adsabs/AstroCLIMB huggingface dataset:
Paper: Discovery of an Accretion Burst in a Free-floating Planetary-mass Object
Authors: Victor Almendros-Abad, Aleks Scholz, Belinda Damian, Ray Jayawardhana, Amelia Bayo, Laura Flagg, Koraljka Mužić, Antonella Natta, Paola Pinilla, Leonardo Testi
Figure caption: Time evolution of the mass accretion rate across all observations seen in Hα (blue circles) and Paschen γ (orange squares). Black diamonds show the average over the following available tracers at each epoch: Paschen δ, Paschen β, Brackett γ, and Hα, each shown with small light gray circles. A small offset in date has been applied for clarity. The Paγ and Hα accretion rate error bars represent the uncertainties from the line flux measurements only. Systematic uncertainties in the accretion rate, primarily from the adopted Lline–Lacc relation and stellar parameters, are not shown; these amount to ∼0.4 dex.
Paper: Discovery of an Accretion Burst in a Free-floating Planetary-mass Object
Authors: Victor Almendros-Abad, Aleks Scholz, Belinda Damian, Ray Jayawardhana, Amelia Bayo, Laura Flagg, Koraljka Mužić, Antonella Natta, Paola Pinilla, Leonardo Testi
Figure caption: Time evolution of the spectral energy distribution for Cha1107-7626, as observed with JWST NIRSpec-PRISM and MIRI-LRS, in 2024 August and 2025 August. Bottom left and right panels show a zoom to the optical and mid-infrared, respectively. The main mid-infrared features discussed in Section 4 are indicated in the bottom right panel.
Paper: Discovery of an Accretion Burst in a Free-floating Planetary-mass Object
Authors: Victor Almendros-Abad, Aleks Scholz, Belinda Damian, Ray Jayawardhana, Amelia Bayo, Laura Flagg, Koraljka Mužić, Antonella Natta, Paola Pinilla, Leonardo Testi
Figure caption: Line profile changes in Hα from VLT/XSHOOTER (left) and Paschen γ from both XSHOOTER and JWST (right) seen in VLT/XSHOOTER observations of Cha1107-7626. The eight epochs are color-coded and clearly show the evolution of the burst.
Paper: Discovery of an Accretion Burst in a Free-floating Planetary-mass Object
Authors: Victor Almendros-Abad, Aleks Scholz, Belinda Damian, Ray Jayawardhana, Amelia Bayo, Laura Flagg, Koraljka Mužić, Antonella Natta, Paola Pinilla, Leonardo Testi
Figure caption: Medium-resolution G140H and G235H JWST NIRSpec spectra from 2025 August 14 (black line) together with the PRISM spectrum from the same epoch (red line). The strongest emission lines detected are identified.
Paper: Numerical Simulations Confirm Wave-induced Shear Mixing in Stellar Interiors
Authors: A. Varghese, R. P. Ratnasingam, L. Ramírez-Galeano, S. Mathis, T. M. Rogers
Figure caption: Diffusion coefficient from the tracer particle simulation (black), along with the theoretical profiles calculated using Equation (3) based on J. P. Zahn (1992; blue) and Equation (5) based on R. J. Garcia Lopez & H. C. Spruit (1991; red) using the horizontal velocity (uh) from the simulation as a function of radius for 3M⊙ (top), 7M⊙(middle), and 20M⊙ (bottom) for the ZAMS models (left, represented with “Z”) and midMS models (right, represented with “m”). The diffusion coefficient calculated from Equation (5) but considering realistic thermal diffusivities from MESA is shown in orange. The green dashed line (Dv,h) is based on Equation (10).
Paper: Numerical Simulations Confirm Wave-induced Shear Mixing in Stellar Interiors
Authors: A. Varghese, R. P. Ratnasingam, L. Ramírez-Galeano, S. Mathis, T. M. Rogers
Figure caption: (a) Initial and (b) final particle distribution from the tracer particle simulations. The color scheme is chosen to visualize the particle motion with time. (c) The trajectory of a single particle over the whole simulation domain at 0.5R⋆ (blue), 0.72R⋆ (red), and 0.7R⋆ (pink). (d) Vorticity from the 2D simulations with positive vorticity (yellow) saturated at 1 × 10−4 s−1 and negative vorticity (blue) saturated at −1 × 10−4 s−1. This is for the 7M⊙ ZAMS model from A. Varghese et al. (2023).
