adsabs/WIESP2022-NER huggingface dataset:
Paper: 2017ApJ...834...44M (fulltext)
calculated using Equation ( 30 ) in the Appendix , which are given bywhere represents a horizontal average. Figure 8. Measurements of horizontally averaged density gradients in layers for two low-Ta* simulations. Also shown are predicted density gradients for systems with given rotation rates and layer heights. Both simulations have Pr = τ = 0.1 and . We indeed find that the density gradients in layers are largest for the simulation with the higher rotation rate (Ta* = 0.1), and that the results from this simulation also fit well with prediction. The case with Ta* = 0.01, on the other hand, shows much greater variability in density gradient, particularly at larger layer heights, making it difficult to determine whether or not the theoretical predictions are valid. From Figure 8 we also see that in both cases the steepening of the density gradient is somewhat greater for smaller layers. However, this dependence of density gradient on layer height is weaker than the theory suggests. In Figure 4 we see that the simulation with Ta* = 0.1 spends a greater amount of time in the three- and two-layered phases than either of the other two simulations, and it consequently takes about twice as long as the nonrotating run (Ta* = 0) to reach the one-layered phase. The root cause is related to the lower supercriticality of convection within the layers combined with the increased stability of the interfaces. Through inspection of the density profiles of our layered simulations, we see that the positions and shapes of the interfaces in rotating simulations have less variability with time compared with nonrotating simulations. We also see in Figure 4 that the fluxes in the rotating layered systems oscillate less than they do in the nonrotating ones. Wood et al. ( 2013 ) already discussed these large-amplitude oscillations in nonrotating simulations and attributed them to the presence of large plumes of fluid punching through interfaces, periodically causing spikes in transport. In our layered rotating simulations (particularly the Ta* = 0.1 case), the absence of large-amplitude oscillations in the layered phase suggests that the motion of these plumes is inhibited, possibly because convection is weaker and the interfacial gradients are stronger. Since the plumes could be key players in the layer merger events, their suppression in the rotating simulations could also explain why the merger timescale is longer. All of these effects also contribute to the reduction of mean fluxes of temperature and chemical composition through density staircases in rotating ODDC compared with the nonrotating case. To show this quantitatively, Table 1 presents measurements of and (with standard deviations) for each value of Ta*. For low Ta* simulations, which clearly form convective layers, the fluxes are measured in the one-layered phase. The simulation with Ta* = 0.01 shows 2% and 7% decreases in thermal and compositional fluxes, respectively, compared to the nonrotating
Paper: 2018MNRAS.476.2117R (fulltext)
giants, as can be seen in Fig. A1 (d). These stars are all flagged with suspect broad lines by the ASPCAP pipeline. 5.4 Carbon rich stars Carbon rich stars (discussed in Section 3.4 ) are also detected as outliers. In Fig. A1 (e), we can see the strong CN compared to OH lines for a few carbon enriched stars. The weirdness score increases with the strength of the CN features. 5.5 Young stellar objects Stars in this group show both H-Br emission lines as well as regular metallic absorption lines. They are mostly young stars included in the INfrared Survey of Young Nebulous Clusters (IN-SYNC, Cottaar et al. 2014 ). We detect stars with both broad and narrow emission, and with absorption that can be broad or narrow as well as double lined (SB2s). SIMBAD classification for stars in this group include ‘Variable star of Orion type’, ‘T Tau-type Star’, ‘Pre-main sequence Star’, and ‘Young stellar object’. 5.6 M dwarfs M dwarf stars are also detected as outliers. This is due to the small number of M dwarf stars in the APOGEE sample. Example spectra are in Fig. A1(g). 