adsabs/FOCAL huggingface dataset:
Sudden outbursts of NH3 simultaneously with H2S detected with the ROSINA-DFMS instrument on the Rosetta S/C point to the presence of abundant ammonium hydrosulphide in or on carbonaceous grains from comet 67P/Churyumov-Gerasimenko. There seems to be a clear distinction between the nucleus ice, where H2S and NH3 exist independently and grains, where they desorb together. S2 is much more abundant on grains compared to water than in the ice of the comet, while S3 is found only in grain impacts. This higher abundance points to radiolysis in these grains, which means they must have been exposed to energetic particles over an extended time. While for operational reasons, S4 could not be measured close to the dust impacts, S4 was clearly identified in periods where the coma was very dusty (Calmonte et al. 2016). Longer sulphur chains very likely are refractory, not sublimating at temperatures reached in the instrument or on grains in the coma. While Sn can also be formed from pure H2S ice by photo processing (Cazaux et al. 2022), the fact that S3 is clearly related to dust and is not found in the normal nucleus ice, where H2S is quite abundant, indicates that S3 is a product of radiolysis of the ammonium salt. In addition, photo processing of H2S results not only in Sn, but also in H2S2 (Cazaux et al. 2022), a species not detected in the DFMS m/z 66 and m/z 65 (HS2) spectra. This exposure rules out a contemporary formation of the salt on the surface or in the interior of the comet or a formation of the salt in the mid-plane of the protoplanetary disc, while the comet accreted. A pre-stellar formation is therefore likely. The salt is semivolatile, less volatile than water and could probably have survived quite high temperatures. It seems that on these grains, acids and ammonia are all locked in salts, be it sulphur, halogens, or carbon bearing acids like HOCN. If indeed, a relatively large part of sulphur and nitrogen is therefore in a semivolatile state in these grains, then the depletion of nitrogen in comets and of sulphur in star-forming regions could probably be explained, primarily because salts escape detection unless they experience temperatures above water sublimation. With the JWST S/C in orbit, there is hopefully the possibility to detect salts, or at least several of the acids in ices, which are supposed to be part of ammonium salt, like HOCN, H2CO, and formamide while looking for ammonium salts in star-forming regions and possibly comets.
Citation: Calmonte et al. 2016
Citation function: Uses
Citation: Cazaux et al. 2022
Citation function: Compare/Contrast; Compare/Contrast
Any massive stellar source from ~8 M⊙ can burn hot enough to completely ionise the surrounding molecular material to form an H II region (Wood & Churchwell 1989; Churchwell 1990; Kurtz 2005) and destroy any complex chemical tracers that are traditionally used to understand the kinematics of their natal environments and their accretion discs. Initial arguments by Walmsley (1995) argued that high accretion rates will cause a very dense H II region that is optically thick to radio emission and thus the H II would not be seen. However, for heavily accreting massive YSOs the onset of the H II region could be delayed via stellar bloating (Palla & Stahler 1992; Hosokawa & Omukai 2009; Hosokawa et al. 2010; Kuiper & Yorke 2013) where the effective temperature of the star is much cooler than it would be considering a main sequence star of the same luminosity. In these models, a halt or considerable reduction in accretion (≪10−3 M⊙ yr−1), or growth beyond ~30−40 M⊙, will result in the YSO contracting to a “main-sequence” configuration, heating significantly, and being able to create an H II region. The fine details of this transition are still somewhat unclear since they depend on the assumed accretion law and the initial conditions chosen for the stellar evolution calculations (Haemmerlé & Peters 2016). Observations in the infra-red (IR) have been made in search of cool stellar atmospheres that point to bloated stars, however these studies remain inconclusive (Linz et al. 2009; Testi et al. 2010). Alternative scenarios are that H II regions can be gravitationally trapped at very early ionisation stages (see Keto 2003, 2007), or flicker due to chaotic shielding of the ionising radiation by an accretion flow, leading to a non-monotonous expansion (Peters et al. 2010b). Hyper Compact (HC) H II regions (0.03 pc) are thought to be the earliest ionisation stage and therefore could relate to the halt of accretion, be a marker of a transition phase.
