adsabs/AstroCLIMB huggingface dataset:
Paper: Multi-passband Observations of a Solar Flare over the He I 10830 Å line
Authors: Yan Xu, Xu Yang, Graham S. Kerr, Vanessa Polito, Viacheslav M. Sadykov, Ju Jing, Wenda Cao, and Haimin Wang
Figure caption: Sample images taken around 17:23 UT for the C3.0 flare that occurred on 2018 May 28. Panel (a): TiO continuum image. Panel (b): He I 10830 red wing +0.8 Å. Panel (c): IRIS SJI 2796 image.
Paper: Multi-passband Observations of a Solar Flare over the He I 10830 Å line
Authors: Yan Xu, Xu Yang, Graham S. Kerr, Vanessa Polito, Viacheslav M. Sadykov, Ju Jing, Wenda Cao, and Haimin Wang
Figure caption: He I 10830 images taken by BBSO/GST at five spectral positions, ±0.8, ±0.4 Å, and line center, in panels (a) to (e). Two areas, marked b1 and b2 in panel (c), are selected as the preflare background. The actual measurements of background intensities are obtained from images around 17:00:57 UT. Boxes of b1 and b2 are shown during the flare time for illustrative purposes only. In panel (e), the target ribbon is outlined by a black box. The small FOV within the black box is displayed in panel (f) for a closer view. The image scale is 0.″062 pix−1 for all panels. Two small boxes in panel (f) indicate two flaring areas. The one close to the ribbon front represents the initial peak of the flare emission, marked as “P.” The rectangular box near the center of panel (f) is located near the back edge of the flare ribbon, representing the trailing or well-developed area of emission marked as “T.” The dotted–dashed box marked as “S” is 10 pixels wide and shows the locations of 10 slits with the lower end covering the ribbon front with negative contrasts. The labels of panel (d) are shared by panels (a) to (e). Panel (f) uses different labels.
Paper: Multi-passband Observations of a Solar Flare over the He I 10830 Å line
Authors: Yan Xu, Xu Yang, Graham S. Kerr, Vanessa Polito, Viacheslav M. Sadykov, Ju Jing, Wenda Cao, and Haimin Wang
Figure caption: Panel (a) shows pseudo line profiles of the He I 10830 Å line in different spatial positions on the flare ribbon, with different colors. Orange: emission profile in the peak flaring area (“P” in panel (f) of Figure 2). Blue: emission profile in the trailing area (“T” in panel (f) of Figure 2). Cyan: absorption profile on the ribbon front. The point of the line center is measured at a different location and therefore highlighted with a red ring. The 1σ uncertainty is calculated from the standard deviation and affected by the area selection. For example, larger areas tend to have larger uncertainties. All colored curves share the same Y-axis on the left side with a unit of percentage of contrast. The black dashed curve links the intensities of the background b, using the right Y-axis. The intensities are normalized to the maximum measured at −0.8 Å and the absolute amplitudes are magnified by a factor of 10 to illustrate the differences among the five spectral positions. Panel (b) shows the averaged contrasts along the 10 slits (dotted–dashed rectangular box marked as “S” in panel (f) of Figure 2), at −0.8 Å.
Paper: Multi-passband Observations of a Solar Flare over the He I 10830 Å line
Authors: Yan Xu, Xu Yang, Graham S. Kerr, Vanessa Polito, Viacheslav M. Sadykov, Ju Jing, Wenda Cao, and Haimin Wang
Figure caption: IRIS SJI and spectra taken at 17:14 UT. Panel (a): SJI 2796 image, in which the dark line indicates the location of slit. Panel (b): Mg II spectrum. Panel (c): Mg II K line profiles for the northern flare ribbon (positive flare), the southern ribbon where enhanced He I 10830 is located (negative flare), and the flaring loops at bottom (flare loop). The positions of spectra are marked using the same color in panel (b).
Paper: The Three-dimensional Collapse of a Rapidly Rotating 16 M⊙ Star
Authors: C. E. Fields
Figure caption: Stellar structure profiles for the 16 M⊙ MESA model at the time of mapping into FLASH. Top panel: convective velocity according to MLT and specific nuclear energy generation rate as a function of mass coordinate. Middle panel: specific angular momentum and angular velocity. Bottom panel: mass-fraction profiles and the specific entropy. The convective O-shell region is shown by the shaded blue region in each panel.
