adsabs/AstroCLIMB huggingface dataset:
Paper: Nature and Scalings of Density Fluctuations of Compressible Magnetohydrodynamic Turbulence with Applications to the Solar Wind
Authors: Xiangrong Fu, Hui Li, 晖 李, Zhaoming Gan, Senbei Du, John Steinberg
Figure caption: Density fluctuation ﹩\overline{\delta \rho }﹩ as a function of (left) turbulent Mach number ﹩\overline{{M}_{t}}﹩ and (right) Alfvén Mach number ﹩\overline{{M}_{A}}﹩ for a series of 3D low-resolution (2563 cells) MHD simulations with different combinations of parameters β, σ, and γ. Circles are data points extracted from simulations, and dash lines are linear fit to these points. ﹩\overline{\delta \rho }﹩ scales as ﹩C\overline{{M}_{t}}﹩, and the coefficient C depends weakly on σ  c  and γ.
Paper: Nature and Scalings of Density Fluctuations of Compressible Magnetohydrodynamic Turbulence with Applications to the Solar Wind
Authors: Xiangrong Fu, Hui Li, 晖 李, Zhaoming Gan, Senbei Du, John Steinberg
Figure caption: Run 1: (a) History of the rms values of the divergence terms that contribute to the growth of density fluctuations, and histograms of the divergence term PDFs at (b) t = 30, (c) t = 60, and (d) t = 90. In the growing phase of turbulence (t < 60), the total divergence ∇ · (ρ  v ) is dominated by the linear term ρ 0∇ ·  v . In the saturation phase, the nonlinear term ∇ · (δ ρ  v ) is also quite strong.
Paper: Nature and Scalings of Density Fluctuations of Compressible Magnetohydrodynamic Turbulence with Applications to the Solar Wind
Authors: Xiangrong Fu, Hui Li, 晖 李, Zhaoming Gan, Senbei Du, John Steinberg
Figure caption: 3D snapshot of density fluctuation at t = 80 for Run 1. The density fluctuations are elongated along the direction of the background magnetic field (x), showing strong anisotropy.
Paper: Nature and Scalings of Density Fluctuations of Compressible Magnetohydrodynamic Turbulence with Applications to the Solar Wind
Authors: Xiangrong Fu, Hui Li, 晖 李, Zhaoming Gan, Senbei Du, John Steinberg
Figure caption: 4D FFT analysis of density fluctuations in Run 4, with t < 60. 97% of the power is in the nonlinear structures, and 3% is in the slow mode. The contribution from the fast mode is negligible. The format is the same as in Figure 6.
Paper:  SpinSpotter : An Automated Algorithm for Identifying Stellar Rotation Periods with Autocorrelation Analysis
Authors: Rae J. Holcomb, Paul Robertson, Patrick Hartigan, Ryan J. Oelkers, Caleb Robinson
Figure caption: TESS  rotators identified by our algorithm (blue) alongside those found by McQuillan et al. (2014; pink) from the Kepler sample. Top: a scatter plot, where each point represents an individual rotating star. Horizontal lines at 9, 13.5, and 45 days represent one-third and one-half of the baseline of a TESS  sector and one-half of the baseline of a Kepler quarter, respectively. Yellow lines show age isochrones from Mamajek & Hillenbrand (2008) at 500 Myr (dotted), 1 Gyr (dashed–dotted), and 4 Gyr (solid). Note the large drop-off in survey sensitivity for stars with periods longer than 9 days; this reflects our limited sensitivity to longer-period rotators due to our short baseline of observations. The increase in rotators with 10 days < P rot < 14 days is dominated by stars flagged as showing half-period aliases of the rotation periods. Bottom: histograms showing the distribution of stars in the SpinSpotter  (bottom left, blue) and McQuillan et al. (2014; bottom right, pink), respectively, intended to more clearly show the features of each distribution. Period bins are equal in log space.
Paper:  SpinSpotter : An Automated Algorithm for Identifying Stellar Rotation Periods with Autocorrelation Analysis
Authors: Rae J. Holcomb, Paul Robertson, Patrick Hartigan, Ryan J. Oelkers, Caleb Robinson
Figure caption: An example of the TESS  LC (top) and ACF (bottom) of a star that displays clear evidence of rotation at 5.8 days. The periodicity causes peaks in the ACF to occur at integer multiples of the lag associated with the rotation period. We fitted parabolas around each of these peaks (red) and used the height of the peak vertex (A; pink), the width of the peak at the zero crossing (B; green), and the R 2 statistic for the parabolic model (red) as parameters to accept or reject the rotation period associated with the location of the peak vertex (P rot; yellow).
