adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016ApJ...824...67Y (fulltext)
( 1969 ) and Yoshimura ( 1975a , 1975b ) succeeded in numerically reproducing the latitudinal evolution of the toroidal magnetic field (butterfly diagram). In order to obtain the equatorward migration of the toroidal magnetic field, it is required that the angular velocity should decrease as the radius increases in the solar convection zone ( for the positive α in the northern hemisphere; : angular velocity). However, the developments of helioseismology revealed that the real radial distribution of the angular velocity in the solar convective zone does not satisfy this requirement. In addition to this point, Parker himself discussed the limitation of the α dynamo in the solar convection zone from the viewpoint of magnetic-field buoyancy (Parker 1975 ). Once the magnitude of a magnetic field has been amplified to some level, the magnetic flux tube starts moving upward from deeper to shallower regions. Since this buoyancy convection timescale, even for a very weak magnetic flux tube, is estimated to be very short as compared with the α -dynamo amplification timescale, it reaches the solar surface without being sufficiently amplified by the dynamo. This is called the magnetic buoyancy dilemma intrinsic to the magnetic flux tube. These difficulties have been considered to be obviated by assuming that the dynamo action is operated in a convectively stable thin layer between the radiative and convection zones (Parker 1975 ). This region is associated with a large differential rotation and is called the tachocline . In the tachocline dynamo picture, the solar magnetic field is generated by a strong differential rotation there and may have enough time to be amplified by the α dynamo before it is raised by the magnetic buoyancy (Spiegel Weiss 1980 ; van Ballegooijen 1982 ). As mentioned above, depending on the location of the field generation, the dynamo models are divided into two types: (i) the flux-transport dynamo, which operates in the overshoot layer below the bottom of the convection zone; and (ii) the mean-field distributed dynamo, which works in the bulk of the convection zone, in particular at the subsurface shear layer. On the surface of the Sun, large-scale meridional circulations with a speed of ∼15 m s −1 in the poleward direction have been observed. The combination of the meridional circulation and the α –Ω mean-field model demonstrated its ability to reproduce the basic features of the solar magnetic activity cycle such as the equatorward migration of the sunspots, the cycle period, and the phase difference between the poloidal and toroidal magnetic fields. In the combination with meridional circulations or convective-diffusive motions such as the Babcock–Leighton-type convection, which produces the east–west tilt of the sunspot pairs under the action of the Coriolis force on the magnetic flux tube during its rise through the convection zone, it is expected that the tachocline dynamo can reproduce several properties of the
Paper: 2018ApJ...854...93H (acknowledgments)
We thank the referee and B. Czerny, C. Done, C. Knigge, J. Matthews, H. Netzer, and D. Starkey for discussions. P.H. and G.S. acknowledge the support of the Natural Sciences and Engineering Research Council of Canada (NSERC), funding reference number 2017-05983. K.H. acknowledges support from STFC grant ST/M001296/1. Software: MATLAB.