Citation: Palla & Stahler 1992
Citation function: Background; Background
The final location of close-in giant planets in our models reflects the strength of the tides that we include in our modeling, which play a very important role in the decay of planetary orbits. These are dynamical tides (e.g., Lai 1997; Ivanov & Papaloizou 2004, 2007, 2011) and in our simulations we used a formulation given by Ivanov & Papaloizou (2007) as described inSect. 2. However, the impulse approximation used in the evaluation of dynamical tides becomes a poor approximation when the circularization proceeds (e.g., Mardling 1995a,b) and the eccentricity becomes low. Equilibrium tides become then effective (e.g., Beaugé & Nesvorný 2012 and references therein) and the tidal evolution may occur on a longer timescale. In short, at the beginning of the orbital evolution that leads to the formation of hot/warm Jupiters, dynamical tides are important in forcing the decay of the orbit. In the last part of the dynamical evolution when the eccentricity has become low, equilibrium tides are more important in determining the location where the planet stops. Unfortunately, at present it is not known when and how the two tides switch. When we change the magnitude of two tides, the final location of the planets can be adjusted (Beaugé & Nesvorný 2012). However, rather than introducing artificial effects, we continue to use dynamical tides in our simulations even for low eccentricities but we stop our simulations when the energy, decreasing from the tide at the pericenter, overcomes the orbital energy leading to a clustering of tidally circularized planets around 0.02 au. However, the final distribution of the inclination of the planets does not depend on this choice and highly misaligned planets would be produced anyway. Since the tidal evolution of planets with arbitrary inclinations is still not well known, we assume that planetary inclination is not significantly changed during tidal evolution (Barker & Ogilvie 2009). Thus, the planets maintain the inclination they have when the circularization begins.
Citation: Beaugé & Nesvorný 2012
Citation function: Uses; Uses
Citation: Beaugé & Nesvorný 2012
Citation function: Uses
The catalogue used in this work is based on a far-IR sample selected in the ∼2 deg2-wide COSMOS field and obtained within the Herschel-PEP survey (Lutz et al. 2011). We consider the latest released blind catalogue selected at 160-μm (DR1, 7047 sources) with >3σ flux density, corresponding to a flux limit of ∼9.8 mJy. The choice of considering as parent sample a far-IR catalogue is guided by the necessity of having several detections at different wavelengths for each system to constrain the dust masses, and yet with a very simple selection function (see Section 3.1). From the original 160-μm selection, we built a multiband catalogue taking advantage of the extensive multiwavelength coverage in the COSMOS field. Concerning the other far-IR PACS band (100 μm) and the mid-IR 24-μm band, we use the association available in the DR1 release and based on the maximum likelihood technique (Sutherland & Saunders 1992; Ciliegi et al. 2001). For the cross-match with the SPIRE far-IR bands (250, 350, and 500 μm), we used the same catalogue considered in previous PACS-based works (i.e. Gruppioni et al. 2013; Delvecchio et al. 2015), the ones provided by the HerMES collaboration (Roseboom et al. 2010) using the Spitzer-MIPS 24-μm positions as priors to extract the SPIRE fluxes. Finally, the IRAC/optical/UV fluxes taken from the COSMOS2015 catalogue (Laigle et al. 2016) were merged to the PACS-160-μm sample by matching the 24-μm counterparts listed in both. The COSMOS2015 is NIR selected, where objects have been detected from the sum of the UltraVISTA-DR2 YJHK and z++ images. By construction, in comparison to the previous i-selected catalogue, this catalogue is missing a fraction of blue, faint, star-forming galaxies (Laigle et al. 2016). For this reason, we decided to cross-match the far-IR sources with no counterparts in the COSMOS2015 catalogue with the Ilbert et al. (2009) i-selected catalogue. Totally, among the 160-μm selected sources (7047), 6002 are with 24-μm counterparts (${\sim }86{{\ \rm per\ cent}}$), of which 5993 with available NIR or optical counterparts (${\sim }99.9{{\ \rm per\ cent}}$, 5783 in the COSMOS2015 and 210 in the Ilbert et al. 2009 catalogue). While the cross-matching with the optical/NIR bands does not involve almost any source loss, a moderate (14${{\ \rm per\ cent}}$) but not negligible fraction of the far-IR sources does not have 24-μm counterparts. A fraction of these sources are likely spurious sources, as shown by the simulations done for the DR1 release PEP catalogue (∼5 per cent at the 3σ flux level). In Section 3.1, we will describe our method to correct for incompleteness and for the presence of spurious systems.
