adsabs/AstroCLIMB huggingface dataset:
Paper: Constraining Cosmology with Machine Learning and Galaxy Clustering: The CAMELS-SAM Suite
Authors: Lucia A. Perez, Shy Genel, Francisco Villaescusa-Navarro, Rachel S. Somerville, Austen Gabrielpillai, Daniel Anglés-Alcázar, Benjamin D. Wandelt, L. Y. Aaron Yung
Figure caption: Clustering statistics for CAMELS-SAM simulations 630 through 649 (each in a unique color). We show the VPF (left), CiC at 27.2 cMpc (center), and real-space 3D 2ptCF (right), for two different selections of z = 0 galaxies (top: stellar mass, bottom: SFR). The shaded black area in the top figure indicates the clustering across five simulations using the “fiducial” cosmological and SAM parameters but generated with different random seeds, to show the effect of cosmic variance for these volumes.
Paper: Constraining Cosmology with Machine Learning and Galaxy Clustering: The CAMELS-SAM Suite
Authors: Lucia A. Perez, Shy Genel, Francisco Villaescusa-Navarro, Rachel S. Somerville, Austen Gabrielpillai, Daniel Anglés-Alcázar, Benjamin D. Wandelt, L. Y. Aaron Yung
Figure caption: Halo mass functions (HMFs) at z = 0 (top) and z = 2 (bottom) for the CV simulations (black stars; fiducial cosmology and unique random seeds). HMFs of the (205 h −1 cMpc)3, N = 12803 “extreme” cosmology volumes, where Ex{0,1,2,3,4} have ΩM = {0.1, 0.5, 0.5, 0.1, 0.3} and σ 8 = {0.6, 1.0, 0.6, 1.0, 0.8}, are shown with {purple, blue, cyan, orange, red} lines.
Paper: Cosmological Simulations of Galaxy Groups and Clusters. III. Constraining Quasar Feedback Models with the Atacama Large Millimeter Array
Authors: Avinanda Chakraborty, Suchetana Chatterjee, Mark Lacy, Soumya Roy, Samrat Roy, Rudrani Kar Chowdhury
Figure caption: An example of “simalma” outputs for one of the BHs. Upper left panel: diagnostic figure of the sky model with pointings. Upper right panel: diagnostic figure of uv coverage and visibilities. Lower left panel: diagnostic figure of clean image and residual. Lower right panel: diagnostic figure of difference and fidelity.
Paper: Rocky Planet or Water World? Observability of Low-density Lava World Atmospheres
Authors: Anjali A. A. Piette, Peter Gao, Kara Brugman, Anat Shahar, Tim Lichtenberg, Francesca Miozzi, Peter Driscoll
Figure caption: The effect of varying the input surface temperature in the VapoRock calculation on the resulting atmospheric temperature profile and emission spectrum. The example shown is for HD 3167 b assuming a 90% rock-vapor fraction, with the volatile component consisting of 80% H2O + 20% CO2. Models are shown for a VapoRock input surface temperature of 3000 K (orange lines) and 2513 K (dark-blue lines, i.e., the substellar temperature of HD 3167 b, as used in Figure 6). The resulting differences in the relative abundances of the rock-vapor species make little difference to the resulting temperature profile and spectrum, especially when compared to the more significant effects of adding volatiles to the atmospheric composition.
Paper: Rocky Planet or Water World? Observability of Low-density Lava World Atmospheres
Authors: Anjali A. A. Piette, Peter Gao, Kara Brugman, Anat Shahar, Tim Lichtenberg, Francesca Miozzi, Peter Driscoll
Figure caption: Retrieved posterior probability distributions for the molecular volume mixing ratios. The posterior distributions correspond to the retrieval cases in Figures 9 and 10 for HD 86226 c (top three rows), HD 3167 b (middle three rows), and 55 Cnc e (bottom three rows). Red squares and error bars show the median retrieved abundances and 68% confidence intervals, respectively. Dashed vertical lines show the nominal H2O and CO2 abundances, without accounting for thermal dissociation effects.