Paper: Numerical Simulations Confirm Wave-induced Shear Mixing in Stellar Interiors
Authors: A. Varghese, R. P. Ratnasingam, L. Ramírez-Galeano, S. Mathis, T. M. Rogers
Figure caption: Initial (blue) and final (orange) particle distributions for the 7M⊙ ZAMS model as a function of grid cells. (a) Distributions centered around 0.8R⋆. The ratio of final to initial widths is σf/σin = 1.001, and the ratio of peak heights is 0.9993. (b) Distributions within the convection zone (cz), normalized for visual comparison. For both panels, the top axis indicates the corresponding stellar radius within which particles are distributed.
Paper: Numerical Simulations Confirm Wave-induced Shear Mixing in Stellar Interiors
Authors: A. Varghese, R. P. Ratnasingam, L. Ramírez-Galeano, S. Mathis, T. M. Rogers
Figure caption: Diffusion coefficient from the tracer particle simulation (black) along with the theoretical profiles calculated using Equation (3) based on J. P. Zahn (1992; blue) and Equation (5) based on R. J. Garcia Lopez & H. C. Spruit (1991; red) using the horizontal velocity (uh) from the simulation as a function of radius for a 1M⊙ model at solar age. The green dashed (Dv,h) line is based on Equation (10).
Paper: Numerical Simulations Confirm Wave-induced Shear Mixing in Stellar Interiors
Authors: A. Varghese, R. P. Ratnasingam, L. Ramírez-Galeano, S. Mathis, T. M. Rogers
Figure caption: Diffusion coefficient from the tracer particle simulation (black) along with the theoretical profiles calculated using Equation (3) based on J. P. Zahn (1992; blue) and Equation (5) based on R. J. Garcia Lopez & H. C. Spruit (1991; red) using the horizontal velocity (uh) from the simulation as a function of radius for the 7M⊙ midMS model at rotation rates (a) 1 × 10−5, (b) 2 × 10−5, (c) 3 × 10−5, (d) 4 × 10−5, and (e) 1 × 10−4 rads−1 (0.13, 0.27, 0.41, 0.55, and 1.37 d−1, respectively).
Paper: Numerical Simulations Confirm Wave-induced Shear Mixing in Stellar Interiors
Authors: A. Varghese, R. P. Ratnasingam, L. Ramírez-Galeano, S. Mathis, T. M. Rogers
Figure caption: The diffusion coefficient from the tracer particle simulation (black) along with the theoretical profile calculated using Equation (C1; pink) using the horizontal velocity (uh) from the simulations for (a) 3M⊙, (b) 7M⊙, and (c) 20M⊙ ZAMS models as a function of radius.
Paper: Numerical Simulations Confirm Wave-induced Shear Mixing in Stellar Interiors
Authors: A. Varghese, R. P. Ratnasingam, L. Ramírez-Galeano, S. Mathis, T. M. Rogers
Figure caption: The diffusion coefficient as a function of total radius at different time differences (τ) for 3M⊙, 7M⊙, and 20M⊙ ZAMS (left, represented with “Z”) and midMS (right, represented with “m”) models. The insets highlight the zoomed-in profiles at the latest τ values.
Paper: Numerical Simulations Confirm Wave-induced Shear Mixing in Stellar Interiors
Authors: A. Varghese, R. P. Ratnasingam, L. Ramírez-Galeano, S. Mathis, T. M. Rogers
Figure caption: The parameter A calculated across masses using Equation (9) as a function of fractional radius.
