adsabs/FOCAL huggingface dataset:
Figure 2 shows the ratios of the relative fluences (element/Mg) measured in Genesis bulk SW (Figure 2(a)) and regime targets (Figure 2(b)) after normalization to the respective solar abundance ratios (Table 7 in Appendix A.5). The bulk SW data for K and Fe in silicon (marked by asterisks in Figure 2(a)) are the preliminary fluences from Rieck (2015), Rieck et al. (2016), and Burnett et al. (2017). If SW had the composition of the solar photosphere, the ordinate value in Figure 2 would be 1 for all elements. The bulk SW shows a fractionation of elemental abundances correlating with FIP, as previously observed by spacecraft (e.g., Geiss 1982; Bochsler 2009). However, the better accuracy and element coverage of Genesis data compared to in situ measurements allows further insights into this FIP-related fractionation. The low-FIP elements (Figure 2(a)) are consistent with a flat pattern, at least below 7 eV, as many pre-Genesis studies suggested. This consistency would make these SW data easier to apply to cosmochemistry, as they would suggest that low-FIP elements are unfractionated relative to each other and hence their SW abundance ratios equal photospheric ratios. However, Figure 2(a) is also consistent with a monotonic increase of the low-FIP elements with decreasing FIP, which could be viewed as a trend continuing that of the high-FIP elements between Ar and C discussed below. The Na abundance given here is about 20% higher than that of Burnett et al. (2017). Na fluences derived from backside depth profiling of diamond-like-C collectors are roughly a factor of 2 lower than ours, which are based on Si collectors (Rieck 2015; Rieck et al. 2016). Jurewicz et al. (2019) propose that diffusion of surface contamination Na might have enhanced fluences from Si collectors. Because the amounts of surface contamination is highly variable, the Na fluences from Si should show significant scatter; however, replicate analyses of bulk and regime samples agree to within 2%–11% (Table 4, Appendix A.1). More analyses are required to resolve this important discrepancy. A ∼20% lower fractionation of K compared to that of Na, if real, would be surprising although within 2σ, both values still overlap.
Citation: Geiss 1982
Citation function: Similarities; Compare/Contrast
Wang et al. (2017), motivated by the fact that previous studies on high-eccentricity migration focused only on the total efficiency of HJs formed, analyzed the efficiency of each high-eccentricity mechanism, trying to understand which mechanism is the dominant one in HJ formation. They considered multiplanetary systems containing from two to five planets with equal masses (1MJ), initially in circular and near-coplanar orbits with a host star of $1\, {\rm M}_\odot$ and $1\, {\rm R}_\odot$. Moreover, various initial semimajor axis conditions were considered, depending on the initial mutual separation between the planets. Their numerical simulations were performed using the classical version of the Mercury code (Chambers 1999), which does not include general relativity (GR) effects and tidal interaction with the central star. They studied how the initial number of planets, the spatial separation between them, and the location of the inner planet influence the efficiencies of high-eccentricity mechanisms. As a result, they found that the Kozai–Lidov mechanism plays the most important role in HJ production. The restriction made in the previous study, regarding not including the contribution of GR and realistic initial mass configurations (unequal masses), makes us wonder about the implications of these contributions on the efficiency of each high-eccentricity mechanism in the activation of the HJ production process. The main effect of GR is to cause the apsidal precession of planetary orbit, where the rate of precession is faster for planets close to the star and with eccentric orbits due to the term a5/2(1 − e2) in the denominator, being a and e the planet’s semimajor axis and eccentricity, respectively (Einstein 1916; Misner, Thorne & Wheeler 1973). Studies related to the influence of GR on the dynamic evolution of systems with more than one massive planet have been carried out, but only considering some particular systems (e.g. Adams & Laughlin 2006; Migaszewski & Goździewski 2009; Veras & Ford 2010; Zhang, Hamilton & Matsumura 2013; Marzari & Nagasawa 2019, 2020). Therefore, the aim of our study is to obtain a broader and more detailed description of the efficiency of each high-eccentricity mechanism in the activation of the HJ production process, including the contribution of GR and different initial planetary mass configurations. The analysis is made by solving the numerical simulation of the exact equations of motion, in the context of general N-body problem. Several initial conditions are considered, changing the initial mass and number of planets, the semimajor axis of the inner planet and the location of the other planets in the system.