Paper: 2021MNRAS.504.4768Z (fulltext)
motion in declination direction. The proper motion of SV Cru, (μ αcos δ, Cep , μ δ, Cep ) = (− 6.43 ± 0.04, 1.31 ± 0.05) mas yr −1 , deviates 10.9σ in declination direction from the proper motion of Ruprecht 97, (μ αcos δ, cl , μ δ, cl ) = (− 6.49 ± 0.07, 0.44 ± 0.07). The mean proper motion is estimated by 23 OC members and the deviation of proper motion may be caused by the strict selection of OC members. Considering the good matching of other constraints, we regard SV Cru as an OC Cepheid. The distance and age of Ruprecht 97 are DM = 12.95 ± 0.53 mag and log ( t yr −1 ) = 8.13 ± 0.21, which are consistent with Cantat-Gaudin et al. ( 2020 ). 3.1.10 UBC 231 and WX Pup WX Pup is located at about 25.9 arcmin from the centre of UBC 231. Its proper motion is 2.2σ from the mean proper motion of UBC 231. The parameters of UBC 231 DM = 12.31 ± 0.17 mag and log ( t yr −1 ) = 7.87 ± 0.30 are in agreement with the reference values DM = 12.31 mag and log ( t yr −1 ) = 8.07, given by Cantat-Gaudin et al. ( 2020 ). 3.1.11 UBC 375 and V438 Cyg V438 Cyg is located at approximately 7.0 arcmin from the centre of UBC 375. V438 Cyg satisfies the constraints of the spatial position and parallax. The proper motion of V438 Cyg is (μ αcos δ, Cep , μ δ, Cep ) = (− 3.28 ± 0.07, −4.27 ± 0.07) mas yr −1 , which is 6.4σ away from the mean proper motion of UBC 375 (μ αcos δ, cl , μ δ, cl ) = (− 3.42 ± 0.07, −5.22 ± 0.11) mas yr −1 . The determined OC’s parameters are DM = 11.08 ± 0.21 mag and log ( t yr −1 ) = 8.03 ± 0.30, which are consistent with Cantat-Gaudin et al. ( 2020 ). Since only 32 OC members are adopted to estimate the mean proper motion, the deviation of proper motion may be attributed to the strict selection of OC members. Considering the good matching of other constraints, we still regard V438 Cyg as an OC Cepheid. 3.1.12 FSR 0158 and GQ Vul GQ Vul is located at approximately 2.4 arcmin from the centre of FSR 0158. GQ Vul agrees well with FSR 0158 in the proper motion (0.6σ) and parallax (0.1σ). This new OC Cepheid is ∼5.3 kpc from us and almost at the detection limit of our method based on proper motions and parallaxes. The far away OC Cepheids in the Milky Way’s disk are helpful to testify the probable bias in Cepheid’s PAR. 3.2 Issue about the known OC Cepheid In this subsection, we discuss known OC Cepheids not included in our catalogue. We compared with Turner ( 2010 ) who established an OC Cepheid catalogue based on optical data and found that 11 of the 24 OC Cepheids in that work are in our list. Nine of the other 13 Cepheids are in small clusters or star associations and they are not included in Cantat-Gaudin’s membership analysis catalogue. Three OC Cepheids have proper-motion differences larger than 10σ. The remaining OC Cepheid is BB Sgr in Collinder 394. 3.2.1 Collinder 394 and BB Sgr BB Sgr was treat as an OC Cepheid for a long time from Turner Pedreros ( 1985 ). However, based on Gaia , the membership is questionable. The proper motion of BB
Paper: 2019ApJ...887L..28H (fulltext)
of [2] and [3], which show FERMI spectra because they provided better statistics. We note that the corresponding RHESSI spectra at those times showed very similar results. The X-ray spectra of background-subtracted data (black solid lines) are plotted together with the fitting results for the isothermal component (red dashed lines), and the nonthermal component (blue dashed lines). The background is represented by the gray solid line. The electron temperature, T , emission measure, EM, electron distribution index, δ , and cutoff energy, E C , as well as the chi-squared of the fitting, χ 2 , are listed. Figure 4. (a)–(d) EM-weighted temperature maps during the M1.3 flare. The dotted lines represent an artificial slit along the trajectory of the jet-like feature, along which we extracted the DEM evolution. (e) Time–distance plot for EM-weighted temperatures along the slit. (f) EM distribution for the position and times marked by the colored circles in panel (e). An animation of panels (a)–(d) is available. The video begins at 21:30:01 and ends at 22:14:49. The realtime duration of the animation is 15 s. (An animation of this figure is available.) In order to infer information on the triggering of the event and its magnetic topology, we initially studied the magnetic field of AR 11944 during the flare under study as well as a few days before. We note that due to the rapidly varying photospheric magnetic field of the AR prior to the occurrence of SOL2014-01-13T21:51M1.3, its closeness to the limb, and the small-scale nature of the event, the magnetograms did not provide us with useful information for the magnetic configuration on the date of the flare occurrence. 3.3. Causal Relationship of Precursors and Flare Features In addition to the bright loop and associated chromospheric brightenings observed for P2 and P3, the impulsive flare phase starts with the appearance of an additional loop (marked by the orange dashed line in the inset of Figure 2 (d)) rooted at a chromospheric brightening (labeled (3)). This was immediately followed by the generation of flare ribbons neighboring L1 and the enhanced emission at L2 as a consequence of nonthermal bremsstrahlung (Figure 2 (e)). No X-ray sources were found at the locations where the ribbons formed encompassing L1 (Figures 3 (d), (e)), possibly due to RHESSI ’s limited dynamic range. However, strong nonthermal emission during the impulsive flare phase was registered at L2, in the form of a compact hard x-ray (HXR) source and a flat nonthermal ( δ ∼ 4.2) component in the corresponding spectrum (Figure 3 (d)). This shows that L2 was the main compact flare region. During this phase, a jet-like feature emerged from L2 (see inset in Figure 2 (e) and Figures 4 (d), (e)), reaching its maximum altitude at . The EM distribution of a position along the jet-like feature before, during and after its occurrence is shown in Figure 4 (f) (see position and corresponding times marked in Figure 4 (e)), showing that it
Paper: 2015ApJ...801...11H (acknowledgments)
This work was supported in part by JSPS KAKENHI Grant Number 24540226.
