adsabs/FOCAL huggingface dataset:
The simulation model Run1-hr of Meyer et al. (2018) begins with the gravitational collapse of $100\, \rm M_{\odot }$ of pre-stellar rotating molecular material on to the stellar embryo. At the end of the free-fall collapse phase, the infalling material ends on a centrifugally-balanced disc, from which the gas is subsequently transferred to the growing protostar. Fig. 1 plots the accretion rate history on to the central massive protostar (solid blue line, in $\rm M_{\odot }\, \rm yr^{-1}$) together with the evolution of the protostellar mass (dashed red line, in $\rm M_{\odot }$) that is calculated as the integrated disc-to-star mass transfer rate through the sink cell. The vertical thin black line indicates the onset of disc formation, when the free-fall collapse of the envelope material on to the protostar stops and the star begins to gain its mass exclusively via accretion from its surrounding disc. After the initial infall of material, the collapse of the parent pre-stellar core material generates an initial increase of the mass flux through the inner boundary at a time ${\approx } 2\, \rm kyr$. The accretion rate then reaches the standard value predicted for MYSOs of $10^{-3}\, \rm M_{\odot }\, \rm yr^{-1}$ (Hosokawa & Omukai 2009) up to the onset of the disc formation happening at ${\approx } 12\, \rm kyr$. Variabilities in the accretion flow begin right after the disc formation and the accretion rate history exhibit numerous peaks of growing intensity as the MYSO becomes heavier. This is not caused by different inner boundary conditions, but simply reflects the time-dependent azimuthal anisotropies induced in the accretion flow by the disc evolution. The efficient gravitational instabilities produce complex substructures in the disc, such as overdense spiral arms in which gaseous clumps of various morphologies form at radii of ${\sim } 100\, \rm au$ and inward-migrate down to the central protostar. This produces luminosity outbursts via the mechanism revealed in Meyer et al. (2017) for massive stars. These outbursts are responsible for step-like increases in the stellar mass evolution (see thick dotted red line) due to the fast accretion of dense circumstellar material inside the sink cell. A more detailed description of the evolution of circumstellar discs around young massive stars irradiating their self-gravitating discs can be found in our precedent study (Meyer et al. 2018). In Fig. 1, several magenta dots mark the time instances of the chosen simulation snapshots considered in this study. Note that the young star becomes, by definition, a massive object when $M_{\star }=8\, \rm M_{\odot }$. Hence, our selected disc models are all in the high-mass regime.
Citation: Hosokawa & Omukai 2009
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A large proportion of systems where one planet or more is orbiting closer to its host star than Mercury to the Sun have been observed. Tidal interactions play a key role in the orbital configuration of these very compact systems since they are likely to circularize orbits, align spins, and synchronize periods (Zahn 1977; Mathis & Remus 2013; Ogilvie 2014). These interactions consists in an exchange of angular momentum between the orbit and the spins of the celestial bodies. This exchange is the consequence of the dissipation of tidal flows. The kinetic energy of these tidal flows is converted into heat through tidal dissipation. Since the planet is synchronized within a timescale of a few thousand years, the stellar tide drives the secular orbital evolution (Guillot et al. 1996; Rasio et al. 1996; Leconte et al. 2010). In this work, we neglect the impact of the dissipation in the radiative zone. In stellar convection zones, there are two kinds of tides and both are dissipated by the turbulent friction applied by convective eddies. On the one hand, the equilibrium tide is the large-scale velocity field associated with tidal deformation, the so-called tidal bulge. This nonwave-like entity corresponds to the hydrostatic adjustment of the star to the gravitational perturbation (Zahn 1966; Remus et al. 2012). The friction applied by convective motions delays the response of the star to the perturbation (e.g., Zahn 1989; Ogilvie & Lesur 2012; Mathis et al. 2016). This results in a lag angle between the axes of the tidal bulge and the line of centers. This angle increases with dissipation magnitude. Hansen (2012) calibrated its value for several stellar masses by constraining the dissipation using observations of planetary systems. Since lower-mass stars have deeper convective envelopes, they dissipate more energy than higher-mass stars. On the other hand, in rotating bodies such as stars, at low tidal frequencies, the Coriolis acceleration acting on this equlibrium tide excites inertial modes (Ogilvie & Lin 2007). Their ensemble, the dynamical tide, constitutes the wavelike part of the tidal response. Its dissipation strongly depends on internal structure since it arises from the reflection of inertial modes on the radiative, stably stratified core (Ogilvie 2013; Mathis 2015). The dynamical tide may also vary over several orders of magnitude with rotation since inertial waves are restored by the Coriolis force. At low frequencies, dissipation of the dynamical tide is several orders of magnitude higher than the dissipation of the equilibrium tide (Ogilvie & Lin 2007).