Paper: The Three-dimensional Collapse of a Rapidly Rotating 16 M⊙ Star
Authors: C. E. Fields
Figure caption: Time evolution of the integrated kinetic (left) and nuclear (right) energy for the duration of the 3D FLASH simulation.
Paper: The Three-dimensional Collapse of a Rapidly Rotating 16 M⊙ Star
Authors: C. E. Fields
Figure caption: Time evolution of the power spectrum for the spherical harmonic decomposition of the turbulent radial velocity field for the convective O-shell region. The spectra are evaluated in the 3D model at eight different times at a radius of rO-shell = 6000 km.
Paper: The Three-dimensional Collapse of a Rapidly Rotating 16 M⊙ Star
Authors: C. E. Fields
Figure caption: Streamlines of the magnitude of the velocity field for the 3D FLASH simulation at a time ≈7 s before iron core collapse. The purple contour shows the iron core at a radius of r ≈ 2400 km shown here. The brighter (orange) stream lines correspond to speeds of ﹩\left|{\boldsymbol{v}}\right|\approx 800﹩ km s−1. The orientation of the rendering is shown with the north vector aligned with the y-axis, the axis of rotation, (with a slight horizontal offset). Large-scale flow is observed with streamlines wrapping about the axis of rotation out to the edge of the convective shell region. The spatial scale of the image corresponds to a spatial extent of ≈40,000 km on a side. This image was produced using Paraview.
Paper: The Three-dimensional Collapse of a Rapidly Rotating 16 M⊙ Star
Authors: C. E. Fields
Figure caption: Top left: angle average of the y-component of the specific angular momentum for the 3D FLASH simulation at five different times include t = 601.5 s (half a second before iron core collapse). The correponding MESA model is also shown by the dashed gray line chosen to match the central density of the 3D model at t = 601.5 s. Turbulent (4πr2Fturb., top right) and advective flux (4πr2Fadv., bottom left) terms for the 3D simulation at three different times. The approximate convective Si- (m ≈ 1.52–1.60 M⊙) and O-shell (m ≈ 1.89–2.80 M⊙) regions at t = 601.5 are denoted by the blue shading.
Paper: The CHIME Fast Radio Burst Population Does Not Track the Star Formation History of the Universe
Authors: Rachel C. Zhang, and Bing Zhang
Figure caption: Probability distribution functions (PDFs) for the SFH model, lognormal delay model with a central value of 10 Gyr, hybrid model with a 20% contribution from the SFH model and 80% from a lognormal delay model with a central value of 13 Gyr, and accumulated SFH model. The top panel shows the PDF of the intrinsic FRB event-rate density ﹩{dN}/({dtdV})﹩, while the bottom panel shows the PDF of the observed FRB event-rate redshift distribution ﹩{dN}/({{dt}}_{\mathrm{obs}}{dz})﹩.
Paper: The CHIME Fast Radio Burst Population Does Not Track the Star Formation History of the Universe
Authors: Rachel C. Zhang, and Bing Zhang
Figure caption: A test of the SFH model against the three observational criteria. The full FRB sample is tested against the model predictions. (a) The ﹩\mathrm{log}{N}_{\gt {{ \mathcal F }}_{\nu }}-\mathrm{log}{{ \mathcal F }}_{\nu }﹩ test, (b) the ﹩\mathrm{log}E﹩ distribution test, (c) the DME distribution test, and (d) the 2D DME − logE distribution for illustrative purpose (not for testing). For panels (a), (b), and (c), the SFH simulations are scaled to the full data. For the energy distribution model, we adopt α = 1.9 and logEc = 41.5.
Paper: The CHIME Fast Radio Burst Population Does Not Track the Star Formation History of the Universe
Authors: Rachel C. Zhang, and Bing Zhang
Figure caption: Similar to Figure 2, but for a test of an accumulated star formation model. All the simulations are scaled to the data. For the energy distribution model, we adopt α = 1.9 and ﹩\mathrm{log}{E}_{c}=41﹩.