5.7 Other outliers Some of the objects detected as outliers did not fall into any of the above classes. These include a brown dwarf, a Wolf–Rayet star, a few AGB stars including an OH-IR star, and known variable stars, as well as three red supergiants observed in the massive stars ancillary program. Also detected as outliers are special non-stellar targets, such as the centre of M32, a few M31 globular clusters, and three planetary nebulae. Two outliers show double peaked H-Br emission lines, as well as absorption lines typical to the APOGEE data set. Both of these objects show RV modulations, suggesting they are multiple star systems. For the first, 2M04052624+5304494, the RV modulation of the absorption lines (determined from APOGEE visit spectra) could be modelled with a period of P = 11.152 ± 0.072 d, and amplitude of K = 78 ± 15 km s −1 . The emission lines in the APOGEE spectra show smaller RV modulation, if any. For the second, 2M06415063−0130177, the absorption lines RV changes by ∼160 km s −1 between two APOGEE visits. The visits are separated by 28 d. For the emission lines we could not get a good estimate on the RVs, as the emission line profiles change significantly between the visits. A CoRoT light-curve is available for this system, showing clear periodic modulation. A period of P = 29.04 d was derived for this light-curve by Affer et al. ( 2012 ). We note that this period does not agree with the RV modulation. For both of these systems, additional work is required to determine their nature. We also detect a group of objects with similar, very broad features. Most of these objects have SIMBAD classifications as contact binaries, mainly W Ursae Majoris. A few objects remain unexplained. We divide these objects into two groups. In the first group, we have objects with spectra that seems to have similar features to typical APOGEE red giants (by means of visual inspection). The
Paper: 2017MNRAS.472.2608S (fulltext)
C iii ] upper limit). Three similar systems from Berg et al. ( 2016 ) with COS detections of all four lines are plotted as blue squares. The ratios of O iii ]/He ii are inconsistent with any AGN models – all have more flux in O iii ] 1666 than in He ii 1640. The deep ESI spectra we obtained also provide constraints on the shape of the ionizing spectrum between 10 and 50 eV. The ratio of flux between the nebular He ii and H recombination lines is closely related to the ratio of He + -ionizing (228 Å) to H-ionizing photons (912 Å). The optical He ii λ4686 and H β lines are ideal for avoiding the uncertainties introduced by dust and aperture corrections. Our ESI spectra easily resolve the nebular and broad components of the He ii line (see Fig. 5 ) and thus allow a robust simultaneous fit to both, to which H β in the same spectrum can be compared (Section 3.4 and Table 10 ). The (dust-corrected) ratio of nebular He ii /H β is plotted as a function of gas-phase metallicity in Fig. 12 . In addition to the sample presented here, we plot the published measurements of this ratio for the local WR galaxies analysed by López-Sánchez Esteban ( 2010 ). The axis on the right displays the ratio Q(He + )/Q(H) assuming ideal Case B recombination proceeds for both H and He + (and 10 4 K, 10 2 cm −3 ; Hummer Storey 1987 ; Draine 2011 ). The Case B predictions should be treated with care as collisional excitation and low ionization parameters can result in deviations from this assumption (e.g. Stasińska Izotov 2001 ; Raiter, Schaerer Fosbury 2010 ), but demonstrate clearly the effect of changing the ionizing spectrum slope at the He + ionizing edge. Figure 12. Measured nebular He ii 4686/H β values (a proxy for the ratio of He + -ionizing to H-ionizing photons) versus gas-phase metallicity. Both our galaxies (blue circles, backed by red squares if nebular C iv was detected) and the WR galaxies presented by López-Sánchez Esteban ( 2010 ) are presented. The right axis indicates the corresponding Q(He + )/Q(H) assuming Case B recombination for both species. Both observational samples show a sharp transition with metallicity, with substantially larger ratios and harder spectra below 12 + log O/H ∼ 8. The nebular He ii /H β ratio reveals a sharp transition in the ionizing spectrum with decreasing metallicity. The maximum value of He ii /H β increases as the gas-phase metallicity decreases, with a marked upturn in this envelope around 12 + log O/H ∼ 8.0–8.2. Systems below 12 + log O/H 8.2 (Z ⊙ /3) reach relative nebular He ii fluxes nearly 10 times higher than those above this metallicity. This implies an order of magnitude increase in Q(He + )/Q(HI) as the metallicity of a young stellar population is decreased from Z ⊙ to Z ⊙ /10. A similar trend towards higher He ii /H β below 12 + log O/H 8.0 was found in the full sample of SDSS spectra analysed by Brinchmann et al. ( 2008b ) and Shirazi Brinchmann ( 2012 ), though we find a far more pronounced transition. This may
Paper: 2016ApJ...820..125T (fulltext)
