adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020A&A...636A..48G (fulltext)
5.2. Comparison with the IRSF monitoring data in the SMC The NIR monitoring data for three stars has recently been published by Ita et al. (2018) using the InfraRed Survey Facility (IRSF, Kato et al. 2007 ). These data were not used in our LC analysis and can therefore serve as an independent test. The K -band data ( H band for 12.627600 −72.858307) were analysed in the same way as the VMC data and literature data. Table 2 compares the periods and amplitudes. For the first two stars the agreement is good. In the last case the periods do not agree within their mutual error bars, but this is a faint source for the IRSF survey ( K ∼ 15.5 mag), while our analysis might be hampered by the relatively small number of data points (18 in total). Table 2. Comparison with three sources in the IRSF survey. Notes. (a) In the K band. 5.3. LPVs with periods over 1000 days Of the likely AGB stars, 34 stars have periods longer than 1000 days 6 . The longest period we find is 3108 d, but this object has alternative periods that are much shorter. Six other stars also have shorter possible periods, and therefore the periods beyond 1000 days are only tentative in these cases. It will be interesting to analyse the IRSF data for the LMC objects when they become available (cf. the SMC data presented in Ita et al. 2018 ). Nine stars with periods over 1000 days in the MCs were previously reported by Wood et al. (1992) and Whitelock et al. (2003) based on NIR observations. Recently, Menzies et al. (2019) presented a compilation of Miras with periods over 1000 days in the Milky Way (17), SMC (3), and LMC (18), while Karambelkar et al. (2019) published a list of 417 luminous IR variables, of which 86 have periods over 1000 days, located in seven galaxies (the majority in M 83 (49 objects), and M 81 (11 objects)). The compilation by Menzies et al. (2019) was based on Wood et al. (1992) and Whitelock et al. (2003) , OGLE data in the SMC and LMC ( Soszyński et al. 2011 , 2009a ), and GS18 (which is partly based on a preliminary analysis of VMC data). Of the 21 MC objects, 18 are confirmed with updated periods. Three objects are not in our list. One object (IRAS 05506−7053) is too red; the period in GS18 is based on WISE data. The two others (HV 888 and HV 11417) are too blue to have passed the selection criteria, and both have been classified as RSGs (GS18, Kraemer et al. 2017 ). The longest period in the study of GS18 of almost 400 AGB stars with Spitzer IRS spectra is 1810 d based on a re-analysis of OGLE-III data for MSX SMC 055 (also known as OGLE-SMC-LPV-08137 and IRAS 00483−7347). No errors on the period are given in GS18 because they are very hard to reliably estimate for these long periods. The period in the OGLE-III release is 1859 d, which gives some indication of the uncertainty. In Groenewegen et al. (2009) such a long-period has been recognised based on OGLE-II data and, it was suggested that the object may be a candidate super-AGB star, based on its
Paper: 2019A&A...631A..64B (fulltext)
the 12 CO line profiles are dominated by infall and outflow motions and we find no significant evidence of rotation in the outflow. As mentioned in Sect. 3.2 , however, CH 3 OH and SO 2 do show clear evidence for rotation within ~16 au ( Fig. 8 ). This emission originates in the central region, where it appears that material is infalling and rotating simultaneously. Even so, the velocity gradients measured from the CH 3 OH and SO 2 lines are consistent with one another, and thus one might expect that the outflow should rotate in the same sense. A 12 CO moment 1 map ( Fig. A.3 ) and line profiles at different positions ( Fig. 4 ), meanwhile, do not show any significant difference in the velocities between northern and southern outflow cavity walls. In the redshifted outflow, however, there is a blueshifted component at small separations (≲30 au) south-west of the protostar ( Fig. 3 ), consistent with the inferred rotation seen in SO 2 and CH 3 OH. Likewise, a redshifted component is observed in the blueshifted outflow on comparable scales, but this component is seen along the jet axis and not the north-western side as would be expected if this component was due to rotation in the outflow. Attributing the inner features to rotation is in any case difficult since the outflow is nearly in the plane of the sky, and expansion of the outflow, deviations of the outflow axis from axisymmetry, or internal shocks (e.g. Fendt 2011 ) can produce similar features. Hints of expansion are indeed seen in the lowermost PV diagram of Fig. 9 , where an extended U-shaped structure is observed at blueshifted velocities. Although this emission structure barely exceeds the 3 σ threshold, itis clear that the velocity is increasing towards the