Paper: The CHIME Fast Radio Burst Population Does Not Track the Star Formation History of the Universe
Authors: Rachel C. Zhang, and Bing Zhang
Figure caption: Similar to Figure 2, but for a test of a lognormal delay model with a central value of 10 Gyr and a standard deviation of 0.8 dex. All the simulations are scaled to the data. For the energy distribution model, we adopt α = 1.9 and ﹩\mathrm{log}{E}_{c}=41﹩.
Paper: The CHIME Fast Radio Burst Population Does Not Track the Star Formation History of the Universe
Authors: Rachel C. Zhang, and Bing Zhang
Figure caption: Similar to Figure 2, but for a test of a hybrid model that includes a young component tracking the SFH (20% of the total population) and an old component (80% of the total population) with lognormal delay with a central value of 13 Gyr and standard deviation of 0.8 dex. All the simulations are scaled to the data. For the energy distribution model, we adopt α = 1.9 and ﹩\mathrm{log}{E}_{c}=41﹩.
Paper: A Significant Detection of X-ray Polarization in Sco X-1 with PolarLight and Constraints on the Corona Geometry
Authors: Xiangyun Long, Hua Feng, Hong Li, Jiahuan Zhu, Qiong Wu, Jiahui Huang, Massimo Minuti, Weichun Jiang, Dongxin Yang, Saverio Citraro, Hikmat Nasimi, Jiandong Yu, Ge Jin, Ming Zeng, Peng An, Jiachen Jiang, Enrico Costa, Luca Baldini, Ronaldo Bellazzini, Alessandro Brez, Luca Latronico, Carmelo Sgrò, Gloria Spandre, Michele Pinchera, Fabio Muleri, and Paolo Soffitta
Figure caption: Intensity of Sco X-1 in each exposure as a function of time measured with PolarLight in the energy range of 2–8 keV and a histogram of the intensity. The dashed line marks the level to separate the low- and high-intensity exposures, with approximately the same number of counts in each segment.
Paper: A Significant Detection of X-ray Polarization in Sco X-1 with PolarLight and Constraints on the Corona Geometry
Authors: Xiangyun Long, Hua Feng, Hong Li, Jiahuan Zhu, Qiong Wu, Jiahui Huang, Massimo Minuti, Weichun Jiang, Dongxin Yang, Saverio Citraro, Hikmat Nasimi, Jiandong Yu, Ge Jin, Ming Zeng, Peng An, Jiachen Jiang, Enrico Costa, Luca Baldini, Ronaldo Bellazzini, Alessandro Brez, Luca Latronico, Carmelo Sgrò, Gloria Spandre, Michele Pinchera, Fabio Muleri, and Paolo Soffitta
Figure caption: Polarimetric modulation curves (red histograms) and corresponding PA versus PF contours of Sco X-1 in the 4–8 keV band, in all and the high-intensity time intervals, respectively. Top: the modulation curves are for visual inspection only, with typical 1σ error bars, and the model curves (gray lines) are derived from the Stokes/Bayesian analysis. Bottom: the red crosses indicate the point estimates and the contours encircle the 1σ, 2σ, and 3σ credible intervals of the Bayesian posterior distribution. The blue point with error bars marks the measurement obtained with OSO-8 around 5.2 keV.
Paper: A Significant Detection of X-ray Polarization in Sco X-1 with PolarLight and Constraints on the Corona Geometry
Authors: Xiangyun Long, Hua Feng, Hong Li, Jiahuan Zhu, Qiong Wu, Jiahui Huang, Massimo Minuti, Weichun Jiang, Dongxin Yang, Saverio Citraro, Hikmat Nasimi, Jiandong Yu, Ge Jin, Ming Zeng, Peng An, Jiachen Jiang, Enrico Costa, Luca Baldini, Ronaldo Bellazzini, Alessandro Brez, Luca Latronico, Carmelo Sgrò, Gloria Spandre, Michele Pinchera, Fabio Muleri, and Paolo Soffitta
Figure caption: Distribution of the PolarLight roll angle during observations of Sco X-1 and those when the source displays low and high intensities.