Citation: Leconte et al. 2010
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GRB 170817A, the short GRB directly associated with the first gravitational wave signal GW170817 from a binary neutron star merger, is a great example presenting afterglow spectra with a beautiful single power-law extending from radio to X-ray (Abbott et al. 2017a, b; Goldstein et al. 2017; Savchenko et al. 2017). A successful launch of a relativistic jet in GRB 170817A is supported by the detection of superluminal motion of a compact radio source (Mooley et al. 2018b; Ghirlanda et al. 2019). The jet was observed from off-axis (i.e. the jet axis and the line of sight was misaligned), leading to a faint gamma-ray emission (Abbott et al. 2017c; Ioka & Nakamura 2018, 2019; Salafia & Ghirlanda 2022, and references therein) and early rising afterglow light curves (Alexander et al. 2017; Hallinan et al. 2017; Margutti et al. 2017; Troja et al. 2017). The slow rising afterglow light curves also rejected a uniform top-hat jet (Mooley et al. 2018a; Troja et al. 2018) and revealed a launch of a structured jet, while various jet structures can explain the observations (see, however, Lamb, Levan & Tanvir (2020) for an alternative scenario with a refreshed shock in an off-axis uniform jet). Often assumed structured jets are a Gaussian jet or a power-law jet (D’Avanzo et al. 2018; Gill & Granot 2018; Lamb & Kobayashi 2018; Lyman et al. 2018; Resmi et al. 2018; Ghirlanda et al. 2019; Lamb et al. 2019; Troja et al. 2019, 2020; Ryan et al. 2020). Other non-trivial jet structures such as hollow-cone jets and spindle jets are also possible candidates as recently discovered by Takahashi & Ioka (2020, 2021), while the consistency of a hollow-cone jet was also confirmed by Nathanail et al. (2020, 2021). As predicted by Troja et al. (2018), the rapid decline of the afterglow after the peak (Mooley et al. 2018c; Lamb et al. 2019; Hajela et al. 2020; Makhathini et al. 2021; Troja et al. 2020; Balasubramanian et al. 2021) revealed a distinctive signature of a successfully launched structured jet, setting it apart from a chocked jet scenario. The observed electron power-law index derived from the afterglow spectrum falls within the range predicted in a theory of particle acceleration at trans-relativistic shock (Keshet & Waxman 2005) as mentioned by Margutti et al. (2018) and is consistent with being constant in time within observational errors (D’Avanzo et al. 2018; Dobie et al. 2018; Margutti et al. 2018; Fong et al. 2019; Kilpatrick et al. 2022).
Citation: Mooley et al. 2018a
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The growth of supermassive BHs is also altered when considering non-Gaussianities. After deriving the merger history of the most massive haloes at z = 6.5 in both the Gaussian and non-Gaussian simulations, we study the evolution of BHs in massive haloes down to z = 6.5. To investigate the cumulative effect over cosmic times on the BHs assembly, we model the growth of BHs in three different ways. Different probabilities for a halo of hosting a seed BH, and different accretion models (either each BH accretes at the Eddington limit for a dynamical time after a major merger or using an accretion rate based on a distribution probability derived from a large-scale hydrodynamical simulation) are adopted. We have not included in our models the effects of ‘kicks’ caused by asymmetric emission of gravitational waves, which have been proposed to be possibly responsible for ejecting BHs from haloes with shallow potential wells, thus halting or reducing the growth of high-redshift BHs hosted in small haloes (e.g. Yoo & Miralda-Escudé 2004; Volonteri & Rees 2006; Tanaka & Haiman 2009). This effect, however, seems to affect less than 10 per cent of binaries and it becomes negligible for BH mergers at z  10 (Volonteri & Rees 2006). We find that non-Gaussianities imply a larger number of massive BHs and also an increase in the mean BH mass (up to 0.36 in the most favourable experiment). A population of supermassive BHs will then grow faster and to higher masses in a universe with scale-dependent non-Gaussian primordial density fluctuations. If the seed masses are similar to those of Pop III star remnants, BHs will not be able to grow above few × 105 M⊙ by z = 6. However, our simulations do not resolve mini-haloes, and we may underestimate the growth of seeds at earlier