Citation: Chambers 1999
Citation function: Motivation; Motivation
Astrochemical models have always dedicated special attention to molecular oxygen. With a cosmic abundance twice that of C, atomic O is the third most abundant element in space. In dense clouds, standard gas phase chemical models therefore suggest a comparable ratio of CO and O2 after times ≥ 3 × 105 yr (e.g. Woodall et al. 2007), where O2 is supposed to be formed especially via OH + O → O2 + H. The OH here can be formed by the dissociative recombination of H3O+, H3O+ +e− → OH + 2H. However, observations with the Submillimeter Wave Astronomy Satellite (SWAS) by Goldsmith et al. (2000) towards Orion and with Odin by Larsson et al. (2007) towards ρ Oph A showed a significant difference between model predictions and measurements. The O2 abundances found were more than 100 times smaller than those predicted by models (Goldsmith et al. 2000). Better agreement with observations can be obtained if freeze-out of O atoms onto dust grains is taken into account in gas-grain chemical models (Bergin et al. 2000; Viti et al. 2001), with consequence surface production of H2O and O2, which may trap a significant fraction of oxygen, leaving only some atomic O and CO in the gas phase. Observations conducted by Liseau et al. (2012) led to a O2 column density of N(O2) = 5.5 × 1015 cm−2 with an upper limit of abundance of N(O2)/N(H2) ∼ 5 × 10−8 in warm gas (T > 50 K) and to N(O2) = 6 × 1015 cm−2 with a little higher abundance in cold gas (T  30 K). Liseau et al. (2012) stated that detecting gas phase O2 might be so difficult because the O2 abundance is transient in ρ Oph A and O2 is no longer detectable after ∼2 × 105 yr. A relatively large amount of O2 has only been found with Herschel in Orion as reported by Goldsmith et al. (2011). This source is quite warm (≥180 K), leading to a grain temperature of ≥100 K. At this temperature the grains are warm enough to desorb H2O ice and keep a large amount of oxygen with a big fraction in the form of O2 in the gas phase. Another explanation for the high O2 abundance found by Goldsmith et al. (2000) is that low-velocity C-shocks might be responsible for the increase of molecular oxygen in the gas phase.
Citation: Bergin et al. 2000
Citation function: Motivation; Motivation
We estimate the DPL parameters in equation (9) via Monte Carlo simulations. We compare the real structure function of an AGN to simulations computed from mock light curves generated using the DPL model of equation (9). We generate ‘mock’ light curves using the Timmer & König (1995) method. To create a single mock light curve, pseudo-random numbers are generated using the Fast Mersenne Twister SFMT19937 generator seeded with hardware-generated random numbers (generated using Intel RDRAND instruction) to ensure that the random number sequences are free of artificial correlations (Coddington 1994) induced by poor random seed choices. At this intermediate stage, the mock light curve is oversampled by a factor of 10 i.e. we generate points at 10 × the required cadence in order to avoid sampling issues. To include low-frequency modes that are not adequately characterized by the length of the observed light curve, the intermediate mock light curve is much longer than is ultimately required to make the final mock; mock light curves generated in this manner are capable of exhibiting low-frequency modes longer than the length of the observed data. Fast Fourier Transforms (FFTs) are most efficient for data sequences that are a power of 2; for this reason, we pick the intermediate (including the oversampling) to be of length 223. This results in the intermediate mock light curve being between 15 × to 45 × the length required for the final mock light curve depending on the actual length of the observed light curve. We pick a uniformly distributed random segment of the intermediate overly long light curve that has the same length as the observed light curve and generate another stream of uncorrelated Gaussian random deviates to simulate the white-noise properties of Kepler instrumentation noise. After adding this ‘measurement noise’, we set data points corresponding to the unobserved cadences in the real light curve to 0. This procedure creates a final mock light curve with identical sampling and noise properties to the real light curve. Fig. 8 shows the true light curve (orange) along with an example mock light curve (light green) for the Sy 1 AGN kplr006932990 illustrating what the mock light curves look like for the best-fitting DPL parameters for this object.
