adsabs/AstroCLIMB huggingface dataset:
Paper: Deep Search for Molecular Oxygen in TW Hya
Authors: Becky J. Williams, L. Ilsedore Cleeves, Christian Eistrup, Jon P. Ramsey
Figure caption: Matched filtering response for the data containing the SO2 and SO transitions, with impulse response on the y-axis and velocity on the x-axis. The two horizontal lines are at 0σ and 3σ. (a) SO2 transition, centered at 235.152 GHz (b) SO transition, centered at 219.949 GHz. The 7σ peak is the C15N detection.
Paper: Transition Probabilities of Near-infrared Ce III Lines from Stellar Spectra: Applications to Kilonovae
Authors: Nanae Domoto, Jae-Joon Lee, Masaomi Tanaka, Ho-Gyu Lee, Wako Aoki, Miho N. Ishigaki, Shinya Wanajo, Daiji Kato, Kenta Hotokezaka
Figure caption: The Ce III lines of HD 190323 (black). In the top panels, the observed spectra are compared with the fitted Gaussian profile (green). In the bottom panels, the observed spectra are compared with the model spectra using the derived gf-values (red). Note that the observed spectrum at ∼15963 Å of HD 190323 (the middle panels) may be affected by unknown lines.
Paper: Transition Probabilities of Near-infrared Ce III Lines from Stellar Spectra: Applications to Kilonovae
Authors: Nanae Domoto, Jae-Joon Lee, Masaomi Tanaka, Ho-Gyu Lee, Wako Aoki, Miho N. Ishigaki, Shinya Wanajo, Daiji Kato, Kenta Hotokezaka
Figure caption: EW of the Ce III λ15961.157 line in the model spectra (colormap) and in the observed spectra (diamonds). Note that the region without color (upper left) is the parameter space where no ATLAS model exists.
Paper: Deep Search for Molecular Oxygen in TW Hya
Authors: Becky J. Williams, L. Ilsedore Cleeves, Christian Eistrup, Jon P. Ramsey
Figure caption: Spectrum of the C15N detection in TW Hya, calculated within a 200 au circular region. The two transitions are N = 2 − 1, J = 5/2 − 3/2 transitions, F = 3 − 2 and F = 2 − 1. The systemic velocity of 2.86 km s−1, centered on the average of the two emission frequencies, is marked with the dashed line.
Paper: Twenty-five Years of Accretion onto the Classical T Tauri Star TW Hya
Authors: Gregory J. Herczeg, 雷歌 沈, Yuguang Chen, 昱光 陈, Jean-Francois Donati, Andrea K. Dupree, Frederick M. Walter, Lynne A. Hillenbrand, Christopher M. Johns-Krull, Carlo F. Manara, Hans Moritz Günther, Min Fang, 敏 房, P. Christian Schneider, Jeff A. Valenti, Silvia H. P. Alencar, Laura Venuti, Juan Manuel Alcalá, Antonio Frasca, Nicole Arulanantham, Jeffrey L. Linsky, Jerome Bouvier, Nancy S. Brickhouse, Nuria Calvet, Catherine C. Espaillat, Justyn Campbell-White, John M. Carpenter, Seok-Jun Chang, Kelle L. Cruz, S. E. Dahm, Jochen Eislöffel, Suzan Edwards, William J. Fischer, Zhen Guo, 震 郭, Thomas Henning, Tao Ji, 涛 纪, Jessy Jose, Joel H. Kastner, Ralf Launhardt, David A. Principe, Connor E. Robinson, Javier Serna, Michal Siwak, Michael F. Sterzik, Shinsuke Takasao
Figure caption: Left: the relative veiling (veiling relative to the low-veiling spectrum of TW Hya) in six wavelength regions (y-axis) compared to the veiling from 5000–5100 Å (x-axis), as measured from ESPaDOnS spectra. Each veiling measurement plotted here is the average of four to eight different 25 Å spectral regions around the labeled wavelength. The line shows a 1:1 relationship and not a fit. The veiling at bluer wavelengths is generally higher than the veiling at red wavelengths, with exceptions. The veiling at 4700 Å is similar to that at 5050 Å because of similar photospheric fluxes. Right: correlated differences between measured and expected veiling at long wavelengths. The inset shows the correlation between veiling at 7000 Å and 5000–5100 Å. The excess veiling is calculated by fitting a line to the r  λ –r 5050 relationship (as can be seen in the left panel) and then subtracting each point from that line. The main plot shows this excess veiling at 7000 Å vs. that at 8700 Å. When the veiling at 7000 Å is higher than expected, the veiling at 8700 Å is also higher than expected. The correlation between excess veiling at 7000 and 8700 Å demonstrates that the scatter in the correlations between veiling measurements is real and not due to S/N or other statistical uncertainties.
