adsabs/AstroCLIMB huggingface dataset:
Paper: Luminous Radio Emission from the Superluminous Supernova 2017ens at 3.3 yr after Explosion
Authors: Raffaella Margutti, J. S. Bright, D. J. Matthews, D. L. Coppejans, K. D. Alexander, E. Berger, M. Bietenholz, R. Chornock, L. DeMarchi, M. R. Drout, T. Eftekhari, W. V. Jacobson-Galán, T. Laskar, D. Milisavljevic, K. Murase, M. Nicholl, C. M. B. Omand, M. Stroh, G. Terreran, B. A. VanderLey
Figure caption: The radio SED at δ t = 1334.7–1374.1 days constrains the location of SN 2017ens (red square) in the phase space of radio observables, spectral peak luminosity L  ν,pk, and ν pk × t pk (where ν pk is the frequency of the spectral peak at time, t pk). Black (orange) dashed lines mark the location of constant outflow velocity (mass-loss rate ﹩\dot{M}﹩, here given for v  w  = 1000 km s−1). GRB-SNe (black circles) show ultrarelativistic outflows and low ﹩\dot{M}﹩, while ordinary H-stripped Type Ib/c SNe at t < 100 days (gray circles) are associated with slower v ∼ 0.1c shocks. The radio properties of SN 2017ens are similar to those of interacting H-stripped SNe at late times t > 550 days (gray squares, Stroh et al. 2021, which were detected by the VLASS, Lacy et al. 2020). If ν pk is not constrained, we use an open square to mark the SN location for an assumed optically thin spectrum F  ν  ∝ ν −1 and ν pk ≈ 0.3 GHz as in Stroh et al. (2021). The location of the other H-poor SLSN with radio detection at the time of radio discovery (i.e., PTF10ghi, Eftekhari et al. 2019) is marked. This plot assumes equipartition ϵ  e  = ϵ  B  = 1/3. References: Soderberg et al. (2012); Stroh et al. (2021).
Paper: Change of Spectral Properties of Magnetic Field Fluctuations across Different Types of Interplanetary Shocks
Authors: Byeongseon Park, Alexander Pitňa, Jana Šafránková, Zdeněk Němeček, Oksana Krupařová, Vratislav Krupař, Lingling Zhao, Ashok Silwal
Figure caption: The distributions of the length scale of proton gyroradius (k  ρ ) normalized to the spectral break scale (k b) as a function of β for the up- (gray dots) and downstream (red dots) of FF (a), FR (c), SF (e), and SR (g) shocks. The up- (black solid line) and downstream (red solid line) histograms of FF (b), FR (d), SF (f), and SR (h) are presented. The black and red horizontal lines in the histograms indicate the median values with the interquartiles.
Paper: Change of Spectral Properties of Magnetic Field Fluctuations across Different Types of Interplanetary Shocks
Authors: Byeongseon Park, Alexander Pitňa, Jana Šafránková, Zdeněk Němeček, Oksana Krupařová, Vratislav Krupař, Lingling Zhao, Ashok Silwal
Figure caption: Profiles of solar wind parameters (B, V p, and N p) as measured across different types of IP shocks: FF (2014 December 21 18:38 UTC, black line), FR (2014 February 17 05:34 UTC, blue line), SF (2013 May 5 17:34 UTC, green line), and SR (2012 December 3 05:31 UTC, red line). Each parameter is shown within [−65 minutes, 65 minutes] from the shock crossing (0 minute, vertical dashed line), and two vertical dotted lines represent ±5 minutes. The blue bars in the top distinguish the intervals of 60 minutes ([−65 minutes, −5 minutes] and [5 minutes, 65 minutes]) where the spectral analysis of the magnetic field fluctuations is applied, while the yellow bars indicate 15 minutes intervals ([−20 minutes, −5 minutes] and [5 minutes, 20 minutes]) for computation of the shock parameters.
Paper: Luminous Radio Emission from the Superluminous Supernova 2017ens at 3.3 yr after Explosion
Authors: Raffaella Margutti, J. S. Bright, D. J. Matthews, D. L. Coppejans, K. D. Alexander, E. Berger, M. Bietenholz, R. Chornock, L. DeMarchi, M. R. Drout, T. Eftekhari, W. V. Jacobson-Galán, T. Laskar, D. Milisavljevic, K. Murase, M. Nicholl, C. M. B. Omand, M. Stroh, G. Terreran, B. A. VanderLey
Figure caption: Temporal evolution of the radio spectral energy distribution (SED) of SN 2017ens in the time range δ t = 55–1602 days. No radio emission is detected in the first year after explosion. Our first radio detection of SN 2017ens occurred at δ t ∼ 1350 days (red). Open symbols mark upper limits.
