adsabs/WIESP2022-NER huggingface dataset:
Paper: 2020MNRAS.493.2834K (fulltext)
The two transonic trajectories with are shown in pink. Green lines show two subsonic trajectories with and and two supersonic trajectories with and . The trajectories that are discontinuous across position are shown in yellow for the same set of constants, , but represent different branches. 3 HAMILTONIAN DESCRIPTION Hamilton’s principle for a conservative system is where the difference of the kinetic and potential energies is the Lagrangian, . This may be generalized for a fluid with internal energy, U (Herivel 1955 ), In non-dimensional units, the Lagrangian for the BP flow is the same as the steady-state equation ( 6 ), where the non-dimensional terms, (log | u | + 2 log | x |), which derive from the dimensional pressure term in equation ( 1 ), are the internal energy. If we define the functional gradient as the directional derivative, for any variation, η, and constant, ε, then the necessary condition of Euler–Lagrange for x to be a critical function of J [ x ] is that the first variation is zero, where the velocity . Since the right-hand side is equivalent to the equation of motion ( 5 ) as can be verified by substitution, then 〈∇ J [ x ]〉 = 0, and the steady-state solution, , is the solution of stationary action toward which the time-dependent solutions evolve. This description is similar to the simple physics problem in which the Lagrangian determines the parabolic trajectory of a particle in a gravitational field and the exact trajectory is determined from the end points as a boundary value problem (BVP) or alternatively from the initial velocities as an IVP. However, the Lagrangian for the BP flow includes the internal energy of the fluid as a function of the density–pressure gradient that is constrained by the conservation equation. Therefore, it is not always possible to find an energy conserving trajectory between two arbitrary velocities. Considered as a BVP, not all boundary conditions are consistent with a steady-state solution. As in Lagrangian mechanics, the energy serves as a Lyapunov function for the stability. The definiteness in sign of the second variation evaluated at the steady state implies at least linear stability (Arnol’d 1965 , 1989 ). Equivalently, the integrand of the second variation, must be definite in sign for every non-zero variation η. Since the middle term with mixed partial derivatives is zero, and This is easily verified for the range x > 0 allowed for BP flows. This implies that the steady-state solution, including all the subsonic, transonic, and supersonic trajectories in the family, is at least linearly stable. The existence of a steady-state solution, a critical point in Lagrangian mechanics, is required for stability. An arbitrary IVP can evolve to the steady state, but the boundary conditions in a BVP must be consistent with the conservation of energy to allow a stable solution. This requirement itself does not determine the steady-state solution or its stability. For example, the derivation of
Paper: 2020ApJ...894...42M (fulltext)
these models demonstrate that banded clouds are easily able to explain the level of polarization measured in Luhman 16A. Considering the disagreement between the measurements and the homogeneous oblate polarization model, we conclude that cloud banding is the dominant source of the measured polarization in Luhman 16A. For Luhman 16B, we can neither rule out nor confirm the presence of banded clouds; the measurements could be explained as due to oblateness alone with clouds thinner than in the f sed = 1 models, or by a combination of oblateness and banded clouds. Although we model the bands and fit the polarization as constant over longitude and time, there is the possibility of variability with amplitudes below our sensitivity limits. Indeed, the bands of Jupiter display variability, yet the bulk features remain relatively constant over time (Ge et al. 2019 ). Similar variability superimposed on the bands discovered here may account for the previously detected photometric variability seen in Luhman 16A. 4.3. Spin Axis Orientations from the Angle of Linear Polarization The angle of polarization for Luhman 16A and B differs by ∼105° (or equivalently, ∼75°, because of a 180° degeneracy in the definition of the angle of polarization). For each object, the angle of polarization could be aligned either parallel or perpendicular to the position angle of the projected on-sky spin axis, as discussed in Section 4.2 . The misalignment of the angle of polarization immediately suggests a misalignment of the two components’ spin axes of either ∼15° or ∼75°, depending on the source of polarization of Luhman 16B. Neither component has an angle of polarization aligned parallel or perpendicular to the longitude of the ascending node