adsabs/AstroCLIMB huggingface dataset:
Paper: Sodium Distribution on the Moon
Authors: S. Narendranath, Netra S. Pillai, Srikar P. Tadepalli, Menelaos Sarantos, K. Vadodariya, A. Sarwade, Radhakrishna V, A. Tyagi
Figure caption: Variation in Na wt% as a function of LST. A significant decrease in Na wt% is measured at 07:00 LST followed by a sinusoidal trend which again tends to increase at 16:00 LST.
Paper: Sodium Distribution on the Moon
Authors: S. Narendranath, Netra S. Pillai, Srikar P. Tadepalli, Menelaos Sarantos, K. Vadodariya, A. Sarwade, Radhakrishna V, A. Tyagi
Figure caption: Mg and Al abundance averaged across latitudes do not show a trend for Na vs. LST. Mg and Al anticorrelate, as expected, for the lunar surface.
Paper: Sodium Distribution on the Moon
Authors: S. Narendranath, Netra S. Pillai, Srikar P. Tadepalli, Menelaos Sarantos, K. Vadodariya, A. Sarwade, Radhakrishna V, A. Tyagi
Figure caption: Mg abundance map (150 km × 12.5 km) showing the known lunar geochemical trend.
Paper: Rest-frame Near-infrared Sizes of Galaxies at Cosmic Noon: Objects in JWST's Mirror Are Smaller than They Appeared
Authors: Katherine A. Suess, Rachel Bezanson, Erica J. Nelson, David J. Setton, Sedona H. Price, Pieter van Dokkum, Gabriel Brammer, Ivo Labbé, Joel Leja, Tim B. Miller, Brant Robertson, Arjen van der Wel, John R. Weaver, Katherine E. Whitaker
Figure caption: Estimate of color gradient strength as probed by the ratio of 4.4–1.5 μm sizes as a function of both stellar mass and rest-frame V − J color in three redshift bins. Marker color indicates UVJ-selected star-forming and quiescent galaxies. We find significant trends with both mass and color, such that redder, more massive galaxies have larger differences between their 4.4 and 4.4 μm sizes. These color gradients indicate that more massive galaxies have increasingly redder centers compared to their outskirts.
Paper: Rest-frame Near-infrared Sizes of Galaxies at Cosmic Noon: Objects in JWST's Mirror Are Smaller than They Appeared
Authors: Katherine A. Suess, Rachel Bezanson, Erica J. Nelson, David J. Setton, Sedona H. Price, Pieter van Dokkum, Gabriel Brammer, Ivo Labbé, Joel Leja, Tim B. Miller, Brant Robertson, Arjen van der Wel, John R. Weaver, Katherine E. Whitaker
Figure caption: Example GALFIT fits for four objects in our sample, ordered from lowest to highest stellar mass. From top to bottom, the objects lie at redshifts of 1.9, 2.2, 1.2, and 1.5. The left column shows three-color images using JWST/NIRCam F150W, F277W, and F444W. The center three panels show the 4.4 μm data, our best-fit model, and the residual. For comparison, the right panel shows the difference between the 4.4 μm data and the best-fit 1.6 μm GALFIT model convolved with the 4.4 μm PSF. The data, model, and residuals are on the same symmetric colorbar for each object, with positive values in gray scale and negative values in red. Our fitting procedure is able to accurately reproduce galaxy cutouts for both isolated and crowded fields, but cannot capture complex morphologies such as spiral arms. The 1.6 μm model tends to underpredict the flux at the center of the galaxy; this is especially true for the higher-mass galaxies in the bottom rows of the plot.
Paper: Rest-frame Near-infrared Sizes of Galaxies at Cosmic Noon: Objects in JWST's Mirror Are Smaller than They Appeared
Authors: Katherine A. Suess, Rachel Bezanson, Erica J. Nelson, David J. Setton, Sedona H. Price, Pieter van Dokkum, Gabriel Brammer, Ivo Labbé, Joel Leja, Tim B. Miller, Brant Robertson, Arjen van der Wel, John R. Weaver, Katherine E. Whitaker
Figure caption: 4.4 μm size as a function of 1.5 μm size in arcseconds, colored by star-forming or quiescent based on rest-frame UVJ colors. Galaxies are slightly smaller in 4.4 μm than in 1.5 μm, the longest-wavelength filter previously accessible with HST. The lower right shows a typical error bar estimated from mock recovery tests (see text).
