adsabs/AstroCLIMB huggingface dataset:
Paper: The Detection of Higher-order Millimeter Hydrogen Recombination Lines in the Large Magellanic Cloud
Authors: Marta Sewiło, Kazuki Tokuda, Stan E. Kurtz, Steven B. Charnley, Thomas Möller, Jennifer Wiseman, C.-H. Rosie Chen, Remy Indebetouw, Álvaro Sánchez-Monge, Kei E. I. Tanaka, Peter Schilke, Toshikazu Onishi, Naoto Harada
Figure caption: From left to right: Band 3 H40α, H50β, and H57γ recombination line integrated intensity images of N 105–1 A. The white contours are the same as in Figure 4.
Paper: The Detection of Higher-order Millimeter Hydrogen Recombination Lines in the Large Magellanic Cloud
Authors: Marta Sewiło, Kazuki Tokuda, Stan E. Kurtz, Steven B. Charnley, Thomas Möller, Jennifer Wiseman, C.-H. Rosie Chen, Remy Indebetouw, Álvaro Sánchez-Monge, Kei E. I. Tanaka, Peter Schilke, Toshikazu Onishi, Naoto Harada
Figure caption: The SED for N 105–1 A, covering a wavelength range from 1.25 μm to 2.7 mm. Filled circles and triangles are valid flux values and flux upper limits, respectively. The flux error bars are smaller than the data points. The photometry and references are provided in Table 2 in Appendix C.
Paper: Uncovering Hidden Massive Black Hole Companions with Tidal Disruption Events
Authors: Brenna Mockler, Denyz Melchor, Smadar Naoz, Enrico Ramirez-Ruiz
Figure caption: Left: the number of TDEs as a function of m 1, with different density profiles represented as different marker shapes (q = 10 in all cases). The points in orange are the lower limits, excluding all stars that move outside r Hill, while the points in blue are the upper limits and include these stars. The blue and orange shaded regions are plotted by determining the slopes between the fractions of TDEs for the α = 1 and 2 runs at the black hole masses used in the simulations and then extrapolating out to higher and lower masses. The black hatched region is the number of TDEs disrupted by m 2 in the same system from chaotic scatterings (the upper limit is from Liu et al. 2013 and the lower limit is scaled down to match Wegg & Nate Bode 2011). The gray shaded region is the number of TDEs predicted from two-body relaxation over 108 yr for a single-SMBH system with mass m 1 (from Stone & Metzger 2016). Right: the number of TDEs as a function of the number of stars within the hierarchical radius. Each m 1 and α combination is plotted as a different color, and the marker shape is determined by the mass ratio q. The upper and lower limits are determined in the same way as for the left panel.
Paper: Ab-initio General-relativistic Neutrino-radiation Hydrodynamics Simulations of Long-lived Neutron Star Merger Remnants to Neutrino Cooling Timescales
Authors: David Radice, Sebastiano Bernuzzi
Figure caption: Angularly averaged profiles of entropy per baryon s; electron fraction Y  e ; convection instability criterion χ; neutrino fractions ﹩{Y}_{{\nu }_{e}}﹩, ﹩{Y}_{{\bar{\nu }}_{e}}﹩, and ﹩{Y}_{{\nu }_{\mu }}﹩; and angular frequency Ω on the equatorial plane for the ℓ mix = 0 SR binary. The profiles are shown as a function of cylindrical radius ϱ and time from merger t − t mrg. The purple contour denotes ρ = 1013 g · cm−3, while the green contours are ρ = 1013.5, 1014, and 1014.5 g · cm−3.
Paper: Ab-initio General-relativistic Neutrino-radiation Hydrodynamics Simulations of Long-lived Neutron Star Merger Remnants to Neutrino Cooling Timescales
Authors: David Radice, Sebastiano Bernuzzi
Figure caption: Baryonic mass and angular momentum at the end of the simulations for the SR binary. The gray shaded area shows the set of all rigidly rotating equilibrium configurations. The green line shows the evolution of the angular momentum due to GW emission, while the red and blue lines show the evolution of the baryonic mass and angular momentum for the disk and the remnant, respectively. The remnant settles to a region corresponding to a stable, rigidly rotating configuration, but it is still rotating differentially at the end of our simulations.