times. We argue that, in a simulation resolving mini-haloes, BHs would have formed earlier through the Pop III remnant scenario, leading to a longer time for them to grow in mass. If we assumed that Pop III remnant seeds with mass 100 M⊙ form at z ∼ 30 in haloes unresolved in our simulations, they would have grown, assuming, optimistically, constant growth at the Eddington rate (but see Johnson & Bromm 2007; Alvarez, Wise & Abel 2009; Milosavljević, Couch & Bromm 2009; Park & Ricotti 2011) to ∼103 M⊙ by z = 18, where we start our analysis. The final BH mass at z = 6 would then be ∼ one order of magnitude larger, a few × 106 M⊙, still short of the ∼109 M⊙ required. The very limited growth obtained for the Pop III remnant case suggests that large seeds or super-Eddington accretion (see Volonteri, Silk & Dubus 2015, and references therein) may be necessary for successful BH growth. We have done the same experiments on BH growth starting with initial 105 M⊙ BH masses (not shown in the paper, but see Section 4). In this case we found that it is much easier for BHs to grow to higher BH masses, but still only to several 108 M⊙. This is not unexpected, because our simulation box does not contain the very rare and biased DM haloes with masses ∼1013 M⊙ believed to be hosting these extreme BHs.
Citation: Yoo & Miralda-Escudé 2004
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The results presented in this work answer a number of questions regarding the character of the GG Tau A system, while raising others and leaving others untouched. First among the questions raised by our results are the questions of whether the detailed morphology of features in the circumbinary torus as seen in our simulations are actually present in the GG Tau A system. The current best resolution observations of the torus provide tantalizing hints that such features do exist, but a definitive statement that such features are present must await higher resolution observations of quality similar to those described by the ALMA Partnership (2015), for the HL Tau circumstellar disk. Another important question concerns planet formation in multiple systems. Whether or not any theoretical mechanism presently exists to explain the formation of giant planets in binary systems, the fact remains that at least a few binary systems, such as γ Cephei (Neuhauser et al. 2007), GI86 (Queloz et al. 2000), and HD 41004 (Zucker et al. 2004), do harbor planets. Therefore, some formation mechanism does in fact exist. While we find that accretion into the circumstellar disks occurs rapidly enough so that they can survive for the comparatively long timescales needed to form planets, other conditions, such as temperatures, remain quite unfavorable. What are the mechanisms still missing from our models that permit such objects to form? Finally, our simulations model the evolution of the “A” component of the full GG Tau system over a time span extending over only a tiny fraction of its formation timescale and in only two spatial dimensions. We neglected the full dynamical effects expected to be present in the system, insofar as our results include neither the distant binary “B” component of the GG Tau system, nor the newly discovered tight binary nature of the GG Tau Ab component. Even so, we find large scale morphological changes even over this short time span, and restricted dimensionality and physical system. Given the vigor of the activity over such a short timescale, we would expect activity of similar scale to occur over longer time spans as well, with correspondingly large consequences on the system morphology. To what extent will 3D effects also play a role in the evolution? What will be the end state configuration of the GG Tau system as a whole? Will the components eventually break apart? Merge? Future investigations extending the work presented here will be required in order to answer these questions.
Citation: Neuhauser et al. 2007
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The Orion Bar is probably the best studied PDR in our Galaxy. It is located between the Orion Molecular Cloud 1 and the H ii region excited by the Trapezium cluster, and is exposed to an FUV field a few 104 times the mean interstellar radiation field. Owing to its proximity (417 pc, Menten et al. 2007) and nearly edge-on orientation, the Bar provides an ideal laboratory for testing PDR models (e.g., Jansen et al. 1995; Gorti & Hollenbach 2002; Andree-Labsch et al. 2017) and a primary target for observational studies of physical and chemical structures of PDRs (e.g., Tielens et al. 1993; Walmsley et al. 2000; van der Wiel et al. 2009; Arab et al. 2012; Peng et al. 2012; Goicoechea et al. 2016; Nagy et al. 2017). Observations of various molecular spectral lines have shown that the emissions could be better interpreted with an inhomogeneous density structure containing an extended and relatively low density (





