adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019MNRAS.488.2473H (fulltext)
by Werner, Rauch Kruk ( 2018 ) in the hot white dwarf in the peculiar binary nucleus of the planetary nebula EGB 6. Ar vii λ = 1065.55 spectral line is observed in ultraviolet spectra of hot PG 1159 and extremely hot helium rich white dwarfs (Werner, Rauch Kruk 2016 , 2017 ). Ar iii and Ar ii spectral lines are used in the determination of argon abundance in B-type star and subdwarf atmospheres (Heber Edelmann 2004 ; Blanchette et al. 2008 ; Lanz et al. 2008 ; Fossati et al. 2009 ). In our recent works, Hamdi et al. ( 2017 ) and Hamdi et al. ( 2018 ), we have reported Stark widths for 300 spectral lines of singly charged argon (Ar ii ). Our Stark widths were compared with experimental results and good agreements were found for most spectral lines. In order to complete the set of Stark broadening parameters of Ar ii spectral lines, we report here Stark shifts. Stark shifts are calculated using semiclassical perturbation approach (SCP) (Sahal-Bréchot 1969a , b ) and a set of atomic data is calculated using Hartree–Fock method with relativistic corrections (Cowan 1981 ). Our results are compared with experimental results for 100 spectral lines. Shifts by collisions with ionic perturbers are also taken into account. We also present results as a function of temperature for 161 spectral lines for collisions with electrons, protons, singly charged helium, and singly charged argon. In previous works (Hamdi et al. 2011 , 2013 , 2014 ), the same method was used and we have determined Stark broadening parameters for 114 spectral lines of Pb IV and 32 spectral lines of Ar iii . 2 THE IMPACT SEMICLASSICAL PERTURBATION METHOD A detailed description of this formalism with all the innovations is given in Sahal-Bréchot ( 1969a , b , 1974 , 1991 ), Fleurier, Sahal-Bréchot Chapelle ( 1977 ), Dimitrijević, Sahal-Bréchot Bommier ( 1991 ), Dimitrijević Sahal-Bréchot ( 1996 ), Sahal-Bréchot, Dimitrijević Ben Nessib ( 2014 ), and earlier papers. The profile F (ω) is Lorentzian for isolated lines: where i and f denote the initial and final states and E i and E f their corresponding energies and w is the half width. The shift ( d ) (in angular frequency units) of an electron-impact shifted spectral line can be expressed as: where N is the electron density, the Maxwellian velocity distribution function for electrons, denotes the impact parameter of the incoming electron. The phase shift due to the polarization potential ( r −4 ), is given in section 3 of chapter 2 in Sahal-Bréchot ( 1969a ) and R D is the Debye radius. The cut-off R 3 is described in section 1 of chapter 3 in Sahal-Bréchot ( 1969b ). Energy levels and oscillator strengths needed for SCP calculations of Stark shifts are obtained using Hartree–Fock method with relativistic corrections (HFR) (Cowan 1981 ). We use as for the width calculations an atomic model including 24 configurations: 3s 2 3p 5 ; 3s 2 3p 4 nl ( nl = 4p, 4f, 5p, 5f, 6p, 6f, 6h, 7p, 7f, 7h) (odd parity) and 3s 3p 6 ; 3s 2 3p 4 n ′ l ′ ( n
Paper: 2019ApJ...873...27L (acknowledgments)
The authors would like to acknowledge Rainer Beck for helpful discussions. J.T.L. and J.N.B. acknowledge the financial support from NASA through grants NNX15AM93G, NAS2-97001, 80NSSC18K0536, and GO7-18073X.