Paper: Modeling the Chromosphere and Transition Region of Planet-hosting Star GJ 436
Authors: Dominik Hintz, Sarah Peacock, Travis Barman, Birgit Fuhrmeister, Evangelos Nagel, Andreas Schweitzer, Sandra V. Jeffers, Ignasi Ribas, Ansgar Reiners, Andreas Quirrenbach, Pedro J. Amado, Victor J. S. Béjar, José A. Caballero, Artie P. Hatzes, David Montes
Figure caption: Top left: temperature structures of a series of models with varying temperature gradients in the upper chromosphere (∇T UC). Top right: changes in number densities of atomic hydrogen of models M2b to M2d (n Mi) relative to the reference model M2 (blue, n ref) at temperatures of 22,900 K (solid line, left axis) and 12,200 K (dotted line, right axis) from where significant contributions to the Lyα and Hα line flux arise, respectively (see Figure 1, left panel). They are plotted against the corresponding ∇T UC gradients. Colors coincide with the models depicted on the left-hand side. Circles depict variations in number density at 12,200 K and crosses those at 22,900 K. Increasing ∇T UC from M2 to M2d leads to a decrease in both n Mi(T = 12,200 K) and n Mi(T = 22,900 K), but the effect is significantly stronger for n Mi(T = 12,200 K). Therefore, Hα is more affected by this temperature profile parameter than Lyα (see also Figure B1). Bottom: corresponding model spectra in the lines of Lyα (right, including the ISM transmission curve) and Hα (left). This figure is continued in Figure B2 in Appendix B for the C IV and He I IRT lines.
Paper: Modeling the Chromosphere and Transition Region of Planet-hosting Star GJ 436
Authors: Dominik Hintz, Sarah Peacock, Travis Barman, Birgit Fuhrmeister, Evangelos Nagel, Andreas Schweitzer, Sandra V. Jeffers, Ignasi Ribas, Ansgar Reiners, Andreas Quirrenbach, Pedro J. Amado, Victor J. S. Béjar, José A. Caballero, Artie P. Hatzes, David Montes
Figure caption: Models M1 and M2 (green and blue lines) obtained from the temperature structures depicted in Figure 1 (with corresponding color-coding) and observations (black lines) of GJ 436 shown in the lines of Lyα together with the ISM transmittance curve for GJ 436 provided by Youngblood et al. (2016, gray dashed line; an exemplary insight into how we apply the ISM transmission to our model M2 is provided in Appendix A) (top panel), C IV (second panel), Hα (third panel), and He I IRT (bottom panel). Red lines correspond to the best linear combinations of the model components M1 and M2. HST observations are scaled by ﹩{d}^{2}/{R}_{\star }^{2}﹩ to the stellar surface. VIS/NIR lines of CARMENES observations and model spectra are continuum normalized. The σ values correspond to the standard deviations of the observed line flux (for Lyα σ is calculated from the line wings only).
Paper: Modeling the Chromosphere and Transition Region of Planet-hosting Star GJ 436
Authors: Dominik Hintz, Sarah Peacock, Travis Barman, Birgit Fuhrmeister, Evangelos Nagel, Andreas Schweitzer, Sandra V. Jeffers, Ignasi Ribas, Ansgar Reiners, Andreas Quirrenbach, Pedro J. Amado, Victor J. S. Béjar, José A. Caballero, Artie P. Hatzes, David Montes
Figure caption: Temperature profiles of models M1 (green) and M2 (blue) as a function of column mass density (right panel). On the left-hand side, we indicate approximate atmospheric temperatures above the temperature minimum where most of the line fluxes of Lyα, C IV, Hα, He I IRT, Ca II IRT, and Na I D2 in model M2 originate from according to the flux contribution function (Magain 1986; Fuhrmeister et al. 2006; Hintz et al. 2019). The corresponding temperature regions of M1 are similar. For the Ca II IRT and Na I D2 lines, they are actually the same, while they are slightly lower for the other lines, but ∼2500 K at most. For M2 we indicate by gray lines the boundaries of the transition region and upper (UC) and lower chromosphere (LC), as well as the photosphere.
Paper: Modeling the Chromosphere and Transition Region of Planet-hosting Star GJ 436
Authors: Dominik Hintz, Sarah Peacock, Travis Barman, Birgit Fuhrmeister, Evangelos Nagel, Andreas Schweitzer, Sandra V. Jeffers, Ignasi Ribas, Ansgar Reiners, Andreas Quirrenbach, Pedro J. Amado, Victor J. S. Béjar, José A. Caballero, Artie P. Hatzes, David Montes
Figure caption: Figure 3 continued for the lines of C IV (upper left), He I IRT (upper right), Na I D2 (lower left), and the bluest Ca II IRT (lower right). The colors coincide with those used in Figure 3.
