adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020A&A...644A.114K (fulltext)
becomes more highly extinct and fainter. Fig. 4 LUCI J -band image of T Tauri from November 2016 displayed on a logarithmic scale. The dotted arc indicates the northern arm of the reflection nebulosity discussed by Stapelfeldt et al. (1998) . The yellow arrows indicate the tadpoles. The red arrows mark bow-shaped features. Extended features around an apparent cavity are enclosed within the black dotted lines. The green, orange, and red dots identify the locations of T Tauri N, T Tauri Sa, and T Tauri Sb, respectively. The inset shows the tadpoles and the cavity with a different gray-scale cut. Fig. 5 Zoom of Fig. 4 on the area southwest of T Tauri N showing the coiling structure and the reflection nebulosity features R2 and R3 (as shown in Fig. 2 of Kasper et al. 2016 ) in the vicinity of T Tauri S. The dotted line connects the inflection points of the coil. The red arrow indicates T Tauri Sa. Fig. 6 T Tauri H 2 line emission map obtained with LBT-LUCI in October 2016. Only features which cannot be explained by image processing artifacts are indicated and discussed in the text. The positions of the stars are indicated by colored dots (T Tauri N, Sa,and Sb are in green, orange, and red, respectively). Fig. 7 IRDIS J -band image of the coiling structure taken in December 2018. The white dotted line indicates the coiling structure’s position in December 2014 (see Kasper et al. 2016 ). The relative distance to the T Tauri stars can be determined from the orbital solution and radial velocity (RV) measurements. According to the best-fit T Tauri N-S orbit ( Schaefer et al. 2020 ), T Tauri S went through the line of nodes in 1936 with an RV difference to T Tauri N of about ± 6.5 km s −1 . The sign of the RV difference cannot be determined from an astrometric orbit solution. Duchêne et al. (2005) found that the heliocentric RVs of T Tauri Sa and Sb appear to be larger than the one of T Tauri N ( Hartmann et al. 1986 ) by 2–3 km s −1 . These values are compatible with the orbit prediction within the error bars if T Tauri S previously went through the ascending node. Since then, T Tauri S has been moving behind T Tauri N for about 80 yr and would now be located roughly 90 AU behind T Tauri N. However, the N-S orbit is still only loosely constrained and orbital solutions may exist for which T Tauri S is located in front of T Tauri N, as suggested by Beck et al. (2020) . Fig. 8 IRDIS J -band image of the R4 structure taken in December 2018. The white dotted line indicates R4’s position, and the “+” symbols indicate the positions of T Tauri Sa and Sb in December 2014 (see Kasper et al. 2016 ). Fig. 9 IRDIS J -band image of the R3 structure taken in December 2018. The white dotted line indicates R3’s position in December 2014 (see Kasper et al. 2016 ). Figure 1 also shows a bow shock at ~1′′.6 west of T Tauri. This bow shock is a signpost of the east–west outflow ( Böhm Solf 1994 ; Herbst et al. 1997 ) with the western part approaching (blue-shifted) at about
Paper: 2019A&A...622A..40K (fulltext)
fully convective (e.g., Spada et al. 2017 ). A radius of ≈0.4 R ⊙ leads to cycle times of about 4–5 years, provided that η T remains unchanged. If η T of cooler stars is reduced with respect to the solar value, as implicitly suggested by the stellar models of Spada et al. (2017) , then the cycle time from Eq. ( 1 ) basically exceeds this limit. Mixing-length estimates of 0.33 u rms l corr for M dwarfs provide about 30% of the solar value, so that the dynamo cycle times of M dwarfs should slightly exceed the duration of the solar cycle. Indeed, the cycle time of one star in Fig. 1 slightly exceeds the 11 yr (≈4000 d) limit but this turned out to be a K2 dwarf star with a rotation period of 22.8 d ( Johnson et al. 2016 ). Figure 1 shows the observed cycle lengths for M dwarfs to be generally shorter than ≈3000 d. GJ270 (spectral type M2) shows the longest known cycle period of any M star with 2700 d (7.4 yr), while the shortest cycle is seen in GJ476 (M4) with 1060 d (2.9 yr) ( Robertson et al. 2013a ). GJ328 as an M0 dwarf with 0.69 M ⊙ shows an activity cycle of ≈2000 d, while a Jupiter-mass planet orbits it with a 4100 d period ( Robertson et al. 2013b ). We note that in the solar system the orbital period of Jupiter coincides with the length of the solar activity cycle, accidental or not. GJ581 possesses an entire planetary system with four or possibly six planets and shows an activity cycle with a length of 1600 d ( Gomes da