 cm−3) medium and a compact and high-density (





 cm−3) component (e.g., Hogerheijde et al. 1995; Young Owl et al. 2000; Leurini et al. 2006, 2010; Goicoechea et al. 2016). However, due to the scarcity of high-resolution observations capable of spatially resolving the density structure, the nature of the high-density clumps or condensations is still not well understood. Lis & Schilke (2003) mapped the Bar in H13CN (1–0) with the Plateau de Bure Interferometer (PdBI) at an angular resolution of about 5″, and detected 10 dense clumps. They proposed that the H13CN clumps are in virial equilibrium and may be collapsing to form stars. Goicoechea et al. (2016) performed Atacama Large Millimeter/submillimeter Array HCO+ (4–3) observations of the Bar and detected over-dense substructures close to the cloud edge, and found that the substructures have masses much lower than the mass needed to make them gravitationally unstable. These two interferometric observations both target molecular spectral lines. A high-resolution map of the dust continuum emission of the Bar, which is highly desirable in constraining the mass and density of the dense condensations, is still lacking. Here we report our Submillimeter Array (SMA) observations of the dust continuum and molecular spectral line observations of the Bar.
Citation: Nagy et al. 2017
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It has also become evident that the normalization of the MS rapidly increases with redshift: distant galaxies form stars at higher paces than in the local Universe, at fixed stellar mass (e.g., Daddi et al. 2007; Elbaz et al. 2007; Whitaker et al. 2012; Speagle et al. 2014; Schreiber et al. 2015). This trend could be explained by the availability of copious molecular gas at high redshift (Daddi et al. 2010a; Tacconi et al. 2010, 2018; Scoville et al. 2017a; Riechers et al. 2019; Decarli et al. 2019; Liu et al. 2019a), ultimately regulated by the larger accretion rates from the cosmic web (Kereš et al. 2005; Dekel et al. 2009a). Moreover, higher SFRs could be induced by an increased efficiency of star formation due to the enhanced fragmentation in gas-rich, turbulent, and gravitationally unstable high-redshift disks (Bournaud et al. 2007, 2010; Dekel et al. 2009b; Ceverino et al. 2010; Dekel & Burkert 2014), reflected on their clumpy morphologies (Elmegreen et al. 2007; Förster Schreiber et al. 2011; Genzel et al. 2011; Guo et al. 2012, 2015; Zanella et al. 2019). IR-bright galaxies with prodigious SFRs well above the level of the MS are observed also in the distant Universe, but their main physical driver is a matter of debate. While a star formation efficiency (SFE = SFR/Mgas) monotonically increasing with the distance from the main sequence (ΔMS = SFR/SFRMS, Genzel et al. 2010, 2015; Magdis et al. 2012; Tacconi et al. 2018, 2020) could naturally explain the existence of these outliers, recent works suggest the concomitant increase of gas masses as the main driver of the starbursting events (Scoville et al. 2016; Elbaz et al. 2018). In addition, if many bright starbursting (sub)millimeter galaxies (SMGs, Smail et al. 1997) are indeed merging systems as in the local Universe (Gómez-Guijarro et al. 2018, and references therein), there are several well documented cases of SMGs hosting orderly rotating disks at high redshift (e.g., Hodge et al. 2016, 2019; Drew et al. 2020), disputing the pure merger scenario. The same definition of starbursts as galaxies deviating from the main sequence has been recently questioned with the advent of high spatial resolution measurements of their dust and gas emission. Compact galaxies with short depletion timescales typical of SBs are now routinely found on the MS, being possibly on their way to leave the sequence (Barro et al. 2017a; Popping et al. 2017; Elbaz et al. 2018; Gómez-Guijarro et al. 2019; Puglisi et al. 2019; Jiménez-Andrade et al. 2019); or galaxies moving within the MS scatter, due to mergers unable to efficiently boost the star formation (Fensch et al. 2017) or owing to gravitational instabilities and gas radial redistribution (Tacchella et al. 2016).
Citation: Elmegreen et al. 2007
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Citation: Ceverino et al. 2010