Paper: Modeling the Chromosphere and Transition Region of Planet-hosting Star GJ 436
Authors: Dominik Hintz, Sarah Peacock, Travis Barman, Birgit Fuhrmeister, Evangelos Nagel, Andreas Schweitzer, Sandra V. Jeffers, Ignasi Ribas, Ansgar Reiners, Andreas Quirrenbach, Pedro J. Amado, Victor J. S. Béjar, José A. Caballero, Artie P. Hatzes, David Montes
Figure caption: Activity measurements for model series M2−M2d as given by the integrated line fluxes (LF; left panel) of Lyα (crosses) and C IV (pluses), as well as the equivalent widths (EW; right panel) of Hα (circles) and the He I IRT (diamonds) for models M2–M2d as a function of ∇T UC. The colors of the models coincide with those in Figures 3 and B2. We also show the corresponding values of M1 (green markers) except for the LF (C IV), which equals to 1251.1 erg s−1 cm−2.
Paper: Modeling the Chromosphere and Transition Region of Planet-hosting Star GJ 436
Authors: Dominik Hintz, Sarah Peacock, Travis Barman, Birgit Fuhrmeister, Evangelos Nagel, Andreas Schweitzer, Sandra V. Jeffers, Ignasi Ribas, Ansgar Reiners, Andreas Quirrenbach, Pedro J. Amado, Victor J. S. Béjar, José A. Caballero, Artie P. Hatzes, David Montes
Figure caption: The linear combination obtained from Section 3.2 for the NUV region compared to an HST STIS observation (G230L grating) of GJ 436. The model spectrum is reduced to the resolving power of 1000 corresponding to the maximum resolution of the STIS G230L grating. Here we also indicate a couple of strong lines and line bands within this spectral region.
Paper: Modeling the Chromosphere and Transition Region of Planet-hosting Star GJ 436
Authors: Dominik Hintz, Sarah Peacock, Travis Barman, Birgit Fuhrmeister, Evangelos Nagel, Andreas Schweitzer, Sandra V. Jeffers, Ignasi Ribas, Ansgar Reiners, Andreas Quirrenbach, Pedro J. Amado, Victor J. S. Béjar, José A. Caballero, Artie P. Hatzes, David Montes
Figure caption: Application of the Lyα ISM transmission curve (green dotted line) for GJ 436 as found by Youngblood et al. (2016) to our model M2 (blue solid line). The resulting modeled Lyα line is shown in red. We also indicate the Lyα line core area we neglected within our modified χ 2-minimization by the two dashed vertical lines.
Paper: Observations and Modeling of Unstable Proton and α Particle Velocity Distributions in Sub-Alfvénic Solar Wind at Parker Solar Probe Perihelia
Authors: Leon Ofman, Scott A. Boardsen, Lan K. Jian, Parisa Mostafavi, Jaye L. Verniero, Roberto Livi, Michael McManus, Ali Rahmati, Davin Larson, Michael L. Stevens
Figure caption: An example of the sub-Alfvénic region observed by PSP during 2022 February 25 to 2022 February 27 UT during Encounter 11. Left: panels, from top to bottom, show M A, proton velocity, T ⊥/T ∥ for proton and α particles, α-to-proton density ratio, and ∣B∣, respectively. Right panels show the proton VDFs at the marked time (green line) on the left panels. The black arrows represent the magnetic field direction in SPAN-I coordinates, where the head is at the SW velocity and the length is the Alfvén speed. The top-right panel shows the proton VDF in the V  r –V  z  plane, the middle-right panel is the V  r –V  y  plane, and the lower-right panel shows lines of energy sweeps at different elevation angles. The obscuring effect of the PSP thermal shield on the observed proton VDFs is evident.
Paper: Observations and Modeling of Unstable Proton and α Particle Velocity Distributions in Sub-Alfvénic Solar Wind at Parker Solar Probe Perihelia
Authors: Leon Ofman, Scott A. Boardsen, Lan K. Jian, Parisa Mostafavi, Jaye L. Verniero, Roberto Livi, Michael McManus, Ali Rahmati, Davin Larson, Michael L. Stevens
Figure caption: The temporal evolution of the α particle VDF for Case 1 with the initial parameters of E11 at 2022 February 25 21:50:50 UT with A  p  = 4.3, A  α  = 1.6, n  α /n  p  = 0.007, and V  d  = 0.32V A. The upper panels show the α particle VDFs in the V  x –V  z  plane, and the lower panels show the cut along V  x  through the corresponding peaks of the α particle VDFs at ﹩t=0,62.5,150,600{{\rm{\Omega }}}_{p}^{-1}﹩.