Paper: Change of Spectral Properties of Magnetic Field Fluctuations across Different Types of Interplanetary Shocks
Authors: Byeongseon Park, Alexander Pitňa, Jana Šafránková, Zdeněk Němeček, Oksana Krupařová, Vratislav Krupař, Lingling Zhao, Ashok Silwal
Figure caption: The distributions of the length scale of proton cyclotron resonance (k c) normalized to the spectral break scale (k b) again as a function of β. The format follows Figure 3.
Paper: Magnetic Fields in Solar Plage Regions: Insights from High-sensitivity Spectropolarimetry
Authors: J. M. da Silva Santos, K. Reardon, G. Cauzzi, T. Schad, V. Martínez Pillet, A. Tritschler, F. Wöger, R. Hofmann, J. Stauffer, H. Uitenbroek
Figure caption: Magnetic field in the photosphere and chromosphere. Panel (a): synchronic HMI LOS magnetogram, matched to the spatiotemporal location of ViSP slit. Panels (b)–(c): LOS component of the field from ViSP λ6302 and λ8542. Panel (d): two-dimensional histogram of LOS magnetic field values derived from ViSP vs. HMI; the solid lines show the HMI = ViSP locus; the cross shows the mean values within the plage mask (panel (h)); r is the Pearson’s correlation coefficient. Panels (e)–(f): transverse components of the field from ViSP; the arrow highlights the fibril-like structure of the field. Panel (g): averaged field strength in the plage as a function of μ for the two ViSP diagnostics. Panels (h)–(i): total field strength from ViSP; the dotted lines show isocontours of μ values; the blue contours show the plage mask (see the main text); the lines F1 and F2 are perpendicular to the fibrils (see Figure 2). The color bars are capped for display purposes. Panels (e), (f), and (h) have been gamma adjusted.
Paper: Magnetic Fields in Solar Plage Regions: Insights from High-sensitivity Spectropolarimetry
Authors: J. M. da Silva Santos, K. Reardon, G. Cauzzi, T. Schad, V. Martínez Pillet, A. Tritschler, F. Wöger, R. Hofmann, J. Stauffer, H. Uitenbroek
Figure caption: Flows and polarization in and around the plage. Panels (a)–(b): zoomed-in views in (λ8542) intensity/Dopplergram of the region delimited by the white box in Figure 2; the color bar range is capped. Panels (c)–(e): Stokes I, U, and V signals at the locations marked by the crosses in the top panels (no spatial binning). The vertical dashed line shows the nominal line center. The open circles in panel (c) show the positions used to determine V LOS. Stokes U and V have been vertically shifted for qualitative comparison; the dotted lines show the ±2 × 10−3 levels for each spectrum; the vertical bar in panel (e) indicates an amplitude of V/I c = 2 × 10−2.
Paper: Magnetic Fields in Solar Plage Regions: Insights from High-sensitivity Spectropolarimetry
Authors: J. M. da Silva Santos, K. Reardon, G. Cauzzi, T. Schad, V. Martínez Pillet, A. Tritschler, F. Wöger, R. Hofmann, J. Stauffer, H. Uitenbroek
Figure caption: Overview of the target on 2022 June 3 at 17:46 UT. Panel (a): SDO/AIA 1700 Å intensities; the dashed black square shows one of the four VBI pointings composed of nine subfields each; the orange (red) rectangle shows the ViSP FOV for arm 6302 Å (8542 Å). Panel (b): VBI G-band intensities. Panel (c): VBI Hβ intensities.
Paper: Magnetic Fields in Solar Plage Regions: Insights from High-sensitivity Spectropolarimetry
Authors: J. M. da Silva Santos, K. Reardon, G. Cauzzi, T. Schad, V. Martínez Pillet, A. Tritschler, F. Wöger, R. Hofmann, J. Stauffer, H. Uitenbroek
Figure caption: Polarization signals arising from the photosphere and chromosphere. Panel (a): continuum-normalized intensity in the core of λ8542; the yellow arrow points to the disk center. Panels (b)/(c): λ8542 TCP/TLP in arbitrary units. Panels (d)/(e): TCP/TLP for λ6302. All color bars are capped for display purposes.