outflow axis. A similar feature is found on the redshifted side at slightly lower S/N, and lower velocities. While an onion-layered velocity structure (whereby velocities close to the outflow axis are high and then fall off with distance) is expected for a rotating, magnetically-launched wind (e.g. Bacciotti et al. 2002 ; Pudritz et al. 2006 ), the symmetry of the PV diagram w.r.t. the zero position offset in the lowermost panel of Fig. 9 is not consistent with this picture. It is, however, consistent with expanding cavity walls. The rotation and expansion velocity at different offsets from the protostar can be estimated using the method presented in Lee et al. (2018) , assuming that only two velocity components are present. While this is an oversimplification of reality, assuming the outflow velocity does not change significantly over the considered distances, it does provide a useful method to deduce any trends in rotation and expansion. We fit 2D Gaussians to several PV diagrams with cuts perpendicular to the outflow axis but with different offsets. One example of such a fit (at zero offset from the continuum peak) is presented in the uppermost panel of Fig. 9 ). We only consider pixels where the emission is above 3 σ ,
Paper: 2017ApJ...851..150H (fulltext)
2009 ), which assumes that all LyC flux converts to unbinding the atmosphere. The LyC flux corresponds to the energy-limited mass-loss rate . In this case, the adiabatic cooling rate is and still has a less than 15% effect in the region of main concern. 8.2. Comparing with Other Hot Jupiter Upper Atmosphere Models Christie et al. ( 2013 ).—The present study agrees with the conclusion in Christie et al. ( 2013 ) that the 2 p occupation is set by radiative excitation and de-excitation, and as an improvement, we include an Ly α radiation transfer inside the atmosphere instead of applying a constant solar Ly α intensity. Because of the ∼30 times stronger stellar Ly α intensity of HD 189733 compared to the Sun and considering the Ly α photons generated inside the atmosphere owing to collisional excitation and recombination cascades, the Ly α mean intensity should be ∼100 times larger for the majority of the atomic layer. In addition, because Christie et al. ( 2013 ) underestimate n 2 p by a factor of 20 owing to a math error, n 2 p should be significantly larger in the whole simulation domain, and thus the atmosphere is optically thick to H α mainly owing to the absorption of H(2 p ). The observed H α absorption width agrees well with an optically thick atmosphere model. Because of the much larger , the photoionization of H(2 l ) is larger compared to photoionization of the ground state in the atomic layer. Hence, n e and n p in this work are ∼10 times larger. Considering the proton collisional l -mixing process with a rate ∼10 times larger than the electron collisional process, as well as the large H(2 p ) population, the creation of 2 s hydrogen is dominated by l -mixing rather than collisional excitation considered in Christie et al. ( 2013 ). As a result, 2 s and 2 p reach collisional equilibrium. In addition, it is shown that the metal lines are crucial in cooling the atmosphere. Assuming solar abundance, lines of Mg and Na can cool the atomic layer by ≃2000–3000 K. Menager et al. ( 2013 ).—Menager et al. ( 2013 ) investigate the Ly α emission and reflection by the atmosphere of HD 189733b. The temperature and electron, hydrogen, and helium number density profiles of HD 189733b from the Koskinen et al. ( 2011 ) unpublished model were applied. The temperature, n e , n p , and are in broad agreement with the profiles presented in this paper in the corresponding pressure range. According to a similar model of HD 209458b presented in Koskinen et al. ( 2013a ), it should be a one-dimensional hydrodynamic model of the upper atmosphere considering hydrogen and helium constructed on top of a full photochemical model of the lower atmosphere. They chose the average solar flux as their stellar spectrum, which is ∼10 times smaller than the synthetic spectrum from MUSCLES. Different from the photoelectron heating efficiency calculated at the fixed ionization fraction x e = 0.1 throughout the model, or a constant η applied in Koskinen et al. ( 2013a ), a
Paper: 2019A&A...622A..61S (fulltext)
