adsabs/AstroCLIMB huggingface dataset:
Paper: Black Hole Dragging: A New Mechanism for Forming Nuclear Star Clusters in Giant Elliptical Galaxies
Authors: Nader Khonji, Alessia Gualandris, Justin I. Read, Walter Dehnen
Figure caption: Surface density profiles and combined inner Sérsic and outer core-Sérsic fits for the five most prominent NSCs: NGC 1600 after kicks with a vk = 0.3 and 0.5 of vesc (top left and top center, respectively), A2147 with vk = 0.3vesc (top right), and NGC 6166 and IC 1101 with vk = 0.5vesc (bottom left and bottom center, respectively). The binned data are shown as black points, the inner Sérsic component of the fit as a blue dotted–dashed line, the outer core-Sérsic component as a green dashed line, and the combined fit as a solid red line. The effective radius for the Sérsic fit is used as the radius of the NSC.
Paper: Black Hole Dragging: A New Mechanism for Forming Nuclear Star Clusters in Giant Elliptical Galaxies
Authors: Nader Khonji, Alessia Gualandris, Justin I. Read, Walter Dehnen
Figure caption: Inner logarithmic slope (γ) as a function of the ratio of kick velocity (vk) to escape speed (vesc) for all six galaxy remnants. A vk/vesc of zero indicates scouring only. If the NSC cannot be excluded from the fit at a vk/vesc of 0.1 or 0.3, there is a spike in γ to ∼0.5 due to its presence. Where the NSC is excluded (unfilled markers), the relation between γ and vesc is restored.
Paper: Black Hole Dragging: A New Mechanism for Forming Nuclear Star Clusters in Giant Elliptical Galaxies
Authors: Nader Khonji, Alessia Gualandris, Justin I. Read, Walter Dehnen
Figure caption: NSC mass (MN) as the proportion of SMBH mass (M•) for the five most prominent NSCs. The galaxies are listed in order of increasing SMBH mass from left to right. Those forming after kicks with a vk/vesc of 0.3 and 0.5 are shown in blue and orange, respectively. MN appears to scale faster than M• and MN/M• may be smaller for kicks with a vk/vesc of 0.5 than 0.3.
Paper: Black Hole Dragging: A New Mechanism for Forming Nuclear Star Clusters in Giant Elliptical Galaxies
Authors: Nader Khonji, Alessia Gualandris, Justin I. Read, Walter Dehnen
Figure caption: Contour plots of NGC 6166 with vk = 0.5vesc. The blue dot is the position of the SMBH remnant, which leaves the area of the plot toward apocenter. The red arrow indicates the direction and magnitude of its velocity. An area of high density is present around the SMBH after the kick, and persists at “Br,” surrounded by a flat core.
Paper: Black Hole Dragging: A New Mechanism for Forming Nuclear Star Clusters in Giant Elliptical Galaxies
Authors: Nader Khonji, Alessia Gualandris, Justin I. Read, Walter Dehnen
Figure caption: NSC radius (rN) for the five most prominent NSCs. Those forming after kicks with a vk/vesc of 0.3 and 0.5 are shown in purple and olive, respectively. The galaxies have increasing SMBH mass (M•) from left to right, and there appears to be a trend for increasing rN with M•.
Paper: Solar Eruption Onset and Particle Acceleration in Nested-null Topologies
Authors: Pankaj Kumar, Judith T. Karpen, Peter F. Wyper, David Lario, Spiro K. Antiochos, C. Richard DeVore
Figure caption: Plasma and magnetic configuration of AR NOAA 12017 on 2014 March 29. (a) AIA 211 Å image of the AR overlaid by HMI magnetogram contours (±200 G). White (black) contours indicate positive (negative) polarities. (b) The zoomed view of the region outlined by the white box in panel (a). N1 represents the approximate position of the lower (inner) null of the nested-null topology, while N2 indicates the approximate position of the higher (outer) null. The arrow points to an small S-shaped filament (outlined by the cyan dashed line). (c) HMI magnetogram (scaled between ±100 G) at 15:06:45 UT. P1 and P2 represent the minority-polarity fluxes (positive) surrounded by opposite polarities. (d) Potential field extrapolation of the AR utilizing the magnetogram shown in panel (c). The cyan field lines illustrate the inner fan-spine topology rooted within the pseudostreamer and surrounding open field (red). N1 and N2 represent the inner and outer nulls, respectively.