Paper: A Significant Detection of X-ray Polarization in Sco X-1 with PolarLight and Constraints on the Corona Geometry
Authors: Xiangyun Long, Hua Feng, Hong Li, Jiahuan Zhu, Qiong Wu, Jiahui Huang, Massimo Minuti, Weichun Jiang, Dongxin Yang, Saverio Citraro, Hikmat Nasimi, Jiandong Yu, Ge Jin, Ming Zeng, Peng An, Jiachen Jiang, Enrico Costa, Luca Baldini, Ronaldo Bellazzini, Alessandro Brez, Luca Latronico, Carmelo Sgrò, Gloria Spandre, Michele Pinchera, Fabio Muleri, and Paolo Soffitta
Figure caption: Schematic drawing of the accretion flow and expected polarization angles of the two models. The radio jet has an orientation of 54° on the sky plane and is likely the symmetry axis of the accretion flow. Given an optically thin corona, if the Comptonization occurs in the accretion disk corona, the PA is expected to be aligned with the disk plane, as indicated by the white arrow; otherwise, if the Comptonization occurs in the transition layer, the PA is expected to be in line with the system axis, as indicated by the black arrow.
Paper: Magnetic Field Amplification by a Plasma Cavitation Instability in Relativistic Shock Precursors
Authors: J. R. Peterson, S. Glenzer, and F. Fiuza
Figure caption: 2D simulation of the transverse dynamics of a dilute, relativistic pair–ion beam with α = 0.1 and γbe = 1000 propagating into a cold pair–ion plasma. The pair loading factor is Z± = 1 and the beam propagates into the page. Magnetic field profiles (a,b) and density lineouts (c,d) are shown at (a,c) ﹩t=500\,{\omega }_{p}^{-1}﹩, the time of saturation of the Weibel instability, and (b,d) ﹩t=5\times {10}^{3}\,{\omega }_{p}^{-1}﹩, the time of saturation of the cavitation instability. The lineouts are taken along the white dashed lines in (a,b). The energy evolution of the system is reported in (e) log scale and (f) linear scale.
Paper: Magnetic Field Amplification by a Plasma Cavitation Instability in Relativistic Shock Precursors
Authors: J. R. Peterson, S. Glenzer, and F. Fiuza
Figure caption: Comparison between analytical growth rates and saturation values for the cavitation instability (solid lines), Weibel instability (dashed lines), and PIC simulations. Results from 2D simulations of the transverse yz-plane are shown for variable pair loading factor Z± ≠ 0 (open circles), pure electron–ion (Z± = 0) beam-plasma (filled circles), and pure electron–positron beam on electron–ion plasma (squares). Simulations capturing the longitudinal dynamics are shown for an electron–positron beam propagating on electron–ion plasma in 3D (triangles) and 2D xy-plane geometry for unmagnetized (×) and magnetized (+) (σ⊥ ∼ 10−6) initial conditions. (a,b) Cavitation instability growth rate. (c,d) Saturation magnetic field wavelength. (e) Saturation magnetization. (f) Beam energy asymmetry. All parameter scans use as fixed parameters α = 0.1 and γbe = 1000.
Paper: Magnetic Field Amplification by a Plasma Cavitation Instability in Relativistic Shock Precursors
Authors: J. R. Peterson, S. Glenzer, and F. Fiuza
Figure caption: Comparison of 2D simulations of the long-term magnetic field amplification for a cold (a,c,e) and a hot beam with comoving temperature Tb = mec2 (b,d,f,h). The dilute (α = 0.1) electron–positron (pair) beam has mean Lorentz factor 〈γbe〉 = 1000 in both cases and propagates in the positive x-direction through an electron–ion plasma. Magnetic field profiles (a,b) and pressure ratio PB/Pth (c,d) are taken at ﹩{t}_{1}=3000\,{\omega }_{p}^{-1}﹩ when the Weibel instability saturates. Magnetic field profiles in (e,f) are taken at ﹩{t}_{2}=12000\,{\omega }_{p}^{-1}﹩ when the cavitation instability saturates. The z-averaged magnetization is shown in (g) at t1 (blue) and t2 (orange) for the cold (dashed) and hot (solid) beam. The evolution of the magnetization and magnetic coherence length are reported in (h) for the hot beam.
