerly referred to as monomict or unbrecciated ureilites; accounting for 95%) and rare polymict ureilites (or brecciated ureilites). The main group ureilites represent the mantle residues after the extraction of feldspar-rich magmas (Cohen et al. 2004; Bischoff et al. 2014; Barrat et al. 2016) and a sulfur-rich iron melt (Goodrich et al. 2004; Warren et al. 2006; Goodrich et al. 2007; Rankenburg et al. 2008; Barrat et al. 2015), whereas polymict ureilites are breccias containing fragments of main group ureilites as well as non-ureilite clasts (Goodrich et al. 2004, 2015). Unlike samples from other differentiated planetary bodies (e.g., terrestrial and lunar samples, Martian meteorites, angrites), some primitive isotopic geochemical features (e.g., O isotope heterogeneities; Greenwood et al. 2005, 2006, 2017) have been preserved in ureilites, which are considered incompatible with a magma ocean event (e.g., Goodrich 1992; Goodrich et al. 2004, 2015). Therefore, the geochemical signatures of ureilite meteorites can provide valuable information about the origin of the ureilite parent body (UPB) and the early evolution of the solar system. For instance, olivine cores from individual ureilite meteorites cover a range of Fe/Mn ratios from 3 to 57 (Goodrich 1992; Goodrich et al. 2004, 2007; Downes et al. 2008) and Δ17O values (mass-independent O isotopic variations) from −2.5‰ to −0.2‰ (Clayton & Mayeda 1988, 1996; Greenwood et al. 2017). The whole rock Δ17O values further correlate with both the Fe/Mn ratios and magnesium number (Mg#: molar ratio of Mg/Mg+Fe) of their olivine cores (Goodrich et al. 2004, 2015). The origin of these covariations is debated and multiple scenarios have been proposed to explain these features, including (1) a “smelting” process which causes redox reactions between C and FeO during UPB differentiation (Singletary & Grove 2003; Goodrich et al. 2007); (2) oxidation of metal due to the presence of water prior to igneous differentiation (Sanders et al. 2017); (3) inherited nebular redox variability from the precursor materials that were not homogenized during partial melting (Warren & Huber 2006; Warren 2012); (4) UPB accretion from the mixing of two different chemical and isotopic reservoirs (Barrat et al. 2017). As such, ureilite genesis still remains unclear. Understanding the origin of these geochemical correlations in ureilites is further complicated by the fact that the timing of formation and early evolution of the UPB is poorly constrained.
Citation: Warren et al. 2006
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In Fig. 1, we investigate the dust grain size distributions at the periastron and apastron of debris disks a belt eccentricity of eb = 0.4 (with dynamical excitation Δeb = 0.1) for different material strengths. Since the radiation pressure has a strong influence on the cutoff in the grain size distribution at smallest sizes, the wavy patterns of the grain size distributions start around the smallest bound grains. We note that the blowout size can be increased with particle porosity and increasing stellar temperature (Burns et al. 1979; Kirchschlager & Wolf 2013; Brunngräber et al. 2017). We find different characteristic wavy patterns in the grain size distributions starting with the depletion of grains with radii below the blowout radius abo, which are caused by the lack of β > 0.5 dust grains (Thébault et al. 2003), depending on different collisional evolutions (Figs. 1 and 2). For grains with radii ~3 μm at the apastron side and grains with radii ~6 μm at the periastron side, a strong wavy pattern develops. This depletion leads to an over-density of slightly larger grains (“first peak”) because grains with radii a abo are depleted and thus can no longer contribute efficiently to the destruction and erosion processes anymore. The overabundance of grains with radii around the first peak (~1.5 abo) in turn induces another depletion of grains with radii around the “second depletion” (~4–40 abo) that is caused by small high-β grains originating inside the disks. Thus, an efficient destruction is responsible for the deep depletion of objects with radii of up to ~4−40 abo. Qualitatively, this first wavy pattern is less pronounced in the size distribution of higher material strengths, where the impact velocities and thus their rate of destructive collisions is significantly lower. A strong depletion for lower material strengths is found for grains with radii ~60 μm, while the depletion in the case of higher material strengths is valid for radii ≤10 μm (Fig. 1). This depletion eventually leads to an over-density of grains with radii around the “second peak”. The overabundance of these grains is shifted from ~100 μm (high material strength) to ~1000 μm (very low material strength).
Citation: Kirchschlager & Wolf 2013
Citation function: Uses
Based on these results, a series of 3D numerical simulations solving the simplified reduced magnetohydrodynamic (RMHD) equations, introduced by Strauss (1976), in Cartesian geometry were performed. The goal of these simulations was twofold: (1) to determine how a coronal loop responds to different photospheric velocity patterns, and (2) to investigate how the electric current sheets are formed, as well as the details of the heating that occurs through the dissipation of magnetic energy (Rappazzo et al. 2007, 2008, 2010; Rappazzo & Velli 2011; Rappazzo & Parker 2013; Rappazzo 2015). The results detailed in these papers confirmed the 2D results that the energy flux entering the corona because of photospheric motions causes a turbulent cascade of energy toward small scales, with electric current sheets continuously being created and destroyed throughout the coronal loop. The turbulence that develops is highly intermittent. The dissipation of energy occurs in current sheets, which are localized structures. Hence, the heating also occurs on small spatial scales. (Note that there is also particle acceleration, although that is beyond the scope of the present model.) When the loop is fully turbulent, it achieves a statistically steady state in which, on average, the Poynting flux induced at the boundaries by footpoint convection is balanced by the dissipation of energy in the electric current sheets. In this steady state, saturation occurs for the fluctuating magnetic energy, kinetic energy, mean square electric current, enstrophy, and other quantities that are then seen to fluctuate around their mean values. When the system becomes nonlinear, its behavior is highly dynamical and chaotic. This state occurs independently of the detailed form of the boundary velocity. A statistically steady state is achieved in which, although the magnetic field footpoints are convected continuously by the boundary flows, the resulting topology of the total magnetic field is not simply a mapping of those flows. By the same token, the turbulent dynamics that occur are due to the inherent nonlinear nature of the system, rather than being a consequence of the complexity of the magnetic field footpoint configuration.
Citation: Rappazzo & Velli 2011
Citation function: Motivation