Paper: Triples as Links between Binary Black Hole Mergers, Their Electromagnetic Counterparts, and Galactic Black Holes
Authors: Smadar Naoz, Zoltán Haiman, Eliot Quataert, Liz Holzknecht
Figure caption: Two examples of the orbital outcome of a recoil kick on an orbit. We consider a system of two low-spin BHs with S1 = S2 = 0.01 and spin–orbit angles of 2°, 1°, Ωs1 = 90°, and Ωs2 = 270°. The mutual inclination is set to 0﹩\mathop{.}\limits^{\unicode{x000b0}}﹩1. Before the kick, the star has a semimajor axis of 20 au and eccentricity e⋆ = 0.9 (right column), and e⋆ = 0.83 (left column). The argument of periapsis of the inner (outer) binary at the time of the merger is 300° (20°). These angles are relevant for the rotation of the various vectors to the invariant plane. See the text for more details. We assume that the kick took place when the star’s phase was at 10° (right column) and 5° (left column). Various post-merger quantities are shown as a function of the mass ratio of the two original BHs, q. The bottom row shows the closest approach, Equation (3). Overplotted are the tidal radius, rcirc, and the Roche radius, rRoche. We expect EM signatures for BHs that approach their stellar companion within these radii. The middle row depicts the post-kick semimajor axis a⋆,n, with the solid red lines, and the pericenter a⋆,n(1 − e⋆,n), with the dashed purple lines. Overplotted are the initial (pre-kick) semimajor axis and pericenter of the star, shown with the solid and dashed gray lines, respectively. Last, Gaia BH1’s semimajor axis and pericenter (aGaia BH1 and rp,Gaia BH1 = aGaia BH1(1 − eGaia BH1)) are also overplotted, shown with the pink and light purple lines (K. El-Badry et al. 2023b). The top row shows the normalized velocity uk = vkick/vr and ﹩\sin \gamma ﹩, where γ is the angle between the radius vector and the post-kick velocity vector. The grid highlights that when uk = 1, γ → 0°, thus Rc → 0. Note that here a lower (higher) pre-kick eccentricity, while keeping all of the other parameters constant, yields a wider, more easily unbound (tighter) post-kick binary. The opposite trend is expected for a kick that takes place at the apocenter; see Equations (5) and (7).
Paper: Triples as Links between Binary Black Hole Mergers, Their Electromagnetic Counterparts, and Galactic Black Holes
Authors: Smadar Naoz, Zoltán Haiman, Eliot Quataert, Liz Holzknecht
Figure caption: An illustration of the system. The three-body system is composed of two stellar-mass BHs and a distant main-sequence stellar tertiary. The angle between the first (second) BH’s spin, S1 (S2) and the inner binary angular momentum LBHs is is1 (is2), and the angle in the plane of the orbit, measured from the semimajor axis, is φ1 (φ2). Finally, the angle between the inner binary angular momentum and the recoil kick vector is α. Post-recoil-kick, there are two possible outcomes: unbound orbits (87% of all systems) and bound orbits (13% of all systems). We highlight that 8% of all systems are bound with a period shorter than 10 yr, which are Gaia-BH-like systems. In 0.2% of all cases, either a bound or an unbound system results in an EM counterpart. These are designated by having the closest approach Rc ≤ rcirc, where 97.6% (2.4%) of all EM-bright sources are on a bound (unbound) orbit when they cross rcirc. About 50% of these systems have a closest approach smaller than rRoche, resulting in a prompt disruption event and yielding an EM counterpart to the GW emission with a time delay of about 10 days. Finally, 50% of the EM signatures have rRoche < Rc ≤ rcirc, which may result in an LMXB. Note that as stars evolve beyond the main sequence and become red giants, their radius expands, yielding a larger fraction of systems undergoing mass transfer events. We discuss these probabilities in detail in Section 4.2.
Paper: Triples as Links between Binary Black Hole Mergers, Their Electromagnetic Counterparts, and Galactic Black Holes
Authors: Smadar Naoz, Zoltán Haiman, Eliot Quataert, Liz Holzknecht
Figure caption: Time-delay distribution for the EM counterpart, assumed to be the time between the merger of the inner BH binary and when the BH remnant–star separation shrinks below rcirc. Both bound (87% of all systems with Rc ≤ rcirc) and unbound systems, as well as those that crossed the star’s Roche limit (∼50% out of all systems with Rc ≤ rcirc), exhibit the same distribution. The average of this distribution is ∼10 days. We note that the fallback time of the bound debris onto the BH is on the order of a day or less; see Equation (19).