). These structures are lying mainly horizontally with respect to z = 0. This behavior is due to the fact that twist introduces local dips and mass tends to accumulate around these dips. However, at the center of the flux rope, even in the case with the strongest twist, the magnetic field lines are essentially horizontal, and therefore the dense material does not reorganize much along these lines, producing the horizontal structuring. Figure 16. Same as in Figure 14 but for N t = 4.3. Figure 17. Same as in Figure 14 but for N t = 8.5. 8. CONCLUSIONS AND DISCUSSION The time evolution of a cold prominence initially embedded in a magnetic flux rope has revealed interesting characteristics of the structure. We have concentrated on sustained prominences, i.e., detached from the photosphere. The different numerical simulations have shown how the magnetic field evolves in response to the dense prominence that is pulled downward by gravity. The various twisted flux ropes considered in this work, together with the embedded prominences, show a different evolution. We have shown that in some cases they are clearly evolving toward a stationary state. These quasi-stationary 3D solutions can be used in the future for other purposes such as the analysis of winking filaments produced by external perturbations like Moreton waves, EIT waves, or nearby flares and CMEs. During the relaxation to the quasi-stationary equilibrium, the system goes through several periodic vertical oscillations. In the linear regime, i.e., when the amplitude of oscillation is small in comparison with the local Alfvén speed, the process of continuum damping transfers the excess of energy in the system to the PCTR. It is in this highly inhomogeneous layer where the process of phase mixing takes place. This energy redistribution explains the attenuation of the position of the CM of the prominence body. Thus, at least in the linear regime the resonant absorption process is crucial in the relaxation of the structure. The energy in the PCTR is eventually dissipated by the numerical scheme used to perform the simulation, and this dissipation can be hardly avoided owing to the continuous generation of small spatial scales with time by the phase-mixing mechanism. Fortunately, although the numerical dissipation produces an energy loss, this does not necessarily mean that the attenuation of the CM is dominated by this artificial dissipation. In fact, it is known that a small dissipation does not change the damping time (Poedts Kerner 1991 ) of the quasi-mode (i.e., the resonantly damped global mode); it affects the dissipation in the layer but not the global attenuation (see Terradas et al. 2006a , for the linear time evolution with dissipation). The different convergence tests performed in the present study point in this direction. Wave leakage is not relevant in our model and does not contribute to the attenuation of the position of the CM. The problem studied here demonstrates the necessity
Paper: 2020A&A...633A..98W (fulltext)
of periapsis ω , and time of periastron t peri ). These can be used to calculate the projected position of both components on the sky at any moment in time. In all previous cases of binary events with significant binary motion, Keplerian orbital motion provided only a small improvement relative to the linear approximation ( Skowron et al. 2011 ; Shin et al. 2012 ). This is not the case here, because, as we show below, the orbital period of the lens is similar to the duration of the event (e.g. Penny et al. 2011 ). Modelling of this event is an iterative process: for given microlensing parameters, we estimated the angular radius and distance to the source, we calculated best-fit microlensing parameters, and we repeated the procedure until all parameters converged. The best-fit microlensing parameters are presented in Table 5 . Uncertainties were calculated using the Markov chain Monte Carlo approach (MCMC; Foreman-Mackey et al. 2013 ) and represent 68% confidence intervals of marginalized posterior distributions. We note that another degenerate solution exists for the microlensing model that differs only by the signs of s z and γ z (( s z , γ z ) → − ( s z , γ z )). The second solution has the same physical parameters (except for Ω → π − Ω and ω → ω − π ) and differs by the sign of the radial velocity. Thus, the degeneracy can be broken with additional radial velocity measurements of the lens ( Skowron et al. 2011 ). Table 5. Best-fit microlensing model parameters of the Gaia16aye binary event. Notes. HJD′=HJD-2450000. We adopt t 0, par = t 0, kep = 7675. 4.3. Source star Spectroscopic observations of the event indicate that the source is a K5-type giant or a super-giant. If the effective temperature of the source were higher than 4250 K, TiO absorption features would be invisible. If the temperature were lower than 3800 K, these features would be stronger than those in the observed spectra. Spectral modelling indicates that the effective temperature of the source is 3933 ± 135 K. According to Houdashelt et al. (2000) , the intrinsic Johnson–Cousins colours of a star of this spectral type and solar metallicity should be , and (error bars correspond to the source of K4- and M0-type, respectively). We used a model-independent regression to calculate the observed colours of the source (we used observations collected in the Bialkow Observatory, which were calibrated to the standard system): V − R = 0.99 ± 0.01 and V − I = 1.91 ± 0.01. Thus, the colour excess is E ( V − I ) = 0.31 and E ( V − R ) = 0.16, consistent with the standard reddening law ( Cardelli et al. 1989 ) and A V = 0.62. According to the best-fitting microlensing model, the amount of light coming from the magnified source is V s = 16.61 ± 0.02 and I s = 14.70 ± 0.02. The V -band brightness of the source after correcting for extinction is therefore V 0 = 15.99 mag. Subsequently, we used the colour–surface brightness relations for giants from Adams et al. (2018) to estimate the