Paper: 2020ApJ...905...32B (fulltext)
the two objects could also exhaust helium shell burning and simply cool into double degenerates. Upon merger, they would likely form a rapidly rotating CO WD, potentially preceded by an R CrB phase if substantial He remains, as expected for hybrid WDs (Perets et al. 2019 ). For further details on these systems, please see Kupfer et al. ( 2020a , 2020b ). 6. Conclusion Using data from the ZTF, we have significantly increased the number of known binary systems with an orbital period of less than one hour. Using high-speed photometric follow-up in combination with spectroscopic follow-up, we have characterized the physical parameters of the systems using model-dependent techniques invoking WD mass–radius relations. In future work, we hope to obtain more model-independent measurements such as orbital decay rates to more robustly characterize the physical parameters of these systems. We will continue to analyze ZTF data as it accumulates more epochs and discovers more such objects. At the current stage of the survey, the sample we have discovered exhibits a remarkably high fraction of pairs of He WDs, likely due to selection effects. As we continue the survey, we anticipate discovering more systems and that our discoveries will begin to include other classes of sources, such as pairs of CO WDs. All of the detached systems in our sample should be undergoing rapid orbital decay due to general relativity, and we expect to detect this orbital decay in every detached system with high significance before LISA launches. The algorithms and analysis techniques we are developing will be applicable to upcoming surveys such as the one that will be conducted by the Vera Rubin Observatory (LSST Science Collaboration et al. 2009 ). These data sets will be accompanied by their own challenges (in particular, lower cadence, split across many filters). Given the clear promise demonstrated by the discoveries emerging from ZTF, we are optimistic that optical time-domain surveys will continue to revolutionize our understanding of binaries in the millihertz regime. These surveys will set the stage for LISA, having discovered dozens, if not hundreds, of sources by the time the instrument begins to operate (Korol et al. 2017 ). In future work, we will characterize our recovery efficiency in greater detail using synthetic lightcurves, in order to better characterize ZTF’s sensitivity to detecting the population of DWDs in the galaxy. Efforts are currently underway to increase sensitivity to the shortest-period binary systems by accounting for their orbital period evolution, which requires algorithms that can accommodate acceleration searches (Katz et al. 2020 ). Binaries with compact object components are transitioning into a golden age, as we enter an era when time-domain surveys such as ZTF enable detection via periodic features. Massively multiplexed spectroscopic surveys acquiring multiple epochs such as SDSS-V (Kollmeier et al. 2017 ) and LAMOST (Cui et al. 2012 )
Paper: 2017A&A...597A..20O (acknowledgments)
We thank the ESO Paranal team for supporting our SPHERE and AMBER observations. We are also grateful to the referee Susanne Höfner for her valuable comments, particularly on the interpretation of the extended H emission. This research made use of the SIMBAD database, operated at the CDS, Strasbourg, France. We acknowledge with thanks the variable star observations from the AAVSO International Database contributed by observers worldwide and used in this research. K.O. acknowledges a grant from Universidad Católical del Norte.
Paper: 2019ApJ...879...22D (fulltext)
source at its given velocity (within 15 km s −1 ), position (within 5 arcmin), and derived a velocity width and size within our cuts used in Section 2.2 . We recovered every injected source at high velocities, where Galactic emission is the weakest, and nearly every source at low velocities and Leo T velocities (97% and 82%, respectively), the latter of which is lower because Leo T’s velocity is often within the region of average Galactic emission >1 K which we purposely ignored. Nevertheless, we can confidently say that SoFiA would successfully detect Leo T at its actual distance in the vast majority of our search area. We next wanted to determine how accurate our detection limit of Leo T-like objects was, based on our observed flux distribution. To do this, we performed the same procedure as before with two alterations; we scaled the peak flux of the injected source to a larger distance (from 0.5 to 1 Mpc), and we reduced the size of the injected source to better correspond to a more distant galaxy. At high velocities, the detection fraction dropped to its lowest value of 86% at 1 Mpc, but the difference at low and Leo T velocities is much more drastic, as shown in the first three rows of Table 3 (with Mass = 1 Leo T). Note, however, that a typical local group galaxy that is 1 Mpc away has a velocity closer to our high-velocity case than either of the lower-velocity cases. In other words, the region of distance-velocity space where our detection fraction is lowest also probably contains the smallest number of dwarf