Paper: Magnetic Fields in Solar Plage Regions: Insights from High-sensitivity Spectropolarimetry
Authors: J. M. da Silva Santos, K. Reardon, G. Cauzzi, T. Schad, V. Martínez Pillet, A. Tritschler, F. Wöger, R. Hofmann, J. Stauffer, H. Uitenbroek
Figure caption: Magnetic field components in the chromosphere. Top panels: LOS component (left) and transverse component (right) of the magnetic field in the region shown in the top panels in Figure 3; B TRV has been gamma-adjusted to enhance the visibility of the fibrils (see Figures 2 and 3). Lower panels: example observed Stokes U and V (black lines) and WFA fit (red lines; within ±120 mÅ from line nominal line center) at the two locations marked in the upper panels.
Paper: Magnetic Fields in Solar Plage Regions: Insights from High-sensitivity Spectropolarimetry
Authors: J. M. da Silva Santos, K. Reardon, G. Cauzzi, T. Schad, V. Martínez Pillet, A. Tritschler, F. Wöger, R. Hofmann, J. Stauffer, H. Uitenbroek
Figure caption: Magnetic field and λ8542 intensity across the fibrils around plage regions. Left panels: enlarged views of the regions around the F1 and F2 lines shown in Figure 4 in λ8542 core intensity and total field strength (chromosphere); the color bars are capped for display purposes. Right panels: intensity, the absolute value of the LOS component, and total field strength along the dashed lines shown on the left panels.
Paper: Magnetic Fields in Solar Plage Regions: Insights from High-sensitivity Spectropolarimetry
Authors: J. M. da Silva Santos, K. Reardon, G. Cauzzi, T. Schad, V. Martínez Pillet, A. Tritschler, F. Wöger, R. Hofmann, J. Stauffer, H. Uitenbroek
Figure caption: Effect of crosstalk removal in the determination of the photospheric magnetic field. Top panels: LOS component of the magnetic field (left), transverse component (middle), and azimuth angle (right) for one of the four ViSP mosaic tiles without any postprocessing (top row) and with crosstalk removal (bottom row). The color bars are capped for display purposes. Bottom panels: an example observed (Level 1.5) full Stokes spectrum (black) and best fit (red) at the location marked by the red cross in the upper panels.
Paper: Magnetic Fields in Solar Plage Regions: Insights from High-sensitivity Spectropolarimetry
Authors: J. M. da Silva Santos, K. Reardon, G. Cauzzi, T. Schad, V. Martínez Pillet, A. Tritschler, F. Wöger, R. Hofmann, J. Stauffer, H. Uitenbroek
Figure caption: Crosstalk correction for the ViSP spectra. Panels (a)–(c): diattenuation matrix parameters along the ViSP slit for one of the four mosaic tiles. Panels (d)–(g): Level 1 (left) and Level 1.5 (right) Stokes Q and V for a given slit position; the dashed lines show the averaged Stokes I around the Fe I lines (arbitrary units) for reference; the color bars are capped for display purposes. Panels (h)–(i): averaged Stokes Q and V along the slit shown in panels (d)–(g), with (red) and without (black) crosstalk correction.
Paper: The Discrepancy between Observed and Predicted Heliospheric Energetic Neutral Atoms below Solar Wind Energy
Authors: André Galli, Igor I. Baliukin, Marc Kornbleuth, Merav Opher, Stephen A. Fuselier, Justyna M. Sokół, Konstantinos Dialynas, Maher A. Dayeh, Vladislav V. Izmodenov, John D. Richardson
Figure caption: Measured vs. modeled total ENA intensities; the same as in Figure 2 but with contributions of the individual ENA populations. Cyan: ISN H contributions; purple: ENAs from solar wind PUI; blue: ENAs from heliosheath PUIs; orange: neutralized solar wind protons; the solid lines refer to the Moscow model populations and the dashed lines refer to Boston model populations.
Paper: The Discrepancy between Observed and Predicted Heliospheric Energetic Neutral Atoms below Solar Wind Energy
Authors: André Galli, Igor I. Baliukin, Marc Kornbleuth, Merav Opher, Stephen A. Fuselier, Justyna M. Sokół, Konstantinos Dialynas, Maher A. Dayeh, Vladislav V. Izmodenov, John D. Richardson
Figure caption: Measured vs. modeled total ENA intensities in the inertial reference frame at 100 au. IBEX observations are depicted by red (IBEX-Lo) and black circles (IBEX-Hi); the results of the Moscow model are shown in dark green and the results of the Boston model are shown in light green.