Paper: 2016A&A...588A.112G (fulltext)
scale height (Klahr Henning 1997; Pinilla et al. 2012; Testi et al. 2014; Johansen et al. 2014) . The regions in the disk midplane that correspond to the snowlines of major volatiles are particularly interesting. These may promote locally efficient grain growth through recondensation across the snowline or by changing the sticking properties of ice-coated grains and, in addition, the local release of volatiles from the ices may induce a local pressure bump that could trap large grains (e.g., Supulver Lin 2000; Wada et al. 2009; Ros Johansen 2013; Gundlach Blum 2015) . The dense regions of disk midplanes can be investigated at submillimeter and millimeter wavelengths, where the dust emission is more optically thin and can now be spatially resolved with present-day facilities, such as the Atacama Large Millimeter/submillimeter Array (ALMA) and the Karl G. Jansky Very Large Array (VLA). These facilities allow us to obtain sensitive and high angular resolution observations of the dust continuum and gas emission lines from protoplanetary disks and therefore study the radial variations of the physical parameters of the disk. Recent results include the study of the dust properties as a function of radius in disks (e.g., Guilloteau et al. 2011; Banzatti et al. 2011; Pérez et al. 2012, 2015; Menu et al. 2014; Tazzari et al. 2016) , and the clear identification of the CO snowline in nearby disks (Mathews et al. 2013; Qi et al. 2013, 2015) , through its effects on DCO and NH abundances. It is thus becoming possible to directly investigate the effect of the CO snowline on grain growth observationally. This snowline is thought to be one of the most important after that of water, because the CO abundance in protostellar ices is found to be about 3040% of HO (Öberg et al. 2011) . In this work we focus on the Herbig Ae star HD 163296, a bright and isolated object at a distance of parsec (van den Ancker et al. 1998) with a relatively massive disk (~0.1 M⊙: Qi et al. 2011; Isella et al. 2007) and an excellent prototype for gas- and dust-rich protoplanetary disks. An estimated age of 5 Myr was obtained from the comparison between H ipparcos astrometric measurements and pre-main sequence evolutionary models by van den Ancker et al. (1998) . The stellar parameters computed by Natta et al. (2004) are , , K. To study the dust properties in the protoplanetary disk around HD 163296, we re-analyse the ALMA science verification observations (de Gregorio-Monsalvo et al. 2013; Mathews et al. 2013; Rosenfeld et al. 2013) and for the first time discuss the continuum emission in Band 6, combined with new VLA observations from the Disks@EVLA collaboration. In Sect. 2 we describe the observational data, in Sect. 3 we report the main new results of our analysis, in Sect. 4 we present the result of our disk modeling, and in Sect. 5 we discuss the main implications for grain properties. Fig. 1 Continuum maps at different wavelengths, from the left : 850 m, 1.3 mm, and 9 mm, the
Paper: 2021A&A...649A..78D (fulltext)
inferred from the shell modeling reported in Sect. 3.1 and Fig. 11 must be considered as strict lower limits. This is because the emerging spectrum is strongly weighted by sightlines with the lowest column densities, given that Ly α escapes through the paths of least resistance (e.g., Eide et al. 2018 ), generally orthogonal to the stream velocities in case of infall ( Gronke et al. 2017 ). Nevertheless, we note that an alternative scenario suggested by recent modeling is that radiation is emitted through the cooling of mixed gas occurring at the boundary between the phases ( Mandelker et al. 2019 ; Cornuault et al. 2018 ). In that case, Ly α radiation would originate from the combined dissipation of the kinetic energy of the stream and the thermal energy from the mixed gas that cools down ( Gronke Oh 2020a ). This scenario, and our calculations of gas mass flows in Sect. 3.1 , also imply a total cold gas mass present at each moment in the diffuse streams of the order of 10 11 M ⊙ (to be compared to the lower limit of ≈10 8 M ⊙ of neutral hydrogen). Is that a reasonable gas mass to be diffused into the intra-group gas? On the one hand, this