galaxies. We also tried reducing the injected sources’ size, which results in a ∼10% drop in the detection fraction at lower velocities compared to the same test without changing the size, meaning the dominant reason for the overall drop is the reduction in flux caused by increasing distance. In the bottom three rows of Table 3 we inject sources with successively larger multiples of Leo T’s mass to determine how complete we are out to 1 Mpc. We note that we would detect ≈50% of galaxies at lower velocities with only 2 times Leo T’s mass, and ≈66% of those with 3 times Leo T’s mass. 6 6 At 4 times Leo T’s mass we naturally recover the same detection fractions as the first row of Table 3 , because multiplying the mass of the third row by 4 is equivalent to halving the third row’s distance in terms of H i flux. Table 3 Detection Fraction of Sources Injected at or beyond a Distance of 500 Note. The H i mass of Leo T is . The last row of 10 times Leo T’s mass is an exception to the previous rows in that the size and w 50 cuts we apply are the dominant cause for reduction in recovered sources in regions significantly contaminated by Galactic emission at low velocities. Without the size and w 50 cuts, the detection fractions are all >92%. The higher level of background noise at low latitude causes SoFiA to chop off the outer regions of the object in its size calculation. Therefore our size cut may lead us to miss some dwarf galaxies
Paper: 2020MNRAS.494.1971C (fulltext)
background magnetic field than those observed in the synthetic observations (slight preference of perpendicular alignment; Fig. 6 ), which were generated from a trans- to super-Alfvénic simulation. Combined with other polarimetric properties, we suggest that the relative alignment between dense cores and pc-scale magnetic fields depends on the magnetic properties of the cloud, and could potentially provide an alternative approach than the traditional Zeeman observations on investigating the magnetization level of the cloud. APPENDIX A: CORES IDENTIFIED IN GAS DATA AND THEIR RELATIVE ALIGNMENTS Figs A1 – A3 show all subregions observed by GAS overlapped with the core boundaries identified by astrodendro . The relative alignment between cores and background magnetic field (i.e. the angle between core major axis and background magnetic field direction) in each subregion is plotted in Fig. A4 . Figure A1. Cores (red contours) identified in each individual region in the Ophiuchus cloud using NH 3 integrated emission (colour maps, in K km s −1 ) observed by GAS. Note that cores that are more rounded (with minor/major ratio >0.8) are removed from the analysis and not shown here. Dendrogram leaves without obvious central peak emission are also excluded. Figure A2. Same as Fig. A1 , but for the Perseus cloud. Figure A3. Same as Fig. A1 , but for the Taurus cloud. Figure A4. The cumulative distributions of the core–background magnetic field alignment for the subregions within the three clouds: Ophiuchus (left), Perseus (middle), and Taurus (right). APPENDIX B: RELATIVE ANGLES IN 2D AND 3D Here, we briefly describe the difference between measuring relative angle between two random vectors in 3D and 2D spaces. Without loss of generality, we assume both vectors are unit vectors and one is pointing straight up (+ z for 3D case and + y for 2D case). The probability for the second vector to have an angle θ with respect to the first one is therefore proportional to the surface area of the stripe of constant latitude θ on the 3D sphere (), or the arc length at θ of the unit circle ( R dθ). For unit vectors, R = 1, and since we only consider angles within [0, 90] deg, the cumulative probability of the relative angle to be smaller than θ is therefore Note that the probability is normalized by the total surface area of the sphere (in 3D) or the circumference of the whole circle (in 2D). This explains why a set of 3D random vectors will show a linear distribution in cos α instead of α, where α is the relative angle between any two vectors from the set. Equations ( B1 ) and ( B2 ) are illustrated in Fig. B1 as black dotted lines, which also shows that two vectors in 3D are more likely to be perpendicular than parallel if measured in degrees. Figure B1. Comparing the relative angles in 3D (top) and 2D (bottom), for both analytic solutions of equations ( B1 ) and ( B2 ) (black dotted lines) and between two sets of random vectors (solid lines), showing the angle in
Paper: 2017ApJ...835...59W (acknowledgments)
We thank the anonymous referee for constructive criticism, particularly regarding estimates of the gas density from the observations. This work was supported by NASA Origins of Solar Systems grant NNX10AH37G.
Paper: 2016ApJ...833..291A (acknowledgments)
We thank the staff at the Nobeyama Radio Observatory and the National Radio Astronomy Observatory for help with the observations. We thank Dr. David Frayer for his help in the observations at the GBT. M.A. thanks the Grant-in-Aid for Scientific Research on Innovative Areas (Grant No. 25108002), Grant-in-Aid for Scientific Research (C) (Grant No. 15K05395), and the Institute for Quantum Chemical Exploration.
