adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019MNRAS.486.4655B (fulltext)
it was about 3 per cent) to about 4.5 per cent at the time of our observations. In this spectral region, we have not detected any change between 2015 and 2018, but data obtained in 1995 were in between the values measured in the two epochs. At 6800 Å, data obtained by Putney ( 1995 ) show the lowest values among all data sets (∼1 per cent). Our BBP measurements in the red (4.0 per cent integrated over a R filter) are consistent with the relatively recent BBP measurement of Butters et al. ( 2009 ), who had detected a signal of 4.06 ± 0.16 per cent. In conclusion, we have clearly detected a change of 1.5 per cent over a 40+ yr interval, but found no evidence of change over a time-scale of the order of three years. 8.3 Linear polarization No large changes of the fraction of linear polarization are detected either on a short (3-yr) or long (40-yr) time-scale. However, over a 40–45 yr interval of time, the position angle has definitely changed by more than 10° both in the blue and the red. In the following, we look at the various linear polarization spectra and their differences in more detail. On the long time-scale sampled by the difference between the 1970s data and our recent ISIS spectra (40 yr), there definitely are changes. In particular, below 4000 Å, the recent level of linear polarization sampled by Q / I seems to have decreased, while that sampled by U / I is larger. This corresponds to a very significant rotation of the position angle of the linear polarization away from the value of ∼20° found in the 1970s to about 33° in recent times. We note that the measurement by West ( 1989 ) obtained 1986 suggests also a decrease of linear polarization compared to the 1970s (3.19 per cent in the B filter, for a position angle of 19.1 ± 0.8) but as discussed in Section 4.3 , these kinds of comparison should be made with caution. It is notable that although the circular polarization is quite different from zero throughout the visible spectrum, the non-zero linear polarization is present essentially below 5500 Å, with possible low-level features suggested by the Steward spectrum, but largely absent from the lower resolution MCSP data. The position angle in the red as measured in the 1970s was around 110°, with a fair amount of scatter due to the small but not quite vanishing amplitude of Q and U in that region. In the ISIS data, both Q and U are also close to zero, therefore the position angle has a large uncertainty, but is consistent with a 10°–30° rotation from the value measured in the 1970s. We conclude that linear polarization measurements are certainly variable over a time-scale of decades, mostly in terms of a rotation, while no significant variation has been detected on time-scale of three years. A rotation constant in wavelength provides a first rough description of the change in linear polarization. 9 CONSTRAINTS ON THE ROTATION PERIOD OF Grw+70° 8247 Measurements of the position angle of the broad-band linear polarization of Grw+70°
Paper: 2021ApJ...912...64Z (fulltext)
also do not show another source within the NACO FoV, indicating that both BHB2007-2AB and VLT 171104.07-272258.8 are more evolved than BHB2007-1. In the analysis we also calculated the photometry in the K S band of BHB2007-2AB and VLT 171104.07-272258.8, listed in Table 2 . Table 1 Astrometric Measurements of the Components Measured with Respect to the Central Source [BHB2007]-1 Table 2 Apparent Magnitudes of the Components Measured with Respect to the Central Source [BHB2007]-1 3.3. Stellar Properties Here we calculate the stellar properties of [BHB2007]-1 and -2 based on the Gaia EDR3 parallax (see Section 3.2 ) and the notion that [BHB2007]-2 is a binary system. We collected the B-to-W4-band photometry of −1 and −2 from Vizier. We investigated the ASAS-SN light curves (Kochanek et al. 2017 ) finding that the visual brightness of both stars is only moderately variable ( mag over years). We adopted a Phoenix model of the stellar photosphere (Hauschildt et al. 1999 ) with the effective temperature constrained by Covey et al. ( 2010 ) ( K and K, for −1 and −2, respectively, with typical uncertainties of ±300 K) and visual extinction calculated from the observed colors between the R , I , and J bands, and stellar luminosity L * from the integrated photospheric model. For [BHB2007]-1, this yielded mag and . However, the edge-on geometry of the disk (see Section 3.4 ) may in principle yield an underestimation of the luminosity. In fact, stars with very inclined disks appear on average older (since less luminous) than the other stars (see Garufi et al. 2018 ). As for [BHB2007]-2, we first had to consider the visually triple nature of the star (see Section 3.1 ). All three stars sit within the 2MASS beam, meaning that the observed photometry is a contribution of all of them. We computed the actual J -, H -, -, and W 1-band magnitudes through the relative flux ratios measured in our NACO images. We then constrained visual extinction and stellar luminosity of [BHB2007]-2A assuming that the constrained by Covey et al. ( 2010 ) refers to this star since it is significantly brighter than the other two. We found that the is comparable with that of [BHB2007]-1 (1.5 ± 0.4 mag) while the luminosity is . Finally, we constrained stellar mass M * and age t of the two sources through multiple sets of pre-main-sequence tracks (Baraffe, MIST, and Parsec; Baraffe et al. 2015 ; Dotter 2016 ; Bressan et al. 2012 ). Bearing in mind the assumptions and the uncertainties described above, we found and Myr for [BHB2007]-1, and and for [BHB2007]-2. The stellar mass of [BHB2007]-1 is significantly smaller than what was constrained by Alves et al. ( 2020 ) from the position–velocity diagram of the CO emission (2.23 ). Such a large discrepancy cannot be explained by an erroneous estimate of the stellar luminosity alone since the evolutionary tracks corresponding to stars with K are nearly vertical, making the stellar mass only marginally impacted by the luminosity.