Paper: Composition of Planetary Debris Around the White Dwarf GD 362
Authors: William T. Reach, Mukremin Kilic, Carey M. Lisse, John H. Debes, Ted von Hippel, Bianca Azartash-Namin, Loïc Albert, Susan E. Mullally, Fergal Mullally, Misty Cracraft, Madison Bernice, Selin L. Erickson
Figure caption: The JWST/NIRSPEC spectrum of GD 362, covering 0.6–5.3 μm. Dots show spectra from each exposure, with the color changing for each of the five nods along the slit. Insets show the spectra in the vicinity of the Hα and Pα photospheric lines of the white dwarf. The black curve shows the median of the five spectra. The black dots show the Spitzer/IRAC and Gaia photometry. For comparison, a white dwarf photosphere model, normalized to the Gaia GRP flux, is shown in light purple.
Paper: Composition of Planetary Debris Around the White Dwarf GD 362
Authors: William T. Reach, Mukremin Kilic, Carey M. Lisse, John H. Debes, Ted von Hippel, Bianca Azartash-Namin, Loïc Albert, Susan E. Mullally, Fergal Mullally, Misty Cracraft, Madison Bernice, Selin L. Erickson
Figure caption: Mineralogical decomposition of GD 362’s mid-infrared spectrum. (Left) The upper-left panel shows the observed spectrum (black with gray error bars) together with the model fit (purple). The lower-left panel shows the contributions from each silicate mineral in the model. (Right) The right panel shows the observed emissivity spectrum after subtracting the modeled silicates (black with gray error bars). Contributions from nonsilicate species are shown, with dashed lines for components that were not included in the final model because their contributions did not significantly improve the fit (as discussed in the text).
Paper: Composition of Planetary Debris Around the White Dwarf GD 362
Authors: William T. Reach, Mukremin Kilic, Carey M. Lisse, John H. Debes, Ted von Hippel, Bianca Azartash-Namin, Loïc Albert, Susan E. Mullally, Fergal Mullally, Misty Cracraft, Madison Bernice, Selin L. Erickson
Figure caption: Radiative transfer models for the spectrum of dust around GD 362. (a) The top panel shows an optically thin model (with a large scale height). (b) The bottom panel shows a moderately optically thick model (with a scale height 3.5 stellar radii at the inner edge). Grains in both models are composed of “Draine” silicates, amorphous carbon, and forsterite. The dust distribution parameters are in Table 3.
Paper: Composition of Planetary Debris Around the White Dwarf GD 362
Authors: William T. Reach, Mukremin Kilic, Carey M. Lisse, John H. Debes, Ted von Hippel, Bianca Azartash-Namin, Loïc Albert, Susan E. Mullally, Fergal Mullally, Misty Cracraft, Madison Bernice, Selin L. Erickson
Figure caption: The JWST/MIRI spectrum of GD 362. The three colored curves show the spectra in MIRI/MRS channels 1, 2, and 3. The channel 4 spectrum is not shown because the star was not detected. The single-channel spike at 7.311 μm is likely an artifact. Black points show the photometry in five bands from Spitzer (M. Jura et al. 2007), and purple dots show the photometry in three bands from JWST.
Paper: Composition of Planetary Debris Around the White Dwarf GD 362
Authors: William T. Reach, Mukremin Kilic, Carey M. Lisse, John H. Debes, Ted von Hippel, Bianca Azartash-Namin, Loïc Albert, Susan E. Mullally, Fergal Mullally, Misty Cracraft, Madison Bernice, Selin L. Erickson
Figure caption: Images of GD 362 in three JWST/MIRI filters at 18 μm (top), 21 μm (middle), and 25.5 μm (bottom). For each filter, the direct images are shown on the left, and the images after subtraction of a point-spread function centered on the star are on the right. In the PSF-subtracted images, a green circle denotes 100 au from the central star.
