adsabs/AstroCLIMB huggingface dataset:
Paper: The Gamma-Ray Pulsar Phenomenology in View of 3D Kinetic Global Magnetosphere Models
Authors: Constantinos Kalapotharakos, Zorawar Wadiasingh, Alice K. Harding, Demosthenes Kazanas
Figure caption: The normalized particle distributions of the encountered voltage, where dN denotes the number of particles within V and V + dV, and N 0 denotes the total number of particles in the simulation. The V values are measured in the voltage across the polar cap. Each panel corresponds to the indicated α values, while the indicated colors denote the different model numbers. The more FF a model, the lower the cutoff of the V distributions. The black-dashed line in the middle panel corresponds to a simulation of w = 0.12 and ﹩{{ \mathcal F }}_{\mathrm{SZ}}=1.5{{ \mathcal F }}_{\mathrm{GJ}}^{0}﹩, which is equivalent to model No. 4, i.e., w = 0.04, and ﹩{{ \mathcal F }}_{\mathrm{SZ}}=1.0{{ \mathcal F }}_{\mathrm{GJ}}^{0}﹩. The magenta line in the top panel denotes a simulation equivalent to model No. 6 in which particle injection occurs in the high E acc region beyond the Y-point (see Section 3.4).
Paper: The Gamma-Ray Pulsar Phenomenology in View of 3D Kinetic Global Magnetosphere Models
Authors: Constantinos Kalapotharakos, Zorawar Wadiasingh, Alice K. Harding, Demosthenes Kazanas
Figure caption: Similar to Figure 8 but for α = 75°.
Paper: The Gamma-Ray Pulsar Phenomenology in View of 3D Kinetic Global Magnetosphere Models
Authors: Constantinos Kalapotharakos, Zorawar Wadiasingh, Alice K. Harding, Demosthenes Kazanas
Figure caption: The physically realistic energy distributions of the high-energy particles for the separatrix zone model and α = 45°. The top (bottom) row shows YP (MP) models. The left-hand column shows distributions for the same ﹩\dot{{ \mathcal E }}﹩ value and different model numbers, i.e., different FF-ness. The right-hand column shows distributions for the same model number and for different ﹩\dot{{ \mathcal E }}﹩ values.
Paper: The Gamma-Ray Pulsar Phenomenology in View of 3D Kinetic Global Magnetosphere Models
Authors: Constantinos Kalapotharakos, Zorawar Wadiasingh, Alice K. Harding, Demosthenes Kazanas
Figure caption: The model SED spectra are denoted by the solid lines for the cases shown in Figure 5. The dashed lines denote the corresponding best fits assuming the model function in Equation (5) for b = 1.
Paper: The Gamma-Ray Pulsar Phenomenology in View of 3D Kinetic Global Magnetosphere Models
Authors: Constantinos Kalapotharakos, Zorawar Wadiasingh, Alice K. Harding, Demosthenes Kazanas
Figure caption: The high accelerating regions for the indicated α and model number values. The color scale denotes the accelerating electric-field components.
Paper: The Gamma-Ray Pulsar Phenomenology in View of 3D Kinetic Global Magnetosphere Models
Authors: Constantinos Kalapotharakos, Zorawar Wadiasingh, Alice K. Harding, Demosthenes Kazanas
Figure caption: Similar to Figure 8 but for α = 45°.
Paper: Hadronic versus Leptonic Origin of Gamma-Ray Emission from Supernova Remnants
Authors: Nicholas J. Corso, Rebecca Diesing, Damiano Caprioli
Figure caption: Cumulative hadronic γ-ray spectra (blue lines) and leptonic γ-ray spectra (red lines) at the transition between the ejecta-dominated and Sedov–Taylor stages of a modeled SNR in a sample of environments. Each panel represents a different ambient medium, with n 0 denoting the matter density normalization and u rad denoting the radiation energy density. As in Figure 1, we consider both homogeneous and wind profiles for the ambient density. Vertical-black dashed lines mark 100 GeV, the dividing energy between our GeV and TeV bands. The top (bottom) panels represent hadronic (leptonic) cases with the extremity of the scenario (i.e., the absolute value of the hadronicity, H) increasing from left-hand to right-hand. Note that the nature of the underlying particle acceleration remains the same across panels. The strong variation in γ-ray emission shown here arises solely from environmental factors.
