adsabs/FOCAL huggingface dataset:
In the study presented here, two solid-phase experiments were run for 3 hours each, to produce cosmic grain analogs from methane (5%) and acetylene (5%) precursors, respectively, seeded in a cold supersonic expansion of argon. Different concentrations of methane were tested in COSmIC, and the 5% concentration was chosen to produce sufficient solid material for analysis in our experimental setup. Methane is a one carbon-atom alkane known to be present in the CSE of carbon-rich stars (Van de Sande & Millar 2019). Acetylene is a two-carbon alkene and the most abundant molecule in the inner region of the CSE where cosmic grains are formed, after H2 and CO (Fonfria et al. 2008), and is a known precursor of aromatic compounds (Frenklach & Feigelson 1989). For each experiment, chemistry was induced by the plasma discharge through fragmentation and ionization of the molecular precursor and plasma-generated chemical reactions between the precursor, its atomic and molecular fragments, radicals, and ions, resulting in the production of larger molecules and eventually solid particles. The resulting carbon grains, analogs of cosmic dust, were deposited on 3 mm diameter transmission electron microscopy (TEM) copper grids with Ultrathin Carbon Film on Lacey Carbon Support Film (Ted Pella, 400 mesh) placed 3 cm away from the slit, far enough away to avoid strong disturbances of the expansion. Ultra-high purity (UHP; 99.9998%) Ar and CH4 gas cylinders were used. The protocol for the addition of C2H2 to the gas mixture using a cold trap has been described in Sciamma-O’Brien et al. (2014). After deposition, the samples were collected under inert argon atmosphere using a custom-made glove box connected to the chamber (see Figure 1), and placed in plastic containers sealed with parafilm to minimize exposure to air before the ex situ analysis of the sample at the SEM facility located at the NASA Ames Advanced Studies Laboratories (Sciamma-O’Brien et al. 2017). The SEM technique was used to characterize the cosmic grain analogs produced in COSmIC. By using the same relative concentration for methane and acetylene (5%), only one experimental parameter (the nature of the precursor) was changed while maintaining all other experimental parameters (energy, temperature, precursor density, and reaction time) identical, thus allowing us to investigate the effect of the molecular precursors (CH4 versus C2H2) on the grain size and morphology under controlled conditions.
Citation: Frenklach & Feigelson 1989
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Radial velocity measurements (Anglada-Escudé et al. 2016) can only give us a minimum planetary mass, as long as no additional information is available to constrain the inclination of the star–planet system with respect to our observational plane. No transit of the planet has been observed so far, leading to the conclusion that the inclination is likely below 90°, and therefore the actual mass of the planet should be larger than 1.27 M⊕. Anglada et al. (2017) proposed the existence of a dust belt at an orbit of about 30 AU from the star, which, if confirmed, would suggest an inclination of the system of 45° with respectto the plane of the sky. Since confirmation has not been possible to date, we assumed four different inclinations spanning the wide range of likely inclinations and hence actual planet masses: 90, 60, 45, and 30°. The corresponding planet masses would then be 1.27, 1.47, 1.8, and 2.54 M⊕ (see also Table 2). The table lists in addition the predicted planet and core radii determined with our interior structure model (Noack et al. 2017) as well as the assumed initial temperature at the core–mantle boundary after planet accretion and solidification based on Stixrude (2014) and Noack & Lasbleis (2020). Both the iron content and the planet mass have a strong effect on the planetary structure and temperature profile. The resulting initial temperature profile in mantle and core, as well as the density, gravitational acceleration, pressure, and electrical conductivity profiles are shown in Fig. 1 for all investigated planet cases. The thermodynamic parameters (density as well as heat capacity and thermal expansivity of the material) were calculated following Noack et al. (2017) and using equations ofstate from Stixrude & Lithgow-Bertelloni (2011) for the mantle and Bouchet et al. (2013) for the metal core (here composed only of pure iron). The electrical conductivity profile is taken to be Earth-like from Xu et al. (2000). Thermal conductivity was calculated with the parametrisations derived in Tosi et al. (2013). For simplicity, the profiles for the different mantle properties were fixed during the evolution and did not take into account the influence of local melting or evolving chemical heterogeneity. The initial temperature profile, which is used in our interior evolution studies (described in Sect. 2.4), resembles the temperature profile depicted in Fig. 1. The temperature at the core–mantle boundary, however, and therefore the adiabatic temperature in the core, was scaled with the core–mantle boundary (CMB) temperature jump, which was varied in our study to investigate its effect on outgassing efficiency. Furthermore, we did not assume the existence of a primary atmosphere (see discussion in Sect. 4.1). During the evolution of Proxima Cen b, changes in atmospheric composition and mass can affect the surface temperature. However, since we did not couple our outgassing model to an atmospheric evolution model, surface temperatures remain constant throughout the evolution.