Paper: Triples as Links between Binary Black Hole Mergers, Their Electromagnetic Counterparts, and Galactic Black Holes
Authors: Smadar Naoz, Zoltán Haiman, Eliot Quataert, Liz Holzknecht
Figure caption: Post-kick orbital configurations of the proof-of-concept population. The top panel presents the post-kick bound population (13% of all systems) and shows the eccentricity (y-axis) and semimajor axes (x-axis) of the star–BH orbits. Overplotted are the orbital parameters of Gaia BH1 (e.g., K. El-Badry et al. 2023b) and the approximate part of the parameter space potentially detectable by Gaia (defined by having a period up to 10 yr and pericenter larger than rcirc). The bottom panel shows the closest approach (which is the pericenter for bound systems) of all of the systems in the simulation as a function of the angle between the recoil kick vector and the star’s initial (pre-kick) orbital velocity, i.e., θ = vr · vkick/(vrvkick). Overplotted is the Roche radius. The color code depicts the star’s initial (pre-kick) semimajor axis. The points with red edges are those that have Rc ≤ rcirc and are therefore likely to produce EM emission.
Paper: A Disk-originated 329 Day Quasiperiodic Oscillation in the Seyfert 1 Galaxy J1626+5120
Authors: Litao Zhu, Zhongxiang Wang, Dong Zheng, Alok C. Gupta, Ju-Jia Zhang
Figure caption: Left: LSP power spectrum (black solid) from the zg LC of J126+5120. The dashed lines indicate the 3σ, 4σ, and 4.53σ significance levels. Right: zg LC (green data points) of J1626+5120 and a sinusoidal fit (red solid line) to the LC, with the latter used to help illustrate the QPO signal. At P = 329 days, the sinusoid has a semiamplitude of 0.26 mag. The inset shows the first three cycles of the QPO variations, which have a waveform (described with the blue solid line; A. Ingram & M. van der Klis 2015) similar to that seen in XRBs.
Paper: A Disk-originated 329 Day Quasiperiodic Oscillation in the Seyfert 1 Galaxy J1626+5120
Authors: Litao Zhu, Zhongxiang Wang, Dong Zheng, Alok C. Gupta, Ju-Jia Zhang
Figure caption: Left: long-term optical and MIR LCs of J1626+5120, where the data were from CRTS, ZTF, ATLAS, Gaia, and Wide-field Infrared Survey Explorer. Vertical dashed lines indicate the epochs of six spectroscopic observations with different telescopes. Right: six spectra of J1626+5120, which are vertically shifted for clarity. The spectrum-taken dates are given.
Paper: A Disk-originated 329 Day Quasiperiodic Oscillation in the Seyfert 1 Galaxy J1626+5120
Authors: Litao Zhu, Zhongxiang Wang, Dong Zheng, Alok C. Gupta, Ju-Jia Zhang
Figure caption: Time lags of the UV and optical LCs with respect to the X-ray LC determined from cross-correlation analysis. A green dashed curve indicates the time delay expected from a disk continuum reprocessing model.
Paper: A Disk-originated 329 Day Quasiperiodic Oscillation in the Seyfert 1 Galaxy J1626+5120
Authors: Litao Zhu, Zhongxiang Wang, Dong Zheng, Alok C. Gupta, Ju-Jia Zhang
Figure caption: Fit to the DESI spectrum of J1626+5120 with PyQSOFit. FWHM(Hβ) = 4046 ± 44 km s−1 and Lλ(5100 Å) = 1043.8±4.3 erg s−1 were obtained from the fitting.
Paper: A Disk-originated 329 Day Quasiperiodic Oscillation in the Seyfert 1 Galaxy J1626+5120
Authors: Litao Zhu, Zhongxiang Wang, Dong Zheng, Alok C. Gupta, Ju-Jia Zhang
Figure caption: Zoomed-in views of the optical spectra of J1626+5120. Left: Hβ and [O III] region. Right: Hα region. The spectra are vertically shifted for clarity.
Paper: A Disk-originated 329 Day Quasiperiodic Oscillation in the Seyfert 1 Galaxy J1626+5120
Authors: Litao Zhu, Zhongxiang Wang, Dong Zheng, Alok C. Gupta, Ju-Jia Zhang
Figure caption: Swift XRT and UVOT LCs of J1626+5120, for which zg and ao LCs are plotted for comparison. Vertical dashed lines mark epochs of LJT (red and green) and HCT (gold) spectroscopic observations.