Paper: Hadronic versus Leptonic Origin of Gamma-Ray Emission from Supernova Remnants
Authors: Nicholas J. Corso, Rebecca Diesing, Damiano Caprioli
Figure caption: Hadronicity as a function of SNR age and ambient density (n 0). The model results are presented as data points, while the contours in the background are generated by 2D interpolation. The left-hand (right-hand) columns correspond to the GeV (TeV) bands, while the top (bottom) rows correspond to homogeneous (wind) density profiles. The black lines represent the Sedov–Taylor times for each density profile. In general, young SNRs tend to be the most hadronic. In all cases, the radiation field (CMB and moderate stellar background) has energy density u rad = 10.3 eV cm−3.
Paper: Hadronic versus Leptonic Origin of Gamma-Ray Emission from Supernova Remnants
Authors: Nicholas J. Corso, Rebecca Diesing, Damiano Caprioli
Figure caption: Cumulative proton spectra (solid lines) and electron spectra (dashed lines) at the transition between the ejecta-dominated and Sedov–Taylor stages of a modeled SNR in a sample of environments. Line colors denote the density normalization used in each model. In the left-hand panel, the ambient medium is taken to be homogeneous; in the right-hand panel, it follows a wind profile. The electron/proton normalization is fixed to K ep = 1.6 × 10−3, as discussed in the text.
Paper: Hadronic versus Leptonic Origin of Gamma-Ray Emission from Supernova Remnants
Authors: Nicholas J. Corso, Rebecca Diesing, Damiano Caprioli
Figure caption: Hadronicity as a function of matter (n 0) and background radiation energy density (u rad). The model results are presented as data points, while the contours in the background are generated by 2D interpolation. The left-hand (right-hand) columns correspond to the GeV (TeV) bands, while the top (bottom) rows correspond to homogeneous (wind) density profiles. Broadly speaking, hadronic emission dominates in environments with high ambient density and low radiation energy density.
Paper: Hadronic versus Leptonic Origin of Gamma-Ray Emission from Supernova Remnants
Authors: Nicholas J. Corso, Rebecca Diesing, Damiano Caprioli
Figure caption: GeV hadronicity as a function of the relative normalization of the electron spectrum, K ep, and ambient density for homogeneous profiles. The total radiation field is fixed as u rad = 10.3 eV cm−3, as was done in Figure 5. Exterior regions with hatch marks bound the region K ep ∼ 10−4 to 10−2, which broadly reflects the observed range of values.
Paper: Hadronic versus Leptonic Origin of Gamma-Ray Emission from Supernova Remnants
Authors: Nicholas J. Corso, Rebecca Diesing, Damiano Caprioli
Figure caption: Time evolution of hadronic and leptonic luminosities for a moderately hadronic homogeneous profile scenario (left-hand) and a moderately leptonic wind profile scenario (right-hand). The environmental parameters are the same as those used in the middle column of Figure 2. The top (bottom) row corresponds to luminosities calculated in the GeV (TeV) band. The vertical-black dashed lines denote the onset of the Sedov–Taylor phase.
Paper: Electron Acceleration and Heating during Magnetic Reconnection in the Earth's Quasi-parallel Bow Shock
Authors: N. Bessho, L.-J. Chen, M. Hesse, J. Ng, L. B. Wilson, J. E. Stawarz
Figure caption: Time evolution of the ion-scale magnetic island, the electric field E  x  and E  y , and the electron trajectory (the same electron as shown in Figure 11). The light blue curves are magnetic field lines, and the grayscale shows E  x  (panels (a), (c), and (e)) and E  y  (panels (b), (d), and (f)). The red curves are the electron trajectory, and yellow dots surrounded by green, blue, and red circles are the positions at each specific time indicated by the numbers in yellow rectangles. While this electron is going around the ion-scale island, it is energized mainly by E  y .
Paper: The Single-pulse Observation of PSR B2111+46 with the Five-hundred-meter Aperture Spherical Radio Telescope
Authors: Q. J. Zhi, J. T. Bai, L. H. Shang, X. Xu, S. J. Dang, D. Li, L. Zhang, P. Wang, X. Y. Xie, R. S. Zhao, A. J. Dong, G. J. Qiao
Figure caption: The single-pulse profiles of PSR B2111+46.
