adsabs/WIESP2022-NER huggingface dataset:
Paper: 2021MNRAS.503.5100M (fulltext)
is inconsistent with a single Rayleigh distribution of inclinations ( i ). Its vertical distribution was better reproduced with two dynamical populations with typical inclination dispersions ( i rms ) of 1° and 9°. Such a scenario is analogous to the classical Kuiper belt, which also has a bimodal inclination distribution (e.g. Brown 2001 ). Similarly, ALMA observations of AU Mic marginally resolved its vertical extent (Daley et al. 2019 ) and constrained as well the distribution of inclinations to i rms ∼ 3°. For moderately inclined and narrow discs, the vertical extent can be constrained by comparing the surface brightness as a function of azimuth if the resolution is comparable or smaller than the disc width (Marino et al. 2016 ). This is because the non-zero vertical thickness of a disc widens the surface brightness distribution near the minor axis, hence reducing slightly its surface brightness compared to a flat disc. This principle has helped to marginally constrain the vertical thickness of HR 4796’s debris ring (Kennedy et al. 2018 ), finding i rms ∼ 3°. Similarly, it is also possible to determine the height of a wide disc if it contains narrow features such as a gap, as done by Marino et al. ( 2019 ) for HD 92945’s disc, finding i rms ∼ 4°. Note that the latter case is very similar to modelling done in protoplanetary discs with gaps to measure the degree of dust settling (e.g. Pinte et al. 2016 ). The width of debris discs provides as well a direct upper limit on the eccentricity levels. This is because for a uniform distribution of longitudes of pericentres, the difference between pericentre and apocentre distance must be equal or smaller than the disc width. This is particularly interesting when applied to narrow debris discs interacting with an eccentric planet since the final width of the disc strongly depends on the forced eccentricity induced by the planet (e.g. Pearce Wyatt 2014 ). Kennedy ( 2020 ) used these arguments to show that the narrowness of the debris disc around Fomalhaut (MacGregor et al. 2017 ) and HD 202628 (Faramaz et al. 2019 ) is inconsistent with the standard secular evolution theory. The dispersion of eccentricities in those discs, constrained by the disc widths, is smaller than expected under the presence of an eccentric planet. This inconsistency could be due to planetesimals being born on eccentric and aligned orbits or their relative velocities being damped reducing the proper eccentricities (e.g. through collisions; Nesvold et al. 2013 ). The multiwavelength analysis of disc sizes can also constrain stirring levels. By combining far-infrared (FIR) and mm spatial information with collisional evolution models, Geiler et al. ( 2019 ) found that the FIR emission of HR8799’s debris disc cannot be reproduced by a dynamically cold wide belt that best matches the mm data. They concluded that to fit both the large extension of the FIR emission and high mm flux, their model required the presence of two
Paper: 2016ApJ...833...90X (fulltext)
between W (CH) and A v than those between W (OH), W (CO) and A v . To better examine these relationships, we adopt the usual definition of “ X -factor”,where N H2 is the column density of H 2 and W is the integrated intensity of the molecular tracers CH, OH, CO. We plot the correlation of “ X -factor” of CH, OH and CO versus the visual extinction in Figure 8 . When the visual extinction A v is in the range between 0.8 and 2.1 mag, the X CH -factor is almost a constant (1.0 ± 0.06) × 10 22 cm −2 K −1 km −1 s. The dispersion of CH X -factor is visibly less than those of OH and CO at 0.8 ≤ A v ≤ 2.1 mag. CH appears to be a better tracer of molecular gas than CO and OH in the transition zone (0.8 ≤ A v ≤ 2.1 mag). Where extinction drops below 1 mag in TBR, the integrated intensities W of the three molecules lie above the overall fitting lines, indicating a stronger intensity of these spectra, which is likely the result of C-shock in the TBR-O (Section 5 ). Figure 8. Correlation plot of the integrated intensity W of CH, OH and CO vs. the visual extinction A v , and the correlation of the “ X -factor” of CH, OH and CO vs. the visual extinction A v . “ r ” means Pearson correlation coefficient. 