Paper: The Discrepancy between Observed and Predicted Heliospheric Energetic Neutral Atoms below Solar Wind Energy
Authors: André Galli, Igor I. Baliukin, Marc Kornbleuth, Merav Opher, Stephen A. Fuselier, Justyna M. Sokół, Konstantinos Dialynas, Maher A. Dayeh, Vladislav V. Izmodenov, John D. Richardson
Figure caption: Ratio of modeled over observed ENA intensities for all six sky regions for the Boston University model (left panel) and the Moscow model (right panel).
Paper: The Discrepancy between Observed and Predicted Heliospheric Energetic Neutral Atoms below Solar Wind Energy
Authors: André Galli, Igor I. Baliukin, Marc Kornbleuth, Merav Opher, Stephen A. Fuselier, Justyna M. Sokół, Konstantinos Dialynas, Maher A. Dayeh, Vladislav V. Izmodenov, John D. Richardson
Figure caption: Discrepancy between average measured (red and black lines) and modeled (green lines) ENA intensities.
Paper: A Model for Nonthermal Particle Acceleration in Relativistic Magnetic Reconnection
Authors: Xiaocan Li, Fan Guo, Yi-Hsin Liu, Hui Li
Figure caption: (a) A schematic diagram illustrating the main processes. The blue and orange regions represent the acceleration and escape regions, respectively. The black curve shows one typical particle trajectory. (b) Two electron trajectories in the kinetic simulations. The orange crosses indicate when particles reach their final energies. The background is at t ≈ 2.44τ  lc , where τ  lc  = L  x /c. (c) Time evolution of the electron energies. The colored lines are the smoothed data that removed the gyromotion effects.
Paper: A Model for Nonthermal Particle Acceleration in Relativistic Magnetic Reconnection
Authors: Xiaocan Li, Fan Guo, Yi-Hsin Liu, Hui Li
Figure caption: (a) The time dependence of the rates within the selected energy bin for runs with different b  g . (b) and (c) The acceleration and escape rates for runs with different σ ic. The dashed lines follow the predicted scalings in Equations (3), (6), and (9) except in (b), where α acc⊥ ∝ t −2 when σ ic = 102.4.
Paper: A Model for Nonthermal Particle Acceleration in Relativistic Magnetic Reconnection
Authors: Xiaocan Li, Fan Guo, Yi-Hsin Liu, Hui Li
Figure caption: (a) Comparing the electron spectra obtained from the model (thick lines) and those in the simulations (thin lines) with b  g  = 0.0 and different σ ic. (b) Comparing the spectra in the runs with σ ic = 102.4 and different b  g . The insets of panels (a) and (b) show the power-law indices and break energies obtained from the models for systems of different sizes. The vertical black line indicates the size of the PIC simulations (L 0). (c) Power-law indices in astrophysically relevant systems with different σ  x  and b  g . The dashed are the fittings. (d) The gyroradius of the electrons with the break energies.
Paper: A Model for Nonthermal Particle Acceleration in Relativistic Magnetic Reconnection
Authors: Xiaocan Li, Fan Guo, Yi-Hsin Liu, Hui Li
Figure caption: The energy dependence of the rates in the runs with σ ic = 102.4. (a) The rates in the run with b  g  = 0. The dashed lines are based on the model scalings. (b) α acc∥ for the runs with different guide fields. The dashed lines follow the scaling. (c) α esc for these runs.
Paper: A Model for Nonthermal Particle Acceleration in Relativistic Magnetic Reconnection
Authors: Xiaocan Li, Fan Guo, Yi-Hsin Liu, Hui Li
Figure caption: Electron energy spectra in the runs with σ ic = 6.4. (a) Time evolution of the spectra when b  g  = 0.0. The insets show the time evolution of the cutoff energy and the power-law index. (b) The spectra in runs with b  g  = 0.0 but different L  x . The inset shows the cutoff energy changing with L  x . L 0 = 4096d  e  for these runs.
Paper: A Model for Nonthermal Particle Acceleration in Relativistic Magnetic Reconnection
Authors: Xiaocan Li, Fan Guo, Yi-Hsin Liu, Hui Li
Figure caption: The time dependence of the rates within the selected energy bin for runs with different b  g . The dashed lines follow the predicted scalings ﹩\sim {L}_{{\rm{O}}}^{-1}\sim {t}^{-1}﹩ in Equations (6), (9), and (7).