Paper: Electron Acceleration and Heating during Magnetic Reconnection in the Earth's Quasi-parallel Bow Shock
Authors: N. Bessho, L.-J. Chen, M. Hesse, J. Ng, L. B. Wilson, J. E. Stawarz
Figure caption: (a)–(c) The trajectory of another electron accelerated in an ion-scale island in the x–y plane. (d) The trajectory in the x–z plane. (e)–(h) Time evolution of γ, W ⊥, W ∥, W total, p ⊥, p ∥, W  x , W  y , and W  z . Depicted with the same format as in Figure 11.
Paper: Electron Acceleration and Heating during Magnetic Reconnection in the Earth's Quasi-parallel Bow Shock
Authors: N. Bessho, L.-J. Chen, M. Hesse, J. Ng, L. B. Wilson, J. E. Stawarz
Figure caption: (a) Time evolution of magnetic field lines in the x–y plane. The red oval regions show where the nonresonant wave grows, while the blue oval regions show where the resonant wave grows. In the plot at Ω i  t = 15.63, the light blue lines represent wave planes of the nonresonant wave. (b) B  y  (red) and B  z  (blue) along y = 20d  i  at Ω i  t = 15.63. Inset: the ion distribution function at x = 55d  i  and y = 20d  i . The white dashed line shows the direction of the magnetic field. (c) 〈B  y 〉 (top) and 〈B  z 〉 (bottom) as functions of x at each time. Inset: the ion distribution function at x = 100d  i  integrated over y.
Paper: Electron Acceleration and Heating during Magnetic Reconnection in the Earth's Quasi-parallel Bow Shock
Authors: N. Bessho, L.-J. Chen, M. Hesse, J. Ng, L. B. Wilson, J. E. Stawarz
Figure caption: (a) (Red) The trajectory of an electron accelerated by the pseudo-potential in the x–y plane. (Black) Magnetic field line. (b) The trajectory in the x–z plane. (c) Time evolution of γ. (d) Time evolution of the work done by the perpendicular electric field (W ⊥ in red), the work done by the parallel electric field (W ∥ in blue), and the sum of these two (W total in green). The parallel work W ∥ is the dominant source of the energy increase.
Paper: Electron Acceleration and Heating during Magnetic Reconnection in the Earth's Quasi-parallel Bow Shock
Authors: N. Bessho, L.-J. Chen, M. Hesse, J. Ng, L. B. Wilson, J. E. Stawarz
Figure caption: Diagrams to explain the maximum energy gain in shock drift acceleration. Panels (a) and (b) show reflected electrons and passing electrons, respectively. See the text in Appendix B for more details.
Paper: The Single-pulse Observation of PSR B2111+46 with the Five-hundred-meter Aperture Spherical Radio Telescope
Authors: Q. J. Zhi, J. T. Bai, L. H. Shang, X. Xu, S. J. Dang, D. Li, L. Zhang, P. Wang, X. Y. Xie, R. S. Zhao, A. J. Dong, G. J. Qiao
Figure caption: The pulse stack of 200 single pulses of PSR B2111+46 obtained from the observation of the FAST. The center frequency is 1250 MHz. From left to right, we show Stokes I, linear polarization, and Stokes V, normalized by their peak values.
Paper: Electron Acceleration and Heating during Magnetic Reconnection in the Earth's Quasi-parallel Bow Shock
Authors: N. Bessho, L.-J. Chen, M. Hesse, J. Ng, L. B. Wilson, J. E. Stawarz
Figure caption: (a) The trajectory of an electron accelerated by shock drift acceleration in the x–y plane. (b) The trajectory in the x–z plane. (c), (d) The time evolution of γ, W ⊥, W ∥, and W total. Depicted with the same format as in Figure 14.
Paper: An Unsupervised Machine Learning-based Algorithm for Detecting Weak Impulsive Narrowband Quiet Sun Emissions and Characterizing Their Morphology
Authors: Shabbir Bawaji, Ujjaini Alam, Surajit Mondal, Divya Oberoi, Ayan Biswas
Figure caption: Left-hand panel: an example 132 MHz radio image is shown in contours, overlaid on an AIA 193 Å  image taken near the middle of the observing duration on 2017 November 27. The contour levels are 0.006, 0.01, 0.04, 0.08, 0.16, 0.32, and 0.64 times the peak in the radio image. The bright compact emission comes the type I noise storm associated with the only active region present on the visible solar disk. Right-hand panel: the contours show an example 133 MHz radio image, overlaid on an AIA 193 Å  image taken near the middle of the observation on 2020 June 20. The contours are at 0.04, 0.08, 0.2, 0.4, 0.6, 0.72, 0.74, 0.76, 0.78, 0.8, and 0.9 times the peak in the radio image.