Paper: Rest-frame Near-infrared Sizes of Galaxies at Cosmic Noon: Objects in JWST's Mirror Are Smaller than They Appeared
Authors: Katherine A. Suess, Rachel Bezanson, Erica J. Nelson, David J. Setton, Sedona H. Price, Pieter van Dokkum, Gabriel Brammer, Ivo Labbé, Joel Leja, Tim B. Miller, Brant Robertson, Arjen van der Wel, John R. Weaver, Katherine E. Whitaker
Figure caption: Galaxies in this paper (gray points) compared to all galaxies in the 3D-HST photometric catalog with ﹩\mathrm{log}{M}_{* }/{M}_{\odot }\gt 9﹩ and 1.0 ≤ z ≤ 2.5 (Skelton et al. 2014). Our targets follow a similar stellar mass and photometric redshifts distribution as the full 3D-HST catalog, showing that the subset is fairly representative of the full survey. Due to the limited area of the current CEERS observation, relatively few quiescent galaxies (indicated by the box in UVJ color–color space) appear in our sample.
Paper: Formation of Moons and Equatorial Ridge around Top-shaped Asteroids after Surface Landslide
Authors: Ryuki Hyodo, Keisuke Sugiura
Figure caption: Comparison of asteroid shapes between the initial figure (t = 0 hr; left panel) and after the disk evolution (t ∼ 700 hr; right panel). The black circle indicates a circle with a radius of 500 m.
Paper: Turbulent Dust-trapping Rings as Efficient Sites for Planetesimal Formation
Authors: Ziyan Xu, Xue-Ning Bai
Figure caption: Snapshots of the projected dust densities in the x–z (upper four panels) and x–y (lower four panels) planes for the runs Z0, Z2H, Z0w, and Z2w at time ﹩T=1380{{\rm{\Omega }}}_{K}^{-1}﹩, when the secondary filament in Z2w is clearly seen. The dust density in the x–z plane is azimuthally averaged over the whole domain, while in the x–y plane, the dust density is vertically averaged around the midplane (z = ±0.05 H), to highlight the dust structures.
Paper: Turbulent Dust-trapping Rings as Efficient Sites for Planetesimal Formation
Authors: Ziyan Xu, Xue-Ning Bai
Figure caption: Upper four panels: radial profiles of the vertically and azimuthally averaged dust density (ρ  d ; the solid curves in the upper two panels), the gas density (ρ  g ; the dashed curves in the upper two panels), and the Keplerian-subtracted azimuthal gas velocity (﹩{v}_{y}^{{\prime} }﹩; the solid curves in the lower two panels), at the midplane (±1 cell from z = 0), averaged between ﹩T=800\mbox{--}1500{{\rm{\Omega }}}_{K}^{-1}﹩. The vertical dashed lines mark the locations of the dust rings (the peaks of dust density), with colors corresponding to those used for the radial profiles of each run. Lower four panels: time evolutions of the radial profiles of ﹩{v}_{y}^{{\prime} }﹩, between ﹩T=800\mbox{--}1500{{\rm{\Omega }}}_{K}^{-1}﹩.
Paper: Turbulent Dust-trapping Rings as Efficient Sites for Planetesimal Formation
Authors: Ziyan Xu, Xue-Ning Bai
Figure caption: Time evolution of the dust maximum density (upper panel), the dust ring width (middle panel), and the dust scale height (lower panel) for all our simulation runs. The dust maximum densities for the no-feedback runs (Z0, Z0H, and Z0w) are scaled to Z = 0.02.
Paper: Turbulent Dust-trapping Rings as Efficient Sites for Planetesimal Formation
Authors: Ziyan Xu, Xue-Ning Bai
Figure caption: A snapshot of a 3D visualization of the dust spatial distribution in our simulation, with a strong gas bump (run Z2H) at the time of ﹩T=1380{{\rm{\Omega }}}_{K}^{-1}﹩ after inserting dust particles. The dust density is mapped to the brightness, where the brighter regions indicate higher dust densities. Also shown in the figure are the vertically averaged gas density, projected to the x–y (R–ϕ) plane, and the vertically and azimuthally averaged radial profiles of the dust and gas densities, with the amplitudes of the profiles being rescaled for better illustration.
Paper: Turbulent Dust-trapping Rings as Efficient Sites for Planetesimal Formation
Authors: Ziyan Xu, Xue-Ning Bai
Figure caption: The azimuthally averaged spatial distribution of the dust density (the upper panels of both plots), the gas rms velocity in the radial direction (the middle panels), and in the vertical direction (the lower panels), time-averaged within ﹩T=800\mbox{--}1500{{\rm{\Omega }}}_{K}^{-1}﹩, for runs Z0 (left) and Z2H (right). Note that this figure is not to the real aspect ratio.