Paper: A Model for Nonthermal Particle Acceleration in Relativistic Magnetic Reconnection
Authors: Xiaocan Li, Fan Guo, Yi-Hsin Liu, Hui Li
Figure caption: Time evolution of ﹩{V}_{x-\max }{\rm{\Gamma }}({V}_{x-\max })﹩ in simulations with σ ic = 102.4, where ﹩{V}_{x-\max }﹩ is the maximum outflow speed at the mid-plane and ﹩{\rm{\Gamma }}({V}_{x-\max })﹩ is the corresponding Lorentz factor. τ  lc  is the light-crossing time.
Paper: NGDEEP Epoch 1: The Faint End of the Luminosity Function at z ∼ 9–12 from Ultradeep JWST Imaging
Authors: Gene C. K. Leung, Micaela B. Bagley, Steven L. Finkelstein, Henry C. Ferguson, Anton M. Koekemoer, Pablo G. Pérez-González, Alexa Morales, Dale D. Kocevski, Guang Yang, 光 杨, Rachel S. Somerville, Stephen M. Wilkins, L. Y. Aaron Yung, Seiji Fujimoto, Rebecca L. Larson, Casey Papovich, Nor Pirzkal, Danielle A. Berg, Jennifer M. Lotz, Marco Castellano, Óscar A. Chávez Ortiz, Yingjie Cheng, Mark Dickinson, Mauro Giavalisco, Nimish P. Hathi, Taylor A. Hutchison, Intae Jung, Jeyhan S. Kartaltepe, Priyamvada Natarajan, Barry Rothberg
Figure caption: Image cutouts for the 31 sources rejected after visual inspection.
Paper: Recipes for Jet Feedback and Spin Evolution of Black Holes with Strongly Magnetized Super-Eddington Accretion Disks
Authors: Angelo Ricarte, Ramesh Narayan, Brandon Curd
Figure caption: Left: magnetic flux parameter ϕ as a function of Eddington ratio f Edd, where color encodes different values of the BH spin a *. For each spin sampled by our simulation library, we plot our fitting function (Equation (3)) in the appropriate color. Right: MHD energy outflow efficiency η MHD as a function of magnetic flux parameter for each of our models. For each spin sampled by our simulation library, we plot the BZ prediction η EM (Equation (5)) as colored lines. The agreement is excellent, implying that a BZ-like EM jet dominates the outflow energy in most of the simulations, except for a * = 0, which features a weaker HD outflow. As a point of reference, we plot the radiative efficiencies of thin disks with a * ∈ {0, 0.68, 0.9, 0.97} as horizontal lines.
Paper: Recipes for Jet Feedback and Spin Evolution of Black Holes with Strongly Magnetized Super-Eddington Accretion Disks
Authors: Angelo Ricarte, Ramesh Narayan, Brandon Curd
Figure caption: Example evolutionary pathways of BH mass and spin computed using the fitting functions derived in this work. In each column, the panels show, from top to bottom, the Eddington ratio f Edd, BH mass M, BH spin a *, and jet power P jet as a function of time. Each pathway is tuned to produce M ≈ 109 M ⊙ at the final time. Left: constant Eddington ratio scenarios, each of which reaches a distinct equilibrium spin from which their recent cosmically averaged f Edd could be inferred. Right: fueling-limited scenarios, where we prescribe ﹩\dot{M}﹩ as a function of time. The constant ﹩\dot{M}﹩ scenario is tuned to match the Wang et al. (2021) quasar found when the Universe was only 670 Myr old (z = 7.642). The power-law ﹩\dot{M}﹩ scenario has a prescribed time-variable accretion rate, ﹩\dot{M}\propto {(1+{(t/{10}^{7}\ \mathrm{yr})}^{2})}^{-1}﹩ (motivated by Hopkins et al. 2006a, 2006b), and might represent the history of a currently low Eddington rate SMBH at the center of a galaxy cluster such as Messier 87.