Paper: Composition of Planetary Debris Around the White Dwarf GD 362
Authors: William T. Reach, Mukremin Kilic, Carey M. Lisse, John H. Debes, Ted von Hippel, Bianca Azartash-Namin, Loïc Albert, Susan E. Mullally, Fergal Mullally, Misty Cracraft, Madison Bernice, Selin L. Erickson
Figure caption: Comparison of the density distributions for models fit to G29-38 (lower two models) and GD 362 (upper two models). The scales and references are listed in Table 4. For G29-38, the model from B22 is a ring (vertically thick but radially narrow), at a radius within the range of the power-law torus of R09. For GD 362, the model from J07 is a flat, opaque disk that warps beyond 50 R⋆.
Paper: Composition of Planetary Debris Around the White Dwarf GD 362
Authors: William T. Reach, Mukremin Kilic, Carey M. Lisse, John H. Debes, Ted von Hippel, Bianca Azartash-Namin, Loïc Albert, Susan E. Mullally, Fergal Mullally, Misty Cracraft, Madison Bernice, Selin L. Erickson
Figure caption: Likelihood distribution for the initial mass of GD 362. The maximum likelihood initial mass is 1.24  ±  0.13 M⊙.
Paper: Dynamical Masses for 23 Pre-main-sequence Stars in Upper Scorpius: A Critical Test of Stellar Evolutionary Models
Authors: A. P. M. Towner, J. A. Eisner, P. D. Sheehan, L. A. Hillenbrand, Y.-L. Wu
Figure caption: As for Figure 8, but showing the full range of dynamical masses for our data. The best-fit line for the full sample is shown as a dotted colored line in each panel, and the best fit for the mass-limited sample is shown as a dashed colored line. As in Figure 8, confirmed or candidate binaries are shown with gray symbols, and the 1:1 relation is shown as a black solid line. Although no PMS model sets show a 1:1 relationship with dynamical mass when using the full sample, the mass values for the G. A. Feiden (2016) magnetic models are still the least discrepant.
Paper: Dynamical Masses for 23 Pre-main-sequence Stars in Upper Scorpius: A Critical Test of Stellar Evolutionary Models
Authors: A. P. M. Towner, J. A. Eisner, P. D. Sheehan, L. A. Hillenbrand, Y.-L. Wu
Figure caption: Stellar mass vs. inclination angle for the CO-derived M⋆ and i in this work (blue), the CO-derived M⋆ and i in S. A. Barenfeld et al. (2017; gray), and the continuum-derived M⋆ and i in S. A. Barenfeld et al. (2017; black). We find that our derived masses and inclination angles show no bias compared to the S. A. Barenfeld et al. (2017) data, but that all samples suffer from small number statistics at the lowest (<30°) and highest (>85°) inclinations.
Paper: Dynamical Masses for 23 Pre-main-sequence Stars in Upper Scorpius: A Critical Test of Stellar Evolutionary Models
Authors: A. P. M. Towner, J. A. Eisner, P. D. Sheehan, L. A. Hillenbrand, Y.-L. Wu
Figure caption: As for Figure 9, but showing the full range of dynamical masses for our data. The CDF using our full sample is shown as a black dotted line. Compare to Figure 9.
Paper: Beyond the Local Group. I. Cosmic Filaments Govern the Spatial Alignments of Galaxy Pairs
Authors: Peng Wang, 鹏 王
Figure caption: Similar to Figure 3, but showing the effect of distance (dfila) between galaxies pairs to the filament spine. The blue curve shows the CDF for pairs with dfila ≤ 0.2 Mpc, the red for dfila ≤ 0.5 Mpc, the green for dfila ≤ 1.0 Mpc, and the orange for dfila ≥ 1.0 Mpc. The colored shaded regions denote the corresponding 95% bootstrap confidence band for each curve, respectively. The legend lists the number of galaxy pairs in each subsample and the confidence levels relative to randomized control samples, respectively. The 1000 gray lines represent CDFs from random realizations, and the black dashed line is the mean CDF of these random samples. The fraction of each subsample with alignment angles smaller than the MW–M31 case, namely ﹩39\mathop{.}\limits^{^\circ }4﹩ shown in vertical red dashed line, is indicated by the corresponding points and annotated text, respectively.