Paper: 2016MNRAS.462.3710C (fulltext)
of 1.5. Convection sets in close to the point where the cold branch of the S-curve ends which leaves this point sensitive to α ml . This effect is even stronger on the critical point where the hot branch starts, as this is where convection is the strongest. Higher values of α ml shift both critical points at which the S-curve bends closer together, leading to a smaller Q Σ crit . This alters the global behaviour of the disc and the outburst properties: the decay from the outburst in a disc with $α ml $ = 6 starts at higher Σ and higher T eff , and less mass is accreted and accumulated in the disc during the outburst cycle. Therefore, more efficient convection produces outbursts that are more frequent and of lower amplitude, and even lack quiescent phases. While it is important to note that this lack of quiescence may be related to uncertainties in the end of the lower branch, there are other ways out. In the discussion below, we examine a few options to restore or modify quiescence. The comparison of DIMRI with $α ml $ = 1.5 and DIMa (Figs 10 and 9 c) highlights the importance of changing the dissipation profile, as this is the only difference between these two models. The main difference is that the outbursts in DIMRI with $α ml $ = 1.5 are wider and the quiescence is longer than in the DIMa light curve. Changing the mixing length parameter is not the only way to increase the outburst amplitude and restore quiescent phases to the DIMRI models. Increasing the mass transfer rate, while all other parameters are fixed, makes the disc hotter and denser. This means that the surface density everywhere in the quiescent disc is closer to the critical value and the disc luminosity is higher. The result is more frequent lower amplitude outbursts for higher mass transfer rate regardless of model, as illustrated by comparing the two mass transfer rates shown in Fig. 9 for each of our three models. Hence, merely reducing the mass transfer rate in the DIMRI model can increase the outburst amplitude and restore quiescent phases. If one sets as low as g s −1 in the DIMRI (with α ml = 6), the elapsed time between two consecutive outbursts starts to be longer and also the outburst amplitude rises (see Table 1 and Fig. 11 ). Note that for such a low , the discs in DIM and DIMa become cold and stable. There is yet another alternative way to restore quiescence in the DIMRI light curves. Observed X-ray fluxes in quiescent dwarf novae are far too large compared to what models predict. A solution to this problem is a truncation of the inner disc. Numerical calculations by Hameury, Lasota Warner ( 2000 ) and Schreiber, Hameury Lasota ( 2003 ) confirm that truncating the inner disc has substantial influence on the dwarf nova light curves and may solve the discrepancy between observations and theory. This truncation may be caused by the magnetic field of the primary white dwarf. A white dwarf magnetic field in the range 10 4 –10 7 G (which translates to a magnetic
Paper: 2019MNRAS.484.5009E (fulltext)
as well, e.g. CALIFA, MaNGA, and SAMI (Croom et al. 2012 ); but the ∼10 times higher spatial resolution of the MAD data boosts the ability to disentangle the different components, in addition to providing information on scales that are still largely unexplored in terms of a systematic study as a function of the location on the SFMS. 4.1 Resolved BPT diagrams We use the well-known [O iii ]/H β versus [N ii ]/H α Baldwin, Phillips Terlevich (hereafter BPT; Baldwin, Phillips Terlevich 1981 ) diagram to identify regions within the galaxies that are ionized by UV radiation from newly born stars (regions of star formation – hereafter, the SF gas) and regions that are ionized by AGN-emitted radiation, shocks, or AGB stars. In these diagrams (e.g. Fig. 1 ), the upper, solid line is the photoionization model of Kewley et al. ( 2001 ), which defines the upper envelope of the region of parameter space occupied by extreme starbursts; this line is a conservative boundary between SF gas and AGN/shocks. The lower, dashed line shows an equivalent but empirical demarcation derived by Kauffmann et al. ( 2003 ). The area of BPT parameter space that lies between these two lines is regarded as identifying regions of gaseous emission with an intermediate ionization spectrum. Figure 1. BPT diagram of all the spaxels from all the galaxies in the sample. The colour scheme has been selected as a continuous variation from purple/blue (SF) to red (AGN/shock dominated), using the η