Paper: Twenty-five Years of Accretion onto the Classical T Tauri Star TW Hya
Authors: Gregory J. Herczeg, 雷歌 沈, Yuguang Chen, 昱光 陈, Jean-Francois Donati, Andrea K. Dupree, Frederick M. Walter, Lynne A. Hillenbrand, Christopher M. Johns-Krull, Carlo F. Manara, Hans Moritz Günther, Min Fang, 敏 房, P. Christian Schneider, Jeff A. Valenti, Silvia H. P. Alencar, Laura Venuti, Juan Manuel Alcalá, Antonio Frasca, Nicole Arulanantham, Jeffrey L. Linsky, Jerome Bouvier, Nancy S. Brickhouse, Nuria Calvet, Catherine C. Espaillat, Justyn Campbell-White, John M. Carpenter, Seok-Jun Chang, Kelle L. Cruz, S. E. Dahm, Jochen Eislöffel, Suzan Edwards, William J. Fischer, Zhen Guo, 震 郭, Thomas Henning, Tao Ji, 涛 纪, Jessy Jose, Joel H. Kastner, Ralf Launhardt, David A. Principe, Connor E. Robinson, Javier Serna, Michal Siwak, Michael F. Sterzik, Shinsuke Takasao
Figure caption: Coadded ESPaDOnS spectra of TW Hya during visits with strong accretion (F high veil, blue) and weak accretion (F low veil, red) over five wavelength ranges. The residual spectrum (F residual, black, scaled by a factor of 3 for visualization) is calculated by veiling the weak accretion spectrum and then subtracting it from the strong accretion spectrum, as F residual = F high veil − (F low veil + r  λ )/(1 + r  λ ), with r  λ  as the veiling at wavelength λ and with all spectra normalized to the continuum level within the spectral region (normalization shifted slightly in this Figure for visualization). The residual spectrum reveals emission lines that would be undetectable in any single spectrum, affecting the spectrum by filling in photospheric absorption lines (see examples in Gahm et al. 2008; Dodin & Lamzin 2012). Many of the narrow emission lines in the residual spectrum are Fe I (e.g., Hamann & Persson 1992; Beristain et al. 1998), although line identification is beyond the scope of this paper. Appendix A describes how the strong, weak, and residual spectra are calculated.(The data used to create this figure are available.)