Paper:  SpinSpotter : An Automated Algorithm for Identifying Stellar Rotation Periods with Autocorrelation Analysis
Authors: Rae J. Holcomb, Paul Robertson, Patrick Hartigan, Ryan J. Oelkers, Caleb Robinson
Figure caption: A comparison of rotation values as derived by our algorithm (y-axis) and literature values (x-axis). Stars in agreement with the literature appear on the line y = x (solid black lines). Stars where the period value disagrees between our algorithm and the literature appear near the lines y = 2x and y = x/2 (gray dashed line). Note the increase in scatter about y = x for stars with P rot > 9 days; this reflects our lack of sensitivity to stars’ periods greater than one-third of the data baseline, since multiple cycles of the rotation period must be present in order to form the characteristic peaks in the ACF that are indicative of stellar rotation.
Paper:  SpinSpotter : An Automated Algorithm for Identifying Stellar Rotation Periods with Autocorrelation Analysis
Authors: Rae J. Holcomb, Paul Robertson, Patrick Hartigan, Ryan J. Oelkers, Caleb Robinson
Figure caption: Examples of LCs and ACFs for illustrative cases of (a) a rotating star; (b) a rotating star with spots in opposite hemispheres, thus creating alias peaks at half values of the rotation period; (c) an EB system; and (d) a star with no clear evidence of rotation. For the case with spots in opposite hemispheres (panel (b)), note the alternating low/high pattern of peaks in the ACF. The even-numbered peaks, which are taller, correspond to the true rotation period.
Paper:  SpinSpotter : An Automated Algorithm for Identifying Stellar Rotation Periods with Autocorrelation Analysis
Authors: Rae J. Holcomb, Paul Robertson, Patrick Hartigan, Ryan J. Oelkers, Caleb Robinson
Figure caption: Stars from our testing and verification sets, plotted to show our selection criteria. A avg is the ACF peak height, B avg is the fractional peak width, and R 2 avg is the goodness-of-fit parameter of the parabolas. Solid black lines show the values of our selection criteria, and dashed lines show suggested values of relaxed criteria that would provide a higher recovery rate of rotators at the expense of a higher false-positive rate.
Paper: Hubble Space Telescope Proper Motion (HSTPROMO) Catalogs of Galactic Globular Clusters. VII. Energy Equipartition
Authors: Laura L. Watkins, Roeland P. van der Marel, Mattia Libralato, Andrea Bellini, Jay Anderson, Mayte Alfaro-Cuello
Figure caption: CMDs for NGC 6752. Here, we show the full catalog with no quality selections or velocity clipping. Top: the data points are colored by proximity to the isochrone, as indicated by the color bar. Bottom: The data points are colored by their estimated stellar mass, as indicated by the color bar. In both panels, the black line shows the adopted isochrone for the cluster.
Paper: Relationship between Total Solar Irradiance and Magnetic Flux during Solar Minima
Authors: Sergey V. Marchenko, Judith L. Lean, Matthew T. DeLand
Figure caption: Upper panel: left—the SOHO/SEM and STEREO A and B EUV fluxes with ±σ measurement uncertainties acquired between 2008 June 1 and October 31; right—the corresponding scatter diagram for the data acquired during the “quiescence” epoch E1 (Table 2) with r = 0.323. Middle panel: left—the Bremen Mg II index compared to the weighted average of the OMI, SORCE, and NOAA 16–18 indices; right—the scatter diagram for the epoch E1, r = 0.512. The projected ≈0.3% point-to-point uncertainties of the Mg II index (https://www.iup.uni-bremen.de/gome/gomemgii.html) are omitted since the good agreement between the two data sets assumes much lower errors. Lower-left panel: the SORCE and VIRGO TSI observations.
Paper: Relationship between Total Solar Irradiance and Magnetic Flux during Solar Minima
Authors: Sergey V. Marchenko, Judith L. Lean, Matthew T. DeLand
Figure caption: Upper panel: the average (20 minute time bin) MDI magnetogram acquired on 2008 September 11. Lower panel: the central part of the magnetogram, with the lighter shades of gray denoting the time progression of the sources.