is comparable to or smaller than the ISM mass of all group galaxies combined, and less than 10% of all baryons expected to be in the group given its DM estimate and assuming a universal baryon fraction. Assuming pressure equilibrium, the flowing gas would need to be confined in clouds with quite a small volume filling factor, of order 10 −3 to 10 −5 , which is qualitatively similar to what is discussed in theoretical works (e.g., Cornuault et al. 2018 ; McCourt et al. 2018 ; Gronke Oh 2020b ). This is required to reach a high-enough density to compensate for the much larger temperature of the hot medium. Such dense gas clouds would produce copious recombination Ly α photons following photoionization, if a source emitting enough ionizing photons towards this medium were to be present. This does not appear to be the case for RO-1001, at least given the evidence presented here. 5. Summary and conclusions The main findings of this work can be summarized as follows: RO-1001 is a group at z = 2.91 defined by four galaxies with stellar masses above 10 11 M ⊙ , originally selected as an overdensity of faint radio sources (12 σ excess) over a small, 10′′ region. The group contains an ongoing SFR of 1200 M ⊙ yr −1 mostly spread among three ALMA-detected, IR-luminous sources characterized by very compact IR sizes, for which we reported CO[3−2] detections with NOEMA. RO-1001 does not contain any detectable evidence of ongoing AGN activity, down to limits that are consistent with a cosmic average co-existence of an AGN and SFR at typical levels. The hosting DM halo mass is estimated to be typical for a group, with ∼4 × 10 13 M ⊙ consistently derived using three methods, including a blind X-ray measurement with significance of 3.1 σ . RO-1001 hosts a giant Ly α halo with a luminosity of 1.3 × 10 44 erg s −1 as
Paper: 2017MNRAS.470...34R (fulltext)
on the low branch. Mass is being fed into the annulus at a greater rate than it is able to process it and so mass accumulates and Σ increases. Eventually, the annulus reaches the low-temperature bifurcation point and leaps to the high hot state. The annulus, now above the Σ-nullcline, begins losing mass because its accretion rate exceeds the supply rate. Once the high-temperature bifurcation point is reached, the annulus drops to the low-temperature state and the process is repeated. The red curve in Fig. 4 (a) corresponds to a numerical simulation of equations ( 15 ) and ( 17 ) and illustrates this basic behaviour quite clearly. In Fig. 4 (b), we plot the evolution of Σ and T c as a function of time. Here, one can observe the two time-scales at work, especially in the (blue) T c curve. In between outbursts there is little change in temperature, but a build up in mass (seen in the red curve) occurring on the slow viscous time-scale t Σ . Interleaving these phases are the abrupt bursts occurring on the short thermal time-scale t th . 3.2.2 Z Cam type simulations Z Cam variables intermittently behave in a similar manner to U Gem variables but occasionally remain on the high-temperature branch in a quasi-steady state for months to years before dropping to quiescence and undergoing a fresh batch of outbursts. This is attributed to a shift in the Σ-nullcline upwards (due to a change in mass transfer, δ ≠ 0), so that the intersection is beyond the saddle node. Of course, the system must be close to criticality if this is to be possible. To demonstrate this in our model, and δ are chosen such that the fixed point oscillates between the stable high-temperature branch and the unstable branch. When the fixed point is on the upper branch the system is stable and will remain in a standstill. A decrease in pushes the fixed point on to the unstable branch and the system undergoes periodic orbits resulting in trains of outbursts. In Fig. 5 , we show an example of standstill–outburst-train cycles. Figure 5. Time evolution of the surface density and central temperature with and b 0 = 0 with mass transfer variation. The variation amplitude is δ = 0.1 and t I = 5 × 10 6 s. The system clearly switches between outburst trains and standstills. 3.3 Results: stochastic With the basics established, we include a fluctuating viscosity (β ≠ 0). The fluctuation time-scale is taken to be 10 t orb , which is motivated by MRI shearing box simulations (Hawley et al. 1995 ). The amplitude is set at b 0 = 0.1. First, we consider the case where there is no mass transfer variability (δ = 0). To obtain a fixed point on the high-temperature branch close to the bifurcation point we choose . A deterministic system with this configuration would remain in the high-temperature state indefinitely. A stochastic system however will see the fixed point jostle about as the two nullclines fluctuate with α, sometimes sitting on the upper branch, sometimes dropping to the middle branch