parameter. To explain how η is defined, we compute the minimum distances ( d 1 , d 2 , and d 3 ) of three general points (P 1 , P 2 , and P 3 , respectively, each one in a different region of the BPT diagram) to the two lines that define the three BPT regions. The spaxels with EW(H α)6 Å lie inside the dash–dot area (see the text). It is also well known that the EW is an extra proxy to constrain the possible ionization source of the ionized gas, restricting the SF regions to those with EW(H α) > 6 Å (e.g. Cid Fernandes et al. 2011 ; Sánchez et al. 2014 ; Barrera-Ballesteros et al. 2016 ; Belfiore et al. 2017 ). The EW(H α) is measured by dividing the H α intensity by the flux density of the underlying continuum of the emission line-free spectra, computed as the mean flux density of two windows of 30 Å wide and centred at ±60 Å from the observed H α wavelength. The spaxels with EW(H α) 6 Å are presented in Fig. 1 as those inside the dash–dotted area. The ratio between all the spaxels with EW(H α) 6 Å and all the spaxels under the empirical SF demarcation of the BPT is 4 per cent. Although this ratio is lower than 2 per cent for 24 out of 38 galaxies, it is significant for some galaxies such as NGC 3521 (23 per cent), IC5273 (11 per cent), or NGC 4941 (8 per cent). To obtain a single-valued continuous parameter that conveys the location of each spaxel on its galaxy’s BPT diagram, we define the variable η as either ( i ) the inward-pointing orthogonal distance from the dashed line, identifying SF
Paper: 2018MNRAS.477..325B (fulltext)
‘truncation’ of the disc could have shrunk the dust disc to its tiny radius. On the other hand, the SED is also compatible with some amount of large dust outside of 0.6 au. For this scenario, radial ‘migration’ coupled with fragmentation interior to the H 2 O snowline (located at 0.6 au) result in a cliff in Σ dust of two orders of magnitude and in a sudden increase of grain sizes. Our modelling shows that this scenario matches the observed SED for outer dust disc radii of R dust, 2 ≲ 45 au. In both scenarios, a dust mass inside of 0.6 au bigger than ∼1 M ⊕ is required to match the observed SED, however due to the optically thick nature of the dust, we do not obtain an upper limit on the dust mass. In the ‘migration’ scenario, the dust mass needs to be fine-tuned to ∼1 M ⊕ inside the H 2 O water snowline in order to match the SED as the dust masses inside and outside of 0.6 au are dependent on each other. In order to distinguish further between these scenarios, the gas radius of the disc can offer important information. If it is truncated at very small radii as well, a dynamical truncation of the disc seems likely. If it is more extended, radial drift in combination with dust physics at the water snowline offer a plausible explanation. We present predictions for CO J = 2–1 fluxes, invoking various M gas – R gas combinations, that will enable future observations to potentially distinguish between these two scenarios. Finally, we note that regardless of the underlying mechanism for the small-scale dust disc, its optically thick nature suggests that several M ⊕ of dust can be located within a small radial extent from the central star. Such a configuration suggests that V410 X-ray 1 may be a potential precursor to a system of tightly packed terrestrial planets, such as the recently discovered TRAPPIST-1 system. APPENDIX A: DUST PROCESSING AT THE WATER SNOWLINE In Section 4.2 , we study whether the truncated dust disc could be related to physical processes happening at the water snowline. The key property that changes across the snowline is frag , the maximum collision velocity of grains for which coagulation, rather than fragmentation, is to be expected. Following Blum Wurm ( 2000 ) and Gundlach Blum ( 2015 ), frag is reduced by an order of magnitude interior to the snowline. If grain growth is limited by fragmentation, the local value of frag is linked to the maximum expected grain size through equating the relative collision velocity Δ for the largest grains at any location with frag . Δ for grains that are coupled to a turbulent flow via drag forces is given by (Voelk et al. 1980 ; Ormel Cuzzi 2007 ): where α is the Shakura–Sunyaev viscosity parameter (Shakura Sunyaev 1973 ), c s the sound speed, and St is the local Stokes number (ratio of drag time to dynamical time) for the largest grains. In the Epstein drag regime, the midplane St for grains of size a max is given bywhere we have assumed a Gaussian vertical structure. When growth is