Paper: 2016A&A...594A.114C (acknowledgments)
F.C. acknowledges support from Swedish National Science Council grant 637-2013-7261. A.L. acknowledges support from the ERC Advanced Grant “ FEEDBACK ”. R.H.I., M.A.P.T. and A.A. acknowledge support from the Spanish MINECO through grants AYA2012-38491-C02-02 and AYA2015-63939-C2-1-P.
Paper: 2016MNRAS.459.3287C (acknowledgments)
We thank the anonymous referee for many comments and suggestions that improved the manuscript. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. This research has also made use of NASA's Astrophysics Data System (ADS), and the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. This research made use of Montage, funded by the National Aeronautics and Space Administration's Earth Science Technology Office, Computation Technologies Project, under Cooperative Agreement Number NCC5-626 between NASA and the California Institute of Technology. Montage is maintained by the NASA/IPAC Infrared Science Archive. This publication makes use of data products from the WISE , which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics and Space Administration. The UKIRT is supported by NASA and operated under an agreement among the University of Hawaii, the University of Arizona, and Lockheed Martin Advanced Technology Center; operations are enabled through the cooperation of the Joint Astronomy Centre of the Science and Technology Facilities Council of the UK. When some of the data reported here were acquired, reported here were acquired, UKIRT was operated by the Joint Astronomy Centre on behalf of the Science and Technology Facilities Council of the UK. Some of the data presented in this paper were obtained from the MAST. STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support for MAST for non- HST data is provided by the NASA Office of Space Science via grant NNX09AF08G and by other grants and contracts. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the US Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS website is http://www.sdss.org/ . RC was supported by NASA Contract 1391817 and NASA ADAP grant 13-ADAP13-0155. MSY acknowledges support from the NASA ADAP grant NNX10AD64G.
Paper: 2015ApJ...808...40W (acknowledgments)
We thank Andrey Kravtsov, Surhud More, Phil Hopkins, Frank van den Bosch, Andrew Hearin, Erik Tollerud, Erwin Lau, and Benedikt Diemer for useful discussions and/or comments on an early draft. We also thank the reviewer for useful comments. A.R.W. acknowledges the hospitality and stimulating environment of the Aspen Center for Physics, supported in part by the NSF. A.R.W. also gratefully acknowledges support from the Moore Center for Theoretical Cosmology and Physics at Caltech. This work was supported in part by NSF grants AST-1412768 1009811, NASA ATP grant NNX11AE07G, NASA Chandra grants GO213004B and TM4-15007X, the Research Corporation, and by the facilities and staff of the Yale University Faculty of Arts and Sciences High Performance Computing Center.