Paper: Relationship between Total Solar Irradiance and Magnetic Flux during Solar Minima
Authors: Sergey V. Marchenko, Judith L. Lean, Matthew T. DeLand
Figure caption: Left panel: SORCE TSI observations and TSI models, along with the scaled and shifted, for clarity, Mg II line index, and the darkening (sunspot) factor from Coddington et al. (2016). In both panels, the epoch under study, 2008 June 1–2008 October 31, is marked by vertical dashed lines. The dotted vertical lines denote the “quiescence” (practically sunspot-free) epoch E1 (2008 June 27–2008 October 7) listed in Table 2. Right panel: the (model–observation) differences normalized by rms(TSI).
Paper: Relationship between Total Solar Irradiance and Magnetic Flux during Solar Minima
Authors: Sergey V. Marchenko, Judith L. Lean, Matthew T. DeLand
Figure caption: Shown in (a) are the time series of the total magnetic flux measured by MDI between 1996 and 2011 (black line) and a model of its variability obtained by regressing the observations against the Bremen Mg II facular-brightening index and the Debrecen sunspot-darkening index (green symbols). Shown in (b) are the individual facular (red line) and sunspot (blue line) contributions to the total magnetic flux variability. Similarly, shown in (c) is the time series of the observed TSI (composite of PMOD and SORCE observations) and a model of its variability. Shown in (d) are the individual facular (red line) and sunspot (blue line) contributions to the TSI variability. The values of r given in (a) and (c) are the correlation coefficients of the model and observations.
Paper: Relationship between Total Solar Irradiance and Magnetic Flux during Solar Minima
Authors: Sergey V. Marchenko, Judith L. Lean, Matthew T. DeLand
Figure caption: Left panel: pdfs of the linear distances between the ∣B∣ ≥ 120 G sources for all the available data (June–October 2008, n = 116 cases), compared to the limited subsample with the enhanced 140–160 pixel scales (23 cases), with ±σ uncertainties. Right panel: same for the 80 < ∣B∣ < 120 G sample. At the Sun’s center, one MDI pixel corresponds to 1.45 Mm.
Paper: Relationship between Total Solar Irradiance and Magnetic Flux during Solar Minima
Authors: Sergey V. Marchenko, Judith L. Lean, Matthew T. DeLand
Figure caption: Left panel: appropriately scaled area fractions of the ∣B∣ > 80 G sources and the normalized TSI measurements with ±σ uncertainties. Right panel: ∣B∣ > 80 G and TSI scatter plot for the data taken during the “quiescence” epoch E1, with r = 0.62 (Table 2).
Paper: Relationship between Total Solar Irradiance and Magnetic Flux during Solar Minima
Authors: Sergey V. Marchenko, Judith L. Lean, Matthew T. DeLand
Figure caption: Distributions of the lifetimes of the ∣B∣ ≥ 120 G sources estimated from the MDI magnetograms acquired on 2008 July 3–2008 July 7.
Paper: Relationship between Total Solar Irradiance and Magnetic Flux during Solar Minima
Authors: Sergey V. Marchenko, Judith L. Lean, Matthew T. DeLand
Figure caption: Surface fractions of the ∣B∣ ≥ 120 G sources in three latitudinal zones: l ≤ −15° (south), ∣l∣ < 15° (equatorial), and l ≥ 15° (north).
Paper: A Rosetta Stone for Eccentric Gravitational Waveform Models
Authors: Alan M. Knee, Isobel M. Romero-Shaw, Paul D. Lasky, Jess McIver, Eric Thrane
Figure caption: The mapping between the eccentricity definitions used by SEOBNRE and TEOBResumS obtained with our mismatch-minimization procedure, for various choices of total (detector-frame) mass, equal-masses, and zero spins. The mismatch calculations assume white noise. All eccentricities are quoted at a reference GW frequency of 10 Hz. The data points are colored by the minimum mismatch between the TEOBResumS and SEOBNRE waveforms before evolving to 10 Hz. The vertical lines indicate where the minimum mismatch crosses certain thresholds: waveforms to the right of the dashed line have ﹩1-{ \mathcal O }\gt 2 \% ﹩, and to the right of the solid line have ﹩1-{ \mathcal O }\gt 5 \% ﹩. The absence of any of these lines means the minimum mismatch never passes that threshold within the explored parameter ranges. The insets show the ratio of the two eccentricities as a function of ﹩{e}_{0}^{\mathrm{SEOB}}﹩. To avoid unphysical model behavior, we have to limit the maximum SEOBNRE eccentricity for heavier systems.