Paper: Recipes for Jet Feedback and Spin Evolution of Black Holes with Strongly Magnetized Super-Eddington Accretion Disks
Authors: Angelo Ricarte, Ramesh Narayan, Brandon Curd
Figure caption: Spin-up parameter s as a function of BH spin a *, with Eddington ratio f Edd encoded in the color. Values computed directly from our GRMHD simulations are plotted in the upper panel, and the predictions of our fitting functions (Equation (13)) are shown in the lower panel. At the lowest accretion rates, the models approximately match the prediction for a razor-thin disk (Equation (B14)), shown as the dotted line. At the highest accretion rates, prograde and zero-spin models approach the curve found for pure GRMHD models (Equation (8)), plotted as a dashed line. We plot our model predictions for s for f Edd = 1 and f Edd → ∞ with blue and red curves, respectively.(The data used to create this figure are available.)
Paper: Recipes for Jet Feedback and Spin Evolution of Black Holes with Strongly Magnetized Super-Eddington Accretion Disks
Authors: Angelo Ricarte, Ramesh Narayan, Brandon Curd
Figure caption: Here we visualize the disk structure and time evolution of the M = 104 M ⊙, a * = 0.9, f Edd = 9.3 model. In the upper two panels, we plot the gas density (color) and magnetic field (black streamlines) within and perpendicular to the disk midplane, respectively. On the right, we plot with a magenta curve the σ ≡ B 2/4π ρ = 1 contour, a common definition of the jet boundary. This snapshot, which corresponds to time t = 28,500t  g , features a flux eruption event, a transient low-density bubble near the horizon. In the lower panels, we plot the Eddington ratio ﹩{f}_{\mathrm{Edd}}=\dot{M}/{\dot{M}}_{\mathrm{Edd}}﹩ and the magnetic flux parameter ϕ as a function of time for this model, demonstrating stability for our period of interest, demarcated by the red horizontal lines. The time corresponding to the snapshot in the upper panels is marked with a blue circle.
Paper: Recipes for Jet Feedback and Spin Evolution of Black Holes with Strongly Magnetized Super-Eddington Accretion Disks
Authors: Angelo Ricarte, Ramesh Narayan, Brandon Curd
Figure caption: Equilibrium spin a *,eq as a function of Eddington ratio f Edd using our model (Equation (13)). Systems with f Edd = 1 reach equilibrium at a *,eq ≈ 0.8, while those with a factor of a few smaller f Edd equilibrate near a *,eq ≈ 1. Systems with both f Edd ≪ 1 and f Edd ≫ 1 reach equilibrium near a *,eq ≈ 0. In the upper x-axis, we plot t Sal/f Edd, the timescale over which both mass and spin evolve and thus the minimum timescale required to reach spin equilibrium.
Paper: Recipes for Jet Feedback and Spin Evolution of Black Holes with Strongly Magnetized Super-Eddington Accretion Disks
Authors: Angelo Ricarte, Ramesh Narayan, Brandon Curd
Figure caption: Pressure scale height measured at a radius of 10r  g  for our simulations, colored by spin. In gray, we plot a linear regression to these data, ﹩h/r=0.21+0.046{\mathrm{log}}_{10}\ {f}_{\mathrm{Edd}}﹩. This correlation suggests that higher Eddington ratios lead to greater pressures that can better confine magnetic flux on event horizon scales.(The data used to create this figure are available.)
Paper: Recipes for Jet Feedback and Spin Evolution of Black Holes with Strongly Magnetized Super-Eddington Accretion Disks
Authors: Angelo Ricarte, Ramesh Narayan, Brandon Curd
Figure caption: The HD spin-up parameter calculated from our simulations (left), fitting function (middle), and residuals (right). The dotted line is the thin disk solution, and the dashed line is the value found by Lowell et al. (2023) for nonradiative simulations.(The data used to create this figure are available.)
Paper: The Influence of Hydrogen Sulfide on the Optical Properties of Planetary Organic Hazes: Implications for Exoplanet Climate Modeling
Authors: Nathan W. Reed, Kevin T. Jansen, Zachary R. Schiffman, Margaret A. Tolbert, Eleanor C. Browne
Figure caption: The calculated transmittance of 405 (blue arrows) and 532 (green arrows) nm light vs. (a) the S:C ratio and (b) the comparative studies. The comparative studies measured the optical properties of haze generated from 0.1% CH4 in N2, except for Khare et al. (1984), which used 1% CH4 in N2. Values from Khare et al. (1984) were taken from comparable wavelengths (413.3 and 563.5 nm). Ugelow et al. (2017) and Khare et al. (1984) produced haze analogs using spark and direct current (D.C.) discharge, respectively, rather than photochemically. Ugelow et al. (2018) only measured at 405 nm and Hasenkopf et al. (2010) only measured at 532 nm.