Paper: Magnetic Field Amplification by a Plasma Cavitation Instability in Relativistic Shock Precursors
Authors: J. R. Peterson, S. Glenzer, and F. Fiuza
Figure caption: Development of the cavitation instability in a 3D simulation of a dilute (α = 0.05) electron–positron (pair) beam propagating in the positive x-direction through a cold electron–ion plasma. The beam has comoving temperature Tb = mec2 and mean Lorentz factor 〈γbe〉 = 1000. The transverse magnetic field amplitude (a), background ion density (b), and beam charge density (c) are reported at ﹩t=12000\,{\omega }_{p}^{-1}﹩, the saturation time of the cavitation instability. The opacity scales linearly with the value.
Paper: Magnetic Field Amplification by a Plasma Cavitation Instability in Relativistic Shock Precursors
Authors: J. R. Peterson, S. Glenzer, and F. Fiuza
Figure caption: Magnetic field amplification from the propagation of a hot pair beam with α = 0.1, 〈γbe〉 = 1000, and ﹩{T}_{b}^{{\prime} }={m}_{e}{c}^{2}﹩ on an electron–ion plasma from 2D fully periodic simulations with longitudinal box size (a,b) 220 c/ωp and (c,d) 1000 c/ωp. Magnetic field profiles are shown at (a,b) ﹩t=1000\,{\omega }_{p}^{-1}﹩, the time of saturation of the Weibel instability, and (c,d) ﹩t=6500\,{\omega }_{p}^{-1}﹩, the time of saturation of the cavitation instability. The evolution of the magnetization in the simulations with small (solid) and large (dashed) longitudinal box size is reported in (e).
Paper: Stochastic Low-frequency Variability in Three-dimensional Radiation Hydrodynamical Models of Massive Star Envelopes
Authors: William C. Schultz, Lars Bildsten, and Yan-Fei Jiang
Figure caption: Spectroscopic HR diagram showing locations of new models (the two Xs) relative to recent TESS observations and stellar evolution tracks (the black lines). The gold circles denote OB star observations from Burssens et al. (2020) while the blue outlines surround stars against which our models are compared in Section 3. The magenta square symbolizes the StarTop model from Jiang et al. (2015): a plane-parallel model of the surface of an 80 M⊙ star. The vertical lines denote locations where the optical depth to the iron-opacity peak, τFe, is equal to (dashed) or 10% of (dotted) of τcrit.
Paper: Stochastic Low-frequency Variability in Three-dimensional Radiation Hydrodynamical Models of Massive Star Envelopes
Authors: William C. Schultz, Lars Bildsten, and Yan-Fei Jiang
Figure caption: Profiles of density, temperature, opacity, and entropy (top to bottom, respectively) from a single temporal snapshot of T42L5.0 (left) and T35L5.2 (right). The pink dashed lines show the radial profile generated from a volume-weighted shell average and the color shows the probability that each quantity has the specified value at that radius. The sum of the color along vertical lines in each panel is 1. The vertical dashed black line shows the location of the iron-opacity peak and the dotted black line represents the photosphere (where 〈τ〉 = 1). The gray shaded region denotes where the average entropy gradient is negative. In the κ panel, the horizontal line represents the opacity above which the model is super-Eddington.
Paper: Stochastic Low-frequency Variability in Three-dimensional Radiation Hydrodynamical Models of Massive Star Envelopes
Authors: William C. Schultz, Lars Bildsten, and Yan-Fei Jiang
Figure caption: Slices of overdensity, radial velocity, opacity, diffusive radiative flux, tangential velocity magnitude, and overtemperature (clockwise from upper left respectively) though a θ–r plane of a single temporal snapshot from T35L5.2. The black dotted line shows the photosphere for the entire snapshot and the scale height at the iron-opacity peak, HFe, is denoted by the scale bar.
Paper: Stochastic Low-frequency Variability in Three-dimensional Radiation Hydrodynamical Models of Massive Star Envelopes
Authors: William C. Schultz, Lars Bildsten, and Yan-Fei Jiang
Figure caption: Velocity profiles for single temporal snapshots of T42L5.0 (top) and T35L5.2 (bottom) vs. radius relative to the location of the iron-opacity peak. The colored shaded regions denote the 95% probability interval of the rms velocities at each radius. Dashed gold and dotted red lines are the volume-weighted average rms tangential and radial velocities, respectively. The solid brown line denotes the convective velocity profile from analogous 1D MESA models. The black scale bar shows the extent of the scale height at the iron-opacity peak. Vertical black lines and gray shaded regions are the same as Figure 2.
