adsabs/FOCAL huggingface dataset:
For the NIR data set, we used standard IRAF11

11
IRAF is distributed by the National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.
 (Tody 1986, 1993) tools to correct the raw images for flat-field and to perform the sky subtraction. A master sky was obtained by combining five sky images of a relatively empty field. For the optical data set, instead, we used images processed, flat-fielded, bias subtracted, and corrected for Charge Transfer Efficiency losses by standard HST pipelines (_flc images). The WFC3 images have also been corrected by the Pixel Area Map using the files available on the HST website. For both NIR and optical data, the photometric reduction was carried out via point-spread function (PSF) fitting techniques in each chip of each image independently, by using DAOPHOTIV (Stetson 1987). The PSF has been modeled by selecting about 200 bright and isolated stars uniformly distributed in each chip, and by using the DAOPHOTIV/PSF routine. We allowed the PSF to vary within each chip following a cubic polynomial spatial variation. The best-fit PSF analytic models obtained for the HST F555W and F814W images, are a Moffat function with β = 1.5 (Moffat 1969) and a Penny function (Penny 1976), respectively. The typical PSF model for the GEMINI images is a Penny function, but in some cases it was necessary to adopt a Lorentz function (Melcher & Gerth 1977) as well. The PSF models thus obtained were then applied to all the star-like sources detected at a 3σ level above the local background by using ALLSTAR. We have derived in this way the stellar instrumental magnitudes. Then, starting from the star lists thus obtained and to fill the gaps among the GSAOI chips, we created a GEMINI master star list containing all the stars measured in at least three Ks images. For HST instead, the master list was made up of all the stars measured in at least four F814W observations. As done in previous works (e.g., Dalessandro et al. 2014 and references therein), the master lists thus created for the two data sets have been used as input for ALLFRAME (Stetson 1994). The files obtained as output have been independently combined to get a GEMINI catalog containing all the stars measured in at least three J and three Ks images and an HST catalog containing stars measured in at least four F555W and four F814W exposures. For every stellar source in each catalog, different magnitude estimates have been homogenized and their mean values and standard deviations have been adopted as the star magnitudes and photometric errors in the final catalog (Ferraro et al. 1991, 1992).
Citation: Moffat 1969
Citation function: Uses
We simulated a laboratory impact experiment into a basaltic rock performed by Nakamura & Fujiwara (1991) to test the implementation of the damage model. In their experiment, a small 0.2 g nylon bullet impacts with an impact speed of 3200 m s-1 into a basaltic rocky sphere with a diameter of 6 cm. We chose to simulate an off-axis impact with an impact angle of 30 ◦ (see Fig. 9 for the definition of the impact angle), which corresponds to an impact parameter of 25% of the diameter. We modeled the basalt target with 523 741 particles and applied the Tillotson equation of state. Owing to the lack of parameters for nylon, we used lucite for the projectile material (values were taken from Benz & Asphaug 1994). The basalt target was modeled with the additional damage model and activation threshold strains were distributed following the Weibull distribution statistically as a preprocessing step (see Sect. 3.2 for a detailed description of the algorithm). Following Benz & Asphaug (1994), we performed several simulations with varying Weibull parameter pairs k and m. Figure 7 shows the damage pattern in the target after 50 μs for the simulation that matched the experimental outcome best. The damage in the region around the impact point is the largest and all particles close to the impact point were immediately damaged upon impact. The damage pattern is radially symmetric around the impact point. Another source of damage is found near the surface of the target. There, damage grows due to spallation stresses that occur when compressive waves get reflected at the free surface. As a result of this nature, the inner core of the target remains intact while the surface is coated with cracks. After 50 μs simulation time, the damage pattern in the target remains unchanged. We identified single fragments using a friend-of-friend algorithm in the following way: first, we removed all fully damaged particles. Then, we clustered all remaining particles that are closer than 1.01 of the original smoothing length to their neighboring particles. Fragments consisting of less than three particles were neglected. Figure 8 shows the fragment mass spectrum obtained with this friend-of-friend clustering algorithm.