Citation: Sutherland & Saunders 1992
Citation function: Uses
Figure 12 shows the results of our tidal evolution calculations. The left panel of Fig. 12 shows the planetary rotational evolution of GJ 1148 b due to star–planet tides. After ~850 Myr, GJ 1148 b reaches a rotation period that is 2∕3 of the orbital period, and remains there with Prot = 27.5 d. During the integration the planetary semi-major axis and eccentricity are mostly unaffected. An asymptotic rotation period that is shorter than synchronous and 2/3 of the orbital period is expected for eb ≳0.24 in the constant Q tidal model (Goldreich & Peale 1966; Cheng et al. 2014). The time for GJ 1148 b to reach asymptotic rotation is inversely proportional to the initial Prot, as long as the initial Prot is much less than 27.5 d, and it depends on the other parameters of GJ 1148 b according to Eq. (3) of Barnes (2017) and Eq. (15) of Cheng et al. (2014). The rotational period of GJ 1148 b is thus very likely much longer than the orbital periods of the hypothetical exomoons, which could be dynamically stable only with orbital periods between 0.7 and 2 d. The right panel of Fig. 12 shows that the longer rotational period of GJ 1148 b (Prot = 27.5 d) leadsto strong orbital decay of the stable exomoon orbits due to tidal interactions with the planet. An exomoon eventually reaches the Roche limit where it is tidally disrupted by the gas giant. Not even one hypothetical “stable” exomoon in the context of Sect. 5.3.1 had survived this test. The maximum time a Mars-like exomoon could survive is ~55 M yr, while for Titan-like moons the maximum survival time is longer, ~255 M yr. The latter is longer by roughly the mass ratio of Mars to Titan, which can be understood from Eq. (2) of Barnes (2017) and Eq. (16) of Cheng et al. (2014). These timescales are optimistic since the orbital decay would start before the planet reaches the asymptotic spin state. In both cases the survival times are much shorter than the age of the system. Therefore, given the relatively fast orbital decay in the small stable region around the planet, we conclude that exomoons around GJ 1148 b are unlikely to exist.
Citation: Barnes (2017)
Citation function: Uses
Citation: Barnes (2017)
Citation function: Uses
Previous observations of AGN that investigate correlations between αOX and Eddington ratio have revealed some similarities with X-ray binary outbursts at high Lbol/LEdd, but these comparisons have not been possible below the critical Lbol/LEdd ≲ 10−2, where an inversion in this correlation is predicted to occur. At higher Eddington ratios of Lbol/LEdd ≳ 10−2, single-epoch X-ray and UV observations of large samples of AGN have previously revealed a hardening of αOX as Lbol/LEdd drops from ∼1 to ∼10−2 (e.g., Vignali et al. 2003; Strateva et al. 2005; Steffen et al. 2006; Just et al. 2007; Grupe et al. 2010; Jin et al. 2012; Wu et al. 2012; Trichas et al. 2013; Vagnetti et al. 2013). This correlation was also observed in multi-epoch UV/X-ray observations of the fading of Mrk 1018 (Noda & Done 2018), which confirms this behavior in an individual AGN. However, the predicted softening of αOX below Lbol/LEdd ≲ 10−2 (thus causing an inversion in the correlation between αOX and Lbol/LEdd) has not been previously observed. This is primarily due to the difficulty of robustly measuring both αOX and Lbol/LEdd for AGN below Lbol/LEdd ≲ 10−2, for three main reasons. First, at low Eddington ratios, AGN are often dust-obscured (Fabian et al. 2008), and thus measuring their intrinsic UV luminosities (and αOX) is difficult. Second, broad emission lines often disappear in low-luminosity AGN below Lbol/LEdd ≲ 10−2, making it difficult to measure MBH (and LEdd). Third, using a sample of AGN with a wide range of Eddington ratios to trace how αOX changes as a function of Lbol/LEdd can be hampered by the 





 scaling of the thin disk temperature with MBH at a fixed Eddington ratio. If the AGN sample has a large range in MBH, this can cause an additional scatter in αOX. Thus, we would ideally use a sample of AGN with a narrow range in MBH, but the difficulty of measuring MBH at Lbol/LEdd ≲ 10−2 also hampers the construction of such a sample. In this paper, we will use a new method to bypass all of these issues, with the goal of extending this spectral comparison between X-ray binaries and AGN to Lbol/LEdd ≲ 10−2.