Paper: Beyond the Local Group. I. Cosmic Filaments Govern the Spatial Alignments of Galaxy Pairs
Authors: Peng Wang, 鹏 王
Figure caption: Cumulative distribution functions (CDFs) of the alignment angles between galaxy pair connecting lines and their corresponding local filament orientations. The blue solid line shows the observed CDF for selected 3087 LG-like galaxy pairs in SDSS. The blue shaded region denotes the 95% bootstrap confidence band by resampling the galaxy pair sample with replacement 1000 times. The gray lines represent CDFs from 1000 random realizations where the pair connecting lines are randomized but filament orientations are kept fixed, illustrating the null hypothesis of no alignment. The black dashed line is the mean CDF of these random samples. The vertical red dashed line marks the angle of ﹩39\mathop{.}\limits^{^\circ }4﹩, corresponding to the MW–M31 pair in the local Universe. Blue and black dots indicate the fraction of observed and random pairs with alignment angles smaller than the MW–M31 case, respectively.
Paper: Beyond the Local Group. I. Cosmic Filaments Govern the Spatial Alignments of Galaxy Pairs
Authors: Peng Wang, 鹏 王
Figure caption: Scatterplots with kernel density estimate (KDE)–based density maps illustrating the alignment angles versus galaxy pair–filament distances. The y-axis represents the alignment angle between the galaxy pair connecting line and the local filament orientation. The x-axis represents the distance of the galaxy pair to the filament spine. Each gray dot corresponds to an individual measurement. The background color gradient and overlaid contour lines represent the number density of points, with the color bar indicating the density scale. The gray dashed horizontal line indicates the angle (﹩57\mathop{.}\limits^{^\circ }3﹩) derived from the random trials.
Paper: Beyond the Local Group. I. Cosmic Filaments Govern the Spatial Alignments of Galaxy Pairs
Authors: Peng Wang, 鹏 王
Figure caption: Distributions of alignment angles for galaxy pairs of fiducial sample. The three panels display histograms of the angle between the galaxy pair connecting line and the filament orientation at the first galaxy (left), the angle between the pair connecting line and the filament orientation at the second galaxy (middle), and the angle between the filament orientations at the two galaxies (right). All angles are restricted to the range 0°–90°. For a random distribution of orientations in three-dimensional space, the expected mean angle is approximately 57﹩\mathop{.}\limits^{^\circ }﹩3. The black solid lines in the left and middle panels indicate the mean alignment angle distribution expected from random trials, where the pair orientations are randomized while keeping the filament orientations fixed.
Paper: AI-powered Reconstruction of Dark Matter Velocity Fields from Redshift-space Halo Distribution
Authors: Xu Xiao, Jiacheng Ding, Xiaolin Luo, Sun Ke Lan, Liang Xiao, Shuai Liu, Xin Wang, Le Zhang, Xiao-Dong Li
Figure caption: Training scheme for reconstructing the DM velocity field in real space. The reconstruction process involves two major steps. The first step is to reconstruct the DM density field, ρDM, from the sparse DM halo density field, ρs, in redshift space. To achieve this, we trained a UNet model to predict ρDM. Subsequently, we trained another UNet model that takes as input the concatenated data from the linear velocity prediction. This model is then used to reconstruct the direction and magnitude fields of velocity. This network structure is also applicable when the final output is the real-space momentum field.
Paper: Beyond the Local Group. I. Cosmic Filaments Govern the Spatial Alignments of Galaxy Pairs
Authors: Peng Wang, 鹏 王
Figure caption: Illustration of the alignment angles between a galaxy pair and the local cosmic filament. The two circles represent the galaxies in the pair, connected by the pair connecting line. The arrow at each galaxy indicates the local filament orientation. The angles θ1 and θ2 denote the alignment angles between the pair connecting line and the filament orientation at each galaxy.