5. CH ABUNDANCE AND C-SHOCK MODEL We have calculated the column density of CH in Section 3.2 . The column density of CH across the TBR is shown in Figure 7 . Mattila ( 1986 ) found a good correlation between the CH column density and visual extinction for diffuse, dark molecular cloud N (CH) = 3.4 × 10 13 ( A v − 03) cm −2 as the dashed line shows in Figure 7 . This correlation is reasonable considering the CH formation threshold, which requires substantial H 2 (Black Dalgarno 1973 ). The derived column density of CH below 1 mag extinction (blue shade in Figure 7 ) in TBR obviously lies above the typical value observed by Mattila ( 1986 ). Furthermore, considering the correction factor for excitation temperature and optical depth (as discussed in Section 3.2 ), the actual column density of CH in TBR-O is likely to be even larger by an additional factor of 3. The true CH column density in TBR-O can thus be one order of magnitude higher than the simple extension from the X -factor for higher extinctions. We compared the CH column density with the prediction of C-shock models. Flower Pineau des Forets ( 1998 ) made a prediction that CH column density would increase from 10 13 to 10 14 cm −2 when extinction increases from 0.1 to 1 mag at a C-shock front, which matches exactly the measured values here. Xu et al. ( 2016 ) also found an overabundance of OH below 1 mag extinction (by a factor of 80). The overabundance of CH and OH suggests that there may be an additional channel of CH and OH production, possibly due to the shock (e.g., Draine Katz 1986a ) produced by the colliding streams (Xu et al. 2016 , Figure 4). When shock waves propagate through the molecular ISM, the gas is compressed, heated, and accelerated. CH can be produced during CH + synthesis
Paper: 2015ApJ...810...60H (acknowledgments)
P.G. and K.H. acknowledge NSF grants AST-1010039 and AST-1412648. This work was also supported by NASA grant HST -GO-12055. K.H. was supported by a NSF Graduate Research Fellowship. D.R.W. is supported by NASA through Hubble Fellowship grant HST-HF-51331.01 awarded by the Space Telescope Science Institute. We thank Bernhard Aringer, Zachary Jennings, and Alex Rudy for helpful conversations. We appreciate the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community. We are most grateful to have had the opportunity to conduct observations from this mountain.
Paper: 2017MNRAS.472.4689R (fulltext)
et al. ( 1996 ) and Giri Chakrabarti ( 2013 ): the cooling rate is ∝ρ 2 T β , where β > 0 is the cooling index. For β = 0.5, the process is bremsstrahlung. In order to ensure the formation of a cold Keplerian disc, we must use higher β in order for the cooling time-scale to be reduced compared with the infall time-scale. Thus we increase cooling gradually inside the would-be Keplerian component. When the region with Keplerian distribution becomes optically thick, the index is raised to 0.6 as used in earlier simulations. This saves the prohibitive time required for Monte Carlo simulations, in order that normal bremsstrahlung radiation may be emitted as blackbody radiation in the optically thick region. Giri et al. ( 2015 ) found that the result is insensitive to the exact cooling process, since blackbody cooling also produced the expected standard disc in the equatorial plane in their simulation. In Fig. 2 (a) and (b), we show the density, velocity and specific angular momentum distribution at t = 19.8 s just before viscosity is turned on. As can be seen clearly, a standing shock at a distance of about 40 r g is formed, with a sudden rise of density and pressure. The colour (online) scale is on the right. The outflow starts in the region between the centrifugal barrier and the funnel wall (Molteni et al. 1996 ). Specific angular momentum remains at its initial value of 1.7. Figure 2. (a) Density, velocity and (b) specific angular momentum distribution at t = 19.8 s just before viscosity is turned on ( t 0 in Fig. 1 a). Here, a standing shock has formed at a distance of about 40 r g , with a sudden rise of density and pressure, and outflow has started between the centrifugal barrier and funnel wall region. Density ρ is chosen to be unity in dimensionless units at the injection radius. In (b), the specific angular momentum distribution is represented by the solid curve for comparison with the Keplerian distribution (dot–dashed curve). At t 0 , the specific angular momentum remains at its initial injected value of 1.7. After viscosity is turned on, as matter accretes, initially the viscous stress near the inner edge remains close to zero. As α increases with time, −α P becomes significant close to the inner edge, due to high pressure, and angular momentum transport starts efficiently. On the equatorial plane, this result is similar to what was seen in Chakrabarti Molteni ( 1995 ). However, at the outer edge, though α itself is high, pressure being very low, the stress is not high enough and thus angular momentum transport remains slower. 