Citation: Benz & Asphaug 1994; Benz & Asphaug (1994)
Citation function: Uses; Uses; Uses
Interestingly, from ∼11:33:40 UT, the PAD of the suprathermal electrons changes to cigar type, in association with a dramatic drop of electron flux (Figure 2(a)). This change of electron PAD and flux is related to a magnetic dip structure that is manifested by a conspicuous decrease of the magnetic field strength (Figure 2(b)). In contrast to the magnetic field enhancement observed at ∼11:33:08 UT, the magnetic dip structure should arise from the local expansion of flux tubes, which are driven by the two opposite flows (see the shaded region in Figure 2(c)). The observed electron cigar distribution and the associated magnetic dip are a strong indication of the betatron cooling effect. This betatron-mediated cigar distribution has been suggested in previous studies (Fu et al. 2011b, 2012c, 2013b; Liu et al. 2017c) but never clearly observed. These observations, for the first time, show a direct link between the cigar distribution and betatron cooling. Note that in the trailing edge of the magnetic dip region, a magnetic hump structure, indicated by the sharp enhancement of magnetic field, is observed at ∼11:33:55 UT (Figure 2(b)). This magnetic hump, reminiscent of the Earthward-propagating dipolarization-front structure typically generated by magnetic reconnection in the midtail (Fu et al. 2012d, 2012e, 2013a, 2015, 2016, 2017; Liu et al. 2013, 2018a, 2018b, 2018c; Cao et al. 2017; Peng et al. 2017; Yao et al. 2017; Chen et al. 2018), was possibly formed due to the local contraction of flux tubes driven by the ion flow with an increasing velocity in the Earthward direction (Figure 2(c)). Associated with this magnetic hump, a very weak pancake distribution of suprathermal electrons is observed (Figure 2(a)), and strong waves near the electron gyrofrequency are also observed (Figures 2(e) and (f)). These waves are whistler-mode; they indicate that the betatron acceleration inside the magnetic hump was accompanied by non-adiabatic effects (Fu et al. 2009, 2010a, 2010b, 2012a, 2012b) and is still ongoing (Fu et al. 2011a; Khotyaintsev et al. 2011; Wang et al. 2017; Yang et al. 2017).
Citation: Fu et al.; 2013b
Citation function: Motivation
Figure 2 gives an overview of particle and solar wind parameters observed during a time span of 8 days before the electron event under consideration. The data seem to show no evidence for a CME and a related interplanetary shock. The solar wind speed, which had a low value of ∼300 km s−1, starts to rise gradually after October 14 and reaches a maximum of ∼650 km s−1 on October 19. At the same time the magnetic field strength and the proton density drop to low values. These variations are a typical signature of the presence of a CIR, i.e., a fast solar wind stream running into a slow solar wind stream and building up a compression region at a radial distance from the Sun of ∼2 au, or between ∼2.7 and 3 au along the field line connecting to the Sun (e.g., Gosling & Pizzo 1999, see Figure 3). The gradual increases in low-energy electron and proton fluxes (see the upper panel of Figure 2) could be caused by particle acceleration at the CIR (Giacalone et al. 2002). We found that on http://www.srl.caltech.edu/ACE/ASC/DATA/level3/ a pair of stream interaction regions (SIRs) is listed for 2002 October 14–15, and that the Harvard–Smithsonian Center for Astrophysics Interplanetary Shock Database (https://www.cfa.harvard.edu/shocks/) does not report any interplanetary shocks during 2002 October 12–20. We therefore conclude that the reflection of electrons in the 2002 October 20 event is likely due to—depending on their pitch angle—scattering or mirroring at a CIR-related compression, and not due to interaction with an interplanetary shock or propagation in a CME-related loop structure. Malandraki et al. (1997) described a technique for estimating the radial distance at which the magnetic compression arises due to stream–stream interaction by mapping the observed time variation of the solar wind speed back to the solar corona. Applying this method using the observed time profile of the solar wind speed (second panel of Figure 3), we estimate the reflectivity to be located at a radial distance of 1.9–2.1 au, corresponding to a distance along the field line connecting to the Sun of ∼2.4–2.7 au.