Citation: Strateva et al. 2005
Citation function: Background
We assume the time-scales we observe here are generated in the corona itself (and note that longer time-scale changes will be driven by the disc outside of the corona) and are made visible by a changing electron temperature and density as a result of local turbulence and coupling to mass accretion rate propagations through the flow from rout to rin. We note that we can discount variations in the seed photon population as the driver for changes in the power spectrum, as the UV emission from the disc is established to be considerably less variable than the corona in NLS1s (Leighly 1999; Smith & Vaughan 2007; Ai et al. 2013; Alston, Vaughan & Uttley 2013; Done & Jin 2016). We assume that the variability generated locally at each radius (rν) is at the viscous frequency (see Churazov et al. 2001) such that
(9)$$\begin{eqnarray}
r_{\nu } = \left[ \frac{2\pi \nu }{\alpha }\left(\frac{H}{R}\right)^{-2}\right]^{-2/3} ,
\end{eqnarray}$$where the frequency is in units of c/Rg (see e.g. Kato, Fukue & Mineshige 1998; Arévalo & Uttley 2006). In the above, α and $\frac{H}{R}$ are the dimensionless viscosity parameter of Shakura & Sunyaev (1973) and scale height of the accretion disc, respectively. We assume that our frequency range of interest, 0.01−1 mHz (i.e. the range over which we can practically fit to the data) corresponds to radii between the ISCO (rin) and some radius within the true outer edge of the corona (i.e. rout ≤ rcorona). The actual frequencies generated in our model therefore depend on the SMBH spin and the combination $\left(\frac{H}{R}\right)^2 \alpha$ (and somewhat on the SMBH mass – although here the range is small). Given the reported high spin values for these bright AGNs (Ogle et al. 2004; Fabian et al. 2013; Done & Jin 2016; Kara et al. 2017; Buisson et al. 2018b), we expect the ISCO to sit at ∼1.25Rg. For the corona at the ISCO to produce variability above our upper frequency limit of 1 mHz requires $\left(\frac{H}{R}\right)^2 \alpha \gtrsim 0.01$. We note that for the mass range subtended by our AGN sample (from 106.00 to 106.63 M⊙, see Table 2), should we instead assume zero spin (rin = 6Rg), the viscous frequency at rin is lower and we have strong curvature in our observed bandpass.
Citation: Done & Jin 2016
Citation function: Uses; Uses
Citation: Done & Jin 2016
Citation function: Uses
We subsequently considered the time evolution of a number of selected cores (based on the requirement that we have sufficient data to follow at least 200 kyr of evolution), finding a remarkably similar chemical evolution in all cores, including one that is strongly stabilized by turbulence and magnetic field and therefore not going through gravitational collapse. Assuming a more or less chemically homogeneous initial condition, the chemical evolution of the cores, particularly regarding deuteration effects, is thus very similar. Of course, environmental effects such as the local cosmic ray ionization rate could induce potential differences, by changing the ionization degree. Similarly, metallicity effects at lower densities may depend on the local conditions. At least within one filament, it is however plausible that the deuteration fraction is indeed indicative of chemical age. We find here that about two free-fall times (as defined for cylindrical systems, see Toalá et al. 2012) are sufficient to reach core deuteration fractions of ≳0.1. We finally investigated also the radial structure of the core, finding overall similar properties as in the isolated collapsing cores studied by Körtgen et al. (2017). The H2 ortho-to-para ratio appears to be approximately flat and only weakly dependent on radius. Both the deuteration fraction and the gas surface density show a peak on scales of about 1000–2000 au, which is particularly pronounced in the case of parallel magnetic fields. As found previously, this peak moves outward with increasing turbulent Mach number, indicating the amount of support against gravity. The difference in the visibility of the peak may result from the difference in the fragmentation mode in both cases, as previously described by Seifried & Walch (2015), thus potentially affecting the structure of the resulting cores. Overall, our results have shown that the observed high deuteration fraction in prestellar cores can be readily reproduced in simulations of turbulent magnetized filaments. We further found that deuteration fractions of order 0.1 can be produced independent of the specific history of the cores, both for high and low virial parameters. The latter suggests that deuteration is potentially very efficient.