Citation: Stixrude (2014)
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We apply our analysis to five open clusters of ages between ∼10 and 100 Myr that were observed within GES (25 Ori, Gamma Vel, NGC 2547, NGC 2451 B, and NGC 2516). These clusters were selected because they cover the age interval in which the effect of radius inflation could be significant, allowing us to investigate how it evolves with age. The 25 Ori cluster is a group of PMS stars that was discovered by Briceño et al. (2005) in the Orion OB1a association, with an estimated age of 6−13 Myr (Downes et al. 2014; Briceño et al. 2019; Kos et al. 2019; Zari et al. 2019); a dispersed, kinematically distinct population was also found in the region using data from the Gaia Second Data Release (DR2; e.g. Zari et al. 2019). Because only a few stars of the secondary population were observed by GES, only the main cluster is considered here. Gamma Vel (age ∼10−20 Myr, Jeffries et al. 2014, 2017) and NGC 2547 (35 ± 3 Myr, Jeffries & Oliveira 2005) are both located in the Vela OB2 association at a relative separation of ∼2°. Both clusters host two kinematically distinct populations (Jeffries et al. 2014; Sacco et al. 2015); the two Gamma Vel populations (Gamma Vel A and B) are also separated by ∼38 pc along the line of sight (Franciosini et al. 2018). NGC 2451 is a double cluster composed of two open clusters of similar age (30−40 Myr, Randich et al. 2018) located at different distances along the same line of sight (Röser & Bastian 1994; Platais et al. 1996). The GES observations cover the background cluster NGC 2451 B and only a few selected regions of the closer and more dispersed NGC 2451 A. For this reason, we considered only NGC 2451 B here. Finally, NGC 2516 is the oldest cluster in our sample, with an age of ∼100 − 140 Myr (e.g. Lyra et al. 2006; Randich et al. 2018), so that most of its members are already close to or at their main-sequence position. All clusters have solar or slightly subsolar metallicities (Biazzo et al. 2011; Jacobson et al. 2016; Spina et al. 2017).
Citation: Jeffries et al. 2014
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Before the detection of GW170817, many associated EM counterparts had already been proposed for NS–NS/BH mergers, and their relative brightness is essentially determined by the properties of the post-merger central remnant object. In general, the merger remnant for an NS–BH merger would be a BH. But an NS–NS merger could lead to a BH or a supramassive NS, depending on the total mass of the NS–NS system and the NS equation of state (Dai et al. 2006; Gao & Fan 2006; Zhang 2013; Lasky et al. 2014; Gao et al. 2016). It is generally believed that, for both cases, sGRBs and their afterglow emission are expected as one of the major EM counterparts of NS–NS/BH mergers (Eichler et al. 1989; Narayan et al. 1992; Berger 2014). On the other hand, for both cases, the merger remnant would be surrounded by a mildly isotropic, sub-relativistic ejecta (which is composed of the tidally ripped and dynamically launched materials during the merger and the matter launched from the neutrino-driven wind from the accretion disk or neutron star surface; Rezzolla et al. 2011; Bauswein et al. 2013; Hotokezaka et al. 2013; Lei et al. 2013; Rosswog et al. 2013; Fernández et al. 2015; Song & Liu 2017). These ejecta are mostly composed of neutron-rich materials, and the radioactivity of these materials and the decay of r-process nuclei would heat the ejecta and then power an optical/IR transient (Li & Paczyński 1998; Metzger et al. 2010). When the merger remnant is a BH, r-process related radioactivity would be the only heating source, so that the luminosity of the optical/IR transient would be roughly ∼103 times that of the nova luminosity (Metzger et al. 2010). But if the merger remnant is a supramassive NS, its magnetic dipole radiation could serve as an additional heating source to the ejecta (Gao et al. 2013; Zhang 2013), which could easily exceed the r-process power, so that the thermal emission from the ejecta would be significantly enhanced (Yu et al. 2013; Metzger & Piro 2014). The luminosity of the optical transient, in this case, would be systematically brighter by more than one order of magnitude than the r-process dominating cases (Gao et al. 2017). Since the thermal emissions from the ejecta are essentially isotropic and also non-relativistic or mildly relativistic (due to the heavy mass loading), they therefore can be detected from any direction if the flux is high enough (see Metzger 2017a for a review).