Paper: Alignments between Galaxies and the Cosmic Web at z ∼ 1–2 in the IllustrisTNG Simulations
Authors: Benjamin Zhang, Khee-Gan Lee, Alex Krolewski, Jingjing Shi, Benjamin Horowitz, Robin Kooistra
Figure caption: Comparison of Horizon-AGN ideal spin alignments from Figure 3 of Codis et al. (2015; black) with the corresponding spin alignments from TNG300-1 (blue) and TNG100-1 (orange). The TNG samples are derived from the samples described in Section 2.1.2, with an additional minimum-mass cut of M * > 108.06 h −1 M ⊙. Error bars for the TNG300-1/TNG100-1 alignments are derived from bootstrap sampling, while errors on Codis et al. (2015) are Poisson, taken from the original figure.
Paper: Alignments between Galaxies and the Cosmic Web at z ∼ 1–2 in the IllustrisTNG Simulations
Authors: Benjamin Zhang, Khee-Gan Lee, Alex Krolewski, Jingjing Shi, Benjamin Horowitz, Robin Kooistra
Figure caption: Shape–cosmic web alignment signal for all galaxies with radius-limited stellar mass M * > 1010 h −1 M ⊙ as a function of total stellar mass. For each redshift, the top panel shows the alignment signal calculated in total stellar mass bins. In the bottom panel, the alignment mean is calculated as a function of a sliding stellar mass bin. Error bars are calculated using bootstrap sampling, with error bars/shaded regions indicating the 16th–84th percentile ranges. The subscript on θ  s  indicates that we are calculating shape alignment.
Paper: Alignments between Galaxies and the Cosmic Web at z ∼ 1–2 in the IllustrisTNG Simulations
Authors: Benjamin Zhang, Khee-Gan Lee, Alex Krolewski, Jingjing Shi, Benjamin Horowitz, Robin Kooistra
Figure caption: Spin–cosmic web alignment signal for all galaxies with more than 50 stellar particles as a function of total stellar mass. For each redshift, the top panel shows the detailed alignment signal calculated in total stellar mass bins. The alignment mean is calculated as a function of a sliding stellar mass bin in the bottom panel. The error bars are calculated using bootstrap sampling, with error bars/shaded regions indicating the 16th–84th percentile ranges. The subscript on θ  j  indicates that we are calculating spin alignment.
Paper: Early Planet Formation in Embedded Disks (eDisk). IX. High-resolution ALMA Observations of the Class 0 Protostar R CrA IRS5N and Its Surroundings
Authors: Rajeeb Sharma, Jes K. Jørgensen, Sacha Gavino, Nagayoshi Ohashi, John J. Tobin, Zhe-Yu Daniel Lin, Zhi-Yun Li, Shigehisa Takakuwa, Chang Won Lee, Jinshi Sai (Insa Choi), Woojin Kwon, Itziar de Gregorio-Monsalvo, Alejandro Santamaría-Miranda, Hsi-Wei Yen, Yuri Aikawa, Yusuke Aso, Shih-Ping Lai, Jeong-Eun Lee, Leslie W. Looney, Nguyen Thi Phuong, Travis J. Thieme, Jonathan P. Williams
Figure caption: Same as Figure A7, but for H2CO (32,2–22,1).
Paper: Early Planet Formation in Embedded Disks (eDisk). IX. High-resolution ALMA Observations of the Class 0 Protostar R CrA IRS5N and Its Surroundings
Authors: Rajeeb Sharma, Jes K. Jørgensen, Sacha Gavino, Nagayoshi Ohashi, John J. Tobin, Zhe-Yu Daniel Lin, Zhi-Yun Li, Shigehisa Takakuwa, Chang Won Lee, Jinshi Sai (Insa Choi), Woojin Kwon, Itziar de Gregorio-Monsalvo, Alejandro Santamaría-Miranda, Hsi-Wei Yen, Yuri Aikawa, Yusuke Aso, Shih-Ping Lai, Jeong-Eun Lee, Leslie W. Looney, Nguyen Thi Phuong, Travis J. Thieme, Jonathan P. Williams
Figure caption: Same as Figure A7, but for H2CO (32,1–22,0).