Paper: Twenty-five Years of Accretion onto the Classical T Tauri Star TW Hya
Authors: Gregory J. Herczeg, 雷歌 沈, Yuguang Chen, 昱光 陈, Jean-Francois Donati, Andrea K. Dupree, Frederick M. Walter, Lynne A. Hillenbrand, Christopher M. Johns-Krull, Carlo F. Manara, Hans Moritz Günther, Min Fang, 敏 房, P. Christian Schneider, Jeff A. Valenti, Silvia H. P. Alencar, Laura Venuti, Juan Manuel Alcalá, Antonio Frasca, Nicole Arulanantham, Jeffrey L. Linsky, Jerome Bouvier, Nancy S. Brickhouse, Nuria Calvet, Catherine C. Espaillat, Justyn Campbell-White, John M. Carpenter, Seok-Jun Chang, Kelle L. Cruz, S. E. Dahm, Jochen Eislöffel, Suzan Edwards, William J. Fischer, Zhen Guo, 震 郭, Thomas Henning, Tao Ji, 涛 纪, Jessy Jose, Joel H. Kastner, Ralf Launhardt, David A. Principe, Connor E. Robinson, Javier Serna, Michal Siwak, Michael F. Sterzik, Shinsuke Takasao
Figure caption: Top left: the relationship between the equivalent width of coadded lines (one set shown here for lines from 5250–5400 Å) vs. veiling, established from ESPaDOnS spectra and with residuals in ESPaDOnS and FEROS spectra shown below. Top right: similar plots as on the left for the flux ratio for a spectral dip around 5205–5215 Å compared with a continuum region. In both cases, we calculate a best-fit relationship between equivalent width (or flux ratios) and veiling with ESPaDOnS. We then apply those relationships to FEROS spectra. The bottom panels on the left and right show the scatter between the veiling calculated from these relationships and the veiling measured by comparing the FEROS spectrum to a low-veiling FEROS template. Bottom: the final comparison of veiling obtained from the combination of line equivalent widths and flux ratios to the veiling from a low-veiling FEROS template. All FEROS spectra are shifted by ∼0.02 to place them on the same scale as the ESPaDOnS veiling measurements. The scatter of 0.013 + 0.045 × r5050 (with a minimum error of 0.02) is applied as an uncertainty to all spectra where veiling is measured from line equivalent widths and flux ratios.
Paper: Twenty-five Years of Accretion onto the Classical T Tauri Star TW Hya
Authors: Gregory J. Herczeg, 雷歌 沈, Yuguang Chen, 昱光 陈, Jean-Francois Donati, Andrea K. Dupree, Frederick M. Walter, Lynne A. Hillenbrand, Christopher M. Johns-Krull, Carlo F. Manara, Hans Moritz Günther, Min Fang, 敏 房, P. Christian Schneider, Jeff A. Valenti, Silvia H. P. Alencar, Laura Venuti, Juan Manuel Alcalá, Antonio Frasca, Nicole Arulanantham, Jeffrey L. Linsky, Jerome Bouvier, Nancy S. Brickhouse, Nuria Calvet, Catherine C. Espaillat, Justyn Campbell-White, John M. Carpenter, Seok-Jun Chang, Kelle L. Cruz, S. E. Dahm, Jochen Eislöffel, Suzan Edwards, William J. Fischer, Zhen Guo, 震 郭, Thomas Henning, Tao Ji, 涛 纪, Jessy Jose, Joel H. Kastner, Ralf Launhardt, David A. Principe, Connor E. Robinson, Javier Serna, Michal Siwak, Michael F. Sterzik, Shinsuke Takasao
Figure caption: Veiling spectra for strong (pink circles, average relative veiling r 5050 = 1.52), above average (blue diamonds, r 5050 = 0.95), below average (red squares, r 5050 = 0.52), and weak (black triangles, r 5050 = 0.21) veiling epochs measured in ESPaDOnS spectra. The left panel shows the relative veiling measurement, with features in the veiling spectrum caused by the shape of the photospheric template, TWA 25. The right panel shows the veiling spectrum multiplied by TWA 25. The spectra with moderate accretion are flat from 4000–9000 Å, while the strongest accretion spectrum has a bluer slope (increasing flux with decreasing wavelength).