of the binary orbit, Ω = 13012 ± 012 (Garcia et al. 2017 ). Luhman 16A’s polarization is at best ∼20° misaligned, and Luhman 16B is ∼40°–60°, indicating that neither object has a spin axis aligned with the orbital plane, regardless of the oblateness or banded interpretation. The misalignment may imply that the two objects did not form alone together within the same disk, but instead may have instead experienced a more raucous dynamical history such as having originally been formed in a triple system (Reipurth Mikkola 2015 ). To the best of the authors’ knowledge, this is the first measurement of the projected spin axis for any brown dwarf (albeit with a 90° degeneracy). Given the low-S/N detection of Luhman 16B, we consider any interpretation of its position angle to be tentative. 5. Total Intensity Variability Interpretation We were unable to measure absolute photometric variability, but the flux ratio of the two objects exhibits quasi-periodic variations, with a period of 1.64 hr. We attribute this signal to Luhman 16B alone ( H -band variability has never been detected in A; e.g., Biller et al. 2013 ; Buenzli et al. 2015b ) and suggest that it may represent a similar three-band scenario to that inferred for other L/T
Paper: 2017ApJ...835..264R (fulltext)
in the F125LP image shown in Figure 1 (top left). The FUV bow shock appears to be asymmetric, being thicker and brighter east of the pulsar. We measure the apex distance (from the pulsar to the leading edge of the bow shock) to be ≈25, consistent with the measurements in H α by Brownsberger Romani ( 2014 ), Gaensler et al. ( 2002 ). To extract the bow-shock flux, we defined the source and two background regions, “inside” and “outside” of the bow shock, shown in Figure 1 (top right). The background regions were selected based on the F125LP image, in which the bow shock was most clearly detected. The regions are close to the bow-shock extraction region to mitigate effects from the inhomogeneous detector background while excluding the pulsar. Because of the larger apparent thickness of the east side of the bow shock, more space was left between the bow-shock extraction region and the background regions. We took the background variations into account by selecting background regions on both sides of the bow shock: one “inside” and one “outside” the bow shock. These regions were used for flux measurements in both F125LP and F140LP. A combined background from the two regions was extracted. The net source count rate was calculated similarly to the pulsar, with the area A b being the combined area of the “outside” and “inside” background regions. The F125LP and F140LP count rates and their uncertainties are given in Table 2 . From the measured count rates we calculated the absorbed fluxes in the F125LP and F140LP wavelength ranges: F (1250–2000 Å) = (8.1 ± 0.4) × 10 −15 erg s −1 cm −2 and F (1350–2000 Å) = (7.0 ± 0.7) × 10 −15 erg s −1 cm −2 , for a flat f λ spectrum. The unabsorbed luminosity can be estimated as L (1250–2000 Å) ≈ 1.9 × 10 29 erg s −1 , for E ( B − V ) = 0.03, and d = 410 pc. The bow shock is not detected in the F475X images (see Figure 8 ). Because of the apparent non-uniformities in the background, which significantly exceed the statistical fluctuations, we chose to measure the upper limits by sampling the background counts from ten different regions in the vicinity of the bow shock and calculating the standard deviation from the mean value. The areas of the background regions were equal to that of the bow-shock region used for FUV flux measurements. The standard deviation can be considered as a conservative 1 σ upper limit. The corresponding 3 σ upper limit on the bow-shock flux in the area of 52 arcsec 2 is F (3750–6000 Å) ≈ 2.9 × 10 −15 erg s −1 cm −2 , for a flat f λ spectrum and E ( B − V ) = 0.03. Figure 8. UVIS/F475X image with the bow-shock extraction regions (the same as in Figure 1 ), used for the data analysis. The pulsar and the direction of its proper motion are shown with a circle and an arrow, respectively. Similar to our analysis of the FUV bow shock around J0437 (Rangelov et al. 2016 ), we measured the ratio of the count rates in F125LP and F140LP filters for the J2124 bow shock. We obtained a ratio of C 125 / C 140
Paper: 2020MNRAS.491.5179M (fulltext)