Silva et al. 2012 ; Robertson et al. 2013a ). The rotation period of the latter star is surprisingly long with 130 ± 2 d ( Robertson et al. 2014 ). Fig. 1. Cycle times vs. rotation period (both in days) for early M stars (rhombic symbols). Among the group of rapid rotators with rotation periods less than a day ( Vida et al. 2014 ) are also the early M dwarfs EY Dra and V405 And ( Vida et al. 2013 ). The remaining data are from Suárez Mascareño et al. (2016) ; Wargelin et al. (2017) , and Distefano et al. (2017) . For GJ476 the cycle time is known but not the rotation period (blue bar). The Schwabe cycle of the Sun is indicated as an asterisk denoted SUN. Gomes da Silva et al. (2012) give further examples of M dwarfs with cycle lengths between three and five years. It is of particular importance to know the relation of the cycle time with the rotation time for a homogeneous sample of stars. Savanov (2012) used light curves of cool M dwarfs to find long-term cycle variations up to timescales of years. His sample is mixed by M stars with outer convection zones and also fully convective stars. There seems to be no systematic dependence of the cycle time on the rotation time even for subsamples with slow and fast rotations. An independent study about the coupling of rotation and cycle times for fast-rotating cool stars has been published by Vida et al. (2014) . Among the analyzed 39 Kepler light curves of late-type stars with rotation periods shorter than one day, nine examples showed cyclic variations with periods between
Paper: 2021ApJ...908L...3R (fulltext)
by comparing the observed spectral energy distribution (SED) to the sum of the SEDs of both models. The reddening is modeled following Cardelli et al. ( 1989 ). From the total luminosity, the fixed light ratio and the other determined parameters, we then derive the individual luminosities. The final spectral fit is shown in Figure 4 . The inferred parameters are given in Table 4 . The spectroscopic mass for the O-type secondary is derived from , where G is the gravitational constant. The spectroscopic mass of the WR primary stems from mass–luminosity relations calculated for chemically homogeneous WR stars by Gräfener et al. ( 2011 ), which for a H-free star depends only on L . The error on the spectroscopic mass of the WR star is propagated from the error in L , and is possibly underestimated in its lower limit given that the WR primary may be chemically inhomogeneous, which would lower M for a given L . Figure 4. Upper panel: comparison between the observed SED (FUSE, IUE, and UBVJHK photometry, blue lines and squares) and the synthetic PoWR SED modeled for WR 133 (red dotted line). The synthetic SED is the sum of a PoWR model for the WR primary (black solid line) and O-type secondary (green dashed line). Lower panels: comparison between the observed and synthetic normalized spectra. See the text for details. Table 4 Inferred Physical Parameters of WR 133 from the PoWR Spectral Analysis Note. a The transformed radius should not be confused with the physical radius of the star; see the text for details. 5. Discussion Sana et al. ( 2012 ) found that about 75% of massive OB stars should have their evolution impacted by a companion. The general assumption used in the analysis of binary populations is that two stars will interact if the system is born with a semimajor axis of less than ∼10 au and the star evolves into a red supergiant (e.g., Sana 2017 ). During a binary interaction, the system initially becomes more compact until the mass ratio is unity and then increases until mass transfer eventually ceases. This has been seen in an active phase for binaries that could become similar to WR 133 with systems such as the LBV-like binaries MWC 314 (Richardson et al. 2016b ) and HDE 326823 (Richardson et al. 2011 ). The supergiant companion to the WR star in WR 133 could have a larger rotation rate than an average field star. The inferred value of of 96 km s −1 is similar to the average value found for O supergiants by Simón-Díaz Herrero ( 2014 ). However, if there were indeed past binary interactions, the O9I star should be rotating in the orbital plane, and thus would likely be a rapid rotator. In fact, if we use our derived value for and the inclination of the system imply a modestly fast rotation rate of 280 km s −1 , which is a reasonably fast rotation for a supergiant O star. However, if the O star had a magnetic field from the interaction, then the rotation rate could be drastically less than critical in a short time (Potter et al. 2012 ).