4.1 Results using viscosity description of Case I In this case, we use . We chose α 0 = 0.12, which is above the critical value of 0.074 required for a Keplerian disc to form with the used injected parameters. In Fig. 3 (a)–(b), we show the distributions at t = 41.58 s of (a) the density and velocity vector and (b) the angular momentum distribution (solid). To compare, we also plot the Keplerian distribution
Paper: 2016A&A...587A.105A (fulltext)
on a relatively short timescale. The surface velocity is divided by a factor of 10 in 200 Myr, until the star spins down enough to enter the unsaturated-dynamo regime and further decreases by another factor of 10 for the following 4.3 Gyr. For clusters older than the Hyades (beyond 700 Myr), the observed distributions of rotation periods strongly narrow. By the time the models reach 1 Gyr, the torque has become much less efficient and the angular velocity profile of the fast rotator is similar to that of the median and slow rotators. The three models later follow the same evolution, which is dictated by the meridional circulation and shear turbulence until the age of the Sun. The surface angular velocity and the total angular momentum content of the slow, medium, and fast rotators are similar beyond 1 Gyr. The angular velocity profiles also share a similar shape and evolution in all three cases. Fig. 4 Angular momentum fluxes carried by meridional circulation (dashed red), shear (dotted blue), and the total flux (solid black) in the fast and slowly rotating models ( top and bottom , respectively). Hatched areas in all plots indicate the convective regions. 4.3. Transport of AM Here we examine the processes that drive the internal transport of AM in our reference models. To do so, we followed the diagnostic procedure proposed by Decressin et al. (2009) , which relies on the analysis of the AM fluxes associated with each transport process that is accounted for. The transport of angular momentum is mainly driven by the meridional circulation as long as the star is rotating relatively fast (above the saturation value), as shown in the left panels of Fig. 4 , and this independently of the initial velocity. After the ZAMS (60 Myr), in all cases meridional circulation weakens because the star does not contract any longer and the surface velocity has already decreased. Angular velocity gradient increases at the base of the convective envelope, leading to a predominance of the shear turbulence over the meridional circulation in this region. For the slower rotators, the shear turbulence takes over the transport of AM in the outer radiative zone below the envelope and efficiently extracts (acting along the angular velocity gradient) the angular momentum, ensuring a continuous decrease of the surface angular velocity from the ZAMS until the age of the Hyades (700 Myr; see lower panel in Fig. 4 ). On the other hand, in the fast rotator, the meridional circulation is very efficient as long as the star contracts (2 orders of magnitude larger than in the slow rotator) and efficiently couples the interior to the convective envelope up to the ZAMS. Beyond this age, its efficiency decreases rapidly, and the shear generated by the torque dominates below the convective envelope, although it is not sufficient to maintain an efficient transfer of AM from the core to the envelope. The overall transport of AM by these processes tends to become less efficient as the
Paper: 2020A&A...641A.172E (fulltext)