spectrum (Figs. 2 and 3 of Ho et al. 2003 ), while atomic hydrogen transitions may produce additional features. That may explain some of the observed multiple lines. In view of the limited significance of the line detections it appears difficult to make specific assignments. Another way to explain multiple lines is to invoke – beyond the radiation from the hot cap at the magnetic dipole – separate hot spots on the stellar surface having higher multipolar magnetic fields. The first strong evidence for this comes from the discovery of strongly phase-variable absorption lines which are visible only for about 20% of the neutron star rotation. The lines were detected in the X-ray spectra of RX J0720.4−3125 and RX J1308.6+2127 ( Borghese et al. 2015 , 2017 , respectively) and have energies a factor of ∼2.5 higher than the broad lines (Table 3 ). Such multipole fields must be quite young compared to the age of the NS. They may be formed by extracting magnetic energy from the toroidal field that resides in deep crustal layers, via Hall drift (e.g., Geppert Viganò 2014 ). We conclude that proton cyclotron lines produced in strong magnetic fields of isolated neutron stars can explain the absorption-type features observed in the soft X-ray spectra of young isolated neutron stars of the M7 class. More accurate measurements in the future in combination with further theoretical work on magnetized atmospheres and the magneto-thermal evolution of neutron stars are required to make further progress in this exciting field. 9. Summary With this contribution we attempt to provide a review on the current state of our knowledge about X-ray sources that show cyclotron line features in their spectra – both for electron cyclotron lines (this is by far the larger part) and for proton cyclotron lines. We concentrate on the observational aspects, providing a number of Tables that contain detailed informations on all sources that we consider reasonably well secured regarding the existence of cyclotron lines. In the case of electron cyclotron lines – observed from binary X-ray pulsars, highly magnetized accreting neutron stars in binaries – we also give a Table with candidate sources, mostly objects for which a detection of a cyclotron line was claimed some time ago, but had not been confirmed in subsequent observations, or sources that were observed rather recently, but only once and with somewhat peculiar properties. On the other hand: for some of the objects listed in the main Tables, the existence of a cyclotron line still needs confirmation (or not) through re-measurements (with a finite probability that they have to be moved into the candidates category). We try to address the basic physics underlying the cyclotron line phenomenon, including a large number of references to theoretical work, with some coverage of specific theoretical modeling. Since the discovery of the first CRSF in Her X-1 in 1976, a new field of Cyclotron Line Research has evolved. With the rather
Paper: 2018ApJ...868..128S (fulltext)
multilayer system. In general, unlike the symmetric case studied, e.g., by Roberts ( 1981b ), the dispersion relation Equation ( 15 ) cannot be decoupled into two equations that correspond to the well-known sausage (oscillations at the slab boundaries in antiphase, that correspond mathematically to the equation containing ) and kink (perturbations of the slab boundaries oscillating in-phase that correspond to the equation with ) MHD waves. Solutions to Equation ( 15 ) will provide modes modified by the density difference at the sides of the magnetic slab similar to a single asymmetric slab investigated by Allcock Erdélyi ( 2017 ). We adopt the notions of quasi-sausage and quasi-kink to describe these eigenmodes, because they have mixed characters. The slab cross-sectional width, that is constant for symmetric kink modes, is affected for asymmetric kink modes and the line of zero perturbation is shifted for asymmetric sausage modes when compared to the symmetric case, namely, the center of the slab. Furthermore, both the sausage and kink modes can be further categorized as body and surface modes. Surface waves exist when , which corresponds to evanescent solutions of Equation ( 5 ), while body waves exist when , which corresponds to spatially oscillatory solutions. This nomenclature for wave classification was introduced by Roberts ( 1981b ) for symmetric and Allcock Erdélyi ( 2017 ) for asymmetric equilibria. Let us now consider some specific cases of an asymmetric multilayered plasma slab structure that are analytically solvable and have potential for solar application. First, by setting p = q = 0, we verify that the dispersion relation reduces to that which governs the asymmetric slab configuration of Allcock Erdélyi ( 2017 ), namely,We will now focus on cases of two layers (one magnetic and one nonmagnetic), and three (a magnetic layer sandwiched between two nonmagnetic) slabs, respectively, and investigate how the multilayered structure affects the eigenmodes and their eigenfunctions. 2.1 Two-slab Case Let us first consider the case of two slabs, represented by p = 1 and q = 0, which is illustrated by Figure 2 . The width of the nonmagnetic and magnetic slabs are denoted as and d 0 , respectively. Such a model could be useful for modeling MHD waves in MBPs of the solar photosphere (as shown on Figure 3 ). MBPs are approximately vertical magnetic structures between supergranule convection cells of different densities and temperatures. Since neighboring granular cells affect the MBP, we can apply a two-slab description to analyze what will change when comparing eigenmodes to MHD waves in a single slab case. Figure 2. Equilibrium configuration for a two-slab case: a magnetic slab, , and a nonmagnetic slab, . Figure 3. Asymmetric multislab approximation for an elongated MBP, based on Figure 11 of Liu et al. ( 2018 ), observed by the New Vacuum Solar Telescope in TiO 7058 Å. From condition ( 15 ), we obtain the dispersion relationNote, that the