Paper: 2017A&A...604A.120N (fulltext)
0.04 and 0.03 mag for and distance modulus, with metallicities from Rey et al. , and 0.06 and 0.05 mag for the and values derived with the metallicities from Sollima et al. At the same time, results from previous studies (e.g., Cacciari Clementini 2003 ; Del Principe et al. 2006 ; Feast et al. 2008 ; Gran et al. 2015 ) suggest that distances derived on the basis of such relations can probably not yet be considered accurate to better than about mag. Thus, we adopt the latter value as a conservative estimate of the systematic error affecting the distance moduli provided in our study. Fig. 8 Fundamental (in purple circles) and first-overtone (in yellow triangles) SX Phe as classified by Olech et al. (2005) , based on their frequency spectrum. The dotted lines are the PL relations derived previously. As can be seen, there are some claimed first-overtone pulsators that also follow the relations defined by fundamental-mode pulsators. Table 4 Adopted distance moduli based on T2Cs and RRL stars. 4.2. Pulsational distance modulus from T2Cs and RRL stars Table 4 lists the weighted-average distance modulus values from T2Cs and RRL stars and sigma as the dispersion around the mean. The weights for each individual distance modulus values were defined as the inverse of the squared of the total error, including statistical and systematic errors (summed up in quadrature). As expected, the error is lower for the PL relations compared to , since the former have lower intrinsic dispersion, and is based on a more extensive dataset. Comparing both variability types, T2Cs lead to a less accurate distance modulus, probably because they are few in number, which increases the statistical error, and they have higher systematic errors since the effect of metallicity and pulsation mode separation in the derivation of the PL relations is not completely established. Considering the distance modulus values found using T2Cs stars as well as RRL stars, both in and , the weighted-average distance modulus to Cen was determined, leading to a final value of mag, where the error bar corresponds to the standard deviation of the mean. Given the discussion of systematic errors in the previous section, our final adopted distance modulus is mag, where the error bars correspond to the adopted statistic and systematic errors, respectively. Adding the errors in quadrature (as recommended, for instance, by Barlow et al. 2002 ), this corresponds to a heliocentric distance of kpc. Table 5 Distance modulus for Cen from the literature. Notes. (a) Most of the error estimates do not consider the contribution of systematic errors, which in the case of T2Cs and RRL stars can reach values of order 0.10 mag. (b) The listed error was determined considering the internal uncertainties of the mean distance modulus value, 0.02 mag, and the observational error, 0.1 mag, as described by McNamara (2000) . (c) Based on the -[Fe/H] relation from Bono et al. (2003) and mean magnitudes from Kaluzny et al.
Paper: 2020ApJ...894...30W (fulltext)
environment theoretically depends on the Mach number or wave amplitude (e.g., Priest 2014 ; Vršnak et al. 2016 ). As the shock wave propagates, its amplitude decreases due to the geometric expansion of its front and the fact that its leading edge is propagating faster than its trailing edge. The amplitude decrease results in the deceleration of the shock wave. Additionally, a power-law fit is a better representation of the distance–time curves of freely propagating shock waves (Warmuth Mann 2011 ) according the the Sedov solution (Sedov 1959 ). On the other hand, we found that the speed of the reflected wave (243.1 ± 30.8 km s −1 ) is much smaller than the speed of the incident wave (which approximately equals that of the unaffected wave at around 05:13 UT: 591.0 ± 25.9 km s −1 ). In other words, the reflected wave was slower than the primary wave. There may be several possibilities for this feature. The primary wave is a nonlinear wave or shock, which is propagating with a Mach number larger than unity, while the reflected wave had changed to be a linear one traveling at the characteristic speed (Liu Ofman 2014 ; Warmuth 2015 ). The reflected wave has to propagate in the rest frame of the flow field associated with the primary wave, which may result in an apparently lower speed (Kienreich et al. 2013 ). Additionally, it may be related to the properties of the medium (magnetic separatrix) causing the shock wave to reflect. The compressible magnetic separatrix may cause the decrease in the reflected wave speed, when the shock wave reacted with the magnetic separatrix. Moreover, it is interesting to find that the refracted and diffracted waves decelerated with a bigger nonconstant deceleration rate. We find that their deceleration rate was decreasing with time, as in a damping motion. These phenomena would help us to further understand the nature of the interaction between magnetohydrodynamic shock waves and the local magnetic field. Appendix