grains (Fujiya et al. 2013 ; Jadhav et al. 2013 ; Liu et al. 2014 ). Stellar evolution theory predicts that, besides the H-ingesting He-shell flash convection on white dwarfs, the i process may also occur in metal-poor low-mass thermally pulsing AGB stars (Lugaro et al. 2012 ), during a He-core flash in metal-poor RGB stars (Campbell, Lugaro Karakas 2010 ), in super-AGB stars (Jones et al. 2016 ), and in Population-III massive stars when their H- and He-burning shells merge (Banerjee Qian Heger 2018 ; Clarkson, Herwig Pignatari 2019 ). In this work, we investigate the robustness of the weak i-process elemental abundances and their ratios that are proposed to contribute to the anomalous chemical composition of the star HD94028. The predicted abundances are affected by uncertainties of the (n,γ) reaction rates for the unstable isotopes involved in this nucleosynthesis. By ‘weak’ with typical neutron exposure τ ∼ 1 we mean an i process that, in spite of having a high neutron density, does not reach a large neutron exposure (time-integrated neutron flux), e.g. because it is terminated by the violent, global, non-radial instability that can be induced by H ingestion into the He-shell convection (Herwig et al. 2014b ), as probably happened in Sakurai’s object Herwig et al. ( 2011 ). As a result of this ‘weakness’ (a low neutron exposure), only the first n-capture peak elements (those with the neutron number close to the magic one N = 50) are produced in a significant amount by the weak i process. We employ the same computational methods and analysis tools that Denissenkov et al. ( 2018 ; hereafter, Paper I ) used in a similar study of i process producing the first peak elements observed in Sakurai’s object. Unlike in Paper I , we do not know the i-process site that contributed to the composition of HD94028. Therefore, we have to use a site-independent one-zone model with the temperature T , density ρ, and initial chemical composition adjusted to mimic the i-process physical conditions appropriate for this case, like it was done by Roederer et al. ( 2016 ). We also deploy constant neutron–density equilibrium models that are independent of the neutron-source reactions. 2 METHODS 2.1 The benchmark simulation For the one-zone simulations of the i-process nucleosynthesis, we use the NuGrid code ppn (Pignatari et al. 2016 ) with the same fixed temperature T = 2 × 10 8 K and density as in Paper I . The initial chemical composition is prepared using the Solar system abundances of Asplund et al. ( 2009 ) scaled to the metallicity of the star HD94028 ([Fe/H] =−1.6), assuming that the α-element abundances are enhanced with the mean value of [α/Fe] equal to +0.4. The initial mass fraction of 12 C is increased to X ( 12 C) = 0.5 to be close to the values found in He-flash convective zones. We also reduce the initial hydrogen abundance to the value of X ( 1 H) = 0.2 adjusted to guarantee that the neutron density will reach the values of N n = 10 13 –10 16 cm −3
Paper: 2021MNRAS.501.4903F (fulltext)
[Ba/Fe]=+0.13 that is significantly above the pristine [Ba/Fe] dex enrichment in Fig. 6 . This suggests that HD 26749 should have a degenerate companion and should have experienced a former epoch of wind accretion. The observable depletion of lithium qualitatively supports this scenario, but for a quantitative argument one would at least need to know the exact amount of the accreted mass. In this context, it is interesting to note that HD 26749 is currently in its subgiant stage of evolution and appears to be 1 or 2 Gyr younger than one might expect for an ancient Population II star. The excess of mass involved amounts to a few hundredths of a solar mass, a reasonable value for wind accretion, but this also depends on a parallax that would need to account for potential orbital motion. While this parallax may be possible with the Gaia DR3 release, it may be worthwhile to observe HD 26749 with high-contrast direct imaging that may disclose the degenerate candidate companion. 4.1.10 HD 34297 Two Hipparcos -based astrometric orbits by Goldin Makarov ( 2006 ) provide an orbital period of about 2 yr for this double-lined spectroscopic binary. The quantitative composite modelling of our Δ v r (Aa–Ab) =+28.9 km s −1 well-resolved 2017 November 28 spectrum of HD 34297 leads to a photometric contrast Δ m V ≃ 0.5 mag, a Population II chemistry [Fe/H]≃ −0.65 and [Fe/Mg]≃ −0.37, and effective temperatures T eff, Aa ≃ 5660 K and T eff, Ab ≃ 5440 K. The surface gravities, log g Aa ≃ 4.43 and log g Ab ≃ 4.53, follow from the Gaia DR2 parallax. Fig. 1 gives an example of the synthetic line formation for the Fe I line λ6200.321. With reference to the bolometric magnitudes M bol, Aa ≃ 4.96 and M bol, Ab ≃ 5.43 and VandenBerg et al. ( 2006 ) evolutionary tracks, the Aa primary, with a stellar radius R Aa ≃ 0.92 R ⊙ and mass M Aa ≃ 0.82 M ⊙ , is already approaching its turnoff position. The Ab secondary is still close to the main sequence at R Ab ≃ 0.80 R ⊙ and M Ab ≃ 0.79 M ⊙ . Our single-epoch radial velocities, and km s −1 , then provide a systemic velocity γ = +27.64 ± 0.27 km s −1 . 4.1.11 HD 51929 This is a known astrometric and spectroscopic binary (Wielen et al. 1999 ; Nordström et al. 2004 ; Makarov Kaplan 2005 ; Frankowski et al. 2007 ); our single-epoch radial velocity, km s −1 (2017 November 28), deviates by 3.73 km s −1 from the Nordström et al. ( 2004 ) velocity. The Ab companion was recently uncovered through direct imaging at ρ = 0.72 arcsec by Kane et al. ( 2019 ), who assess this to be a faint red dwarf, which we do confirm, since the Aa primary displays the pristine [Ba/Fe] dex enrichment (cf. Table 3 and Fig. 6 below). 4.1.12 HD 65721 For this astrometric and spectroscopic binary (Wielen et al. 1999 ; Nordström et al. 2004 ), Hipparcos -based orbital solutions were first presented in Goldin Makarov ( 2006 , 2007 ). In spite of their preliminary character with d and , the eccentricity appears indeed high, for HD 65721 presents an