Citation: Gosling & Pizzo 1999
Citation function: Compare/Contrast
This paper presents a new 3D RT code, RAdiation SCattering in Astrophysical Simulations (RASCAS), which was designed to construct accurate multi-wavelength mock observations (spectra, images, or datacubes) from high-resolution simulations. RASCAS deploys a general two-step methodology (e.g. Hummels et al. 2017; Barrow et al. 2017). The first step consists in extracting all relevant information from the simulation outputs: (1) the information concerning the medium through which light will propagate, and (2) the information concerning the sources of radiation. Point (1), for example, determines the number density of H I atoms, their thermal velocity dispersion and bulk velocity, and the dust density, everywhere in a chosen volume. Fully coupled radiation-hydrodynamic simulations would naturally provide the information about ionised states, but retrieving this information in pure hydrodynamic simulations may be tricky. In such case, it is necessary to process the simulation outputs with additional software and models, typically with CLOUDY (Ferland et al. 2013), as in Hummels et al. (2017) and Barrow et al. (2017), or with independent codes which solve for H and He ionisation by propagating ionising radiation in post-processing (e.g. Li et al. 2008; Yajima et al. 2012). Point (2) is also involved to a greater or lesser extent, depending on the sources. Computing the continuum emission from star particles is relatively straightforward, using spectro-photometric models of stellar populations (e.g. Bruzual & Charlot 2003; Eldridge et al. 2017). However, computing the emission lines from the gas (e.g. in the Lyman-α line or in other nebular lines) again requires a detailed knowledge of the ionisation and thermal state of the emitting species. This may be provided by the simulation code, as is the case for H and He with RAMSES-RT (Rosdahl et al. 2013), or it may be necessary to post-process the simulation to estimate the emissivities of the gas. This first step is very simulation- and model-dependent, and RASCAS chooses to encapsulate it in a simulation-plugin module and to implement two stand-alone pre-processing codes which generate an adaptive mesh with all the needed physical information about the gaseous medium, and the initial conditions for light emission in the form of lists of photon packets. These datasets, which could easily be generated from other simulations with any post-processing code, serve as inputs to RASCAS to perform the radiative transfer computation.
Citation: Hummels et al. (2017)
Citation function: Uses; Uses; Uses
Since the manifold spirals arise in a system of reference which corotates with the bar, the manifold theory in its basic form predicts that the spiral arms should have the same pattern speed as the bar. This remark seems to come in conflict with observations both in our Galaxy (as reviewed e.g. in Bland-Hawthorn & Gerhard 2016; see also Antoja et al. 2014; Junqueira et al. 2015 and references therein) and in other galaxies (e.g. Vera-Villamizar et al. 2001; Boonyasait, Patsis & Gottesman 2005; Patsis, Kaufmann & Gottesman 2009; Meidt, Rand & Merrifield 2009; Speights & Westpfahl 2012; Speights & Rooke 2016). Considering again, galactic disc simulations, the leading paradigm over the years refers to simulations showing the coexistence of multiple pattern speeds (Sellwood & Sparke 1988; Little & Carlberg 1991; Rautiainen & Salo 1999; Quillen 2003; Minchev & Quillen 2006; Dubinski, Berentzen & Shlosman 2009; Quillen et al. 2011; Minchev et al. 2012; Baba, Saitoh & Wada 2013; Roca-Fabrega et al. 2013; Font et al. 2014; Baba 2015; but see also a noticeable exception in Roca-Fabrega et al. 2013), possibly connected also to the phenomenon of nonlinear coupling of multiple disc modes (Tagger et al. 1987; Tagger & Athanassoula 1991; Sellwood & Wilkinson 1993; Masset & Tagger 1997). On the other hand, it is well known that even isolated barred galaxies undergo substantial secular evolution (see Athanassoula 2013; Binney 2013; Kormendy 2013 in the tutorial volume Falcon-Barroso & Knapen 2013). The tendency to transfer angular momentum outwards (e.g. towards the halo or across the disc, Tremaine & Weinberg 1984; Debattista & Sellwood 1998; Debattista & Sellwood 2000; Athanassoula 2002; Athanassoula & Misiriotis 2002; Athanassoula 2003; OŃeill & Dubinski 2003; Holley-Bockelmann, Weinberg & Katz 2005; Berentzen, Shlosman & Jogee 2006; Martinez-Valpuesta, Shlosman & Heller 2006) leads the bar to slow down and grow in size at a rate which produces non-negligible change in dynamics at time-scales comparable even to a few bar periods. This process becomes complex, and even partially reversed due to the growth of ‘pseudo-bulges’ or peanuts (Kormendy & Kennicutt 2004), caused by dynamical instabilities such as chaos or the ‘buckling instability’ (Combes & Sanders 1981; Combes et al. 1990; Pfenniger & Friedli 1991; Raha et al. 1991; Bureau & Athanassoula 1999; Martinez-Valpuesta & Shlosman 2004; Bureau & Athanassoula 2005; Debattista et al. 2006). The reduction in size of the bar by the transfer of angular momentum under constant pattern speed is discussed in Weinberg & Katz (2007). Spiral activity acts as an additional factor of outwards transfer of angular momentum (Lynden-Bell & Kalnajs 1972), while a radial re-distribution of matter can take place even under a nearly preserved distribution of angular momentum (Hohl 1971; Sellwood & Binney 2002; Avila-Reese et al. 2005). Radial migration is enhanced by the amplification of chaos due to the overlapping of resonances among the various patterns (Quillen 2003; Minchev & Quillen 2006; Quillen et al. 2011).