Citation: Toalá et al. 2012
Citation function: Uses; Uses
The broad-band spectral energy distribution (SED) of blazars is characterized by two broad humps, one at optical/UV/X-ray bands and the other in the HE γ-ray band (see Padovani et al. 2017 for a recent review). It is believed that the first peak (low-energy component) is mostly due to synchrotron emission from relativistic electrons, whereas the origin of the second component is highly debatable. Within conventional leptonic scenarios, this component is produced when the synchrotron-emitting electrons inverse Compton up scatter the photons of internal (synchrotron self-Compton, SSC; Ghisellini, Maraschi & Treves 1985; Maraschi, Ghisellini & Celotti 1992; Bloom & Marscher 1996) or external (external inverse Compton, EIC; Sikora, Begelman & Rees 1994; Błażejowski et al. 2000; Ghisellini & Tavecchio 2009) origin. The nature of the external photon fields depends on the distance of the emitting region from the central black hole (Sikora et al. 2009) and can be dominated either by the photons directly emitted from the accretion disc (Dermer, Schlickeiser & Mastichiadis 1992; Dermer & Schlickeiser 1993) or disc photons reflected from the broad-line region (BLR; Sikora et al. 1994) or IR photons emitted from the dusty torus (Błażejowski et al. 2000). Recently, after associating TXS 0506+056 with the IceCube-170922A neutrino event (IceCube Collaboration et al. 2018a,b; Padovani et al. 2018), it is more evident that the HE component could be initiated by the interaction of energetic protons when they are effectively accelerated in the blazar jets. The HE component can be either from proton synchrotron emission (Mücke & Protheroe 2001) or from secondary particles from pion decay (Mannheim & Biermann 1989; Mannheim 1993; Mücke & Protheroe 2001; Mücke et al. 2003; Böttcher et al. 2013). In the latter case, blazars are also sources of very high energy neutrinos (Ansoldi et al. 2018; Keivani et al. 2018; Murase, Oikonomou & Petropoulou 2018; Padovani et al. 2018; Sahakyan 2018; Cerruti et al. 2019; Gao et al. 2019; Righi, Tavecchio & Pacciani 2019; Sahakyan 2019).
Citation: Maraschi, Ghisellini & Celotti 1992
Citation function: Compare/Contrast
In the following, all measurements and maps derived from snapshots prior to t = 800 Myr are therefore using the original simulation (R+13), with the others extracted from the simulation at 1 pc resolution (for the gaseous component). Note that star formation was turned on at t ∼ 745 Myr in the simulation as to avoid gas being prematurely consumed. Details on the implemented recipes for star formation, stellar feedback (photoionization, radiative pressure, supernova explosions) are described in R+13. Adding on AGN feedback would significantly impact on the distribution, kinematics and physical status of the gas, specifically for the close environment of the black hole (Ciotti & Ostriker 1997; Haehnelt, Natarajan & Rees 1998; Silk & Rees 1998; Kauffmann & Haehnelt 2000). For the present simulations, however, we do not include the potential feedback from an AGN, thus focusing on a time window (t ∼ 750–830 Myr) when we consider that the AGN itself is quiet (or in an ‘off-state’). This is partly justified by the assumption that AGN have low duty cycle at low redshift and for black holes of a few 106 M⊙ (Haehnelt & Rees 1993; Wang et al. 2009; Shankar et al. 2010; Shankar, Weinberg & Miralda-Escudé 2013) and by the short time range we are considering. More importantly, it allows us to narrow down our study to probe the interplay between the dynamical evolution and the effect of star formation (similarly to e.g. Levine et al. 2008; Hopkins & Quataert 2010a). Turning on the AGN in such a simulation would be paramount to understand any potential fuelling cycle starting from the large-scale down to the vicinity of the black hole, and such an implementation has already been included in ramses by a direct calculation of the Bondi accretion rate (see Teyssier et al. 2011; Gabor & Bournaud 2014, and references therein). It would nevertheless require to probe various feedback schemes, and triggering mechanisms, which is beyond the scope of the present paper.
Citation: Kauffmann & Haehnelt 2000
Citation function: Future Work