Paper: Twenty-five Years of Accretion onto the Classical T Tauri Star TW Hya
Authors: Gregory J. Herczeg, 雷歌 沈, Yuguang Chen, 昱光 陈, Jean-Francois Donati, Andrea K. Dupree, Frederick M. Walter, Lynne A. Hillenbrand, Christopher M. Johns-Krull, Carlo F. Manara, Hans Moritz Günther, Min Fang, 敏 房, P. Christian Schneider, Jeff A. Valenti, Silvia H. P. Alencar, Laura Venuti, Juan Manuel Alcalá, Antonio Frasca, Nicole Arulanantham, Jeffrey L. Linsky, Jerome Bouvier, Nancy S. Brickhouse, Nuria Calvet, Catherine C. Espaillat, Justyn Campbell-White, John M. Carpenter, Seok-Jun Chang, Kelle L. Cruz, S. E. Dahm, Jochen Eislöffel, Suzan Edwards, William J. Fischer, Zhen Guo, 震 郭, Thomas Henning, Tao Ji, 涛 纪, Jessy Jose, Joel H. Kastner, Ralf Launhardt, David A. Principe, Connor E. Robinson, Javier Serna, Michal Siwak, Michael F. Sterzik, Shinsuke Takasao
Figure caption: The accretion luminosity measured from flux-calibrated spectra of TW Hya. The left panel shows the spectrum of TW Hya (black), modeled (light blue) as the sum of a stellar photosphere (red) and accretion continuum (blue), with a residual accretion spectrum (pink) calculated by subtracting the scaled photosphere from the TW Hya spectrum. The right panel shows the TW Hya spectrum and the residual accretion spectrum on an epoch of strong accretion (2019 July 7) and weak accretion (2021 April 8). The weak accretion spectra tend to have larger Balmer jumps, while the strong accretion spectra have a bluer slope (increasing flux to shorter wavelengths).
Paper: Twenty-five Years of Accretion onto the Classical T Tauri Star TW Hya
Authors: Gregory J. Herczeg, 雷歌 沈, Yuguang Chen, 昱光 陈, Jean-Francois Donati, Andrea K. Dupree, Frederick M. Walter, Lynne A. Hillenbrand, Christopher M. Johns-Krull, Carlo F. Manara, Hans Moritz Günther, Min Fang, 敏 房, P. Christian Schneider, Jeff A. Valenti, Silvia H. P. Alencar, Laura Venuti, Juan Manuel Alcalá, Antonio Frasca, Nicole Arulanantham, Jeffrey L. Linsky, Jerome Bouvier, Nancy S. Brickhouse, Nuria Calvet, Catherine C. Espaillat, Justyn Campbell-White, John M. Carpenter, Seok-Jun Chang, Kelle L. Cruz, S. E. Dahm, Jochen Eislöffel, Suzan Edwards, William J. Fischer, Zhen Guo, 震 郭, Thomas Henning, Tao Ji, 涛 纪, Jessy Jose, Joel H. Kastner, Ralf Launhardt, David A. Principe, Connor E. Robinson, Javier Serna, Michal Siwak, Michael F. Sterzik, Shinsuke Takasao
Figure caption: Examples of the correlation between absolute veiling measured from WTTS templates and relative veiling measured from weak-accreting spectra of TW Hya itself. Colored dots indicate the measured r abs and r rel in each sextile with different colors representing different template stars. The solid line and shaded region in the same color as dots denote the best-fit linear relation between r abs and r rel and its 1σ error for each template. The yellow stars at r rel = 0 is the best-fit r zpt determined as the inverse-variance weighted mean of the y-intercepts. The r zpt measured from different templates are consistent with each other.
