adsabs/FOCAL huggingface dataset:
In active galactic nuclei (AGNs), the gaseous accretion disc around the SMBH may facilitate binary formation and mergers of stellar-mass compact objects (McKernan et al. 2014; Bellovary et al. 2016; Bartos et al. 2017; Stone, Metzger & Haiman 2017; Tagawa, Haiman & Kocsis 2020a). In this scenario, BHs form in situ in the vicinity of a GN and sink to the inner region due to mass segregation or they are delivered to these regions by infalling globular clusters (Morris 1993; Miralda-Escudé & Gould 2000; Freitag, Amaro-Seoane & Kalogera 2006; Hopman & Alexander 2006; O’Leary et al. 2009; Antonini 2014), then get captured in the disc by hydrodynamic drag as they cross the disc (e.g. Goldreich, Lithwick & Sari 2002; Bartos et al. 2017; Yang et al. 2019b; Tagawa et al. 2020a). Alternatively, some BHs may have formed in the disc itself (Levin 2007; Stone et al. 2017). Once in the disc, BHs get transported to the inner regions by exchanging angular momentum with the surrounding gas (Goldreich & Tremaine 1979). In certain regions, the BHs open an annular gap in the accretion disc and accumulate in a narrow range of radii, the so-called migration traps (Bellovary et al. 2016; Secunda et al. 2019, 2020, 2021). Bellovary et al. (2016) argue that migration traps may be expected to be close to the SMBH from ∼20 to ∼300 Schwarzschild radii (rS = 2GMSMBH/c2) from the central SMBH of mass MSMBH.6 However, they may exist in slim discs near the innermost stable circular orbit (ISCO; Peng & Chen 2021) or near the boundary of a gap region if a gap opens due to a heavy stellar-mass BH or an intermediate-mass BH (McKernan et al. 2014). Dynamical encounters frequently happen in migration traps leading to the formation and subsequent merger of binary black holes (BBHs) on short time-scales (Secunda et al. 2019, 2020; Yang et al. 2019a), where the binary separation is efficiently reduced by gas dynamical friction (Escala et al. 2004; Kim & Kim 2007; Baruteau, Cuadra & Lin 2011) to the point where GW emission drives the binaries together. Alternatively, BBHs may also form and merge in the disc outside migration traps (Tagawa et al. 2020a). Because of the deep potential barrier of the SMBH, the merger remnant BH remains near the migration trap and may undergo subsequent mergers with additional BHs, which leads to high BH masses, characteristic spin properties, and possibly non-zero eccentricity identifiable via GW observations (Yang et al. 2019a; Secunda et al. 2020; Tagawa et al. 2020b, 2021a,b; Samsing et al. 2022).
Citation: Tagawa et al. 2020a
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Samples of dust particles from comet 81P/Wild 2 were returned to Earth by the Stardust mission in 2006. They were collected in aerogel with entrance velocity of 6 km s−1 and some of them made craters in the aluminium foils separating the aerogel tiles. Analyses of 81P/Wild 2 samples trapped in aerogel and of residues in the craters in the Al foils have shown that this comet contains mineral components similar to those found in primitive meteorites, including fragments of refractory minerals found in Ca–Al-rich inclusions (CAIs) (Zolensky et al. 2006; Simon et al. 2008; Schmitz et al. 2009), and chondrule fragments (Leroux et al. 2008a; Nakamura et al. 2008). However, during the capture process, the impacting 81P/Wild 2 dust particles were dispersed along the impact tracks and mixed with the aerogel (Leroux et al. 2008b; Stodolna et al. 2012). The mineralogy of terminal particles (at the end of the tracks) is dominated by olivine and pyroxene minerals, the number of pyroxene grains compared to that of olivine being higher (∼1) than that currently observed in carbonaceous chondrites, where olivine grains dominate the silicate mineralogy (number of pyroxenes/olivines 0.1) (Zolensky et al. 2006; Tomeoka et al. 2008). Apart from these single minerals present at the end of the impact tracks, fine-grained particles mixed with melted aerogel have partially survived the impact and are present along the tracks, together with significant amounts of crystalline silica (Leroux 2012; Roskosz & Leroux 2015). The oxygen isotopic composition of the minerals exhibits 18O/16O and 17O/16O ratios within the range observed in carbonaceous chondrites (Nakashima et al. 2012; Ogliore et al. 2015). No evidence for 26Mg excess was found in the 81P/Wild 2 minerals, suggesting that short-lived 26Al radionuclide was not incorporated at the time of their formation, implying a formation a few million years after the formation of the meteoritic CAIs (Ogliore et al. 2012; Nakashima et al. 2015). The abundance of isotopically anomalous presolar grains in 81P/Wild 2 is estimated to be less than 1000 ppm (Floss et al. 2013).