Paper: 2017ApJ...837..156O (fulltext)
been identified by IE methodology, where the IE is less than 1.5%. c ( N IE /Total) × 100%. The same study explained in the above paragraph was performed with two other identifiers of MCs (see Huttunen et al. 2005 , their Table 2, and Wu et al. 2003 , /WIND list). In Table 3 , the results are shown in the last two columns. The low successfulness (55.5%) with Huttunen et al. ( 2005 ) has an explanation. The IE methodology was unsuccessful in four events: one event was identified as a cloud candidate (cl) with a time extension of eight hours; two events were classified as MCs, however, on that date ICMEs were not reported; and the last MC was also identified by Wu et al. ( 2003 , /WIND list, event 42 (September 21) with poor quality (3) and with a time extension of seven hours. 4.1.5. False Alarm Cases Furthermore, we saw false alarm cases. False alarm cases are grouped into three categories: (1) non-MC ICMEs according to the catalog of Richardson Cane ( 2010 , Table 3) (20 of these cases were identified); (2) stream interaction regions (SIRs) according to the catalog of Jian et al. ( 2006 , p. 367–371), where 36 SIRs were identified; (3) other cases (non-MC ICMEs and non-SIRs). In Table 4 , the previous categories are shown. In the first category, 15% of false alarms are reported. This is important and can save time, e.g., when a specialist wants to identify new MCs in a large data period. These three cases could be MC-ICME; however, they are not studied in this work. Table 4 Summary of Events That Are Not MCs Reported in 1999 and False Alarm Cases in the IE Analysis Notes. a Number of false alarm cases, i.e., where . b Percent of false alarm cases /Total) × 100%). c Adapted from Richardson Cane ( 2010 , Table 3), from 33 ICMEs: 13 were MCs-ICMEs and 20 were non-MC ICMEs. d Jian et al. ( 2006 , p. 367–371). In the second category, there is only one false positive case. In summary, the IE methodology is not useful to identify SIRs. Finally, five cases are in the third category. These cases are interesting because disturbances in the IMF can be observed. In addition, we are left with two open questions. What really are these events? Why do they have low IE values, similar to an MC? We think that these cases could be flux ropes. However, further study should be done to investigate the above cases. In summary, the IE methodology identified 9 from 13 MCs, and nine false alarm cases are emitted. In a year (1999), a total of 788 windows of 2500 values exist, which means that the percent of false alarms was 1.14%, which is a good result. However, we prefer to propose it as a preliminary result because one year is considered poor statistics to accept the result as a robust identifier of MCs. Despite the difficulties, we can conclude that the IE methodology is useful for studying MCs, and in the next section, the IE methodology is used to identify two MCs. 4.2. Results: Part II After tests on efficiency, the methodology is applied to study some SW
Paper: 2020MNRAS.492.1154B (fulltext)
an order of magnitude lower just by increasing the disc viscosity by a factor of 3–6 (Stehle 1999 ). Therefore, while increasing the disc temperatures increases the opening angle of the resulting spiral arms (e.g. Makita, Miyawaki Matsuda 2000 , and references therein), increasing the disc viscosity seems to have a similar (and seemingly stronger) effect. The point we make here is that wide open spiral arms may be more a signature of a high viscosity disc than of a hot disc, and that a large disc viscosity might lead to detectable, wide open tidally induced spiral density waves even when the disc is relatively cool and small. Thus, the detection of spiral arms in V2051 Oph during quiescence suggests that its quiescent accretion disc has high viscosity. This is in line with the results of Baptista Bortoletto ( 2004 ) – which inferred a high viscosity α ≃ 0.1−0.2 for the quiescent accretion disc of V2051 Oph – and with those of Baptista et al. ( 2007 ) – which suggest that the outbursts of V2051 Oph are the result of bursts of enhanced mass transfer on to a high-viscosity accretion disc. 5 SUMMARY Infrared JHK s quiescent light curves of V2051 Oph