Paper: A Disk-originated 329 Day Quasiperiodic Oscillation in the Seyfert 1 Galaxy J1626+5120
Authors: Litao Zhu, Zhongxiang Wang, Dong Zheng, Alok C. Gupta, Ju-Jia Zhang
Figure caption: Periodicity analysis of the zg LC of J1626+5120. Left: color-scaled WWZ power spectrum, with the 3σ and 4σ confidence contours (gray and black dashed lines, respectively) overlaid. The gray dotted region indicates the cone of influence (COI; C. Torrence & G. P. Compo 1998; M. Tarnopolski et al. 2020), within which the results are affected by edge effects. The main QPO signal at ∼329 days is located outside the COI, confirming the detection is reliable. Right: same as the left panel of Figure 3, with the time-averaged WWZ power spectrum (red histogram) plotted.
Paper: A Disk-originated 329 Day Quasiperiodic Oscillation in the Seyfert 1 Galaxy J1626+5120
Authors: Litao Zhu, Zhongxiang Wang, Dong Zheng, Alok C. Gupta, Ju-Jia Zhang
Figure caption: Photon index (Γ) vs. unabsorbed 0.3–10 keV X-ray flux (﹩{F}_{{\rm{X}}}^{u}﹩) for J1626+5120 obtained from Swift XRT observations. The blue solid line represents the best linear fit, which shows a weak softer-when-brighter trend (with a 3.3σ significance): ﹩{F}_{{\rm{X}}}^{u}﹩/10−12 ∼ (0.10 ± 0.03) × Γ + (1.63 ± 0.08). The black dashed line marks the weighted mean of the Γ values (with their uncertainties as the weights), which is ≃1.91.
Paper: A Disk-originated 329 Day Quasiperiodic Oscillation in the Seyfert 1 Galaxy J1626+5120
Authors: Litao Zhu, Zhongxiang Wang, Dong Zheng, Alok C. Gupta, Ju-Jia Zhang
Figure caption: Cross-correlation analysis of different LC pairs, which were used to estimate interband time lags. Each set of panels displays the two LCs being cross-correlated, the resulting cross correlation function, and the CCCD/CCPD lag distributions from PyCCF.
Paper: Observational Evidence for Spin Alignment between Galaxy Groups and Their Central Galaxies
Authors: Peng Wang, 鹏 王
Figure caption: Distributions of key properties for the sample of galaxy groups and their central galaxies. Top right: redshift distribution of galaxy groups. Top middle: distribution of group richness (Ngal). Top left: correlation between group richness and redshift. The red line indicates the best-fit linear relation, with the correlation coefficient r quantifying the strength of the correlation. Bottom left: distribution of group masses and the stellar masses of central galaxies (﹩{\mathrm{log}}\,{M}_{\odot }﹩). Bottom middle: distribution of central galaxy colors (g − r). Bottom right: relation between the central galaxy color (g − r) and the group mass ﹩{\mathrm{log}}\,{M}_{\odot }﹩. The red line shows the best-fit linear relation, while the correlation coefficient r reflects the strength of the linear correlation.
Paper: Observational Evidence for Spin Alignment between Galaxy Groups and Their Central Galaxies
Authors: Peng Wang, 鹏 王
Figure caption: Probability density of the projected alignment angle between the spin of galaxy groups and that of their central galaxies. The solid curves labeled “Method 1-All” and “Method 2-All” show results for all groups. The dotted curves labeled “Method 1-Merger” and “Method 2-Merger” show the corresponding results for merging groups. For clarity, only the 99.73% (3σ) bootstrap confidence intervals for all samples are shown. The “Isotropic” curve denotes the expectation for random spin axes in projection. In the legend, we provide the sample size N, the mean alignment angle 〈θ〉 with its corresponding uncertainty, and, in parentheses, the detection significance relative to isotropy.
Paper: Observational Evidence for Spin Alignment between Galaxy Groups and Their Central Galaxies
Authors: Peng Wang, 鹏 王
Figure caption: Left: same as Figure 2, but for groups with 6 ≤ Ngal < 10 in different redshift bins. Shaded regions indicate the 3σ confidence intervals. The legend lists the redshift range, sample size, mean alignment angle with its 1σ error, and the detection significance relative to random samples. The black dashed line marks the isotropic expectation. Middle: alignment angle as a function of group richness Ngal for several redshift intervals, with error bars showing the 1σ uncertainties. In the legend, N gives the sample sizes for the four richness bins shown by the points on each curve. Right: alignment angle as a function of redshift for different richness bins Ngal, with 1σ error bars. N gives the sample sizes for the four redshift bins shown by the points on each curve.