Paper: Twenty-five Years of Accretion onto the Classical T Tauri Star TW Hya
Authors: Gregory J. Herczeg, 雷歌 沈, Yuguang Chen, 昱光 陈, Jean-Francois Donati, Andrea K. Dupree, Frederick M. Walter, Lynne A. Hillenbrand, Christopher M. Johns-Krull, Carlo F. Manara, Hans Moritz Günther, Min Fang, 敏 房, P. Christian Schneider, Jeff A. Valenti, Silvia H. P. Alencar, Laura Venuti, Juan Manuel Alcalá, Antonio Frasca, Nicole Arulanantham, Jeffrey L. Linsky, Jerome Bouvier, Nancy S. Brickhouse, Nuria Calvet, Catherine C. Espaillat, Justyn Campbell-White, John M. Carpenter, Seok-Jun Chang, Kelle L. Cruz, S. E. Dahm, Jochen Eislöffel, Suzan Edwards, William J. Fischer, Zhen Guo, 震 郭, Thomas Henning, Tao Ji, 涛 纪, Jessy Jose, Joel H. Kastner, Ralf Launhardt, David A. Principe, Connor E. Robinson, Javier Serna, Michal Siwak, Michael F. Sterzik, Shinsuke Takasao
Figure caption: Left: the size of the Balmer jump (here defined as the flux ratio of 3600–4200 Å) vs. the veiling estimated from low-resolution spectra. The Balmer jump tends to be higher when the veiling is lower. When the Balmer jump is larger, more emission escapes at short wavelengths, so the bolometric correction is larger. Right: the bolometric correction, used to convert the measured accretion flux at 5050 Å to the total accretion luminosity, vs. the veiling, as measured from low-resolution spectra.
Paper: Twenty-five Years of Accretion onto the Classical T Tauri Star TW Hya
Authors: Gregory J. Herczeg, 雷歌 沈, Yuguang Chen, 昱光 陈, Jean-Francois Donati, Andrea K. Dupree, Frederick M. Walter, Lynne A. Hillenbrand, Christopher M. Johns-Krull, Carlo F. Manara, Hans Moritz Günther, Min Fang, 敏 房, P. Christian Schneider, Jeff A. Valenti, Silvia H. P. Alencar, Laura Venuti, Juan Manuel Alcalá, Antonio Frasca, Nicole Arulanantham, Jeffrey L. Linsky, Jerome Bouvier, Nancy S. Brickhouse, Nuria Calvet, Catherine C. Espaillat, Justyn Campbell-White, John M. Carpenter, Seok-Jun Chang, Kelle L. Cruz, S. E. Dahm, Jochen Eislöffel, Suzan Edwards, William J. Fischer, Zhen Guo, 震 郭, Thomas Henning, Tao Ji, 涛 纪, Jessy Jose, Joel H. Kastner, Ralf Launhardt, David A. Principe, Connor E. Robinson, Javier Serna, Michal Siwak, Michael F. Sterzik, Shinsuke Takasao
Figure caption: The zero-point of relative veiling as a function of wavelength. Yellow stars represent the final r zpt determined by averaging the individually measured r zpt from different stellar templates in colored dots in the background. The shape resembles the relation between relative veiling and wavelength.
Paper: Nitrogen as a Tracer of Giant Planet Formation. II. Comprehensive Study of Nitrogen Photochemistry and Implications for Observing NH3 and HCN in Transmission and Emission Spectra
Authors: Kazumasa Ohno, Jonathan J. Fortney
Figure caption: Vertical distributions of NH3 (solid lines) and HCN (dashed lines) volume mixing ratios. Different colored lines show the distributions for different planetary equilibrium temperatures. The gray bold lines denote the abundance equal to the bulk nitrogen abundances, while the green bold lines denote the NH3 abundance predicted by our semianalytical model (Equation (1)). We have assumed a solar composition atmosphere in this figure. The top left panel shows the results for a 1 Gyr old Jupiter mass planet with K zz = 108 cm2 s−1, the top right panel instead sets the age of 0.1 Gyr to demonstrate the effect of planetary age, the left bottom panel sets M p = 0.1M J to demonstrate the planetary mass dependence, and the bottom right panel sets K zz = 1011 cm2 s−1 to be compared with the bottom left panel for showing the K zz dependence.