Paper: Observations and Modeling of Unstable Proton and α Particle Velocity Distributions in Sub-Alfvénic Solar Wind at Parker Solar Probe Perihelia
Authors: Leon Ofman, Scott A. Boardsen, Lan K. Jian, Parisa Mostafavi, Jaye L. Verniero, Roberto Livi, Michael McManus, Ali Rahmati, Davin Larson, Michael L. Stevens
Figure caption: The temporal evolution of the proton VDF for Case 1 with the initial parameters of E11 at 2022 February 25 21:50:50 UT with A  p  = 4.3, A  α  = 1.6, n  α /n  p  = 0.007, and V  d  = 0.32V A. The upper panels show the proton VDFs in the V  x –V  z  plane, and the lower panels show the cut along V  x  through the corresponding peaks of the proton VDFs at ﹩t=0,62.5,150,600{{\rm{\Omega }}}_{p}^{-1}﹩.
Paper: Observations and Modeling of Unstable Proton and α Particle Velocity Distributions in Sub-Alfvénic Solar Wind at Parker Solar Probe Perihelia
Authors: Leon Ofman, Scott A. Boardsen, Lan K. Jian, Parisa Mostafavi, Jaye L. Verniero, Roberto Livi, Michael McManus, Ali Rahmati, Davin Larson, Michael L. Stevens
Figure caption: An example of kinetic wave activity analysis of the magnetic field observed by the PSP FIELDS instrument during 2021 November 22 00:00 UT to 04:00 UT. Left panels show (a) M A, (b) magnetic compressibility ﹩{(\delta {B}_{\parallel }/\delta B)}^{2}﹩, (c) magnetic PSD, (d) DOP, (e) ellipticity, and (f) WNA in degrees. The two white lines (b)–(f) are at the proton and α cyclotron frequencies. Right panels show the (g) linearly detrended magnetic fluctuations δ B time series in field-aligned coordinates for 20 s, and the (h) power spectrum with the compressional (red) and transverse (blue) polarization. The proton-cyclotron frequency F cp is indicated by the vertical line.
Paper: Observations and Modeling of Unstable Proton and α Particle Velocity Distributions in Sub-Alfvénic Solar Wind at Parker Solar Probe Perihelia
Authors: Leon Ofman, Scott A. Boardsen, Lan K. Jian, Parisa Mostafavi, Jaye L. Verniero, Roberto Livi, Michael McManus, Ali Rahmati, Davin Larson, Michael L. Stevens
Figure caption: An example of the sub-Alfvénic region observed by PSP on 2022 June 3 03:30–5:00 UT during Encounter 12. The format of the panels is similar to Figure 1.
Paper: Observations and Modeling of Unstable Proton and α Particle Velocity Distributions in Sub-Alfvénic Solar Wind at Parker Solar Probe Perihelia
Authors: Leon Ofman, Scott A. Boardsen, Lan K. Jian, Parisa Mostafavi, Jaye L. Verniero, Roberto Livi, Michael McManus, Ali Rahmati, Davin Larson, Michael L. Stevens
Figure caption: An example of the sub-Alfvénic region observed by PSP during 2021 November 21 to 23 during Encounter 10. The format of the panels is similar to Figure 1.
Paper: Scaling of Magnetic Reconnection Electron Bulk Heating in the High-Alfvén-speed and Low-β Regime of Earth’s Magnetotail
Authors: M. Øieroset, T. D. Phan, M. Oka, J. F. Drake, S. A. Fuselier, D. J. Gershman, K. Maheshwari, B. L. Giles, Q. Zhang, F. Guo, J. L. Burch, R. B. Torbert, R. J. Strangeway
Figure caption: Electron temperature change ΔT e vs. inflow parameters. Events with lobe inflow are in blue and events with plasma sheet inflow are in black. (a) ΔT e vs. inflow Alfvén speed based on the reconnecting field component B  X . (b) ΔT e vs. available magnetic energy per particle. (c) ΔT e vs. inflow density. (d) ΔT e vs. inflow ∣B  X ∣. (e) ΔT e vs. inflow electron temperature. (f) ΔT e vs. inflow electron β. (g) Outflow ΔT e∣∣ vs. outflow ΔT e⊥. (h) Inflow T e∣∣ vs. inflow T e⊥. The black solid lines in (g) and (h) mark T e∣∣ = T e⊥, while the dashed lines in (b), (g), and (h) are linear least-squares fits to the data.