Paper: 2016MNRAS.457..822B (fulltext)
line profiles. Calculations of lines formed in disc winds from white-dwarf stars have been limited to resonance lines tend to produce absorption profiles for inclinations i ≳ 70° (Proga et al. 2002 ). However, disc-wind models for T Tauri stars can produce similar H α line profiles to those seen in RS Oph (Kurosawa, Harries Symington 2006 ). Polarization measurements may be able to distinguish between inflow and outflow models, since, in the inflow models, the fraction of scattered light contributing to the blue wing is larger. 4 SODIUM ABSORPTION LINES 4.1 Na <sc>i</sc> <sc>D lines in RS Oph</sc> Sodium absorption lines have been suggested as a useful diagnostic of the circumstellar environment in both supernovae and their progenitor systems (Patat et al. 2007 ; Chugai 2008 ; Patat et al. 2011 ). The Na i D lines, together with other strong metal lines such as the Ca ii H K lines, probe the density and ionization conditions in the circumstellar environment. We modelled the line profiles expected from the simulations, with the aim of understanding where the observed features in RS Oph form in the CSM and whether the same circumstellar components can be responsible for metal line variations observed in some SN Ia. We have modelled Na i D absorption line profiles directly from the simulations, using the results from the ionization calculation described above. All neutral sodium atoms are assumed to be in the ground state and contribute to the absorption lines. We present line profiles from calculations in which the mass-loss rate has been scaled to 5 × 10 −7 M ⊙ yr −1 , since this gives better agreement with observations both for the ionization structure and the line profiles. We did not include a detailed treatment for the H ii region in the Na i D line ionization calculation; instead, we neglected the contribution from all material within 3 × 10 13 cm. The results are not particularly sensitive to this choice of radius. We assume that the continuum flux around the Na i D lines is dominated by the RG. The presence of an emission component in RS Oph limits the extent to which this is true (Patat et al. 2011 ). The emission component has a width v ∼ 100 km s −1 , higher than the width of the absorption lines, and is likely associated with emission from the ionized region. The line profiles were derived by calculating the optical depth to the photosphere along lines of sight through the CSM. The photosphere was subdivided into lines of sight to ensure that the substructure along the line of sight is resolved. Finally, the flux from each line of sight was combined and weighted by area. Fig. 6 shows the resulting synthetic line profiles. For inclinations i ≳ 70°, the line profile shows both red and blueshifted components. The redshifted components arise in the wind that is falling back on to the white dwarf and accretion disc, and is therefore strongest when the white dwarf is closest to the observer and the line of sight intersects the most
Paper: 2020A&A...642L..18K (fulltext)
the satellite at two specific times, and where its spin is indicated by a curled arrow. When cold rotationally supported disc galaxies merge they produce complex tidal tails ( Toomre Toomre 1972 ; Eneev et al. 1973 ; Quinn 1984 ; Barnes 1988 ). These tidal tails can be quite different from those originating in spherical dispersion-supported systems. This is also what we find in our simulations, as shown in Fig. 4 . The bottom row of this figure zooms in on the discy satellite and reveals that some of the stars in green (i.e. the ridge in the bottom panel of Fig. 3 ), are lost first. These are followed by some stars in cyan with lower final eccentricity and slightly less retrograde motions since they are part of the ridge in the middle panel of Fig. 3 . We also see that there is some alternation in green and cyan stars amongst the material that becomes unbound slightly later. Since a discy satellite has internal and orbital angular momentum, mass loss depends on the exact configuration of the merger (e.g. inclination, spin) and not only on the internal binding energy of the stars as in the case of a spheroidal system. For example, the presence of a single extended counter-tail can be explained by the alignment of the spins of the two galaxies ( Struck Smith 2012 ), i.e. although the secondary moves on a retrograde orbit, its sense of rotation is prograde with respect to the primary. A difference in eccentricity also arises from these additional degrees of