Error Estimations In this study, we determine the dynamic characteristics of EUV waves by using the method in Liu et al. ( 2010 ). The track technique ignores the heights of EUV waves by assuming that all emission originates from the spherical solar surface, and this distance measurement was corrected by the sphericity of the solar photosphere. This method would cause an overestimation in the measurement of some parameters (e.g., distance, velocity). We will supply some error estimations in subsequent sections. As Figure A1 (a) shows, a coronal structure B at a height of Δ h above the solar surface is projected onto the solar surface A observed by an observer on Earth (E). We set a new “north pole” at the disk center ( D ) and make AC = d , AC ⊥ OE, OA = OD = R ⊙ , OB = OF = R ⊙ + Δ h , and DE = 1 au. Defining θ = ∠AEO, we can get We make OG ⊥ EG and get Then, we getAccording to the law of sines, in △BOE, we can get Based on the above equations, we can get whereFirst, if the EUV wave originally
Paper: 2017ApJ...843...64C (fulltext)
between 3D and 1.5D calculations might be quite influenced by ingredients of the models that require improvement. We show the relevance of these ideas commenting on two of the few papers existing on this topic. The key is that the influence of the 3D radiation field on the Hanle polarization has been posed since the beginning in terms of inhomogeneities, and not in terms of velocity gradients. The initial conclusions in this regard were obtained by Manso Sainz Trujillo Bueno ( 2011 ) with an analytical/numerical study approaching the problem in the static regime. In more recent numerical studies the action of 3D velocities in the polarization is included, but the contributions of vertical and horizontal velocity gradients are not separated. As the effects of horizontal velocity gradients and inhomogeneities are in turn blended, it is also unknown how large their relative contributions are to the differences between 1.5D and 3D calculations. Consider, for instance, Figure 3 of Štěpán Trujillo Bueno ( 2016 ), which points out that the Hanle effect at disk center tends to depolarize when it is accompanied by other symmetry breaking effect, as happens in the well-known case of the solar limb. The figure also confirms that macroscopic motions are yet an efficient polarizing source for the spectral line considered (as known from 1D simulations of the Ca ii IR triplet; Carlin et al. 2013 ). However, it is unknown whether the polarization introduced by velocity in that figure is due to the horizontal or to the vertical velocity gradients already considered in 1D calculations. This matters not only for comparing with 1D calculations but also because horizontal velocity gradients can potentially compensate and exceed the polarization created by horizontal inhomogeneities (as explained in previous paragraphs). Thus, the dominant symmetry breaking and polarizing effect at disk center could not be the inhomogeneities, as always affirmed, but the horizontal velocity gradients. Hence, we think that is key quantifying the sensitivity of such results to the model horizontal distribution of velocity gradients , which also can help detach the conclusions from the eventual lack of realism in them. Finally, it should also be noted that when those radiative transfer simulations do not add microturbulent velocity, they use the temperature distribution of MHD models that are known to represent a chromosphere cooler than the solar one. Thus, as temperature influences significantly the effect of the velocities in the LP (see Section 4.2 ), the relative strength of the polarizing effects discussed here might change accordingly. In summary, horizontal velocity gradients compete with other physical mechanisms, such as inhomogeneities and the Hanle effect, for polarizing and depolarizing the light; hence, we need to explicitly quantify whether they are relevant for inferring magnetic fields with the Hanle effect. We also need to find theoretical and observational
Paper: 2017MNRAS.470.2339L (fulltext)
ratio moving in circular orbits, the filling of the Roche lobe is followed by violent mass transfer by the primary to the secondary and the formation of a common envelope. On the other hand, an equal-mass binary may evolve into a fully contact system of two stars, both of which fill their Roche lobes simultaneously. In the limiting case, violent mass transfer may occur on the hydrodynamic time-scale, with an amount of released energy of the order of the thermal energy of the components or the common envelope: GM 2 / a 510 48 ( M /M ) 2 ( a /R ), where M is the mass of the component or of the common envelope and a is the radius of the common envelope. For 3M common envelope mass and energy from Table 1 , the envelope radius is about 10R . The high energy of the studied explosion is a consequence of compactness and the relatively high mass of the parent system. From this point of view, it is very important to obtain observational evidence for such