that BSF galaxies both in the cocoon and in the supercluster have significantly lower stellar masses than galaxies from other populations. This is what could be expected from galaxies which are still forming their stellar content and increasing their stellar mass. Small differences in stellar masses of galaxies from other populations are not statistically significant and we can conclude that galaxies from these populations in our sample have statistically similar stellar masses, in a range of M * ≈ 3 × 10 9 h −1 M ⊙ − 3 × 10 11 h −1 M ⊙ . Figure 10 (upper left panel) shows that BSF galaxies have the highest values of the concentration index, C , and VO galaxies have the lowest C values among our galaxy populations, and they are similar in the supercluster and in the cocoon. In this respect, RQ galaxies differ from others. In the supercluster they have lower values of C than in the cocoon, these differences are statistically significant at a very high level, with a p -value of p ≤ 0.01 (Table 4 ). The distribution of probabilities for galaxies to be of late-type, P late , are plotted in Fig. 10 (upper right panel). This probability is the highest for BSF galaxies and for RQ galaxies in the cocoon, and the lowest for VO galaxies (in both environments). However, KS test shows that these differences are not statistically significant (Table 4 ). BSF galaxies also have the lowest values of stellar velocity dispersions and VO galaxies have the highest (Fig. 10 , lower left panel). This agrees with earlier results which have shown that galaxies with young stellar populations have lower stellar velocity dispersions than galaxies with old stellar populations ( Shankar et al. 2009 ; Bezanson et al. 2012 ; Einasto et al. 2018b ). The stellar velocity dispersions of RQ galaxies are similar in the supercluster and in the cocoon, indicating, in accordance with the results for P late , that the overall morphological type of RQ and RSF galaxies is the same in both environments. As the stellar mass distribution for VO, RQ, and RSF samples are similar, this suggests that VO galaxies are mainly non-rotating early-type galaxies, but RQ and RSF galaxies (and BSF galaxies) may have also a significant rotation component corresponding to the dynamics of disky galaxies. The metallicities Z of BSF, RQ, and RSF galaxies (Fig. 10 , lower right panel) are statistically similar in both environments, BSF and RQ galaxies having the lowest metallicity values. VO galaxies in the supercluster have higher metallicities than galaxies from other populations in our study. In the cocoon, they are similar to the metallicities for galaxies from other populations. We additionally checked whether the galaxy content of long filaments with length ≥20 h −1 Mpc depends on the distance of filament member galaxies from the supercluster centre, D c . The values of D n (4000) index versus D c for galaxies in long filaments are plotted in Fig. 11 . In this figure, background colours show the
Paper: 2016A&A...585A..87S (fulltext)
properties across luminosity and redshift indicates very similar underlying physical mechanisms that can be understood in terms of Eigenvector 1 ( Boroson Green 1992 ; Francis et al. 1992 ) where the Eddington ratio drives the relative strength of the lines, and orientation effects influences the observed kinematics of the lines ( Shen Ho 2014 ). The specific local physical conditions that regulate the accretion rate then account for the majority of the inter-quasar variation. In this section, we will investigate whether the targets selected are simply scaled-up version of average quasars, or remarkable in some sense. 6.1. Comparison to global population The QSOs selected for this study lie amongst the brightest values of compared to the spectroscopically confirmed SDSS QSOs ( Shen et al. 2011 ), which are uniformly selected to 90 percent completeness ( Richards et al. 2002a ; Vanden Berk et al. 2005 ). It is clear that the sample presented here represents some of the brightest quasars existing between redshift 1 and 2. It is therefore potentially a biased subset of the global quasar population. Quasars have previously been shown to exhibit a high degree of homogeneity over many orders of magnitude ( Dietrich et al. 2002 ), but this is only true for quasars resembling average quasars and not exotic sub-types. To ensure that the targets selected for this sample are not outliers in color, the spectrophotometric colors are compared to those of the quasar population presented in Pâris et al. (2014) 3 . We show the comparison in Fig. . From the magnitudes, we first confirm that the quasars selected here are among the intrinsically brightest visible in -band with a 2.7 sigma magnitude distance from the mean absolute -band magnitude, normalized at redshift 2, of the full SDSS quasar catalog for DR10. Despite their luminosities the colors are representative of the subset