Paper: Nitrogen as a Tracer of Giant Planet Formation. II. Comprehensive Study of Nitrogen Photochemistry and Implications for Observing NH3 and HCN in Transmission and Emission Spectra
Authors: Kazumasa Ohno, Jonathan J. Fortney
Figure caption: Atmospheric P–T profiles (top panel) and vertical distributions of CH4 abundances (bottom panel) in a Neptune-mass planet (0.11M J) with the age of 1 Gyr. Different colored lines show the profiles for different equilibrium temperatures. The solid and dashed lines show the profiles for 100× solar metallicity and solar composition atmospheres, respectively. The green dashed and solid lines in the top panel denote the equal equilibrium abundance P–T relation of N2 and NH3, computed by the law of mass action as described in Zahnle & Marley (2014) for NH3 = N2, for solar and 100× solar metallicity atmospheres, respectively. The gray solid and dotted lines in the top panel also denote the condensation temperature of NH3 and HCN when assuming the volume mixing ratio of 10−4, typical peak abundances seen in our photochemical calculations (see Figures 3 and 6), where we have used the vapor pressure of Fray & Schmitt (2009).
Paper: Nitrogen as a Tracer of Giant Planet Formation. II. Comprehensive Study of Nitrogen Photochemistry and Implications for Observing NH3 and HCN in Transmission and Emission Spectra
Authors: Kazumasa Ohno, Jonathan J. Fortney
Figure caption: The quenched NH3 abundance predicted by the semi-analytic model of Paper I (Equation (1)) as a function of planetary mass and age, where g and T int are extracted from the thermal evolution tracks of Fortney et al. (2007). The black line denotes the abundance contours of ﹩{f}_{{\mathrm{NH}}_{3}}={10}^{-5}﹩ and 10−6, and yellow dashed and dotted lines show the contours corresponding to 90% and 50% of the bulk nitrogen abundance. The left and right columns show the results for solar metallicity and 10× solar metallicity atmosphere, respectively, where we have assumed that the N/H is scaled by the metallicity.
Paper: Nitrogen as a Tracer of Giant Planet Formation. II. Comprehensive Study of Nitrogen Photochemistry and Implications for Observing NH3 and HCN in Transmission and Emission Spectra
Authors: Kazumasa Ohno, Jonathan J. Fortney
Figure caption: Same as Figure 3, but for different atmospheric compositions. The upper left and right panels show the NH3 (solid curves) and HCN (dashed curves) distributions for the N/H of 0.1× and 10× solar values, the lower left panel shows the results for C/O = 1.1, and the lower right panel shows the results for the atmospheric metallicity of 100× the solar value. We assume a Jupiter-mass planet at the age of 1 Gyr for the variation of N/H and C/O, while we assume a Neptune-mass planet for the 100× solar atmospheres. We set K zz = 108 cm2 s−1 for all calculations.
Paper: Nitrogen as a Tracer of Giant Planet Formation. II. Comprehensive Study of Nitrogen Photochemistry and Implications for Observing NH3 and HCN in Transmission and Emission Spectra
Authors: Kazumasa Ohno, Jonathan J. Fortney
Figure caption: NH3 abundance at pressure levels of 1 mbar (red lines) and 100 mbar (blue lines). The solid, dashed, and dashed–dotted lines show the abundances for K zz = 1011, 108, and 105 cm2 s−1, respectively. From left to right, each panel shows the abundances for planets at the age of 0.1, 1, and 10 Gyr. The top and bottom panels show the results for Jupiter-mass (1M J) and Neptune-mass (0.11M J) planets, respectively. The gray lines denote the bulk nitrogen abundance, and the green lines denote the quenched nitrogen abundance predicted by Equation (1). We have assumed solar composition atmospheres.
Paper: Nitrogen as a Tracer of Giant Planet Formation. II. Comprehensive Study of Nitrogen Photochemistry and Implications for Observing NH3 and HCN in Transmission and Emission Spectra
Authors: Kazumasa Ohno, Jonathan J. Fortney
Figure caption: Photospheric pressure level at each wavelength across the UV. The thick gray line shows the pressure level of the vertical optical depth of unity, and other colored lines show the contribution of each molecule. We assume a solar composition atmosphere, M p = 1M J, age of 1 Gyr, T eq = 800 K, and K zz = 108 cm2 s−1.