Citation: Zolensky et al. 2006; Zolensky et al. 2006
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Intensive and extensive efforts in numerical simulations have revealed so far that no successful explosion is obtained in spherical symmetry (e.g., Liebendörfer et al. 2001, 2005; Rampp & Janka 2002; Sumiyoshi et al. 2005), and that multidimensional effects are crucially important (Burrows et al. 2006; Bruenn et al. 2009; Marek & Janka 2009; Suwa et al. 2010; Müller et al. 2012; Takiwaki et al. 2012; Couch 2013; Couch & Ott 2013; Hanke et al. 2013; Murphy et al. 2013; Lentz et al. 2015; Melson et al. 2015; Nakamura et al. 2015; Bruenn et al. 2016; Roberts et al. 2016; O’Connor & Couch 2018). Among them are rotation (Fryer & Heger 2000; Kotake et al. 2003; Thompson et al. 2005; Marek & Janka 2009; Iwakami et al. 2014a; Nakamura et al. 2014; Takiwaki et al. 2016; Summa et al. 2018), a magnetic field (Akiyama et al. 2003; Kotake et al. 2004; Yamada & Sawai 2004; Sawai et al. 2005; Obergaulinger et al. 2006, 2014, 2018; Burrows et al. 2007; Takiwaki et al. 2009; Sawai & Yamada 2014, 2016; Mösta et al. 2015), non-spherical structures of the progenitor (Couch & Ott 2013; Takahashi & Yamada 2014; Couch et al. 2015; Takahashi et al. 2016), turbulence (Murphy & Burrows 2008; Murphy & Meakin 2011; Murphy et al. 2013; Couch & Ott 2015; Mabanta & Murphy 2018), (magneto)hydrodynamical instabilities (Blondin et al. 2003; Scheck et al. 2006; Blondin & Mezzacappa 2007; Iwakami et al. 2008; Guilet et al. 2010; Wongwathanarat et al. 2010; Fernández et al. 2014; Takiwaki et al. 2014; Fernández 2015), general relativistic gravity (Dimmelmeier et al. 2002; Shibata & Sekiguchi 2004, 2005; Kuroda et al. 2012, 2016; Müller et al. 2012; Ott et al. 2012), and neutrino transport (Nagakura et al. 2014, 2017, 2018; Dolence et al. 2015; Pan et al. 2016). It is true that large-scale dynamical simulations have played a crucial role in recent progresses in our understanding of these ingredients, but we believe that a more phenomenological approach that employs toy models still plays an indispensable and complementary role to understand each effect more deeply.
Citation: Takiwaki et al. 2016
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The many unexpected surprises of comet 67P/Churyumov-Gerasimenko (hereafter 67P) revealed by the historic Rosetta mission highlight the importance of observing the evolution of comets throughout their orbits. One of the surprises was the drastic heterogeneity in both the major and minor volatile species in the coma that was observed early on in the mission (Hässig et al. 2015; Luspay-Kuti et al. 2015, hereafter ALK15). When Rosetta first arrived at comet 67P in August 2014, the Rosetta Orbiter Mass Spectrometer for Ion and Neutral Analysis/Double Focusing Mass Spectrometer (ROSINA/DFMS; Balsiger et al. 2007) detected large diurnal variations in the intensity profiles of various species in the coma from distances to the comet as far as 250 km. At this time, 67P was still at a distance of about 3 AU and inbound from the Sun. The intensity variations in the major and minor volatile species were found to be periodic, and were dependent on both the observing sub-spacecraft latitude and longitude (Hässig et al. 2015; Luspay-Kuti et al. (2015)). As reported in Hässig et al. (2015), the intensity of H2O in the coma dominated the overall signal, with maxima in the H2O signal every ~6 h, about twice during a rotation. Interestingly, however, CO2 and CO displayed a separate additional maximum when the H2O signal was near its minimum. This independent maximum in CO2 and CO only occurred at negative observing latitudes that are associated with a particular “view” of Rosetta at 67P, with the larger lobe blocking out the neck and head. At this time, 67P had not yet reached its first equinox (10 May 2015), and the poorly illuminated southern hemisphere was experiencing winter. In addition, the largest H2O activity was localized at the well-illuminated neck region, as also seen by the Microwave Instrument on the Rostta Orbiter (MIRO; Gulkis et al. 2015; Biver et al. 2015; Lee et al. 2015) and by the Visible InfraRed Thermal Imaging Spectrometer (VIRTIS; Bockelée-Morvan et al. 2015; Migliorini et al. 2016). VIRTIS also measured weak H2O production in regions with low solar illumination, while CO2 was outgassing from both illuminated and non-illuminated regions pre-inbound equinox (Bockelée-Morvan et al. 2015; Migliorini et al. 2016; Fink et al. 2016). The observed outgassing pattern of the major cometary species suggested that CO and CO2 may be sublimating from a depth below the diurnal skin depth.