processes in binary systems and to determine the corresponding physical properties: the released energy, ejected mass and angular momentum. Ultraluminous red novae may be observational manifestations of violent mass transfer in the binary (Blagorodnova etal. 2017 ). We already pointed out that observation of the progenitor of V1309 Sco in the form of a contact binary with a period of 1.7d (Tylenda etal. 2011 ) provides strong evidence for the merger model of LRNe. This conclusion is also corroborated by the approximate occurrence rate of such events one event in 2040yr per galaxy like the Milky Way or M31. To show this, let us use the distribution of binary semimajor axes a (Masevich Tutukov 1981 ). This distribution has the form d N 0.2dlog( a /R ). Systems with 10 a /R 1000 can pass through the common-envelope stage during the second mass exchange stage. According to the initial distribution of binaries over semimajor axes of Masevich Tutukov ( 1981 ), N 0.2dlog( a /R ), the frequency of formation of common envelopes with radii 610R is about 0.04 per year. Systems with larger component separations are less energetic. 5 CONCLUSIONS In this article, we described the discovery of the nova M31LRN and long-term observations of its light curve with the MASTER network of robotic telescopes. It is an important point that the entire observational part of the study was performed on identical telescopes equipped with identical photometers. The resulting light curve agrees fairly well with the independent light curve published by Kurtenkov etal. ( 2015b ). An analysis of the photometric observations and velocity data for the nova M31LRN made using radiative hydrodynamic modelling showed that the rather long observable plateau (50d) could be achieved only with a higher merged stellar mass M 3M and common envelope radius R 10R . The corresponding explosion energy should be lower, 310 50 erg, whereas the total kinetic energy of the ejected envelope is lower by three orders of magnitude. The proposed
Paper: 2019MNRAS.488.1760S (fulltext)
the literature value of 5.2 kpc, but consistent within the uncertainty, and so this star is still marginally consistent with being an LBV on the S Dor instability strip. The higher precision expected in future Gaia data releases is needed to say more. WS 1: WS 1 was classified as a bona-fide LBV by Kniazev, Gvaramadze Berdnikov ( 2015 ) based on observations of significant photometric and spectroscopic variability. The Gaia DR2 distance of 2.48 kpc is significantly lower than the baseline literature value of ∼11 kpc, although this object has a negative parallax and like the previous two, is dominated by the prior. However, the literature luminosity is also highly uncertain because it is not based on a measurement – it was derived by assuming that the star fell on the S Dor instability strip Gvaramadze et al. ( 2012 ), so perhaps a large revision is not so surprising. Gaia data suggests that the star is significantly lower in luminosity, although again DR3 data are needed for a confident result. MN48: MN48 was classified as a bona fide LBV by Kniazev et al. ( 2016 ) who identified spectroscopic and photometric variability typical of LBVs. There is large excess noise in the Gaia DR2 data for this star, and the resulting 68 per cent HDI spans 1.44–5.18 kpc. The upper end of this range is marginally consistent with the previous literature value of ∼5 kpc, although again, the literature value was highly uncertain as well. 3.2 Candidate LBVs HD 80077: This blue hypergiant was included as an LBV candidate by van Genderen ( 2001 ) based mainly on its high luminosity and spectrum. Its new Gaia DR2 distance is significantly closer than the baseline literature value, reducing its nominal luminosity by about 40 per cent, although its luminosity is still above the S Dor strip. The Gaia DR2 distance and the original cluster-fitting distance of Steemers van Genderen ( 1986 ) (2.8 ± 0.4) agree within their quoted errors. SBW1: As in the cases of Sher 25 and HD 168625, this blue supergiant is an LBV candidate because of its circumstellar nebula that bears a remarkable resemblance to the ring nebula around SN 1987A. In fact, it has been argued that SBW1 is in some respects the best Galactic analogue to the progenitor of SN 1987A, in terms of both its nebula and the properties of the weak-winded central star (Smith et al. 2013 , 2017 ). Based on various considerations, Smith et al. ( 2017 ) favoured a distance of about 7 (±1) kpc for SBW1. The Gaia DR2 parallax indicates a distance of 5.2 kpc (Table 3 ). This distance is closer, but the range of 4.0–6.8 kpc overlaps with the uncertainty of 7 (±1) kpc for the previous distance. This new distance corresponds to a somewhat lower effective single-star initial mass of 13–15 M ⊙ , as compared to ∼18 M ⊙ for the progenitor of SN 1987A (Arnett 1989 ). This confirms the notion that SBW1 is an analogue of the progenitor of SN 1987A and its nebula (Smith et al. 2013 ), although at somewhat lower initial mass.