Paper: Giant Impact Events for Protoplanets: Energetics of Atmospheric Erosion by Head-on Collision
Authors: Kenji Kurosaki, Shu-ichiro Inutsuka
Figure caption: Snapshots of impact simulations. The leftmost panels are color-coded diagrams of lost and bound particles. Red indicates lost atmospheric particles. Green indicates bound atmospheric particles. Black indicates escaping rock particles, and blue indicates bound rock particles. The top five panels contain simulation snapshots for the target mass as 10 M ⊕ with 30% of the atmosphere by mass, while the impactor mass is 2.1 M ⊕ without the atmosphere (run (31) C3+D3). Those snapshots shows the time when t = 0 × τ ff (panel (1)), t = 2.5 × τ ff (panel (2)), t = 4.5 × τ ff (panel (3)), and t = 30 × τ ff (panel (4)) from left to right, respectively. Run (31) simulation is start from t = − 2 × τ ff, where τ ff = 1.8  hr. The bottom eight panels show detail simulation snapshots for run (31) C3+D3. Those snapshots shows the time when t = 2.0 × τ ff, t = 2.5 × τ ff, t = 3.0 × τ ff, t = 3.5 × τ ff, t = 4.0 × τ ff, t = 4.5 × τ ff, t = 5.0 × τ ff, and t = 5.5 × τ ff from left top to right bottom, respectively. Colors represent the target and the impactor compositions: Green represents the atmosphere particles of the target and the impactor. Red and blue are the core particles of the target and the impactor, respectively.
Paper: Giant Impact Events for Protoplanets: Energetics of Atmospheric Erosion by Head-on Collision
Authors: Kenji Kurosaki, Shu-ichiro Inutsuka
Figure caption: Snapshots of impact simulations. Top five panels show simulation central cross sections reflecting the target mass as 1 M ⊕ with 20% of the atmosphere by mass, while the impactor mass is 1 M ⊕ with 20% of the atmosphere by mass (run (3) B2+B2). The top leftmost panel is color-coded diagrams of lost and bound particles. Red indicates lost atmospheric particles. Green indicates bound atmospheric particles. Black indicates escaping rock particles, and blue indicates bound rock particles. Those snapshots shows the time when t = 0 × τ ff (panel (1)), t = 1 × τ ff (panel (2)), t = 3 × τ ff (panel (3)), and t = 30 × τ ff (panel (4)) from left to right, respectively. Run (3) simulation is start from t = − 1 × τ ff, where τ ff = 3.0  hr. The bottom five panels show simulation snapshots indicating the target mass as 10 M ⊕ with 20% of the atmosphere by mass, and impactor mass is 1 M ⊕ with 20% of the atmosphere by mass (run (6) B4+B2). The bottom leftmost panel is color-coded diagrams of lost and bound particles, whose colors are the same as the top leftmost panel. The other bottom four snapshots show the time when t = 0 × τ ff (panel (1)), t = 1.25 × τ ff (panel (2)), t = 3.25 × τ ff (panel (3)), and t = 30 × τ ff (panel (4)) from left to right, respectively. Run (6) simulation is start from t = −1 × τ ff, where τ ff = 1.5  hr. The colors are the same as in Figure 1.
Paper: Giant Impact Events for Protoplanets: Energetics of Atmospheric Erosion by Head-on Collision
Authors: Kenji Kurosaki, Shu-ichiro Inutsuka
Figure caption: The interior structure at t = 50 τ ff. The radius is normalized by the target radius. The left panel shows the target mass as 10 M ⊕ with 30% of the atmosphere by mass. The impactor mass is 2.1 M ⊕ without the atmosphere. The right panel shows the target whose mass is 1 M ⊕ with 30% of the atmosphere, while the impactor whose mass is 0.21 M ⊕ without the atmosphere. The red line represents the density structure normalized by the central density. The purple line reflects the internal energy normalized by 1012 erg g−1. The blue line indicates the hydrogen mass fraction.
Paper: Giant Impact Events for Protoplanets: Energetics of Atmospheric Erosion by Head-on Collision
Authors: Kenji Kurosaki, Shu-ichiro Inutsuka
Figure caption: The interior structure at t = 50 τ ff. The radius is normalized by the target radius. The left panel shows the target whose mass is 1 M ⊕ with 20% of the atmosphere by mass while the impactor is the same as the target. The right panel shows the target whose mass is 10 M ⊕ with 20% of the atmosphere by mass while the impactor whose mass of 1 M ⊕ with 20% of the atmosphere by mass. Colored lines are the same as those in Figure 4.