Paper: Observational Evidence for Spin Alignment between Galaxy Groups and Their Central Galaxies
Authors: Peng Wang, 鹏 王
Figure caption: Dependence of the mean projected alignment angle between group spin and central galaxy spin on various physical properties. Top left: mean angle as a function of group mass (log M⊙). Top right: mean angle as a function of central galaxy stellar mass (log M⊙). The red dashed line in each panel indicates the expected value for random (isotropic). Bottom left: mean angle for elliptical and spiral central galaxies. Bottom right: mean angle as a function of central galaxy color (g − r). All error bars represent the standard error of the mean in each bin or group.
Paper: Tidal Disruption Event Demographics in Supermassive Black Hole Binaries over Cosmic Times
Authors: Denyz Melchor, Smadar Naoz, Suvi Gezari, Brenna Mockler
Figure caption: Projected TDE rate per SMBH mass. The pink curve shows the predicted upper-limit TDE rate as a function of disrupting SMBH mass for a binary with fixed mass ratio (q = 100) and a core-like stellar density profile. The solid and dashed black lines show single-SMBH TDE rates from two-body relaxation assuming core-like and cusp-like profiles, respectively. The orange line represents the TDE rate estimated by B. Rom & R. Sari (2025). The hatched gray band shows the observed TDE rate in PSB galaxies (K. D. French et al. 2020a), while the solid gray band denotes the average observed TDE rate (S. van Velzen et al. 2020). The green line marks the TDE rate in green-valley galaxies estimated by Y. Yao et al. (2023). Disruptions of Sun-like stars are not expected beyond MBH ∼ 108M⊙ (e.g., N. C. Stone et al. 2018; T. Wevers et al. 2019a; S. Gezari 2021; M. Polkas et al. 2024), so a decline in the TDE rate is expected at higher masses.
Paper: Tidal Disruption Event Demographics in Supermassive Black Hole Binaries over Cosmic Times
Authors: Denyz Melchor, Smadar Naoz, Suvi Gezari, Brenna Mockler
Figure caption: Correlation between TDE light curves and black hole mass. Scatter points show τdecline, measured from the light-curve decay timescale of individual TDEs, plotted against the estimated black hole mass (MBH) for each event. The blue shaded region indicates the parameter space where SMBH candidates may form TDEs. Left panel: TDE light-curve properties adapted from E. Hammerstein et al. (2022), and TDE black hole masses adapted from A. Mummery et al. (2024). Right panel: TDE light-curve properties adapted from S. van Velzen et al. (2019a), and TDE black hole masses adapted from A. Mummery et al. (2024). For both panels, the mass of the black hole is derived from MBH–σ⋆ when available, and from MBH–Mgal; derivations can be found in A. Mummery et al. (2024). The orange line in both panels indicates the best-fit line with slope matching the tfb theoretical line. The yellow range represents a 1σ deviation. In both plots, the systematic errors as stated in A. Mummery et al. (2024) are shown in the corners: approximately ∼0.3 dex and ∼0.8 dex for MBH–σ⋆ (upside-down triangles) and MBH–Mgal (circles), respectively.
Paper: Tidal Disruption Event Demographics in Supermassive Black Hole Binaries over Cosmic Times
Authors: Denyz Melchor, Smadar Naoz, Suvi Gezari, Brenna Mockler
Figure caption: Upper-limit TDE redshift evolution. Top panels: comoving volumetric TDE rate as a function of redshift z. Left: predicted rate from SMBH binaries (thick solid line) across different masses: 105M⊙ (pink), 106M⊙ (green), 107M⊙ (blue), and 108M⊙ (orange). Middle: single-SMBH rates with core-like (solid) and cusp-like (dashed) stellar profiles, across the same mass bins. Right: rates integrated over SMBH masses from 105 to 108M⊙, comparing the different models. Bottom panels: cumulative number of TDEs per unit redshift for the same models: left (binaries), middle (single-SMBH core/cusp), right (mass-integrated comparison).