Citation: Coddington 1994
Citation function: Uses; Uses
On the other hand, the line-shift-absence conundrum of ZP74 is easily explained through geometrical considerations. Essentially, the arguments leading to this conundrum assume that the collapse is roughly spherically symmetric and monolithic, so that the infall motions are coherent, and directed towards a single collapse centre, at the geometrical centre of the cloud. This assumption is inconsistent with our current understanding of the structure of MCs, which are known to be far from spherically symmetric, and instead consist of an intricate and inhomogeneous network of filaments and clumps within them (e.g. Bally et al. 1987; Feitzinger et al. 1987; Gutermuth et al. 2008; Juvela, Pelkonen & Porceddu 2009; Myers 2009; André et al. 2010; Henning et al. 2010; Men’shchikov et al. 2010; Molinari et al. 2010; Arzoumanian et al. 2011; Busquet et al. 2013). The central clumps (‘hubs’) appear to accrete from the filaments, while in turn the filaments seem to accrete radially from their surroundings (Schneider et al. 2010; Kirk et al. 2013; Peretto et al. 2014; Gong et al. 2018; Lu et al. 2018; Williams et al. 2018; Shimajiri et al. 2019). Thus, the geometry is far from being spherically symmetric, and therefore the accreting gas is not isotropically distributed around the collapse centres (the hubs). In addition, the velocity field is highly complex and chaotic (e.g. Gómez & Vázquez-Semadeni 2014; Zamora-Avilés, Ballesteros-Paredes & Hartmann 2017; Gómez, Vázquez-Semadeni & Zamora-Avilés 2018), so there is no reason to expect a systematic redshift of the absorption lines produced in the gas surrounding the hubs. Instead, the accretion flow is most directly observed as velocity-centroid gradients along the filaments, directed towards the hubs. Indeed, synthetic CO observations of simulations of the regime often show only marginal or no evidence for infall profiles, due to the chaotic motions and perhaps velocity crowding effects, although the line profiles do look similar to observed ones (e.g. Heitsch et al. 2009; Heiner, Vázquez-Semadeni & Ballesteros-Paredes 2015; Clarke et al. 2018). Nevertheless, recent dedicated searches for evidence of infall signatures in CO lines from GMCs have met with success. For example, Schneider et al. (2015) have found the classical combination of self-absorbed and blue-skewed optically thick lines (12CO (3–2)) together with centrally peaked optically thin (13CO (1–0)) lines, indicating collapse in the molecular gas surrounding IRDC G28.37+0.07, while Barnes et al. (2018) have measured shifts between the lines of 12CO (tracing gas in the outer parts of the cloud) and 13CO (tracing gas deeper into the cloud) in the pc-scale, massive clumps of the CHaMP survey, finding systematic velocity differentials between the two lines that imply an average mass accretion time-scale of ∼16 Myr, consistent with the time-scales we discuss in this paper (cf. Section 7.1 and Fig. 13).
Citation: Myers 2009
Citation function: Differences
The close approach of a star to a supermassive black hole (SMBH) can lead to the destruction of the stellar body in a process known as a tidal disruption event (TDE; Hills 1975). Gravitationally bound material returns to the black hole and is accreted, giving rise to a flare whose electromagnetic signature peaks in the extreme ultraviolet (EUV) band (Rees 1988; Ulmer 1999). These flares were first detected in the soft X-ray band by ROSAT (Komossa et al. 2004; Bade et al. 1996; Komossa & Greiner 1999; Komossa & Bade 1999), later by XMM-Newton and Chandra (Esquej et al. 2007; Saxton et al. 2012a, 2017; Maksym et al. 2010; Lin et al. 2015) (see review by Komossa 2017), and also in the UV band by GALEX (Gezari et al. 2006, 2008, 2009). In recent years, large-area optical surveys have detected candidate TDEs emitting at temperatures of a few ×104 K (van Velzen et al. 2011; Cenko et al. 2012a; Gezari et al. 2012; Arcavi et al. 2014; Holoien et al. 2016), ostensibly too cool to be coming from an accretion disc (e.g. Bonning et al. 2007). This optical radiation has been interpreted as being due to reprocessing of the accretion radiation by an optically thick screen (Metzger & Stone 2016; Roth & Kasen 2018; Dai et al. 2018) or to emission from shocks (Piran et al. 2015). Super-Eddington accretion in the initial phase of the disruption causes a large-scale, radiation-driven outflow of material from the central engine (Strubbe & Quataert 2009), which Metzger & Stone (2016) showed would initially completely absorb the radiation from the central engine and convert it into optical/UV photons with an effective temperature similar to that observed. In this model, the screen density is expected to drop after a few months to the point where the inner thermal radiation would become visible, with the delay time and ratio of X-ray to optical/UV flux depending on the line of sight (Metzger & Stone 2016; Dai et al. 2018). Observationally, the evidence for differences in the X-ray and UV/optical timescales is mixed. The X-rays may have lagged the UV by ∼32 days in ASASSN-14li (Pasham et al. 2017), but broadly fell on the same timescale (Brown et al. 2017), as they did in 2MASX 0740-85 (Saxton et al. 2017). In SDSS J1201+30 the UV flux did not change, while the X-rays dropped by a factor of 100 (Saxton et al. 2012a), whereas in ASASSN-15oi the X-ray luminosity was quite low (LX ∼ 1041 ergs s−1) at the peak of the optical flare, but 200–300 days later had increased to LX ∼ 1042ergs s−1 (Holoien et al. 2018; Gezari et al. 2017).