Citation: Roca-Fabrega et al. 2013
Citation function: Compare/Contrast
The study of the variations in chromospheric flux using Ca-K spectral line at all time-scales is important to study the solar dynamo because of relation between the Ca-K line and the magnetic fields on the Sun (Leighton 1964; Skumanich, Smythe & Frazier 1975; Sivaraman & Livingston 1982; Ortiz & Rast 2005). Short time-scale variations provide the information about the dynamics and energy transport in the chromosphere whereas the long period variations provide valuable knowledge about the dynamo process. The observed Ca-K line irradiance variation with solar cycle is due to the contribution of various features such as plages, enhanced, active and quiet network to the flux. Also, the flux modulation occurs because of varying contribution with latitude due to limb darkening. Irradiance variations in Ca-K line have been studied by many researchers using long series of spectroheliograms obtained at Kodaikanal, Mount Wilson, Sac Peak and other observatories (Ermolli et al. 2009; Foukal et al. 2009; Tlatov, Pevtsov & Singh 2009; Priyal et al. 2013, and references therein) using plage areas as a proxy to the chromospheric irradiance variation. Worden, White & Woods (1998) analysed 1400 spectroheliograms obtained at the National Solar Observatory, Sac Peak spread between the period of 1980–1996 to identify the different features such as plages, enhanced, active and quiet network using the intensity contrast technique to study the rotational modulation and solar cycle variations in the chromosphere. Walton, Preminger & Chapman (2003) concluded that faculae are responsible for 80 per cent of the solar cycle variations whereas Ermolli, Berrilli & Florio (2003) estimated that during the ascending phase of the solar cycle (1996–2002) the network contributed 40–50 per cent to the solar cycle variations. Following this work Priyal et al. (2013) digitized the spectroheliograms obtained at Kodaikanal Observatory and analysed the data for three solar cycles (1955–1985) and found good correlation with the Mount Wilson data and sunspot number.
Citation: Ermolli et al. 2009
Citation function: Background
In this paper, we report on spectroscopic CH3CN, CH3OH (methanol), and dust continuum observations with the Atacama Large Millimeter/submillimeter Array (ALMA) at 349 GHz with an angular resolution of 0′′.1. We exploit the CH3CN (19K–18K) K-ladder, with excitation energies ranging from 168 K (for K = 0) to 881 K (for K = 10), to probe, at different radii, the physical conditions in the accretion disk of an early-type young star. We targeted the star-forming region G023.01−00.41, at a trigonometric distance of 4.59

$^{+0.38}_{-0.33}$





−0.33

+0.38



 kpc from the Sun (Brunthaler et al. 2009), where we recently revealed the accretion disk around a young star of 104.6 L⊙, corresponding to a ZAMS star of 20 M⊙ (Sanna et al. 2019, their Fig. 1); the disk was imaged by means of spectroscopic ALMA observations of both CH3CN and CH3OH lines at 0′′.2 resolution inthe 230 GHz band. The disk extends up to radii of 3000 au from the central star where it warps above the midplane; here, we resolve the outer disk regions in two apparent spirals projected onto the plane of the sky. We showed that molecular gas is falling in and slowly rotating with sub-Keplerian velocities down to radii of 500 au from the central star, where we measured a mass infall rate of 6 × 10−4 M⊙ yr−1 (Sanna et al. 2019, their Fig. 5). The disk and star system drives a radio continuum jet and a molecular outflow aligned along a position angle of 57°, measured east of north (Sanna et al. 2016, their Fig. 2); their projected axis is oriented perpendicular to the disk midplane whose inclination with respect to the line-of-sight was estimated to be less than 30° (namely, the disk is seen approximately edge-on; Sanna et al. 2014, 2019). Previously, we also measured the average gas conditions over the same extent of the whole disk, by means of Submillimeter Array (SMA) observations of the CH3CN (12K–11K) emission, and we estimated a kinetic temperature of 195 K and CH3CN column density of 5.1 × 1016 cm−2 (Sanna et al. 2014, their Fig. 2 and Table 4).