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The Pleiades open cluster does not only offer us a beautiful spectacle during the fall, it is one of the best-studied stellar associations and one of the cornerstones to understand stellar properties and evolution. In fact, the literature includes more than one thousand refereed papers dealing with the Pleiades, most of which use the Pleiades as a reference, just in the last ten years. In spite of this, the Pleiades cluster still has many secret and basic parameters, such as its distance and age, that are not clearly established. Even the census of this cluster is incomplete, although the recent works by Stauffer et al. (2007), Lodieu et al. (2012), and Bouy et al. (2013) have considerably improved the membership list, In fact, Bouy et al. (2015) has increased the number of known members by 50% with public archival data, very accurate proper motions, and multiwavelength photometry (see additional details in Sarro et al. 2014). Regarding its distance, there are currently two different methodologies based on parallaxes from Hipparcos (Perryman et al. 1997) and isochrone fitting, respectively. Pre-Hipparcos distances for the Pleiades range between 125 and 130 pc (see, for instance, Soderblom et al. 1993d), whereas the initial distance derived by Hipparcos is much closer, about 119 pc (van Leeuwen 1999). This last value is significantly different from the distance derived by Pinsonneault et al. (1998), who used color-magnitude diagrams and fitting isochrones and obtained 133.5 ± 1.2 pc. More recently, van Leeuwen (2009), by reanalyzing Hipparcos data, derived a distance of 120.2 ± 1.9 pc. These values should correspond to the distance to the cluster center, whose core radius should be around 3 degrees, which corresponds to 5–6 pc. Another trigonometric parallax, based on Hubble Space Telescope data and three members, was found by Soderblom et al. (2005), who obtained 134.6 ± 3.1 pc. Recently, Melis et al. (2014) derived a distance of 136.2 ± 1.2 pc based on an accurate parallax for four bona fide members obtained with the VLBI. This value also agrees with that derived by Galli et al. (in prep.) using accurate proper motions and the convergence point method (137.7 ± 2.5 pc).
Citation: Perryman et al. 1997
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The majority of the particles collected by Stardust are olivine and pyroxene silicates with solar isotopic compositions, which suggests an origin in our solar system rather than an interstellar provenance. These polymineralic particles dominate those made of a single mineral even down to sizes smaller than 100 nm, indicating that the dust composition is surprisingly consistent at different scales and that the smallest subunits of the dust may be as small as tens of nanometers (Hörz et al. 2006; Zolensky et al. 2006). The sizes of these smallest single mineral impactors are similar to those of the nanocrystals determined by Rietmeijer (1993). As discussed above, they might also be existing in MIDAS dust particles and might be fused into the 100 nm features. Price et al. (2010) and Wozniakiewicz et al. (2012) investigated the sizes of particles smaller than 10 μm that impacted the aluminum foils of the Stardust probe. The distribution peaks at about 175 nm, but if we assume that the particles areagglomerates of smaller subunits, as indicated by their common polymineralic nature, then the subunit size distribution would peak at sizes below 100 nm (Price et al. 2010). A study of over 450 particles that do not seem to be agglomerates, that is, those that show single mineral impactors of silicate or sulfide, found geometric mean sizes of 

$532^{741}_{-310}$




532

−310

+741



 nm for the silicate particles and 

$406^{491}_{-222}$




406

−222

+491



 nm for the sulfides (Wozniakiewicz et al. 2013). These sizes are notably larger than the 175 nm (or less) found for the whole dataset. This large spread of subunit sizes could indicate a size distribution with a large width. No fits of these sizedistributions are available, but the figures in Wozniakiewicz et al. (2012) and Price et al. (2010) indicate that the differential sizes may follow a log-normal distribution. When we assume that the smallest subunit sizes are possibly between tens and hundreds of nanometers, the subunit size range found for MIDAS smallest features would be encompassed. The determination of the size distributions for the small Stardust particles and a detailed comparison to the distributions obtained for comet 67P could be the work of an interesting future project.
Citation: Wozniakiewicz et al. (2012)
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Citation: Wozniakiewicz et al. (2012)
Citation function: Uses; Uses