Citation: Cenko et al. 2012a
Citation function: Background
The simplest explanation for these metal abundances, i.e., that they reflect the yields of normal core-collapse SNe (averaged over the stellar initial mass function (IMF)), fails to predict anything like the observed stellar abundances of the extremely metal-poor stars ([Fe/H]  −3). Of course, at these low metallicities, the number of progenitor SNe enriching the ISM may be small, so models typically allow for metal-poor or metal-free progenitor stars, and an arbitrary mix of progenitor stellar masses (i.e., assuming the abundances might come from just one or at most a few SNe with individual explosion and progenitor parameters fitted to the observations). However, even with these degrees of freedom, the models still often fail to explain the abundances of certain individual species at the order-of-magnitude level (see, e.g., Nomoto et al. 2006; Heger & Woosley 2010; Lee et al. 2014; Placco et al. 2015), although they undoubtedly explain many of the observed abundance ratios. For the lowest metallicity stars observed ([Fe/H]  −4), and in particular for the CEMP stars, these remaining discrepancies have led to more “exotic” models with a number of free parameters, invoking a mix of normal/faint SNe and hypernovae (with variable explosion energies of ∼1051–1054 erg); jets, prior “failed explosions,” and fallback episodes; rotation and adjustable mixing layers allowing for a tunable stellar abundance profile in the progenitor stars; and pollution of the stars via companions (e.g., Tominaga et al. 2007; Ishigaki et al. 2014; Takahashi et al. 2014; Abate et al. 2015). These additions can improve the agreement with observations; however, there is still no consistent theoretical scenario that simultaneously explains most of the observed stars, and even the best-fit models for many individual stars can still have order-of-magnitude discrepancies with certain outlier elements (see Tominaga et al. 2014; Frebel et al. 2015; Placco et al. 2015, and references therein).
Citation: Lee et al. 2014
Citation function: Differences; Similarities
The dynamical importance of CRs is even more uncertain. This is in part because most early work on this question focused only on galactic conditions similar to those found locally (Jokipii 1976; Badhwar & Stephens 1977; Ghosh & Ptuskin 1983; Chevalier & Fransson 1984; Boulares & Cox 1990; Ko, Dougherty & McKenzie 1991; Ptuskin 2001), and/or focused largely on the question of how and whether CRs can drive galactic winds originating in the ionized, low-density medium found several scale heights above galactic planes (Ipavich 1975; Breitschwerdt, McKenzie & Voelk 1991; Zirakashvili et al. 1996; Ptuskin et al. 1997; Zirakashvili & Völk 2006; however, for an exception see Breitschwerdt, McKenzie & Voelk 1993). More recent numerical and analytic models have continued in this vein (e.g. Everett et al. 2008; Jubelgas et al. 2008; Samui, Subramanian & Srianand 2010; Wadepuhl & Springel 2011; Uhlig et al. 2012; Booth et al. 2013; Pakmor et al. 2016; Simpson et al. 2016; Recchia, Blasi & Morlino 2016, 2017; Ruszkowski, Yang & Zweibel 2017; Pfrommer et al. 2017; Buck et al. 2019), rather than address the question of whether CRs represent a significant contribution to the support of the neutral material that dominates the total mass budget and occupies at least $\sim 50{{\ \rm per\ cent}}$ of the volume (e.g. Dekel et al. 2019) near the mid-plane. Indeed, the vast majority of published simulations that include CR transport do not resolve the neutral phase or galactic scale heights (∼100 pc), and those that do (e.g. Hanasz et al. 2013; Salem & Bryan 2014; Salem, Bryan & Corlies 2016; Chan et al. 2019) generally assume that CR transport in the neutral ISM is identical to that in the ionized ISM (though see Farber et al. 2018), an assumption that is almost certainly incorrect (e.g. Zweibel 2017; Xu & Lazarian 2017; Krumholz et al. 2020). Only a few published models attempt to address the question of CR pressure support in the neutral ISM for non-Solar neighbourhood (mostly starburst or Galactic Centre) conditions (e.g. Thompson et al. 2006; Socrates, Davis & Ramirez-Ruiz 2008; Lacki, Thompson & Quataert 2010; Lacki et al. 2011; Crocker et al. 2011; Crocker 2012; Lacki 2013; Yoast-Hull, Gallagher & Zweibel 2016; Yoast-Hull & Murray 2019; Krumholz et al. 2020).