Paper: A Rosetta Stone for Eccentric Gravitational Waveform Models
Authors: Alan M. Knee, Isobel M. Romero-Shaw, Paul D. Lasky, Jess McIver, Eric Thrane
Figure caption: Detectability of the residual differences between SEOBNRE and TEOBResumS after minimizing their mismatch. The dark blue (red) curve shows the minimum mismatch between the waveforms, assuming LVK design sensitivity, as a function of SEOBNRE eccentricity at 10 Hz for a M = 100 M ⊙ (M = 250 M ⊙) BBH at a distance of 1 Gpc (5 Gpc). The light red/blue curves show the inverse squared optimal SNR of the difference, δ h = h TEOB − h SEOB, between the mismatch-minimized waveforms, again computed with LVK sensitivity. For reference, we have drawn horizontal dashed lines for certain values of optimal SNR. The differences are detectable wherever the darker-shaded curve crosses above the lighter curve of the same base color. The differences between the waveforms for a 100 M ⊙ BBH at 1 Gpc become detectable above roughly ﹩{e}_{0}^{\mathrm{SEOB}}=0.25﹩, whereas they are never detectable for the M = 250 M ⊙ system at 5 Gpc in the parameter range we explored.
Paper: A Rosetta Stone for Eccentric Gravitational Waveform Models
Authors: Alan M. Knee, Isobel M. Romero-Shaw, Paul D. Lasky, Jess McIver, Eric Thrane
Figure caption: Marginal posterior probability distributions of eccentricity for ten selected GW events, obtained with either SEOBNRE (red) or TEOBResumS (blue). Note that GW190909 has since been downgraded to a non-astrophysical trigger (Abbott et al. 2021b). In many cases the peaks in the posterior distributions rail against the upper prior limit. The upper prior limit used for TEOBResumS is at the right edges of the plots. The upper prior limit of SEOBNRE is shown by the dashed vertical lines. The upper prior limit of TEOBResumS is ﹩{e}_{0}^{\mathrm{TEOB}}=0.3﹩ at 10 Hz.
Paper: A Rosetta Stone for Eccentric Gravitational Waveform Models
Authors: Alan M. Knee, Isobel M. Romero-Shaw, Paul D. Lasky, Jess McIver, Eric Thrane
Figure caption: Eccentricities in our chosen set of numerical relativity simulations measured with TEOBResumS and SEOBNRE using our mismatch-minimization procedure. The eccentricities are evolved to a frequency of 10 Hz. The simulations were scaled to a total mass of 150 M ⊙ in the detector frame, and the mismatches are computed with respect to either white (circles) or LVK (crosses) noise, integrating from 10 Hz. Note that three simulations (SXS:BBH:1136, SXS:BBH:1149, and SXS:BBH:1169) lack an SEOBNRE measurement due to having spins outside the allowed range of this model, ∣χ 1,2∣ < 0.6.
Paper: A Rosetta Stone for Eccentric Gravitational Waveform Models
Authors: Alan M. Knee, Isobel M. Romero-Shaw, Paul D. Lasky, Jess McIver, Eric Thrane
Figure caption: Eccentricity maps calculated with LVK sensitivity instead of white noise. This plot shows the same information as in Figure 5, except that we replace the flat PSDs with detector PSDs in the calculation of the mismatch.
Paper: A Rosetta Stone for Eccentric Gravitational Waveform Models
Authors: Alan M. Knee, Isobel M. Romero-Shaw, Paul D. Lasky, Jess McIver, Eric Thrane
Figure caption: Effect of varying the mass ratio and component spins on the eccentricity mapping between TEOBResumS and SEOBNRE (at 10 Hz). The total mass here is fixed at 150 M ⊙. The top left panel is repeated from Figure 5. Again, the insets show the ratio of the mapped eccentricities, and the vertical lines denote the same mismatch thresholds from Figure 5.