Paper: Stochastic Low-frequency Variability in Three-dimensional Radiation Hydrodynamical Models of Massive Star Envelopes
Authors: William C. Schultz, Lars Bildsten, and Yan-Fei Jiang
Figure caption: Top: lightcurve of the T35L5.2 model (dark blue points) for 8 days after the model has reached a steady-state equilibrium in the outer regions. The gold dashed line is a first-order polynomial fit of the lightcurve used to zero-mean the lightcurve before taking the power spectrum. Middle: power spectrum of the lightcurve from the top panel (solid dark blue line) compared to the normalized SLF fit, α(ν), from observed OB stars calculated by Bowman et al. (2020; dashed blue lines of different shades). The colors of the dashed lines match the outlines of the observed locations in Figure 1. The solid gold line shows the result of removing the periodic signals and linearly smoothing the power spectrum. The pink dotted–dashed line denotes the α(ν) fit from T35L5.2 and the pink shaded region shows the 95% confidence interval. Bottom: cumulative power spectrum of all normalized power spectra on the middle panel normalize to be zero at the left limit (ν = 0.5 day−1) and one at the right edge.
Paper: Could the Magnetic Star HD 135348 Possess a Rigidly Rotating Magnetosphere?
Authors: Rahul Jayaraman, Swetlana Hubrig, Daniel L. Holdsworth, Markus Schöller, Silva Järvinen, Donald W. Kurtz, and George R. Ricker
Figure caption: The top two panels show TESS light curves for HD 135348 from Sectors 11 and 38. The bottom panel contains a discrete Fourier transform of the Sector 38 data, showing a period of 2.0593 days. The data gap in the middle of the light curves is caused by the lack of observations during data downlink. Approximately 5% of the data points in the Sector 11 light curve were clipped. The removed points arise from scattered light entering the aperture; the deviations of these points from unity were found to be much greater than the semi-amplitude of the periodic flux variations (∼5 ppt).
Paper: Could the Magnetic Star HD 135348 Possess a Rigidly Rotating Magnetosphere?
Authors: Rahul Jayaraman, Swetlana Hubrig, Daniel L. Holdsworth, Markus Schöller, Silva Järvinen, Donald W. Kurtz, and George R. Ricker
Figure caption: A comparison of binned, phase-folded TESS light curves for various magnetic stars with that of HD 135348 in the top panel. Its light curve most closely resembles that of the known RRM star HD 23478, shown in the second panel. The light curves of two other confirmed RRM stars, σ Ori E and HD 345439, are shown in the third and fourth panels. The fifth panel shows the light curve of HD 66765, a star that shows Hα in emission (a strong indicator that an RRM exists), and the bottom panel shows the light curve of a CM star, ALS 3694. The markers on the plot for HD 135348 represent the rotational phases at which the magnetic field measurements were made (see Section 2.2 and Table 1). We emphasize the significant differences in the light-curve morphologies for CM and RRM stars. These light curves are binned to 0.002 units of phase. Note that the y-axis scales vary in each panel, in order to more clearly show the individual stars’ flux variations. The scatter visible in the light curve of ALS 3694 is likely not astrophysical in origin, and may simply be rms scatter.
Paper: Could the Magnetic Star HD 135348 Possess a Rigidly Rotating Magnetosphere?
Authors: Rahul Jayaraman, Swetlana Hubrig, Daniel L. Holdsworth, Markus Schöller, Silva Järvinen, Donald W. Kurtz, and George R. Ricker
Figure caption: A plot of four Stokes I spectra taken with the RSS over four different epochs with insets showing the variability of the line profiles of the Si III triplet at 4552.62 Å, 4567.84 Å, and 4574.76 Å, along with the profiles of He I at 4921.93 Å. These wavelengths have been labeled with dotted vertical lines in both inset plots. Three of the four spectra have been offset for better visibility. We have labeled two other key spectral lines (another He I line, and also Hβ) visible in the plot.