Citation: Rezzolla et al. 2011
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The basic idea, summarized in Dekel et al. (2009), is that during VDI, the high surface density of gas and ‘cold’ young stars, Σ, drives the Toomre Q parameter below unity, Q ∼ σΩ/(πGΣ) ≲ 1, where σ is the 1D velocity dispersion and Ω is the angular frequency, a proxy to the epicyclic frequency κ, which is related to the potential well (Toomre 1964). It has been established that under such conditions, the disc will fragment and produce large star-forming clumps. This has been shown using idealized simulations of isolated galaxies (Noguchi 1999; Gammie 2001; Immeli et al. 2004a,b; Bournaud, Elmegreen & Elmegreen 2007; Elmegreen, Bournaud & Elmegreen 2008; Bournaud & Elmegreen 2009; Hopkins et al. 2012b), as well as cosmological simulations (Agertz, Teyssier & Moore 2009; Ceverino et al. 2010; Ceverino et al. 2012; Genel et al. 2012; Mandelker et al. 2014; Oklopcic et al. 2016). The ratio of clump mass to the mass of the cold disc scales as Mc/Md ∝ δ2, where δ = Md/Mtot is the ratio of the cold disc mass to the total mass within the disc radius, which includes the bulge and dark matter halo (e.g. Dekel et al. 2009). This leads to much larger clumps at z ∼ 2 than the low-redshift giant molecular clouds (GMCs). Gravitational interactions in the perturbed disc drive turbulence causing the disc to self-regulate in a marginally stable state with Q ≲ 1 (Dekel et al. 2009; Ceverino et al. 2010; Krumholz & Burkert 2010; Cacciato, Dekel & Genel 2012; Forbes, Krumholz & Burkert 2012; Forbes et al. 2014) that can last for more than a Gyr so long as the accretion is not interrupted. Some recent works have called into question the validity of linear Toomre analysis in the context of these highly non-linear galaxies (Behrendt, Burkert & Schartmann 2015; Tamburello et al. 2015; Inoue et al. 2016) and others have suggested alternate fragmentation mechanisms related to turbulence (e.g. Hopkins 2013). However, since clump formation is largely determined by the balance between self-gravity, turbulent pressure and the centrifugal force, the largest clumps are always roughly at the Toomre scale. Larger clumps would be disrupted due to the shear and/or tidal forces within the disc, or would not collapse in the first place due to the centrifugal force. Therefore, regardless of the full validity of linear Toomre analysis, it is plausible that the Toomre Q parameter can serve as a crude criterion for instability, possibly with a critical value that is larger than unity.
Citation: Dekel et al. (2009)
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Citation: Dekel et al. 2009
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The relationship established in our study between the occurrence rates of major flares and FHCMEs and the annual average latitude of sunspot groups may indicate the role of long-distance magnetic connectivity in the dependence of occurrence rates on the SC phase. In particular, the annual average latitude of sunspot groups determines the average length of trans-equatorial loops (TLs). A TL is a coronal loop connecting two separate ARs located in the opposite hemispheres (e.g., Tsuneta 1996; Pevtsov 2000). It has been found to be a common phenomenon, with up to one-third of all ARs exhibiting TLs in soft X-ray images (Pevtsov 2000). Therefore, our results may be a statistical indication that large-scale magnetic connectivities highlighted by TLs are important for producing solar activity such as major flares and FHCMEs. There have been many studies on the topology of the large-scale solar magnetic field structure (e.g., Babcock 1961; Hansen & Hansen 1975) and the relationship between the occurrence of solar flares / CMEs and large-scale magnetic connectivity (e.g., Antiochos et al. 1999; Khan & Hudson 2000; Wang & Sheeley 2003; Wang et al. 2007; Zhou et al. 2007; Fu & Welsch 2016). In particular, Dalla et al. (2007) established that the flaring rate (weakly) increases with the appearance of a new AR within 12° of the pre-existing flaring AR. Also, it has been reported that the interaction of ARs increases as sunspots move toward the equator (McIntosh & Leamon 2014; McIntosh et al. 2014). Such interactions may be related to the reconnection of magnetic fields that belong to different ARs, and the subsequent formation of large-scale structures such as TLs, which play an important role in triggering flares and CMEs (Khan & Hudson 2000). Moon et al. (2002) showed that sympathetic flares are more frequent in TLs than in loops connecting two ARs in the same hemisphere. Chen et al. (2006) examined the relation between the number of soft X-ray TLs and that of ARs each year from 1991 to 2001. They found good correlations of the TL numbers with SC indices, with the increase in TL number during the descending phase of SC 22 in comparison to other phases of SC 22. The effect of the dependence of flaring activity on AR separation and connectivity may be understood in terms of the interaction of magnetic energy between these ARs (Fu & Welsch 2016).