Citation: Hässig et al. 2015
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We use the N-body simulation package—MERCURY (Chambers 1999)—to numerically investigate the effects of photo-evaporation on the dynamical evolution of planet–satellite systems. We choose the Bulirsch–Stoer integration algorithm, which can handle close encounter accurately. It is important in the simulations, as we will see below, that many close encounters among moons and the planet are expected to happen. Collisions among moons, the planet, and the central star are also considered in simulations and treated simply as inelastic collisions without fragmentations. Each simulation consists of a central star, a planet, and some moons orbiting around the planet. The photo-evaporation is simply modeled as a slow (adiabatic) and isotropic mass-loss process of the planet. In reality, the photo-evaporation is a very slow process on a timescale of the order of 107–108 year (Owen & Wu 2013). However, it is impractical and unnecessary to perform a simulation on such a long timescale. Instead, we model the mass-loss process on a timescale of τevap, and each simulation typically lasts for several τevap. As long as the adiabatic requirement is met, i.e., the mass-loss timescale is much longer than the dynamical timescale of the system (τevap ≫ Pp, where Pp is the orbital period of the planet), one could study the dynamical effects of the mass-loss process equivalently. As we discussed in Section 3.3, the results converge if τevap > 102–103 Pp, indicating the adiabatic condition is met. Therefore, in all other simulations, we set τevap = 104 Pp. Other parameters are set to represent the typical values of Kepler planets. In particular, we consider a planet–satellite system orbiting a star of solar mass (M⋆ = M⊙) in a circular orbit (ep = 0.0) with semimajor axis of ap = 0.1 au. The orbit has a period of ∼10 days (typical value of Kepler planets), and it is sufficiently close to the central star to be subject to significant photo-evaporation effect (Owen & Wu 2013), which removes massive hydrogen envelopes of the planet. The planet has an initial mass of Mpi and a final mass of Mpf after photo-evaporation. In this paper, we adopt Mpi = 20 M⊕ and Mpf = 10 M⊕ nominally (close to the standard model adopted in Owen & Wu 2013). The mean density of the planet is set to the same as that of Neptune (1.66 g cm−3). The effect of changing the planetary density is discussed in Section 3.3. We performed a number of sets of simulations by considering different planet–satellite configurations. Similar to the definition in MERCURY, hereafter, we define “small moons” as test particles (TPs) whose mutual gravity and corresponding effects on the planet and the star are ignored, while “big moons” are gravitationally important enough that their gravitational effects are fully considered. Table 1 lists the initial setups and parameters of various simulations, whose results are presented in the following subsections.
Citation: Owen & Wu 2013
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Summarizing the discussion above: we have a number of suggested r-process sites, but only one of them is proven by a direct observation of the explosive event. Observations of low metallicity stars show essentially three types of patterns, a weak or limited r-process, a strong solar-type r-process, and an actinide-boosted r-process (in some publications, this is also referred to as a weak, main, and strong r-process). Whether the latter two types are produced in different sites or a result of variations within the same site (e.g., neutron star mergers) is still debated. The question is now how such observations can point back to the r-process sites, and whether it is possible to identify features which can provide additional insight. A promising approach is to look for correlations among different elements, which might directly identify the nucleosynthesis of a specific site (see e.g., Barklem et al. 2005; François et al. 2007; Mashonkina et al. 2007; Kratz et al. 2008). Cowan et al. (2005) compared the abundances of Fe, Ge, Zr, and r-process Eu in low metallicity stars. They found a strong correlation of Ge with Fe, indicating the same nucleosynthesis origin (core-collapse supernovae), a weak correlation of Zr with Fe – indicating that other sites than regular core-collapse supernovae (without or low Fe-ejection) contribute as well –, and no correlation between Eu and Fe, essentially pointing to a pure r-process origin with negligible Fe-ejection. More recent data from the SAGA and JINA databases (Suda et al. 2008; Abohalima & Frebel 2018) permit a correlation between Eu and Fe for [Eu/Fe]  0.3, that is for stars with lower than average r-process enrichment. If interpreted in a straightforward way, this would point to a negligible Fe/Eu ratio (in comparison to solar ratios) in the major r-process sources, while a noticeable coproduction of Fe with Eu is possible in less strong r-process sources, for example with a weak r-process. Such cases could again be identified with the limited-r entry in observations (Hansen et al. 2018).