Paper: Ab-initio General-relativistic Neutrino-radiation Hydrodynamics Simulations of Long-lived Neutron Star Merger Remnants to Neutrino Cooling Timescales
Authors: David Radice, Sebastiano Bernuzzi
Figure caption: Angularly averaged profiles of entropy per baryon s; electron fraction Y  e ; convection instability criterion χ; neutrino fractions ﹩{Y}_{{\nu }_{e}}﹩, ﹩{Y}_{{\bar{\nu }}_{e}}﹩, and ﹩{Y}_{{\nu }_{\mu }}﹩; and angular frequency Ω on the equatorial plane for the ℓ mix(ρ) SR binary. The profiles are shown as a function of cylindrical radius ϱ and time from merger t − t mrg. The purple contour denotes ρ = 1013 g · cm−3, while the green contours are ρ = 1013.5, 1014, and 1014.5 g · cm−3. This figure should be contrasted with Figure 3, which shows the same quantities for the ℓ mix = 0 binary.
Paper: Ab-initio General-relativistic Neutrino-radiation Hydrodynamics Simulations of Long-lived Neutron Star Merger Remnants to Neutrino Cooling Timescales
Authors: David Radice, Sebastiano Bernuzzi
Figure caption: Profiles of entropy per baryon s, electron fraction Y  e , convection instability criterion χ, and neutrino fractions ﹩{Y}_{{\nu }_{e}}﹩, ﹩{Y}_{{\bar{\nu }}_{e}}﹩, and ﹩{Y}_{{\nu }_{\mu }}﹩ along the z-axis for the ℓ mix = 0 SR binary. The profiles are shown as a function of time from merger t − t mrg. The purple contour denotes ρ = 1013 g · cm−3, while the green contours are ρ = 1013.5, 1014, and 1014.5 g · cm−3. This figure should be contrasted with Figure 3, which shows the same quantities in the equatorial plane.
Paper: Ab-initio General-relativistic Neutrino-radiation Hydrodynamics Simulations of Long-lived Neutron Star Merger Remnants to Neutrino Cooling Timescales
Authors: David Radice, Sebastiano Bernuzzi
Figure caption: GW luminosity and neutrino luminosity and average energies. The data are smoothed using a running average with window width Δt = 0.5 ms. The GW luminosity is shown for the three most dominant modes: (ℓ, m) = (2, 1), (2, 2), and (3, 3). ﹩{L}_{{\nu }_{\mu }}﹩ is the luminosity of one of the heavy-lepton neutrino species. We do not track separately ν  μ , ν  τ , and their antiparticles, so the overall neutrino luminosity of the four species combined is ﹩4\times {L}_{{\nu }_{\mu }}﹩.
Paper: Ab-initio General-relativistic Neutrino-radiation Hydrodynamics Simulations of Long-lived Neutron Star Merger Remnants to Neutrino Cooling Timescales
Authors: David Radice, Sebastiano Bernuzzi
Figure caption: GW and neutrino cooling timescales. The data are smoothed using a running average with window width Δt = 0.5 ms. Neutrino radiation becomes the dominant mechanism for the evolution of the remnant ∼10 ms after merger.
Paper: Ab-initio General-relativistic Neutrino-radiation Hydrodynamics Simulations of Long-lived Neutron Star Merger Remnants to Neutrino Cooling Timescales
Authors: David Radice, Sebastiano Bernuzzi
Figure caption: Time evolution of diagnostic quantities in the simulations: minimum of the lapse function, ﹩{\alpha }_{\min }﹩; maximum rest-mass density, ﹩{\rho }_{\max }﹩; and maximum temperature T max. The dynamical phase of evolution terminates ∼50 ms after merger.