fulfilling the selection criteria. As a consistency check, the composite generated in this work is used to fit for the dust content of quasars with varying degree of extinction and consistent results were obtained as compared with composites generated from other subsamples of the parent population. For a comparison between SDSS and quasars detected in GALEX, UKIDSS, WISE, 2MASS and Spitzer , see Krawczyk et al. (2013) . Fig. 6 Quasar color, , as a function of redshift, . The full quasar sample from SDSS DR10 ( Pâris et al. 2014 ) is shown in gray color. Bounding contours show the 0.5, 1, 1.5, 2 contours of the density of points. The olive points are individual quasars from the SDSS sample which satisfy . Overplotted in pink are the quasars contributing to the composite presented here. The quasar colors have been marginalized over in the right side where the tall pink lines are at the position of the composite quasars and the olive shaded area is the kernel density estimation of the distribution of quasars fulfilling the selection criteria. The gray-shaded area is the projection of the
Paper: 2020ApJ...889L..38K (fulltext)
September 20, this was an 1833 (spatial) by 15 (slit width) box, 2019 October 1 had 15575 by 10, 2019 October 9 had 260 by 15, and 2019 October 10 had 18′′ by 24. The imaging data were extracted using circular aperture photometry with a radius of 250 and largely followed the reduction procedure outlined in Farnham et al. ( 2000 ) using the dust reflectance slope measured from our 2019 October 10 spectrum. The standard star used (HD 81809) is noted in that work to be less than ideal for blue/UV observations, so we have overestimated errors when possible. We note that for the 2019 September 20 and October 1 data sets, insufficient data were collected to allow a median combination of spectra, making the two data sets “look” even noisier than their later counterparts even though cosmic rays were cleaned from them. 3. Analysis, Modeling, and Results In order to search for molecular emission features, it is necessary to produce a flux-calibrated solar spectrum fit to subtract from the flux-calibrated data. This process creates a dust reflectance spectrum for the comet, which can be used to look for a spectral slope. The solar spectrum was taken from Chance Kurucz ( 2010 ) and binned to the resolution of each data set. The original flux-calibrated spectra, best-fit reflectance, and dust-subtracted comet spectra are shown in Figure 1 . The noisiness of the Bok data prevents a high-quality solar reflectance subtraction, while the later data become very linear after subtraction indicating a relatively good fit of the solar spectrum to the continuum points. Figure 1. Spectral observations of 2I/Borisov. The flux-calibrated and dust-subtracted spectra of all three nights of spectroscopy from 2019 October. The heliocentric and geocentric distances are listed underneath the telescope names and dates on the figure. The same background and background-subtraction process was applied to all data sets. The CN B–X (0–0) band centered at 388.3 nm can be seen clearly in the MMT and LBT data. (The data used to create this figure are available .) We consider the spectral slope between the blue, green, and red continuum regions defined for the Hale–Bopp filter set (Farnham et al. 2000 ) to compare reflectance slopes with other workers. These filters are centered at 4450, 5260, and 7128 Å, respectively. If we normalize our highest-quality data from 2019 October 10 at the green continuum point, then we find the blue-to-green slope to be S ′ = 22+/−4%/micron, and the green-to-red slope to be S ′ = 11+/−3%/micron. Fitzsimmons et al. ( 2019 ) suggested that the dust might be “redder” at shorter wavelengths, and our best data support that conclusion. After the reflected sunlight from dust was subtracted, we searched for emission from molecules in the coma of 2I/Borisov (Table 2 ). For each possible feature we subtracted a linear estimate of the background as determined by a fit to data on either side of the feature and then integrated using a trapezoidal method all flux
Paper: 2019MNRAS.482.2374A (acknowledgments)
ACKNOWLEDGEMENTS We acknowledge financial support by the Deutsche Forschungsgemeinschaft within project Ur 32/27-2. Simulations were performed at the High Performance Cluster Elwetritsch (Regionales Hochschulrechenzentrum, TU Kaiserslautern, Germany).