Paper: Beyond the Local Group. I. Cosmic Filaments Govern the Spatial Alignments of Galaxy Pairs
Authors: Peng Wang, 鹏 王
Figure caption: Histogram showing the distribution of the mean distances of galaxy pairs to their nearest filament spine. The x-axis represents the mean pair–filament distance, (dfila), in megaparsecs (Mpc). The distribution peaks at approximately 0.4 Mpc and exhibits a long tail toward larger distances, indicating that most galaxy pairs reside close to the filament.
Paper: Dissecting the Mass Quenching in TNG50: Galaxy Size Determines the Quenching Mode
Authors: Haochen Jiang, Enci Wang, Kai Wang, Chengyu Ma, Xu Kong
Figure caption: The stacked SFR surface density of three types of quiescent galaxies at the star-forming time (tSF, left), SMBH growth flattened time (tBH, middle), and quiescent time (tQ, right). The top panels show the stacked face-on image with a color bar encoding the SFR surface density. The bottom panels show the SFR surface density as a function of physical galaxy-centric distance and Re-normalized galaxy-centric distance. In each panel, the thin lines show the individual result, while the thick lines show the average of each type of quiescent galaxies. This figure clearly demonstrates that the triggering of kinetic AGN feedback at tBH can effectively quench star formation activities within ≈1–2 kpc, which causes those compact and small galaxies to be classified as quenched, while those extended galaxies are still counted as SFGs.
Paper: Dissecting the Mass Quenching in TNG50: Galaxy Size Determines the Quenching Mode
Authors: Haochen Jiang, Enci Wang, Kai Wang, Chengyu Ma, Xu Kong
Figure caption: The evolution of galaxy properties around tBH for three types of quiescent galaxies selected at z = 0 (red: Fast-QG; green: Reju-QG; blue: Slow-QG) and the evolution of reference SFGs in black. The background faded lines show the evolution trajectory of individual galaxies, and the solid lines show the median of each type of galaxy.
Paper: Dissecting the Mass Quenching in TNG50: Galaxy Size Determines the Quenching Mode
Authors: Haochen Jiang, Enci Wang, Kai Wang, Chengyu Ma, Xu Kong
Figure caption: The relationship between the deviation from the SFMS and the effective radius (left) and kinematics (right), which is quantified by the fraction of stellar particles with orbital circularity above 0.7, measured at tBH for three categories of quiescent central galaxies in TNG50. The blue solid lines show the linear fitting, and the shaded regions show the 1​​​​​σ range of bootstrap samples. The surrounding three small panels show the distribution of corresponding properties.
Paper: Dissecting the Mass Quenching in TNG50: Galaxy Size Determines the Quenching Mode
Authors: Haochen Jiang, Enci Wang, Kai Wang, Chengyu Ma, Xu Kong
Figure caption: The 46 samples from the three classes (orange circle: Fast-QG; green triangle: Reju-QG; magenta square: Slow-QG) on the tQ versus (tBH–tQ) plane and the effective radius versus tBH plane. Many galaxies across all three classes have similar tBH values but different sizes. Galaxies with smaller sizes tend to quench rapidly, while larger galaxies quench more slowly or incompletely.
Paper: Dissecting the Mass Quenching in TNG50: Galaxy Size Determines the Quenching Mode
Authors: Haochen Jiang, Enci Wang, Kai Wang, Chengyu Ma, Xu Kong
Figure caption: Similar to Figure 5 but for the gas surface density. This figure also demonstrates that the triggering of kinetic AGN feedback expels the gas content within ≈1–2 physical kpc, and this causes those small and compact galaxies to be gas deficient and easily quenched.
Paper: Dissecting the Mass Quenching in TNG50: Galaxy Size Determines the Quenching Mode
Authors: Haochen Jiang, Enci Wang, Kai Wang, Chengyu Ma, Xu Kong
Figure caption: The evolution of total SFR in blue solid lines and the gas inflow rate in golden solid lines around tBH and tQ for three types of quiescent galaxies. The dashed lines are for the star-forming reference sample.
Paper: Probing Evolution of Long GRB Properties through their Cosmic Formation History Aided by Machine Learning–predicted Redshifts
Authors: Dhruv S. Bal, Aditya Narendra, Maria Giovanna Dainotti, Nikita S. Khatiya, Aleksander Ł. Lenart, Dieter H. Hartmann
Figure caption: Left: the τ–k distribution for the PS. We select the k value corresponding to τ = 0 for making the luminosity and 1 + z independent of each other. Right: the τ–k distribution for the CS. We select the k value corresponding to τ = 0 to make the luminosity and 1 + z independent of each other.