Paper: 2017A&A...606A...8D (fulltext)
CO(3–2) and HCO(4–3) lines with the ASTE telescope, aiming at studying in detail the spatial distribution and physical characteristics of the molecular gas associated with the nebula. To analyze the properties of the dust in the nebula and its surroundings we use images from the ATLASGAL and Herschel databases. To account for the interaction between the ionized and molecular gas and to determine the evolutionary status of the H ii region, we use radio continuum data from the NVSS archives. This analysis also allows us to identify spots of dense molecular gas and dust where the massive star formation is developing, particularly the central region corresponding to the source IRAS 18139-1842, where star formation processes seem to be very active. In addition, we look for signatures of star formation in the surroundings using archival photometric data and we investigate a probable scenario of triggered star formation at the border of the nebula. We also study the characteristics of an identified stellar cluster projected onto the border of the bubble using new IR data. This provides some additional information that helps us to better understand the star formation activity in the nebula. N5 is an open IR bubble of about 8 in diameter centered at RA, Dec(J2000) = 181702, 184019, coincident with the H ii region Sh2-39, and located at a distance of 4.1 kpc ( Bronfman et al. 1996 ; Faúndez et al. 2004 ). In Fig. 1 we present a 8.28 m (MSX-A) emission image of the nebula, which has a typical arc-shaped morphology in the mid-IR wavelengths. At wavelengths of about 8 m most of the emission originates in strong features of PAH molecules, which are considered to be good tracers of warm UV-irradiated photodissociation regions (PDR; Hollenbach Tielens 1997 ). Since these complex molecules are destroyed inside the ionized gas of an H ii region, they indicate the limits of the ionization front and delineate the boundaries of the bubble nebula, giving a glimmer of the location of the parental molecular gas against which N5 is likely expanding. The open morphology of the IR nebula suggests that Sh2-39 is the classical H ii region density bounded toward the outer side (eastward) and ionization bounded toward the molecular cloud (westward). The nebula is believed to be linked to the IR sources IRAS 18139-1842 and IRAS 18139-1839, which both have characteristics of massive young stellar objects (see Sect. 6). The CS(2–1)line was detected by Bronfman et al. (1996) at a velocity 1 of 40.2 km s in the direction of IRAS 18139-1842. NH line and HO maser emission, which are believed to be excellent indicators of early stages of massive star formation, were also detected toward this IRAS source ( Sunada et al. 2007 ; Urquhart et al. 2011 ), while radio recombination line observations (H87 and H88) show emission at the same position at a velocity of 41 km s (Lockman, 1989 , see Fig. 1 ), which indicates the existence of high density ionized gas at the same radial velocity
Paper: 2019ApJ...877L..21Y (fulltext)
of the ejecta, the luminosity of this radioactivity-powered AIC optical emission cannot be much higher than ∼10 41 erg s −1 , even though the synthesization of 56 Ni in the ejecta is very efficient. In principle, the AIC NSs could sometimes (but may not always) be converted into magnetars, as these NSs can initially rotate very rapidly and differentially. Following this consideration, it has been suggested that the rapid spin-down of the magnetars can substantially affect the AIC optical transient, e.g., by providing an extra energy source to heat the AIC ejecta (Yu et al. 2015 , 2019 ). In more detail, the energy released from the magnetars can be initially in the form of a Poynting flux, which could eventually transform into a relativistic electron-positron wind. Then, questions arise as to what kind of emission can be produced by the pulsar wind (PW) and how energy is transferred from the PW to the AIC ejecta. It is expected that the PW emission, if detected, can help to identify AIC candidates in the future. Basically, PW emission could be mainly determined by the collisions of the PW with the AIC ejecta and, in particular, with the companion