Paper: Asteroid Measurements at Millimeter Wavelengths with the South Pole Telescope
Authors: P. M. Chichura, A. Foster, C. Patel, N. Ossa-Jaen, P. A. R. Ade, Z. Ahmed, A. J. Anderson, M. Archipley, J. E. Austermann, J. S. Avva, L. Balkenhol, P. S. Barry, R. Basu Thakur, J. A. Beall, K. Benabed, A. N. Bender, B. A. Benson, F. Bianchini, L. E. Bleem, F. R. Bouchet, L. Bryant, K. Byrum, J. E. Carlstrom, F. W. Carter, T. W. Cecil, C. L. Chang, P. Chaubal, G. Chen, H. C. Chiang, H.-M. Cho, T-L. Chou, R. Citron, J.-F. Cliche, T. M. Crawford, A. T. Crites, A. Cukierman, C. M. Daley, E. V. Denison, K. Dibert, J. Ding, M. A. Dobbs, D. Dutcher, W. Everett, C. Feng, K. R. Ferguson, J. Fu, S. Galli, J. Gallicchio, A. E. Gambrel, R. W. Gardner, E. M. George, N. Goeckner-Wald, R. Gualtieri, S. Guns, N. Gupta, R. Guyser, T. de Haan, N. W. Halverson, A. H. Harke-Hosemann, N. L. Harrington, J. W. Henning, G. C. Hilton, E. Hivon, G. P. Holder, W. L. Holzapfel, J. C. Hood, D. Howe, J. D. Hrubes, N. Huang, J. Hubmayr, K. D. Irwin, O. B. Jeong, M. Jonas, A. Jones, T. S. Khaire, L. Knox, A. M. Kofman, M. Korman, D. L. Kubik, S. Kuhlmann, C.-L. Kuo, A. T. Lee, E. M. Leitch, D. Li, A. Lowitz, C. Lu, D. P. Marrone, J. J. McMahon, S. S. Meyer, D. Michalik, M. Millea, L. M. Mocanu, J. Montgomery, C. Corbett Moran, A. Nadolski, T. Natoli, H. Nguyen, J. P. Nibarger, G. Noble, V. Novosad, Y. Omori, S. Padin, Z. Pan, P. Paschos, S. Patil, J. Pearson, K. A. Phadke, C. M. Posada, K. Prabhu, C. Pryke, W. Quan, A. Rahlin, C. L. Reichardt, D. Riebel, B. Riedel, M. Rouble, J. E. Ruhl, B. R. Saliwanchik, J. T. Sayre, K. K. Schaffer, E. Schiappucci, E. Shirokoff, C. Sievers, G. Smecher, J. A. Sobrin, A. Springmann, A. A. Stark, J. Stephen, K. T. Story, A. Suzuki, C. Tandoi, K. L. Thompson, B. Thorne, C. Tucker, C. Umilta, L. R. Vale, T. Veach, J. D. Vieira, G. Wang, N. Whitehorn, W. L. K. Wu, V. Yefremenko, K. W. Yoon, M. R. Young
Figure caption: Predicted asteroid S/N at 2.0 mm in the RA13HDEC-25 and RA23HDEC-25 fields. We base these predictions on the number of observations in which the asteroid is present in the field as well as the mean asteroid flux modeled by STM at 2.0 mm during those observations. Each point in the plot represents an asteroid present in the field during at least one observation. The dotted lines trace out levels of constant S/N at values equal to 1, 5, and 10.
Paper: Asteroid Measurements at Millimeter Wavelengths with the South Pole Telescope
Authors: P. M. Chichura, A. Foster, C. Patel, N. Ossa-Jaen, P. A. R. Ade, Z. Ahmed, A. J. Anderson, M. Archipley, J. E. Austermann, J. S. Avva, L. Balkenhol, P. S. Barry, R. Basu Thakur, J. A. Beall, K. Benabed, A. N. Bender, B. A. Benson, F. Bianchini, L. E. Bleem, F. R. Bouchet, L. Bryant, K. Byrum, J. E. Carlstrom, F. W. Carter, T. W. Cecil, C. L. Chang, P. Chaubal, G. Chen, H. C. Chiang, H.-M. Cho, T-L. Chou, R. Citron, J.-F. Cliche, T. M. Crawford, A. T. Crites, A. Cukierman, C. M. Daley, E. V. Denison, K. Dibert, J. Ding, M. A. Dobbs, D. Dutcher, W. Everett, C. Feng, K. R. Ferguson, J. Fu, S. Galli, J. Gallicchio, A. E. Gambrel, R. W. Gardner, E. M. George, N. Goeckner-Wald, R. Gualtieri, S. Guns, N. Gupta, R. Guyser, T. de Haan, N. W. Halverson, A. H. Harke-Hosemann, N. L. Harrington, J. W. Henning, G. C. Hilton, E. Hivon, G. P. Holder, W. L. Holzapfel, J. C. Hood, D. Howe, J. D. Hrubes, N. Huang, J. Hubmayr, K. D. Irwin, O. B. Jeong, M. Jonas, A. Jones, T. S. Khaire, L. Knox, A. M. Kofman, M. Korman, D. L. Kubik, S. Kuhlmann, C.-L. Kuo, A. T. Lee, E. M. Leitch, D. Li, A. Lowitz, C. Lu, D. P. Marrone, J. J. McMahon, S. S. Meyer, D. Michalik, M. Millea, L. M. Mocanu, J. Montgomery, C. Corbett Moran, A. Nadolski, T. Natoli, H. Nguyen, J. P. Nibarger, G. Noble, V. Novosad, Y. Omori, S. Padin, Z. Pan, P. Paschos, S. Patil, J. Pearson, K. A. Phadke, C. M. Posada, K. Prabhu, C. Pryke, W. Quan, A. Rahlin, C. L. Reichardt, D. Riebel, B. Riedel, M. Rouble, J. E. Ruhl, B. R. Saliwanchik, J. T. Sayre, K. K. Schaffer, E. Schiappucci, E. Shirokoff, C. Sievers, G. Smecher, J. A. Sobrin, A. Springmann, A. A. Stark, J. Stephen, K. T. Story, A. Suzuki, C. Tandoi, K. L. Thompson, B. Thorne, C. Tucker, C. Umilta, L. R. Vale, T. Veach, J. D. Vieira, G. Wang, N. Whitehorn, W. L. K. Wu, V. Yefremenko, K. W. Yoon, M. R. Young