Paper: Scaling of Magnetic Reconnection Electron Bulk Heating in the High-Alfvén-speed and Low-β Regime of Earth’s Magnetotail
Authors: M. Øieroset, T. D. Phan, M. Oka, J. F. Drake, S. A. Fuselier, D. J. Gershman, K. Maheshwari, B. L. Giles, Q. Zhang, F. Guo, J. L. Burch, R. B. Torbert, R. J. Strangeway
Figure caption: MMS3 observations of a magnetotail high-speed flow reversal event with lobe inflow. (a) Magnetic field. (b) Proton spectrogram. (c) Electron spectrogram. (d) Proton velocity. (e) Proton V  X  component perpendicular to B. (f) Proton density. (g) Parallel (red) and perpendicular (black) electron temperature. (h) Alfvén speed based on proton density and B  X . (i) and (j) Maxwellian fit (red curve) to the average electron distribution function measurements in the two lobe (inflow) intervals marked by pairs of blue vertical lines in (a)–(h). Only the red data points have been used for the Maxwellian fitting (see the text for details). The two Maxwellian fits in (i) and (j) result in inflow electron temperatures of 89 eV and 83 eV, respectively. The solid black curve in (j) and (k) is the one-count level for the electron instrument. The measurement uncertainty for each data point is displayed with black vertical lines. (k) Schematic of the reconnection geometry in the magnetotail and approximate spacecraft trajectory (blue dashed curve).
Paper: Scaling of Magnetic Reconnection Electron Bulk Heating in the High-Alfvén-speed and Low-β Regime of Earth’s Magnetotail
Authors: M. Øieroset, T. D. Phan, M. Oka, J. F. Drake, S. A. Fuselier, D. J. Gershman, K. Maheshwari, B. L. Giles, Q. Zhang, F. Guo, J. L. Burch, R. B. Torbert, R. J. Strangeway
Figure caption: MMS3 observations of a magnetotail high-speed flow reversal event with plasma sheet inflow. (a) Magnetic field components in GSM coordinates. (b) Proton spectrogram in differential energy flux. (c) Electron spectrogram. (d) Proton velocity. (e) Proton V  X  component perpendicular to B. (f) Proton density. (g) Parallel (red) and perpendicular (black) electron temperature. (h) Alfvén speed based on proton density and B  X . (i) Schematic of the reconnection geometry in the magnetotail and the approximate MMS spacecraft trajectory (the blue dashed curve). The pairs of red vertical lines in (a)–(h) mark outflow intervals and the pair of blue vertical lines mark an inflow interval.
Paper: Scaling of Magnetic Reconnection Electron Bulk Heating in the High-Alfvén-speed and Low-β Regime of Earth’s Magnetotail
Authors: M. Øieroset, T. D. Phan, M. Oka, J. F. Drake, S. A. Fuselier, D. J. Gershman, K. Maheshwari, B. L. Giles, Q. Zhang, F. Guo, J. L. Burch, R. B. Torbert, R. J. Strangeway
Figure caption: Electron temperature change normalized to inflow parameters. Events with lobe inflow are in blue and events with plasma sheet inflow are in black. (a) ΔT e normalized to m i                      ﹩{V}_{\mathrm{Ax},\mathrm{in}}^{2}﹩ vs. inflow electron temperature. (b) ΔT e and available magnetic energy per particle normalized to the inflow electron temperature. (c) ΔT e normalized to the inflow electron temperature vs. the inflow electron β. (d) ΔT e vs. the exhaust density divided by inflow density. The dashed line in (b) is the linear least-squares fit to the data.