Paper: Giant Impact Events for Protoplanets: Energetics of Atmospheric Erosion by Head-on Collision
Authors: Kenji Kurosaki, Shu-ichiro Inutsuka
Figure caption: Time evolution of the escaping mass fractions for the atmosphere and core. The left panel shows that the target mass is 10 M ⊕ with 30% of the atmosphere by mass. The impactor mass is 2.1 M ⊕ without the atmosphere. The right panel shows the target mass as 1 M ⊕ with 30% of the atmosphere by mass, while the impactor mass is 0.21 M ⊕ without the atmosphere. The purple line is the atmosphere loss fraction described by Equation (16), and the green line is the core loss fraction described by Equation (17).
Paper: Giant Impact Events for Protoplanets: Energetics of Atmospheric Erosion by Head-on Collision
Authors: Kenji Kurosaki, Shu-ichiro Inutsuka
Figure caption: The relationship between the total kinetic energy of escaping particles (K esc) and the reference kinetic energy for gravitational escape velocity (E esc,ref) defined by Equation (27). The bottom panel shows the value of β = K esc/E esc,ref for all numerical results. The black line shows a linear regression, K esc = β E esc,ref, where β = 0.42. Plot colors are the same as those in Figure 8.
Paper: Giant Impact Events for Protoplanets: Energetics of Atmospheric Erosion by Head-on Collision
Authors: Kenji Kurosaki, Shu-ichiro Inutsuka
Figure caption: Snapshots of impact simulations. The leftmost panels are color-coded diagrams of lost and bound particles. Red indicates lost atmospheric particles. Green indicates bound atmospheric particles. Black indicates escaping rock particles, and blue indicates bound rock particles. The other four panels contain simulation central cross sections showing the target mass as 1 M ⊕ with 30% of the atmosphere by mass, while the impactor mass is 0.21 M ⊕ without the atmosphere (run (28) C1+D1). Those snapshots shows the time when t = 0 × τ ff (panel (1)), t = 1 × τ ff (panel (2)), t = 2.25 × τ ff (panel (3)), and t = 30 × τ ff (panel (4)) from left to right, respectively. The colors are the same as in Figure 1.
Paper: Giant Impact Events for Protoplanets: Energetics of Atmospheric Erosion by Head-on Collision
Authors: Kenji Kurosaki, Shu-ichiro Inutsuka
Figure caption: The relationship between the initial total energy calculated by Equation (23) and the approximate total energy after a collision calculated by the right-hand side of Equation (24). Purple circles are the impact simulations when the impactor has an atmosphere, while red squares are the bare-rock impactor. Filled marks reflect that the planet retains a layered structure after the impact, while open marks represent the planet has a region wherein the atmosphere and core are mixed. The black dotted line shows ﹩{E}_{\mathrm{final}}^{\mathrm{app}}={E}_{\mathrm{tot}}﹩.
Paper: Giant Impact Events for Protoplanets: Energetics of Atmospheric Erosion by Head-on Collision
Authors: Kenji Kurosaki, Shu-ichiro Inutsuka
Figure caption: Time evolution of the escaping mass fraction after the impact event. The left panel shows the target whose mass is 1 M ⊕ with 20% of the atmosphere by mass while the impactor is the same as the target. The right panel shows the target whose mass is 10 M ⊕ with 20% of the atmosphere by mass while the impactor whose mass of 1 M ⊕ with 20% of the atmosphere by mass. Colored lines are same as those in Figure 3.
Paper: The Gamma-Ray Pulsar Phenomenology in View of 3D Kinetic Global Magnetosphere Models
Authors: Constantinos Kalapotharakos, Zorawar Wadiasingh, Alice K. Harding, Demosthenes Kazanas
Figure caption: The ratio of the emitting power over the accelerating power as a function of the rate of radiation reaction losses. All the panels correspond to the separatrix zone model number 4 for α = 45°. Each panel shows the simulation particles for the indicated ﹩\dot{{ \mathcal E }}﹩ values of YPs. The red lines indicate the fraction of the total γ-ray luminosity that is emitted by all the particles with ﹩\geqslant {\dot{\gamma }}_{{\rm{E}}}﹩ as depicted in the right-hand vertical axes. The high-emitting particles do not reach the RRL regime, i.e., P E/P A ≈ 1 for ﹩\dot{{ \mathcal E }}\lesssim {10}^{34}\,\mathrm{erg}\,{{\rm{s}}}^{-1}﹩.