Citation: Wang et al. 2007
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From analytical and semi-analytical work (see for example Lovelace & Romanova 2014) we know that the local conditions where the RWI grows are what trigger the growth of one mode over another and in turn the strengths of each mode directly impacts the PDS. In the following section we focus on only one spin (a = 0.995) and position (rc = 2.7rg) where the RWI is triggered but look at different local conditions (see Eq. 3 for the parametrization), meaning the (unstable) zone where the RWI is active will be different. We can then see how the local condition affect the PDS, especially focusing on which modes are selected and how they relate to one another when present. We saw in previous studies of the RWI (Tagger & Varnière 2006; Vincent et al. 2013; Casse et al. 2017; Casse & Varniere 2018) that every simulation triggers a lot of modes, with up to m = 19 being detectable in Fourier space; though lower modes tend to be a lot stronger, especially once the instability reaches saturation. Because of the timescale near the last stable orbit of a black-hole of ten solar masses, we mostly observe the saturated state of the instability in a standard observation. We have therefore first performed several simulations with the same aforementioned corotation radius rc = 2.7rg but with various ε and δ parameters leading to different geometries of the unstable zone. A typical GRHD simulation example and the related ray-tracing synthetic observation can be found in Fig. 1. Once a simulation has reached the saturation stage, namely when the amplitude of the fluctuations remains approximately constant, we stop the simulation and proceed to use GYOTO in order to compute the light curve of the simulation as perceived by a remote observer. We stress here that the combination of the GR-AMRVAC code and the GYOTO code encompasses all general relativistic effects at work upon both the disc and its radiative emission, here providing reliable synthetic observations.
Citation: Tagger & Varnière 2006
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Substantial effort has been made by research groups in this field to improve the fibre link performance by employing non-circular fibres, double scramblers, fibre shakers, and fibre stretchers (Pepe et al. 2000; Chazelas et al. 2010; Perruchot et al. 2011; Avila 2012; Chazelas, Pepe & Wildi 2012; Bouchy et al. 2013; Reynolds & Kost 2014; Halverson et al. 2015a; Spronck et al. 2015; Ishizuka et al. 2016; Sutherland et al. 2016). Recent test results on the performance of MM optical fibres with non-circular cores show improved scrambling performance over that of standard circular MM fibres (Chazelas et al. 2012; Feger et al. 2012; Sturmer et al. 2016). However, a native non-circular fibre has also been shown to fail to provide the high level of stability and homogeneity in both the near-field and far-field of the fibre that is required for the next generation of PRV instruments operating below the 1 m s−1 radial velocity precision level. Much work has been done to improve the degree of stability and homogeneity that the fibre link can deliver by investigating and testing various combinations of fibre core geometries combined with a double scrambler. Optical double scramblers placed in the fibre link act to exchange the near and far-field light distributions of the fibre by means of a lens rely (Barnes & MacQueen 2010). This method has been shown to improve the level of scrambling in both the near and far-fields of the fibre output (Halverson et al. 2015a; Seifahrt, Sturmer & Bean 2016). Although these methods improve the fibre scrambling they do not necessarily reduce modal noise. Modal noise mitigation/suppression requires a time variable correction which is typically achieved by means of shaking and/or stretching the fibre (Reynolds & Kost 2014; Roy et al. 2014; Halverson et al. 2015a). In this paper, we present a new fibre technology (multicore fibre and photonic lanterns) that has the possibility to provide improved fibre scrambling and modal noise characteristics.