Paper: 2018A&A...609A.132C (fulltext)
GASP2D photometric decomposition of the mock images of the simulated galaxies is shown in Fig. 3 for the lenticular remnant resulting from the merger experiment gS0dE0o100 seen at an inclination . We listed the more relevant best-fitting structural parameters (i.e., the effective radius , Sérsic index , and axial ratio of the bulge, the scale-length and axial ratio of the disk, and the difference between the position angles of bulge and disk ) of the mock images of the simulated lenticular remnants seen at different inclinations in Table A.1 for the giant-giant galaxy mergers, Table A.2 for the giant-dwarf galaxy mergers, and Table A.3 for the giant S0-dwarf E0 galaxy mergers with different orbital parameters. Following Méndez-Abreu et al. (2017) , we adopted and as uncertainties on the axial ratio and position angle of both bulge and disk, respectively. 4.3. Bulge intrinsic shape of the simulated lenticular remnants seen at different inclinations We made use of our statistical method to retrieve from the mock images the probability distribution of the intrinsic axial ratios and of the bulges of the simulated lenticular remnants seen at different inclinations (, and ). We excluded from the analysis the edge-on configurations () because they do not allow us to constrain and due to the unknown orientation of the triaxial bulge in the disk plane. The probability distribution of and for the bulge of the simulated lenticular remnants resulting from the merger experiment gS0dE0o100 and seen at an inclination is shown as an example in Fig. 4 . The face-on configurations of the simulated lenticular remnants ( and ) provided the same result in terms of the probability distribution of and for all the simulated bulges, as expected if no observational and theoretical bias affected the adopted statistical method. Furthermore, we confirmed that the tightest constraints for are given when galaxies are seen face on, whereas remains unconstrained for these galaxies. We also found consistent probability distributions of and for the same simulated lenticular remnant seen at intermediate inclinations (, , and ). Because of this, we were confident of having correctly recovered the bulge intrinsic shape and suggested us to set a limit on the inclination of real galaxies () to robustly apply our statistical method to their bulges (see Sect. 2 ). In general, the probability distribution of and is tighter at with respect to or . Therefore, we considered this inclination as the ideal viewing angle for future analyses of the bulge intrinsic shape in real galaxies. The probability distributions of and of the bulge of the simulated lenticular remnant resulting from the merger experiment gS0dE0o100 seen at different inclinations are shown in Fig. 5 , while the remaining galaxies are in Fig. B.1 . The values of and derived for all the simulated lenticular remnants seen at different inclinations are listed in Table 1 . We realized that the accuracy of the photometric
Paper: 2017A&A...597A..87S (acknowledgments)
This paper is dedicated to the memory of Janusz Kaluzny, founder and long-time leader of CASE, who prematurely passed away in March 2015. We thank an anonymous referee for very insightful remarks that significantly improved the final shape of the paper. A.A.P. acknowledges partial support from the Polish National Science Center (NCN) grant 2015/17/B/ST9/02082. M.R. was supported by the NCN grant DEC-2012/05/B/ST9/03931.
Paper: 2019ApJ...879...85V (fulltext)
caused by the erupting filament plasma and the bright ribbon fronts on the outer border of the growing dimming region that come into the pixel’s FOV. Figure 7. Time evolution of the relative density changes, i.e., actual density divided by its pre-event value, derived from AIA DEM maps for pixels corresponding to the EIS spectra in Figure 6 . The segments in red indicate the phase of the impulsive density drop: its start is defined by the data steeply decreasing below 1, and its end is calculated from the time derivative of the curve (defined by the first positive increase). The green vertical lines indicate the time at which the corresponding pixel was covered by the EIS slit. To relate the evolution of the densities reconstructed from the AIA DEM maps (Figure 7 ) in the 4 pixels indicated in Figure 5 with the corresponding spectra and flows determined from the EIS scan (Figure 6 ), we need to relate the dynamics observed in AIA with the time at which the EIS slit covered the particular pixels. We first discuss the behavior of the top three pixels plotted in Figure 7 , and thereafter the bottom-most one, which is related to different structures as we will show. In the density evolution retrieved from the DEM maps (Figure 7 ), we can see that the sharp density drop due to the dimming occurs first in the leftmost (easternmost) pixel, and