Paper: Ab-initio General-relativistic Neutrino-radiation Hydrodynamics Simulations of Long-lived Neutron Star Merger Remnants to Neutrino Cooling Timescales
Authors: David Radice, Sebastiano Bernuzzi
Figure caption: Mixing length prescription used for the viscous simulations. The data points are obtained by combining the values of α reported by Kiuchi et al. (2018, 2023) with angular velocities and sound speed values of Radice (2020). The solid line is the fitting function adopted here.
Paper: Ab-initio General-relativistic Neutrino-radiation Hydrodynamics Simulations of Long-lived Neutron Star Merger Remnants to Neutrino Cooling Timescales
Authors: David Radice, Sebastiano Bernuzzi
Figure caption: Baryonic mass and angular momentum at the end of the simulations for the ℓ mix(ρ) SR binary. The gray shaded area shows the set of all rigidly rotating equilibrium configurations. The green line shows the evolution of the angular momentum due to GW emission, while the red and blue lines show the evolution of the baryonic mass and angular momentum for the disk and the remnant, respectively. This figure should be contrasted with Figure 4, which shows the same quantities for the ℓ mix = 0 binary.
Paper: Monster Radiative Shocks in the Perturbed Magnetospheres of Neutron Stars
Authors: Andrei M. Beloborodov
Figure caption: Left: evolution of the wave profile E(ξ) and κ(ξ), where ξ ≡ t − r/c. The wave has frequency ν = ω/2π = 10  kHz and initial power L = 1043  erg s−1; the magnetosphere has magnetic dipole moment μ = 1033 G cm3 and density parameter ﹩{ \mathcal N }={10}^{39}﹩. Five snapshots are shown, when the wave packet reaches r/R × = 0.9 (black), 1 (red), 1.1 (green), 1.5 (blue), and 3.7 (magenta). The electric field E is normalized to E 0, which would be the wave amplitude if it propagated in vacuum. The plasma Lorentz factor γ is related to the proper compression κ by γ = (1 + κ 2)/2κ. Black dotted curves show the analytical result for κ(ξ) (Equation (78)). The simulation neglected the shock precursor effect, which can reduce γ in the interval 3π/2ω < ξ < ξ sh (Section 5.5). Right: same wave but now launched into the magnetosphere with ﹩{ \mathcal N }={10}^{37}﹩. Here, κ becomes extremely small, breaking the MHD description and transitioning to a two-fluid regime. We argue in Section 5.3 that the two-fluid calculation will likely give the same evolution of E(ξ) and κ(ξ) as found in the MHD model.
Paper: Monster Radiative Shocks in the Perturbed Magnetospheres of Neutron Stars
Authors: Andrei M. Beloborodov
Figure caption: A magnetosonic wave with wavevector  k ⊥ B  bg at two oscillation phases: when the wave field  B  w =  B  −  B  bg is aligned with  B  bg (left) and antialigned with  B  bg (right). The wave electric field  E  w =  E  is always orthogonal to  B  bg. The quantity B 2 − E 2 reaches zero when  B  w = −  B  bg/2. At this moment, the plasma drift speed  v  = c  E  ×  B /B 2 approaches c, so it experiences ultrarelativistic acceleration in the direction opposite to the wave propagation direction.
Paper: Monster Radiative Shocks in the Perturbed Magnetospheres of Neutron Stars
Authors: Andrei M. Beloborodov
Figure caption: Excitation of magnetospheric waves by a star quake with ω ≫ c/R ⋆ and a vertical (radial) wavevector  k . The wave polarization is set by the direction of the electric field  E  =  B  ×  v /c in the crustal shear oscillation with a horizontal velocity  v . Left: crustal motion with  v ∥ n  ≡  k  ×  B  bg excites an Alfvén wave  E ⊥ n . Right:  v ⊥ n  excites a magnetosonic wave  E ∥ n . The wave amplitude is small, so  B  ≈  B  bg near the star.