Paper: The Evaporating Massive Embedded Stellar Cluster IRS 13 Close to Sgr A*. I. Detection of a Rich Population of Dusty Objects in the IRS 13 Cluster
Authors: Florian Peißker, Michal Zajaček, Lauritz Thomkins, Andreas Eckart, Lucas Labadie, Vladimír Karas, Nadeen B. Sabha, Lukas Steiniger, Maria Melamed
Figure caption:  H-band observation of IRS 13 in 2004.
Paper: Interaction of Inner Heliosheath Ions with Electromagnetic Ion Cyclotron Waves
Authors: Dae-Young Lee, Sung-Jun Noh
Figure caption: Threshold temperature anisotropy of the hot protons (A p) required for ion cyclotron instability shown as a function of the anisotropy of one of two helium ions (A He+, A He2+) with the other set to 0. (a) Results for four different sets of background conditions (plasma density n  e  in per cubic centimeter and B  o  in nanotesla), which were chosen to represent the IHS conditions as available from the observations of Voyager and previous reports (the representative cases described in Section 2 in the text). (b) Results for different parametric choices while the other parameters remain the same as for the blue dashed curve in (a): the black curve is the case with reduced hot proton density by 50% (and increased cold protons accordingly), the magenta curve is the case with reduced temperatures by 1 order for all of the hot populations, and the olive and orange curves are the cases with κ = 4 and 5, respectively. f is defined by the ratio between the electron plasma and gyrofrequencies, ω  pe /Ω e .
Paper: Interaction of Inner Heliosheath Ions with Electromagnetic Ion Cyclotron Waves
Authors: Dae-Young Lee, Sung-Jun Noh
Figure caption: Expanded view of the unstable EMIC wave solution taken from Figure 2 and the corresponding minimum resonant energies E min for interaction with protons and helium ions.
Paper: Interaction of Inner Heliosheath Ions with Electromagnetic Ion Cyclotron Waves
Authors: Dae-Young Lee, Sung-Jun Noh
Figure caption: Demonstration of pitch-angle scattering effects for 48 protons launched with the same initial energy (KEo = 5 keV) and pitch angle (α o = 135°) but different gyro-phases. (a) Local pitch angles over time shown only for six protons for visual clarity. The average pitch angle ﹩\left\langle {\boldsymbol{\alpha }}\right\rangle ﹩ estimated for all 48 protons is also shown as a thick line. (b) The diffusive (red) and advective (blue) pitch-angle changes. The average pitch angle (black) is repeated for reference. The scales of the diffusive and advective changes are shown on the right axis, and the scale of the average pitch angle is shown on the left axis.
Paper: Interaction of Inner Heliosheath Ions with Electromagnetic Ion Cyclotron Waves
Authors: Dae-Young Lee, Sung-Jun Noh
Figure caption: Example of the ion cyclotron wave solution of the kinetic dispersion equation.
Paper: Resonant Excitation of Planetary Eccentricity due to a Dispersing Eccentric Protoplanetary Disk: A New Mechanism of Generating Large Planetary Eccentricities
Authors: Jiaru Li, Dong Lai
Figure caption: Result of the long-term time evolution in the the fiducial system (a p = 0.2r 0, M p = 3 × 10−4 M ⊙, and τ d = 105 P 0). The disk mass M d, which is used as the horizontal axes in both panels, is a function of time (see Equation (16)). The upper panel shows the eccentricity evolution of the planet and the disk. The lower panel shows the behaviors of the precession rates and the eccentricity coupling rates for the secular dynamics model (Equations (8) and (9)). The vertical dashed gray lines mark the M d value when ω d = ω d,m + ω d,p = ω p,d.