Paper: Morphological Regimes of Rotating Moist Convection
Authors: Whitney T. Powers, Adrian E. Fraser, Evan H. Anders, Jeffrey S. Oishi, Benjamin P. Brown
Figure caption: Panel (a): horizontally- and time-averaged vertical profiles of moist static energy fluxes for the simulation shown in Figures 6–7. The turbulent heat flux (dashed orange line) and the turbulent latent heat flux (dashed blue line) sum to form the moist static energy flux (solid green line). We also show the diffusive thermal flux (purple) and the diffusive moisture flux (red). The total energy flux (brown line) is given by the sum of the solid lines. Panel (b): horizontally- and time-averaged profiles of potential temperature θ, latent heat density Λ, and moist static energy density m. The solid orange line shows the potential temperature of the evolved state, and the dashed orange line shows the potential temperature of the initial atmosphere. The solid blue line shows the latent heat of the evolved state, the dotted blue line shows the latent heat at saturation of the evolved state, and the dashed blue line shows the latent heat of the initial atmosphere. Note that the latent heat of the initial atmosphere is identical to the latent heat at saturation and the moist static energy of the initial atmosphere.
Paper: ALMA-IMF. XX. Core Fragmentation in the W51 High-mass Star-forming Region
Authors: T. Yoo, A. Ginsburg, J. Braine, N. Budaiev, F. Louvet, F. Motte, A. M. Stutz, B. Thomasson, M. Armante, M. Bonfand, S. Bontemps, L. Bronfman, G. Busquet, T. Csengeri, N. Cunningham, J. Di Francesco, D. J. Díaz-González, M. Fernández-Lopez, R. Galván-Madrid, C. Goddi, A. Gusdorf, N. Kessler, A. Koley, H.-L Liu, T. Nony, F. Olguin, P. Sanhueza, M. Valeille-Manet, L. A. Zapata, Q. Zhang
Figure caption: An overview of the W51-A region. The image at the center is the ALMA 3 mm continuum image of ALMA-IMF combining two subregions, W51-E and W51-IRS2. Each continuum image is truncated and merged in the middle to look seamless. Three regions marked with the blue boxes are magnified with the 1.3 mm continuum images from the ALMA-IMF data (left) and high-res data (right). We annotate massive young stellar objects (YSOs) studied in previous studies (e.g., C. Goddi et al. 2020), W51e2 (#39) and W51e8 (#32) in W51-E and W51-North (#11) in W51-IRS2.
Paper: ALMA-IMF. XX. Core Fragmentation in the W51 High-mass Star-forming Region
Authors: T. Yoo, A. Ginsburg, J. Braine, N. Budaiev, F. Louvet, F. Motte, A. M. Stutz, B. Thomasson, M. Armante, M. Bonfand, S. Bontemps, L. Bronfman, G. Busquet, T. Csengeri, N. Cunningham, J. Di Francesco, D. J. Díaz-González, M. Fernández-Lopez, R. Galván-Madrid, C. Goddi, A. Gusdorf, N. Kessler, A. Koley, H.-L Liu, T. Nony, F. Olguin, P. Sanhueza, M. Valeille-Manet, L. A. Zapata, Q. Zhang
Figure caption: The spectral index distribution of high-res sources detected in both bands. Left: 3 mm peak flux vs. 1.3 mm peak flux diagram. Here, the 1.3 mm peak fluxes are measured from the convolved continuum images with the common beam. The error bars denote the background noise level measured in each continuum image. Typical calibration errors, 5% in Band 3 and 10% in Band 6, are marked as gray crosses at the lower right corner. The gray dashed lines indicate a spectral index = [−1, 0, 1, 2, 3, 4] from bottom to top. Right column: histograms of the spectral indices in W51-E and W51-IRS2. High-res sources in each region are grouped depending on whether the spectral index indicates a certain dusty source or certain H II regions, or possibly free–free emission contamination.