Paper: Wobbling Jets in Common Envelope Evolution
Authors: Noam Dori, Ealeal Bear, Noam Soker
Figure caption: The ratio of the convective velocity to the Keplerian velocity, v conv(r)/v Kep(r), as a function of radius inside the envelope. The thin-black line in each panel depicts this quantity when we just start to remove mass at a high rate. The blue line represents the envelope after we removed an envelope mass of ΔM env = 9.5M ⊙ to have M env(0) = 1.9M ⊙; we set t = 0 when mass removal ends. At t = 0 we start to inject energy from radius R in(t) = 0.07R 1(t) to R 1(t), where R 1(t) is the stellar radius. The green, red, and cyan lines denote t = 0.1, 0.25, and 0.35 yr, respectively. The insets give the mass removal rate and the power of energy injection ﹩(\dot{M},{\dot{E}}_{\mathrm{in}})﹩. The last panel does not have the line at t = 0.35 yr because the expansion at that time is on a shorter timescale than the dynamical timescale, and the results are unreliable. Note the different scaling of the axes in the different panels.
Paper: Wobbling Jets in Common Envelope Evolution
Authors: Noam Dori, Ealeal Bear, Noam Soker
Figure caption: Similar to Figure 1, but the energy injection is to a shell R in(t) = 0.25R 1(t) < r < R 1(t) (rather than into a shell R in(t) = 0.07R 1(t) < r < R 1(t)).
Paper: Near-Sun In Situ and Remote-sensing Observations of a Coronal Mass Ejection and its Effect on the Heliospheric Current Sheet
Authors: O. M. Romeo, C. R. Braga, S. T. Badman, D. E. Larson, M. L. Stevens, J. Huang, T. Phan, A. Rahmati, R. Livi, S. T. Alnussirat, P. L. Whittlesey, A. Szabo, K. G. Klein, T. Niembro-Hernandez, K. Paulson, J. L. Verniero, D. Lario, N. E. Raouafi, T. Ervin, J. Kasper, M. Pulupa, S. D. Bale, M. G. Linton
Figure caption: Illustrations of (a) a single ICME model based on coronal magnetograms and in situ data, and (b) a simplified two-part ICME model based on white-light coronagraphs and in situ observations. PSP travels along a circular orbit around the Sun from top to bottom in each figure. The magnetic field lines are represented by the black curves connected to the solar surface, with arrows depicting their respective polarity. The purple dashed lines depict the magnetic inversion lines observed by PSP at times t 4a  and t 7a . The single ICME in (a) is portrayed by the blue region, while CME1 and CME2 in (b) are shown by the blue and orange regions, respectively. The active region AR 13102 is shown on the solar surface in (a) with positive (red) and negative (blue) polarity to agree with the ADAPT-GONG magnetograms. The gray area in (b) represents the region of space outside of CME1 that PSP travels through between times t 4a  and t 6. CME1 could be concave in three dimensions, allowing PSP to exit and re-enter the structure.
Paper: Near-Sun In Situ and Remote-sensing Observations of a Coronal Mass Ejection and its Effect on the Heliospheric Current Sheet
Authors: O. M. Romeo, C. R. Braga, S. T. Badman, D. E. Larson, M. L. Stevens, J. Huang, T. Phan, A. Rahmati, R. Livi, S. T. Alnussirat, P. L. Whittlesey, A. Szabo, K. G. Klein, T. Niembro-Hernandez, K. Paulson, J. L. Verniero, D. Lario, N. E. Raouafi, T. Ervin, J. Kasper, M. Pulupa, S. D. Bale, M. G. Linton
Figure caption: Encounter 13 PSP orbit in HEEQ coordinates, where the trajectory follows a counterclockwise motion around the Sun (red, not to scale) with time. The CME1 (blue region) model at time 22-09-05/17:27:20 UT, when PSP observes a fast-forward shock, is shown in the (a) XY plane (solar equatorial plane) and (b) XZ plane. The dark blue outline corresponds to the CME cross section within the orbital plane of PSP. The shock plane and normal are represented by a red line segment and arrow, respectively. At time 22-09-05/20:00 UT when PSP briefly enters the flank, CME1 is illustrated in the (c) XY plane and (d) XZ plane. Both CME1 and CME2 (orange region) at time 22-09-06/02:30 UT, when PSP enters the leg of CME1, are shown in the (e) XY plane and (f) XZ plane. The dark red outline along the wire frame of CME2 corresponds to the cross-sectional area within the orbital plane of PSP. At time 22-09-06/16:40 UT, when PSP passes the second magnetic inversion line, CME1 and CME2 are modeled in the (a) XY plane and (b) XZ plane. In addition, the orbital position of STA and Earth are depicted by black arrows, with the coronagraph COR2 FOV extent portrayed by two green lines in the solar equatorial plane.