Citation: Feger et al. 2012
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The diffusion time-scale of the field lines in the disc is comparable to the viscous time-scale which is much longer than the interaction time-scale of the field lines and the inner disc, tint = |ΩK − Ω*|−1 (Fromang & Stone 2009), where Ω* is the angular velocity of the neutron star. The field lines cannot slip through the disc. Theoretical studies and numerical simulations show that the field lines interacting with the inner disc in a narrow boundary inflate and open up within the interaction time-scale (Aly 1985; Lovelace, Romanova & Bisnovatyi-Kogan 1995; Hayashi, Shibata & Matsumoto 1996; Miller & Stone 1997; Uzdensky, Königl & Litwin 2002; Uzdensky 2004). If the system is in the SP phase, the matter flowing in to the boundary at the innermost region of the disc is expelled from the system along the open field lines. The open lines reconnect on the dynamical time-scale, $\Omega _{\mathrm{K}}^{-1}$, and continue to apply torque on the matter until they open up again (Lovelace, Romanova & Bisnovatyi-Kogan 1999; Ustyugova et al. 2006). The simulations show that the field lines outside the interaction boundary are decoupled from the disc. The magnetosphere inside the boundary is the region where the closed field lines and the plasma can rotate together at radius rm ≃ rin. These results indicate that a strong-propeller mechanism could be sustained if the field lines expel the matter from the inner boundary at the same rate as that of the mass-flow from the outer disc. Furthermore, for this strong-propeller phase to be steady, the matter should be accelerated to speeds greater than the escape speed, υesc, within tint. Ertan (2017) showed that the maximum radius at which the strong-propeller condition is satisfied is much smaller than rA, while the accretion rates estimated for the WP/SP transition in this model seem to be in agreement with the transition properties of tMSPs (Ertan 2017, 2018). The inner boundary in the spin-up phase is also continuously evacuated, as in the case of strong-propeller, but now due to accretion on to the star. At rin the field lines decelerate and bring the matter into co-rotation within tint in the spin-up phase. In other words, the same equation could determine the inner disc radius in both the strong propeller phase (for rin > rco) and the spin-up phase (for rin rco). However, this is not the whole story, since the conditions in the spin-up (SU) phase are rather different, in particular rA enters the picture (Sections 2.1–2.4).
Citation: Uzdensky 2004
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We note that there has recently been some debate about potential observational effects (beam-smearing effect) on the dynamical state of (field) high-z galaxies (e.g., Di Teodoro et al. 2016; Kohandel et al. 2020). More specifically, when studying the gas kinematics with low spectral and spatial resolution data, the unresolved rotations within the PSF can artificially increase the value of σ0 and decrease the value of Vmax, leading to a severe systematic underestimation of the Vmax/σ0 ratio. The beam-smearing effect has been widely explored in local galaxies (e.g., see Fig. 6 in Di Teodoro & Fraternali 2015) and high-z simulated galaxies (e.g., Kohandel et al. 2020). Subsequently, given the high spatial and spectral resolution of ALMA observations, which have been able to reach subkpc spatial resolution in high-z galaxies, the beam-smearing effect is less significant. Compared to previous data, ALMA allows for a more robust way to measure intrinsic velocity dispersions. However, in our case, comparing two cluster SBs (with a spatial resolution of ∼0.3″ and a channel width of 30 km s−1) with other field SBs at similar redshifts and similar resolutions (with spatial resolutions of 0.1″–0.7″ and channel widths of 20–50 km s−1, except for one with a channel width of ∼100 km s−1), the σ0 values of the two cluster SBs are about three times lower than that of those field SBs (blue triangles in Fig. 6; Calistro Rivera et al. 2018; Barro et al. 2017; Swinbank et al. 2011; Tadaki et al. 2017). In addition, even at the same spatial and spectroscopical resolution, the σ0 values of the two cluster SBs are still more than three times lower than those of the two cluster MSs (green stars in Fig. 6). Furthermore, the simulations have confirmed the robustness of σ0 and Vmax for the two SBs and two MSs in our measurements (see Sect. 3.6). Therefore, we argue that the higher Vmax/σ0 and lower σ0 of the two cluster SBs with respect to field galaxies is a real physical difference, rather than being driven by observational effects.
Citation: Di Teodoro & Fraternali 2015
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