after removal of the ellipsoidal contribution from its mass donor star show an out-of-eclipse double-wave modulation the amplitude of which increases with wavelength. Eclipse maps derived from these light curves reveal two asymmetric arcs extended in azimuth at the intermediate/outer disc regions, which are reminiscent of those seen in eclipse maps of the dwarf nova IP Peg in outburst and are similarly interpreted as tidally induced spiral density waves in the V2051 Oph accretion disc. In agreement with this interpretation, the maxima of the double-wave modulation coincide with the binary phases where the spiral arms are seen face-on. The spiral arms are located at different distances from disc centre; the inner arm is at an average radius of (at a corresponding Keplerian velocity of ), while the outer arms is at . We estimate the opening angle of these spirals to be θ s = 21° ± 4°. The detection of wide open spiral arms in V2051 Oph during quiescence suggests that its quiescent accretion disc has high viscosity. Broad-band JHK s disc spectra were extracted for concentric annular disc regions. The inner disc atmosphere () is made of hot and tenuous gas ( and surface densities ), suggesting that either an extended hot and tenuous chromosphere veils the emission from an underlying, colder accretion disc or that the inner disc itself has evaporated into a hot and tenuous corona. The disc gas becomes cold and opaque in its outer regions (). The spiral arms are hotter and have lower column densities than the disc at same radial range, with differences within the 2σ confidence level. The JHK s best-fitting hydrogen continuum emission models are in very good agreement with the B -band observations of Baptista Bortoletto ( 2004 ), suggesting that the inner disc regions were also optically thin and with
Paper: 2016ApJ...816...41Y (fulltext)
the elongated EUV channel-like structure (Zhang et al. 2012 ; Chen et al. 2014 ; Cheng et al. 2014 ) in the high-temperature passbands (e.g., 94 and 131 Å) of the Atmospheric Imaging Assembly (AIA; Lemen et al. 2012 ). More recently, the detailed structure and evolution of flux ropes in the low solar atmosphere were presented by Wang et al. ( 2015 ) and Kumar et al. ( 2015 ). The flux ropes in the low solar atmosphere show some similarities in the morphology of the filaments but have not been confirmed. The relationship between the filaments and the sigmoids and the elongated EUV channel-like structure are still unknown. However, they show some common characteristics in the formation process. For instance, in the framework of magnetic field reconnection, an increasing number of observations suggested that magnetic flux ropes are created either via flux cancellation in the photosphere (Green Kliem 2009 ; Tripathi et al. 2009b ; Green et al. 2011 ; Savcheva et al. 2012 ; Wang et al. 2015 ) or the tether-cutting reconnection in the corona (Moore et al. 2001 ; Kim et al. 2008 ; Liu et al. 2010 ; Chen et al. 2014 ; Cheng et al. 2015 ). Therefore, the study of the evolution and formation of the filament may shed new light on the evolution and formation of the flux rope. At the chromospheric level, observations always show that the dark, loop-like threadlike structures are the most prominent features observed in H α images. Those threadlike structures that mark the emergence of the new magnetic flux are called an arch filament system (Bruzek 1969 ), and those that span and transverse to older PILs of active region are called field transition arches (FTAs; Prata 1971 ). In addition, other kinds of loop-like threadlike structures are nearly parallel to the PILs. However, both of them connect with opposite polarities and trace out the magnetic line of force. In this study, employing the excellent observations from the AIA and the Helioseismic and Magnetic Imager (HMI; Schou et al. 2012 ) on board the Solar Dynamics Observatory ( SDO ; Pesnell et al. 2012 ), we present the rapid formation of a filament via magnetic reconnection between two sets of threadlike structures on 2013 May 29. In the photosphere, at a region where some ends of the two sets of threadlike structures are close to each other, magnetic flux convergence and cancellation occur. This provided us with a good opportunity to detail the evolution and formation of the filament. Detailed observations and methods are presented in Section 2 . The results are shown in Section 3 . The conclusion and the discussion are given in Section 4 . 2. OBSERVATIONS AND METHODS The event was well-observed by the SDO /AIA and HMI, as well as the Global Oscillation Network Group (GONG) at the National Solar Observatory (NSO). AIA takes full-disk images in 10 ultraviolet (UV) and EUV wavelengths covering a wide temperature range, with a pixel size of 06 at a high cadence of up to 12 s. We mainly used the Level