thereafter in the two neighboring pixels to the right. This is consistent with the expansion of the dimming region toward the west, i.e., from left to right (see Figure 8 and the associated movie). On the other hand, EIS was in scanning mode, and we may expect to see the largest upflow components in those pixels that have been activated as dimming pixels (i.e., the associated fields were “opened”) closest in time to when the EIS slit was mapping them. Figure 8. Evolution of the western dimming region in SDO /AIA 211 Å filtergrams. We show the same zoomed FOV as for the EIS maps in Figure 5 . The yellow boxes mark pixels for which we show the EIS spectral profiles in Figure 6 . The still figure shows the event at two time steps (annotated in each panel). The animated figure shows the event evolution from 2012 September 27, 23:50 UT to September 30, 01:16 UT. (An animation of this figure is available.) The EIS scan was from west to east, and the strongest upflow component is observed at the pixel ( x = 563′′, y = 44′′), which was covered by the slit at 00:04:16 UT. This is the pixel where we find that the total amount of upflowing plasma is larger than the stationary plasma component. The AIA density evolution shows that in this pixel, the strong drop in the density, indicative of the impulsive dimming onset, starts exactly at that time when it was mapped by the EIS slit. In the neighboring EIS pixel to the left, centered at x = 560′′, the impulsive density drop occurred during 23:58 to 00:06 UT, and the EIS slit covered it during this time range at 00:04:58 UT. This finding is consistent with EIS still showing a
Paper: 2020ApJ...905...36W (fulltext)
in the ionization balance in photoionization equilibrium (Ferland et al. 1998b ; Ferland 2003 ). In order to provide reliable DR rate coefficients for use in plasma modeling, the storage ring DR experiments combined with the new calculations with the AUTOSTRUCTURE code and the FAC code have made huge progress in the past few decades. Storage ring DR experiments of highly charged iron have been performed with almost all of the charge states for Fe 7+ to Fe 22+ (Savin 2007a ; Schippers 2012 ). Following the first paper to describe the calculation goals and methodology (Badnell et al. 2003 ), a series of the papers have been published based on the calculations by the AUTOSTRUCTURE code, and the calculated data are collected into the recent database CHIANTI (Dere et al. 2009 , 2019 ). By comparison of these calculations and the storage ring experiments, DR rate coefficients for carbon-like and oxygen-like ions have been identified as the most urgent needs for astrophysical applications, and therefore precision measurements of the DR rate coefficients of these ions that are abundant in astrophysical plasmas to benchmark the theories are required. Calcium is an abundant element in the solar corona and also in high-temperature astrophysical plasmas. Emission lines from highly charged calcium ions play an important role for diagnostics and modeling in astrophysical plasmas as well as fusion plasmas (Del Zanna Mason 2018 ; Träbert et al. 2018 ). For a C-like Ca 14+ ion, the previous observations were mostly from solar flare and solar wind based on the lines that are largely in the EUV range. These EUV lines are usually from the transitions within the 2 s 2 2 p 2 –2s 2 2 p 2 and 2 s 2 2 p 2 –2 s 2 p 3 configurations. The spectral lines of 555 Å and 1098 Å from 2 s 2 2 p 2 –2s 2 2 p 2 transitions of Ca xv were measured and identified by the Solar Ultraviolet Measurements of Emitted Radiation spectrometer on the Solar and Heliospheric Observatory (SOHO). With the Skylab observation, the ratios of the 181.90 or 215.37 Å lines versus the resonance EUV line at 200.98 Å are excellent density diagnostics at high densities of solar flares (Dere et al. 1979 ). This Skylab observation of Ca xv emission lines to derive the electron density was later on confirmed by a new data analysis and a measurement from the TEXT tokamak plasma (Keenan et al. 1988 , 2003 ). The Ca xv 181.90/200.98 Å ratio was also investigated by the observations of solar coronal emission lines with the Extreme-Ultraviolet Imaging Spectrometer (Culhane et al. 2007 ) on Hinode spacecraft (Kosugi et al. 2007 ), as discussed by Warren et al. ( 2008 ) and Del Zanna ( 2008 ), and several Ca xv emission lines at EUV range are formed at about 4MK and the electron density measurements at these temperatures are potentially important for understanding the coronal heating mechanism. In addition to the EUV emission lines, transitions of C-like Ca xv within the 2 s 2 2 p 2 –2 s 2 2 p 3 d configurations
Paper: 2019ApJ...877..130S (acknowledgments)
M.B.S. acknowledges a PhD student fellowship from IISER, Pune. We acknowledge insightful comments from the anonymous referee that has helped us improve the paper.