Citation: Mashonkina et al. 2007
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Figure 24 shows that the median Re increases for increasing mass (from ~5 Kpc to ~20 Kpc) and, given a stellar mass, our passive ETGs have median sizes smaller than those of the parent sample even by ~15% (at the highest masses). Furthermore, the entire Re distribution is extended to smaller radii in the case of passive ETGs, especially for log (M/M⊙) ≳ 11.5. Remembering that small differences in galaxy size imply large differences in stellar mass density, the derived trends suggest that the ETGs analyzed in this work should have formed from higher density progenitors which, in the hypothesis of no coeval mergers, do not increase their mass during the evolution. On the other hand, we cannot exclude the possibility that galaxies in the parent sample have experienced more dry mergers than our massive and passive galaxies, which have increased their size across cosmic time (Naab et al. 2009; Johansson et al. 2012a). The observed trend is also in agreement with some literature studies at low redshift which also suggest that more compact galaxies contain older stellar populations than larger ones (e.g. Saracco et al. 2009; Shankar & Bernardi 2009; Williams et al. 2010; Poggianti et al. 2013; McDermid et al. 2015). In general, this is qualitatively consistent with the high zF inferred by our analysis (z> 5) since, according to the cosmological evolution of the baryonic matter density, gas was denser at these cosmic epochs. These findings also agree with recent observations of compact systems (with various level of SF) at z ~ 2−3, which could be identified as local massive ETGs progenitors (e.g. Daddi et al. 2004; Cattaneo et al. 2013; Finkelstein et al. 2013; Marchesini et al. 2014; Nelson et al. 2014; Williams et al. 2015). Moreover, recent models of elliptical galaxy formation also predict the formation of compact and dense progenitors at high redshifts (e.g. Johansson et al. 2012a; Naab et al. 2014 and references therein).
Citation: Johansson et al. 2012a
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Citation: Johansson et al. 2012a
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Gravitational-wave astronomy has recently become a reality with the first detection of gravitational waves (GW hereafter) by the Laser Interferometer Gravitational-Wave Observatory (LIGO) and Virgo ground-based interferemeters. To date, eleven events have been reported by the LIGO and Virgo detectors (Abbott 2018), and this number will quickly increase to tens of events in the coming years. Some of these events may correspond to gravitationally lensed events with magnification factors ranging from a few tens to a few hundreds (Dai et al. 2017; Ng et al. 2018; Li et al. 2018a; Smith et al. 2018a,b; Broadhurst et al. 2019). Recent works have studied lensing effects in the existing LIGO/Virgo O1 and O2 events (Hannuksela et al. 2019; Broadhurst et al. 2019), while Smith et al. (2018c) searched for candidate galaxy cluster lenses for the GW170814 event. The most likely lenses for such events would be massive galaxies or galaxy clusters (Ng et al. 2018; Dai et al. 2017; Smith et al. 2018b; Broadhurst et al. 2019). On the other extreme of the lens mass regime, compact objects with masses of a few hundreds to a few tens M⊙ can also act as lenses (Lai et al. 2018). In this case, the geometric optics limit is not valid since the Schwarszchild radius of the lens is comparable to the wavelength of the wave. For these relatively low masses, the lensing effect has a modest impact on the average magnification, but it can introduce a frequency dependence on the magnification (see for instance, Jung & Shin 2019; Lai et al. 2018). An even smaller mass regime was considered in Christian et al. (2018) where the authors find that lenses with a mass as low as 30 M⊙ could be detected with current experiments. They also consider future, higher-sensitivity experiments and show how they can push the limit to even smaller masses of order 1 M⊙. These conclusions are, however, obtained assuming isolated microlenses and without accounting for the effect of the macromodel, or other nearby microlenses. In the small mass regime, microlenses such as neutron stars have been also considered as scattering sources of GWs, and it is found that a GW can be focussed at a focal point near the neutron star surface (Halder et al. 2019; Stratton & Dolan 2019).