star. An illustration of these interactions is presented in Figure 1 . Considering the fact that the spin-down luminosity of a magnetar can be as high as ∼10 45 erg s −1 , the companion could be completely evaporated by the PW in a few days, because the gravitational binding energy of the companion is probably not much higher than ∼10 48 erg. The evaporated companion material (ECM) would finally form a torus due to the orbital motion, which surrounds the magnetar and blocks the PW in the orbital plane. The formation of such an ECM torus and its interaction with the PW are investigated in Section 2 . Meanwhile, the termination shock (TS) of the PW due to its interaction with the AIC ejecta is calculated in Section 3 . Both of these PW emission components can penetrate the AIC ejecta and become detectable as long as the ejecta becomes transparent for X-rays. A summary of our calculations is given in Section 4 . In Section 5 , we discuss the possible application of the AIC model in explaining the multi-wavelength emission of a recently discovered fast-evolving luminous optical transient: AT 2018cow. Figure 1. Illustration of a post-AIC system (not to scale), including the stages before and after the breakdown of an ECM torus. 2. The Interaction of a PW with a Companion 2.1. The Evaporation of a Companion Some NSs formed from AIC events are considered to rotate very rapidly and possibly differentially. In this case, the magnetic fields of the NSs can be significantly amplified through magnetorotational instability (Duncan Thompson 1992 ; Dessart et al. 2007 ; Cheng Yu 2014 ). However, in practice, it is still difficult to assign a fiducial value to the spin periods and the magnetic fields of the NSs, because these parameter values are highly dependent on the properties of progenitor WDs. In principle,
Paper: 2016A&A...588A..90L (fulltext)
resulting in the gray eclipse. Also, no significant pulsations in maximum light are present in the historical light curve of MASTER OT J095310.04+335352.8, which also speaks in favor of the cold dust disk. Of course, the edge-on disk orientation is not the only possibility to explain this system, and further study of this system will clarify this question. The spectral energy distribution (SED) of the system was very similar in the all our spectra despite the great difference in continuum fluxes. The main contributor to the SED is a red giant in eclipse as well as after eclipse, and the TiO bands of the red giant are the most prominent absorption features in the spectra. The spectrum obtained in eclipse also exhibits some faint emission lines against the TiO bands (see Sect. 2.2). We neglect the interstellar reddening because the system is above the Galactic plane and there is not any sufficient continuum distortion. The spectrum obtained in eclipse (February 10, 2013) includes the red region, allowing us to estimate the spectral type of the partially eclipsed red giant by means of [TiO] spectral index ( Kenyon Fernandez-Castro 1987 ). The star is an M1-type giant (0.5 subclass). Differences between the observed spectrum and standard SEDs of M0III and M1III do not exceed in the red part of spectrum, but for M2III the difference are about . There are not any signs of CaH absorption bands in any of our spectra, therefore, the luminosity class of the star cannot be (MS star). Another luminosity criterion is a relative intensity of FeI 5250 Å and CaI 5269 Å. According to this criterion the star is a giant (III class) rather than supergiant ( class). We note that the SED in the red region is nearly the same as that of a standard M1-type giant, but in the blue region the situation is quite different; see Figs. 2 and 5. The continuum fluxes are nearly 20% higher than the standard M1III SED. The second spectrum was obtained after eclipse and does not include the red region and, therefore, we cannot trace the emission lines evolution. The main difference between two spectra is that after eclipse the observed SED is similar to a standard M1 SED even in the blue region; see Fig. 4 . A difference between the SED observed in eclipse and a standard M1III SED is shown in Fig. 5 . A SED of the difference characterizes a sufficiently hot source because its continuum flux grows from the visual to blue region. The emission of this source was visible only during eclipse (when the contribution of the red giant to the total flux was greatly diminished). Without the eclipse, the SED is the pure contribution from the M1III-type star. The SED of the hot component corresponds poorly to a SED of any standard hot star with spectral class earlier than F5 and does not correspond to SEDs of later stars (thick line in Fig. 5 is a standard B6V SED); we tried to approximate the observed hot component spectrum by means of SEDs for MK standard stars for spectral classes