Paper: Monster Radiative Shocks in the Perturbed Magnetospheres of Neutron Stars
Authors: Andrei M. Beloborodov
Figure caption: Shock structure solution in the approximation of steady e ± streams (viewed in the shock rest frame). The solution was calculated for χ = 0, 103, 105, and 107. Left: for each case, we show g(s) (black), b(s) (blue), and ﹩{\gamma }_{{\rm{D}}}(s)=B/\sqrt{{B}^{2}-{E}^{2}}﹩ normalized to ﹩\sqrt{{\sigma }_{{\rm{u}}}}﹩ (red). Flows with χ ≫ 1 radiate most of their energy before developing gyration; this results in γ ≪ γ u. Each solution ends where the downstream begins, i.e., where the e ± streams develop gyration (then they become unstable, thermalizing their gyration energy in the drift frame). The drift Lorentz factor γ D at the end point represents the downstream speed relative to the shock. Right: relation between g and b. The approximate analytical solution (Equation (B19)) is shown by the dotted curve and compared with accurate numerical integration (solid curve). Both numerical and analytical solutions are plotted until the e ± streams reach ﹩| {\tilde{\beta }}_{x}| =1/2﹩, which corresponds to 60° gyration in the drift frame.
Paper: Monster Radiative Shocks in the Perturbed Magnetospheres of Neutron Stars
Authors: Andrei M. Beloborodov
Figure caption: Schematic illustration of the flow of characteristics C + in FFE and MHD, shown in the t–ξ plane. Each C + has its initial position ξ i and carries a value of ﹩K({\xi }_{{\rm{i}}})={rE}/\mu ﹩ determined by the initial wave profile. In the FFE limit, the characteristics are vertical straight lines; they propagate with speed β + = 1, which corresponds to d ξ +/dt = 0. The MHD correction to FFE implies d ξ +/dt = 1 − β + ≠ 0, so the C + characteristics are no longer static in ξ; they become significantly bent in the region where E 2 approaches B 2, leading to the formation of a shock (red curve).
Paper: Monster Radiative Shocks in the Perturbed Magnetospheres of Neutron Stars
Authors: Andrei M. Beloborodov
Figure caption: Expansion of a spherical wave front through the magnetosphere. The front thickness (wave duration) is much smaller than r/c. Its shock-free part (﹩\sin \theta \lt \sin {\theta }_{\times }﹩) is plotted in green and the part carrying the shock is in red. The region swept by the shock is shaded in pink; its boundary θ(r) is defined by Equation (110). The wave front is shown at two moments: (1) when it crosses radius R × (the shock has formed in the equatorial plane, θ = π/2); and (2) when it expands to a larger radius (the shock now occupies a range of θ). The shock radiates gamma rays (red arrows) beamed along  β  D (Equation (111)); most of them convert to e ± pairs.
Paper: Monster Radiative Shocks in the Perturbed Magnetospheres of Neutron Stars
Authors: Andrei M. Beloborodov
Figure caption: Evolution of shock parameters in the simulation with ﹩{ \mathcal N }={10}^{39}﹩: upstream Lorentz factor γ u, shock compression factor q ≡ κ d/κ u, and radiative parameter χ. Red dotted curves show the analytical results for γ u(r), q(r), χ(r) derived in Section 5 (Equations (75), (92), and (94)). The simulation neglected the precursor effect (Section 5.5).
Paper: Monster Radiative Shocks in the Perturbed Magnetospheres of Neutron Stars
Authors: Andrei M. Beloborodov
Figure caption: Lorentz factor of the shock relative to the downstream plasma as a function of the radiative parameter χ (solid curve, calculated as explained in Appendix B). The dashed line shows the analytical result at χ ≫ 1 (Equation (54)).
Paper: Exploring the Origin of Stars on Bound and Unbound Orbits Causing Tidal Disruption Events
Authors: Shiyan Zhong, Kimitake Hayasaki, Shuo Li, Peter Berczik, Rainer Spurzem
Figure caption: Distribution of the stars that can cause TDEs on the e–β plane (left: M BH = 0.01; right: M BH = 0.05). The solid magenta curve denotes ﹩{e}_{\max }﹩. The ﹩{e}_{\min }﹩ (solid) and e lcb (dashed) lines are plotted with two different Q boost values that are indicated by the colors, respectively (see also Equations (22)–(24)).