Citation: Broadhurst et al. 2019
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Classic examples of this approach to investigating solar eruptions have been the many simulations of the magnetic breakout model, which uses reconnection as the eruption mechanism (Antiochos 1998; Antiochos et al. 1999). In typical breakout simulations magnetic energy is built up in the corona by prescribed surface motions at the lower boundary, either large-scale shearing flows (e.g., Karpen et al. 2012) or small-scale randomized motions leading to helicity condensation (e.g., Dahlin et al. 2022). The normal component of the coronal field is held fixed at this lower boundary, which is presumed to represent the high-beta photosphere, and the coronal flux is assumed to be multipolar so that a separatrix surface and null point are present. Note that even a bipolar AR implies the presence of a coronal separatrix surface and null when the field of the bipole is embedded in the solar global dipole (e.g., Lynch et al. 2008). The initial conditions in most breakout simulations are taken to be the minimum-energy potential field (
J
 ≡ ∇ × 
B
/μ
0 = 0 where 
B
 is the magnetic field), but the imposed boundary motions inject magnetic free energy into the corona, leading to the formation of a sheared-arcade filament channel with current sheets at the overlying separatrix surface and null between the filament flux system and neighboring flux systems. In the breakout model, magnetic reconnection (often called breakout reconnection) at the separatrix current sheet is necessary so that the sheared filament flux can expand upward sufficiently for a vertical current sheet to form below/inside this flux system. The onset of magnetic reconnection in this vertical current sheet (often loosely described as flare reconnection) greatly accelerates the eruption, leading to the ejection of the shear and the relaxation of the corona back down to a near-potential state (e.g., Karpen et al. 2012; Wyper et al. 2017). A related but physically distinct reconnection-driven eruption model is the so-called tether cutting, which posits that with sufficient boundary shear a vertical current sheet can form spontaneously in the filament channel, without the need for breakout reconnection, leading to an eruption (e.g., Moore et al. 2001). Recent simulations using the same approach as the breakout ones just described appear to support this scenario (Jiang et al. 2021).
Citation: Lynch et al. 2008
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Variable stars have a long tradition of being tracers of different stellar populations. The most well known of the pulsating variable stars in dwarf galaxies are RR Lyrae stars, which unequivocally trace an old population (>10 Gyr, Smith 1995). RR Lyrae stars have been found in almost all of the satellites of the Milky Way (recent compilation in Martínez-Vázquez et al. 2019), even in low-luminosity systems such as Segue I (Simon et al. 2011). This is a clear indication that an old population is ubiquitous among dwarf galaxies. On the other hand, anomalous Cepheids, which are pulsating stars above the horizontal branch, may trace an intermediate age population. Their progenitors are likely massive (1–2 M⊙) stars that have evolved off the main sequence (Fiorentino & Monelli 2012). These stars have been commonly found in the classical satellites of the Milky Way but are rare in globular clusters (Clement et al. 2001), which contain exclusively old populations. Since they have been observed also in dwarf spheroidal (dSph) galaxies with predominantly old stellar populations like Draco and Ursa Minor (Nemec, Wehlau & Mendes de Oliveira 1988; Kinemuchi et al. 2008), other formation channels are needed for these stars. Binary evolution can also bring stars to the region of the instability strip where anomalous Cepheids live (Gautschy & Saio 2017). Thus, old progenitors are also possible. Dwarf Cepheid stars (a.k.a. SX Phoenicis and/or δ Scuti stars) are also pulsating stars in the instability strip, but they are found below the horizontal branch (Breger 2000). Their very short periods (just a few hours) and small amplitudes make them very hard to detect, especially in distant systems. They may trace either intermediate-age populations (massive main sequence stars in the instability strip) or old populations (pulsating blue stragglers). Finally, Classical Cepheids indicate the presence of young populations (0.03–0.7 Gyr). They are rare in dSph galaxies, with Leo I being an exception (Fiorentino et al. 2012).
Citation: Fiorentino & Monelli 2012
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