freedom. For a given internal binding energy stars have a range of internal angular momenta, which translates in debris with high and low orbital angular momentum, resulting respectively in more circular and more radial final orbits (the latter being typical of shell-like features). In addition to these effects, dynamical friction causes the satellite’s core to spiral inwards. This enhances the prominence of the only tidal arm apparent in Fig. 4 , which trails behind the core. 3.2. Some implications We have just shown that the stars in the velocity arch constitute debris lost early on in the merger. This probably also has implications for the chemical abundances of its stars since dwarf galaxies are known to possess metallicity gradients. For M ⋆ ∼ 10 9.6 M ⊙ the gradients have an amplitude ∼ − 0.064 dex kpc −1 according to Ho et al. (2015) . Thus, if Gaia -Enceladus had a physical extent of 5 kpc, this could imply a metallicity difference of 0.32 dex between its centre and outskirts. Myeong et al. (2019) studied some of the stars in the velocity arch, and argued that these stars, together with other retrograde stars, were part of a different accreted galaxy that they named Sequoia. These authors have shown that the stars are more metal poor on average than the debris of Gaia -Enceladus by ∼0.30 dex, which suggests they formed in an object that is ten times smaller ( Matsuno et al. 2019 ). However, an alternative explanation supported by our simulations is that these stars (or some of them)
Paper: 2020A&A...640A..19T (fulltext)
the Orion disks dust content, as measured with ALMA, are comparable with those measured for Perseus with VLA observations ( Fig. 5 bottom panel): 16 and 45% for Class 0 and Class I, respectively. Figure 6 presents a different visualization of the distribution of exoplanet systems dust masses compared to the range of disk masses observed in Class 0 and Class I disks in Perseus. The conclusions of our work show that for the complete sample of disks in Perseus and with a large sample of known exoplanets, there is enough solid material in the Class 0 stage to explain the solid content in observed exoplanetary systems. This conclusion is consistent with Najita Kenyon (2014) and Greaves Rice (2011) but now with much more robust statistics. In recent years, several studies, specifically with the use of microlensing observations, have estimated that nearly all stars could host a planet of at least 10 M ⊕ ( Cassan et al. 2012 ; Suzuki et al. 2016 ). Very few cases and irreproducibility of microlensing observations suggest caution should be taken when extrapolating the results to all systems. It should be kept in mind that a large population of wide-orbit planets could pose a challenge to efficiency of planet formation even at the Class 0/I stage ( Najita Kenyon 2014 ). Actual timescales for the different phases of low-mass star formation are uncertain. Based on a statistical analysis of a population of protostars, it is estimated that the Class 0 evolutionary phase lasts for ~0.1 Myr after the beginning of the collapse, and the Class I phase ends when the protostar is ~ 0.5 Myr old ( Dunham et al. 2014b ). More recent estimates of the half-lives of the protostellar phases give Class 0 half-life values of ~0.05 Myr and ~0.08 Myr for Class I ( Kristensen Dunham 2018 ). Therefore, our results indicate the start of the planet formation begins less than 0.1 Myr after the beginning of the cloud collapse. This is consistent with the ages of the oldest meteorites in our Solar System ( Connelly et al. 2012 ). 4.3 The context of planet formation models It is worthwhile to put our empirical constraints on the solid mass reservoir in the context of planet formation models. Broadly speaking, planets can either form bottom-up through the assembly of smaller building blocks in the so-called core accretion scenario, or top-down in the gravitational instability scenario via direct gravitational collapse of the disk material. In the latter case (see Kratter Lodato 2016 for a review), planets need not contain rocky cores and we may therefore have overestimated the solid mass locked in planets. In addition, in this view planets form at the very beginning of the disk lifetime when the disk is gravitationally unstable, possibly at even earlier stages than we probe here. If this is the case, our observations do not put constraints on the mass budget required for planet formation since planets would already be formed in the disks we are considering. Our results are instead