Citation: Recchia, Blasi & Morlino 2016
Citation function: Background; Background; Background
Solar-like oscillations are excited stochastically by motions in the convective envelope of stars. For low-mass stars, they are found in all evolutionary states, between the main sequence and horizontal branch of helium-core burning stars (e.g. Leighton et al. 1962; Frandsen et al. 2002; Carrier et al. 2003; Hekker et al. 2009; Chaplin et al. 2011; Huber et al. 2011; Kallinger et al. 2012; Mosser et al. 2013) and were even detected in the M5 super giant α Her (Moravveji et al. 2013). These very characteristic oscillations lead to a nearly regular spaced comb-like pattern in the power spectrum. It was shown by Deubner (1975) that this ridge structure is governed by the degree of oscillation modes and resembles the predictions made by Ando & Osaki (1975). Empirically, the frequency patterns of solar-like oscillations were described through scaling relations by Kjeldsen & Bedding (1995). Since then, these relations have been tested and revised from large sample studies based on high-precision space photometry of red giants in clusters and in eclipsing binaries besides single stars (e.g. Corsaro et al. 2012b; Frandsen et al. 2013; Kallinger et al. 2010). Indications of non-radial oscillation modes were found in the variations of the absorption lines of bright red giants (Hekker et al. 2006; Hekker & Aerts 2010), but were firmly established in a large set of red giants observed with the CoRoT satellite (De Ridder et al. 2009). The identification of dipole mixed modes extended the sensitivity of the seismic analyses towards the core of evolved stars also (Beck et al. 2011; Bedding et al. 2011; Mosser et al. 2011). The analysis of solar-like oscillations enabled us to unravel many open questions on stellar structure and evolution, such as constraining the internal rotational gradient (Elsworth et al. 1995; Beck et al. 2012; Deheuvels et al. 2012) or determining the evolutionary status in terms of nuclear burning of a given red giant star (Bedding et al. 2011; Mosser et al. 2011).
Citation: Deubner (1975)
Citation function: Background; Background
In this context, massive (ℳ>1011 M⊙) PGs (MPGs) deserve particular attention. These systems are expected to evolve mainly through (dry) mergers (e.g. Hopkins et al. 2009; De Lucia & Blaizot 2007). If this is the case, in this mass range we should detect a stronger signal of the size-growth with respect to a lower mass range. So far, because MPGs are extremely rare, there have been very few studies that have investigated the combined evolution of the number density and of the age of MPGs as a function of their compactness (Carollo et al. 2013; Fagioli et al. 2016). Carollo et al. (2013) found that the number density of massive quiescent and elliptical galaxies with Re 2.5 kpc decreases by about 30% from z~1 to z~0.2 and that their U−V colours are consistent with passive evolution. They concluded that the driving mechanism for the average size-growth of the whole population is the appearance at later epochs of larger quiescent galaxies. More recently, Fagioli et al. (2016, hereafter F16) analysed the spectroscopic properties of ~500 MPGs (defined as galaxies with absent or very weak emission lines and no MIPS detections) at 0.2 z 0.8 in the zCOSMOS-bright 20 K catalogue (Lilly et al. 2007). From the analysis of stacked spectra of small and large MPGs, they dated the stellar content of these groups and found that the two sub-populations have similar ages. The authors concluded that, in this mass regime, the size growth of individual galaxies through dry mergers is the most likely explanation for the increase in the mean effective radius of the whole population. A recent analysis by Zahid et al. (2016) on the physical properties of compact post starburst galaxies at 0.2z0.8 with ℳ>1011 M⊙ provides new insights. On the basis of both their number density and of their ages, which have been found to be 1 Gyr, the authors suggest that this class of objects are the progenitors of compact quiescent galaxies. They conclude that a substantial fraction of dense quiescent galaxies at z0.8 are newly formed.
Citation: Carollo et al. (2013)
Citation function: Background