Paper: 2020ApJS..250...16A (fulltext)
H i Ly α (lower left), low-TZ Si iii 1206 Å (upper left), low-TZ C iii 1175 Å (lower right), and coronal soft X-rays (upper right). These flux–flux diagrams are more populated given the better FUV coverage of the two stars, although the C iii comparison for B is still missing points for epochs when E140H exposures were taken: these cut off just below Si iii 1206 Å, so C iii 1175 Å was excluded. As tentatively seen in the Mg ii panels previously, the A correlations line up nicely with the solar equivalents, although A’s X-rays still are shifted downward, and noisier than the others perhaps partly because of non-simultaneous coverage. Again, there is somewhat disjoint behavior in several of the panels between the A and B relations. Notice also that A’s Ly α slope closely matches that of the solar correlation, but B’s trend is somewhat shallower. Similarly, the Si iii correlation against C ii is noticeably shallower for B, and curiously shifted downward, compared to A or the Sun. At the same time, the C iii behavior of B appears to be steeper. Finally, A’s X-rays (though scattered) seem to show a steeper rise with increasing C ii , as in the lower branch of the Sun’s bent power law, whereas B displays a flatter relation (although still with α ∼ 2). Figure 16 (c) displays correlations between low-TZ Si iii 1206 Å and low-TZ C iii 1175 Å (lower left), mid-TZ Si iv 1400 Å (upper left), mid-TZ C iv 1550 Å (lower right), and mid-TZ N v 1240 Å (upper right). Again, the C iii comparison is lacking a few points for α Cen B, most conspicuously on the low-flux side. Despite that, all four panels have the same general appearance: the A points closely track the translated solar power laws, while B shows universally steeper slopes. Curiously, the B trends have “retreated” backward to the extent that the A and B origins of the four correlations are similar (again taking into account the abbreviated C iii series for B). This is a consequence of the behavior seen in the upper left panel of Figure 16 (b), wherein B’s Si iii points are seen to parallel those of A, but also approach a similar minimum. This behavior is unlike the other diagnostics shown earlier, for which the minimum B value generally was comparable to, or exceeded, that of A’s maximum. The contrary behavior emphasizes the unique position of Si iii 1206 Å in both the Sun and α Cen A, as a “super-active” species, showing the steepest power laws relative to the other FUV features, even including hot N v . In α Cen B, in contrast, the Si iii feature has declined in strength, and shows the expected “stellar” behavior, namely Si iv , C iv , and especially N v all exhibit power laws comparable or steeper than unity against Si iii . The final figure, Figure 16 (d), illustrates a number of miscellaneous, but nonetheless illuminating, comparisons. The lower left panel compares the two Mg ii resonance lines, which are both affected by ISM absorptions, but the k-line more so than the weaker
Paper: 2015ApJ...801..104H (acknowledgments)
We thank Joseph Trollo for assisting with observations and Dianne Harmer for invaluable technical support at Kitt Peak. This research was supported by NASA through the Spitzer Exploration Science Program Weather on Other Worlds (program GO 80179) and ADAP award NNX11AB18G. This publication makes use of the SIMBAD online database, operated at CDS, Strasbourg, France, and the VizieR online database (see Ochsenbein et al. 2000 ). This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. This research has benefited from the M, L, T, and Y dwarf compendium housed at DwarfArchives.org. We have also made extensive use of information and code from Press et al. ( 1992 ). We have used digitized images from the Palomar Sky Survey (available from http://stdatu.stsci.edu/cgi-bin/dss_form ), which were produced at the Space Telescope Science Institute under U.S. Government grant NAG W-2166. The images of these surveys are based on photographic data obtained using the Oschin Schmidt Telescope on Palomar Mountain and the UK Schmidt Telescope. Facility : KPNO 2.1 m
Paper: 2019A&A...629A..49O (acknowledgments)
GMO acknowledges support of the Research Foundation – Flanders under contract G075916N and under grant number V434818N. HvW acknowledges support from the Research Council of the KU Leuven under grant number C14/17/082. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France. This research has made use of NASA’s Astrophysics Data System Bibliographic Services. We thank the anonymous referee for the useful comments that improved the manuscript.
Paper: 2018ApJ...856...48C (acknowledgments)
We thank the referee Durgesh Tripathi for useful suggestions. We also thank Linhua Deng for providing us with the NVST data, Rui Liu for helpful discussions, and the ISSI workshop on “Decoding the Pre-Eruptive Magnetic Configuration of Coronal Mass Ejections” led by S. Patsourakos and A. Vourlidas. SDO is a mission of NASAs Living With a Star Program. X.C. and M.D.D. are supported by NSFC under grants 11722325, 11733003, and 11790303 and by Jiangsu NSF under grant BK20170011. X.C. is also supported by the “Dengfeng B” program of Nanjing University. B.K. acknowledges support by NASA under Grants NNX16AH87G and 80NSSC17K0016, the NSFC, and the DFG.