Citation: Sanna et al. 2019
Citation function: Extends; Extends; Extends
White dwarfs (WDs) are the end state of stars with mass ≲8M⊙. A WD attains its stable equilibrium configuration by balancing the outward force due to the degenerate electron gas with the inward force of gravity. If the WD has a binary companion, it pulls out matter from the companion, and as a result, the mass of WD increases. Once the mass of the WD reaches the Chandrasekhar mass limit (Chandrasekhar 1931; ∼1.4M⊙ for carbon–oxygen nonrotating, nonmagnetized WDs), the pressure balance no longer sustains, and the WD bursts out to produce a Type Ia supernova (SN Ia; Choudhuri 2010). The similarity in peak luminosities of SNe Ia is used as one of the standard candles to estimate the luminosity distances for various astronomical and cosmological objects (Lieb & Yau 1987; Nomoto et al. 1997). However, recent discoveries of various under- and overluminous SNe Ia question the complete validity of considering luminosities of SNe Ia as standard candles. SNe Ia such as SN 1991bg (Filippenko et al. 1992; Mazzali et al. 1997), SN 1997cn (Turatto et al. 1998), SN 1998de (Modjaz et al. 2001), SN 1999by (Garnavich et al. 2004), and SN 2005bl (Taubenberger et al. 2008) were discovered with extremely low luminosities, which were produced from WDs with 56Ni mass content as low as ∼0.1M⊙ (Stritzinger et al. 2006). On the other hand, a different class of SNe Ia, such as SN 2003fg (Howell et al. 2006), SN 2006gz (Hicken et al. 2007), SN 2009dc (Yamanaka et al. 2009; Tanaka et al. 2010; Silverman et al. 2011; Taubenberger et al. 2011; Kamiya et al. 2012), SN 2007if (Scalzo et al. 2010; Yuan et al. 2010; Scalzo et al. 2012), SN 2013cv (Cao et al. 2016), and many more was discovered with an excess luminosity, with the observed mass of 56Ni as high as ∼1.8M⊙ (Kamiya et al. 2012), violating the Khokhlov pure detonation limit (Khokhlov et al. 1993). It was inferred that these underluminous SNe Ia were produced from WDs with a mass ∼0.6M⊙ (Mazzali et al. 1997; Turatto et al. 1998), while the same for overluminous SNe Ia could be ∼2.8M⊙ (Scalzo et al. 2010; Kamiya et al. 2012). Hence, these progenitor WDs of peculiar SNe Ia violate the Chandrasekhar mass limit: the underluminous SNe Ia were produced from sub-Chandrasekhar limiting-mass WDs (WDs burst before reaching the mass ∼1.4M⊙), and the overluminous SNe Ia were produced from super-Chandrasekhar limiting-mass WDs (WDs burst well above the mass ∼1.4M⊙). These new mass limits are important, as they may lead to modifying the standard candle.
Citation: Garnavich et al. 2004
Citation function: Background; Background; Background
Blaauw (1961) hypothesized that runaway stars are created when a massive component within a binary system explodes as a Type II supernova. The secondary is then ejected with a velocity comparable to the orbital velocity at the time of the supernova event. The explosion may not disrupt the system (Hills 1983), which will be observed as a OB-neutron star/black-hole system and eventually become a high-mass X-ray binary. Given that D2–EB is a double-lined O-type binary, it is unlikely that such a mechanism took place. Alternatively, D2–EB could have been part of a triple system including a very massive component (≥85 M⊙) with a lifetime of less than 4.5 Myr. However, the inferred radial velocity curve lacks evidence of a low-mass third companion, unless it orbits with a high eccentricity and/or long period. Observations of D2–EB over a longer time span than the data presented here are required to investigate the possibility of a third body in the system. Poveda et al. (1967) suggested that runaway stars are dynamically ejected due to encounters of collapsing protostars in/near the core of young clusters. Interactions in clusters that contain initial hard binaries increase the number of escapees via binary-binary interactions (Mikkola 1983). The less massive binary system reaches peculiar space velocities up to ~200 km s-1, while the more massive system travels at a speed less than ~100 km s-1 (Leonard & Duncan 1988). The binary frequency for runaways with V∞> 30 km s-1 is predicted to be 10% and it is striking that systems hosting components of M ~ 20 M⊙ as the most massive members are not predicted to escape with more than 50 km s-1 (Leonard & Duncan 1990). This does not conflict with the observed lower limit for the space velocity of D2–EB. A young, dense cluster core increases the possibility of dynamical ejection, with the consequence that the kinematic age of the runaways is similar to the age of the cluster (Gualandris et al. 2004). In this case, the unreasonably low estimated value for the tangential velocity of D2–EB makes the possibility of the dynamical ejection through binary-binary interaction unlikely.
Citation: Poveda et al. (1967)
Citation function: Background