Figure caption: SPT-SZ, SPTpol, and SPT-3G observation fields plotted on a Mollweide projection of the sky using the equatorial coordinate system. The yellow, orange, and cyan boxes denote the sky regions observed using the SPT, and thicker boxes denote regions analyzed in this paper. The blue-green-yellow color scale represents galactic dust emission at 545 GHz from the Planck Public Data Release 2, with yellow indicating higher emission (Planck Collaboration et al. 2016). The dashed gray line represents the ecliptic plane.
Paper: The Molecular Composition of Shadowed Proto-solar Disk Midplanes Beyond the Water Snowline
Authors: Shota Notsu, Kazumasa Ohno, Takahiro Ueda, Catherine Walsh, Christian Eistrup, Hideko Nomura
Figure caption: The fractional molecular ice abundances with respect to water ice n X, ice/﹩{n}_{{{\rm{H}}}_{2}{\rm{O}},\mathrm{ice}}﹩ for the coma of comet 67P/Churyumov-Gerasimenko and our standard disk model calculations at t =106 yr. The comet data (green filled circles and lines) are taken from Rubin et al. (2019, 2020) and originally derived from in situ measurements by the Rosetta Orbiter Spectrometer for Ion and Neutral Analysis for the coma of comet 67P in 2015 May/June before perihelion. The blue crosses show the results of our standard model calculations at r =5.3 au and for f =0.03 (in the shadowed region), the red open squares show those at r =5.3 au and for f =1.0 (between CO2 and CH4 snowlines in the non-shadowed disk), and the black open circles show those at r =30 au and for f =1.0 (outside CO and N2 snowlines in the non-shadowed disk). We note that the points that are not shown in this Figure indicate low ice abundances (<10−7 with respect to water ice) of those species. In our standard model, we assume molecular initial abundances and ξ CR(r) =1.0 × 10−17 [s−1] (see Sections 3.1 and 3.2).
Paper: The Molecular Composition of Shadowed Proto-solar Disk Midplanes Beyond the Water Snowline
Authors: Shota Notsu, Kazumasa Ohno, Takahiro Ueda, Catherine Walsh, Christian Eistrup, Hideko Nomura
Figure caption: The radial profiles of fractional abundances with respect to total hydrogen nuclei densities n X/n H at t =106 yr for H2O (top panels), CO (second-row panels), CO2 (third-row panels), and O2 (bottom panels). The left panels show the results for the disk midplane with the monotonically decreasing density and temperature profile (f =1.0), and the right panels show the results for the shadowed disk midplane (f =0.03). In each panel, the dashed and solid lines show the profiles for gaseous and icy molecules, respectively. The green and blue lines show the results for molecular initial abundances (the “inheritance” scenario), and the red and black lines show the results for atomic initial abundances (the “reset” scenario). The green and red lines show the results for the higher cosmic-ray ionization rate ξ CR(r) =1.0 × 10−17 [s−1], whereas the blue and black lines show the results for the lower cosmic-ray ionization rate ξ CR(r) =1.0 × 10−18 [s−1].
Paper: The Molecular Composition of Shadowed Proto-solar Disk Midplanes Beyond the Water Snowline
Authors: Shota Notsu, Kazumasa Ohno, Takahiro Ueda, Catherine Walsh, Christian Eistrup, Hideko Nomura
Figure caption: Same as Figure 10, but for atomic initial abundances (“reset” scenario).