Paper: 2017ApJ...850...43R (acknowledgments)
The authors acknowledge Chuck Horowitz, Evan O’Connor, and Todd Thompson for productive conversations, concerning insight into and help with the microphysics, and Ernazar Abdikamalov, Sean M. Couch, Luke F. Roberts, and Christian D. Ott for fruitful discussions on the nature of core-collapse supernovae. Support was provided by the Max-Planck/Princeton Center (MPPC) for Plasma Physics (NSF PHY-1523261). D.R. gratefully acknowledges support from the Schmidt Fellowship. A.B. acknowledges support from the NSF under award number AST-1714267. J.D. acknowledges support from a Laboratory Directed Research and Development Early Career Research award at LANL. The authors employed computational resources provided by the TIGRESS high performance computer center at Princeton University, which is jointly supported by the Princeton Institute for Computational Science and Engineering (PICSciE) and the Princeton University Office of Information Technology and by the National Energy Research Scientific Computing Center (NERSC), which is supported by the Office of Science of the US Department of Energy (DOE) under contract DE-AC03-76SF00098. The authors express their gratitude to Ted Barnes of the DOE Office of Nuclear Physics for facilitating their use of NERSC. This paper has been assigned a LANL preprint # LA-UR-17-20973.
Paper: 2017MNRAS.472.3256L (acknowledgments)
Acknowledgements We thank the referee for many useful comments and suggestions that helped to improve the paper. We would like to thank all those who are, or have been, associated with BiSON. BiSON is funded by the Science and Technology Facilities Council (STFC). MNL acknowledges the support of The Danish Council for Independent Research | Natural Science (Grant DFF-4181-00415). WJC, SJH, GRD, YPE and RH acknowledge the support of the UK Science and Technology Facilities Council (STFC). Funding for the Stellar Astrophysics Centre (SAC) is provided by The Danish National Research Foundation (Grant DNRF106).
Paper: 2020MNRAS.494.5682S (fulltext)
of ε Cet are mostly not able to be individually separated, and the use of assumptions are required to untangle them (Docobo Andrade 2013 ). The angular separation of the components was 0.1 arcsec in 2015 (Mason et al. 2001 ). We obtained sixteen polarized spectra for ε Cet over a 6.0 h interval on 2016 September 21, corresponding to 2.5 per cent of the estimated rotational period of 10.2 d. The mean LSD profile (Fig. B4 ) combined 6693 lines, and a magnetic detection was obtained from the mean LSD Stokes V profile. We are unable to determine whether the magnetic detection was on the F5V or F6V components because the two stars could not be separated spectroscopically. We estimated B l at +0.6 ± 0.1 G. Our radial velocity determination +15.81 ± 0.26 km s −1 (Table 5 ) is similar to the value of +15.49 ± 0.19 km s −1 in Anderson Francis ( 2012 ), and our S-index 0.186 ± 0.003 (Table 5 ) was less than 0.219 in Pace ( 2013 ). A4 α Com α Com (HIP 64241) is a triple star consisting of a double-lined spectroscopic binary (Duquennoy, Mayor Halbwachs 1991 ) of F5V and F6V components (Edwards 1976 ; Soubiran et al. 2018 ), and a V magnitude 10.2 component with an angular separation of 85 arcsec in 2014 based on data from the continuously updated Washington Double Star Catalogue (WDS) (Mason et al. 2001 ). The system was first observed by Struve ( 1875 ) who reported that occultations were observed (Muterspaugh et al. 2015 ). Roman ( 1950 ) describes α Com as a close visual binary, with equal F5 components, and showing weak lines. Hartkopf et al. ( 1989 ) describes α Com as a ‘lovely edge-on system’ with an orbital inclination close to 90 deg, and suggested a partial eclipse was possibly missed in early 1990. The eclipsing binary has the third longest period known ( P = 25.85 yr) (Lipunov et al. 2016 ), with only eight eclipsing stars having periods longer than 10 yr. We obtained twenty-seven polarized spectra for α Com over a 3.8 h interval on 2017 February 16, corresponding to 5.3 per cent of the estimated rotational period of 3.0 d. The mean LSD profile (Fig. B5 ) combined 7628 lines, and a magnetic detection is visible in the Stokes V profile. We are unable to determine whether the magnetic detection was on the F5V or F6V components because the two stars could not be separated spectroscopically. We estimated B l at +3.1 ± 0.3 G. Our radial velocity determination −18.02 ± 0.27 km s −1 (Table 5 ) is larger the value of −16.05 ± 0.21 km s −1 in Anderson Francis ( 2012 ), and our S-index 0.272 ± 0.008 (Table 5 ) is smaller than the 0.308 in Pace ( 2013 ). A5 ω Dra ω Dra (HIP 86201) is a triple star system consisting of double-lined spectroscopic binary (Abt Levy 1976 ; Konacki et al. 2010 ), and a V magnitude 13.2 common proper motion companion with an angular separation of 70 arcsec (Abt Levy 1976 ; Fekel, Tomkin Williamson 2009 ). The mases of the primary component are 1.46 M ⊙ and secondary 1.18 M ⊙ (Konacki et al. 2010 ), which corresponds to