Paper: Turbulent Dust-trapping Rings as Efficient Sites for Planetesimal Formation
Authors: Ziyan Xu, Xue-Ning Bai
Figure caption: Upper panel: time evolution of the dust ring width w  d  as a function of time ﹩T=720{{\rm{\Omega }}}_{K}^{-1}﹩ to ﹩2880{{\rm{\Omega }}}_{K}^{-1}﹩, featuring later stages of the runs Z1w and Z2w. Lower two panels: time evolution of the azimuthally and vertically averaged radial profiles of the dust density at the midplane (±1 cell from z = 0), for runs Z1w and Z2w, for the same time interval as the upper panel.
Paper: Turbulent Dust-trapping Rings as Efficient Sites for Planetesimal Formation
Authors: Ziyan Xu, Xue-Ning Bai
Figure caption: CDFs of the dust density for runs Z2H and Z2w, as labeled in the titles of each panel. The upper panels show CDFs for the whole simulation domain, whereas the bottom two panels measure the CDFs of run Z2w for dust particles residing in the x < −0.5 H and x > −0.5 H regions separately, in order to investigate the clump properties in the primary and secondary rings. The CDFs are measured by averaging within different time intervals in units of ﹩{{\rm{\Omega }}}_{K}^{-1}﹩, as marked by the different colors (see the legend). The black curves in each panel shows corresponding runs with no feedback, i.e., Z0 and Z0w, averaged from ﹩T=720{{\rm{\Omega }}}_{K}^{-1}﹩ to ﹩1520{{\rm{\Omega }}}_{K}^{-1}﹩.
Paper: Turbulent Dust-trapping Rings as Efficient Sites for Planetesimal Formation
Authors: Ziyan Xu, Xue-Ning Bai
Figure caption: The same as Figure C1, but for the runs Z0w and Z2w.
Paper: Cooling Envelope Model for Tidal Disruption Events
Authors: Brian D. Metzger
Figure caption: Evolution of TDE cooling envelope model with M ⋆ = M ⊙, M • = 2 × 106  M ⊙, β = 1, shown as a function of time since the circularization/envelope assembly time t = t fb. Solid lines show a model neglecting mass accretion onto the SMBH (i.e., ﹩{\dot{M}}_{\bullet }={\dot{E}}_{\bullet }=0﹩), while dotted lines show the effects of including SMBH accretion mass loss (Equation (21) for α = 0.01, H/r = 0.3) but not energy feedback. The top panel shows the characteristic envelope radius R  v , mass M  e /M ⋆, and photosphere radius R ph, while the bottom panel shows the envelope luminosity L rad, optical luminosity ν L  ν  at frequency ν = 6 × 1014 Hz (g band), effective temperature T eff, and ﹩{L}_{{\rm{X}}}\equiv 0.01{\dot{M}}_{\bullet }{c}^{2}﹩, taken as a proxy for the X-ray luminosity from the inner disk (potentially observable only through a narrow polar region). The envelope evolution concludes roughly once R  v  decreases to R circ (horizontal dashed line), as occurs roughly at the times t disk (Equation (33)) and ﹩{t}_{\mathrm{acc}}^{\star }﹩ (Equation (34)), respectively, in the two models. A dashed black line illustrates ∝ t −3/2 decay.
Paper: Cooling Envelope Model for Tidal Disruption Events
Authors: Brian D. Metzger
Figure caption: Schematic illustration of the model. Rapid circularization of the TDE debris forms a quasi-circular pressure-supported envelope around the SMBH of characteristic radius R  v  and photosphere radius R ph ∼ 10R  v  that powers the early optical emission. The envelope luminosity L rad primarily derives from the gravitational energy released from its cooling-driven contraction at close to the Eddington limit, though fallback accretion (which deposits its energy at a radius R acc ≳ R  v , where the stream dissolves inside the envelope after passing through pericenter) may contribute significantly at early times. As the envelope cools and contracts, roughly as R  v  ∝ R ph ∝ t −1 initially, the effective temperature rises T eff ∝ t 1/2 while the optical luminosity ﹩\nu {L}_{\nu }\propto {R}_{\mathrm{ph}}^{2}{T}_{\mathrm{eff}}\propto {t}^{-3/2}﹩ drops. The SMBH accretion rate rises in tandem, as controlled by the envelope density and viscous time near the circularization radius, potentially powering thermal X-ray and/or jetted activity along the initially narrow polar funnel. The inner accretion flow also acts as a source of energy to the envelope, which can delay the envelope’s contraction in a regulated manner, flattening the late-time optical and X-ray light-curve decay.
Paper: Cooling Envelope Model for Tidal Disruption Events
Authors: Brian D. Metzger
Figure caption: Optical-band light curves ν L  ν  at ν = 6 × 1014 Hz and proxy X-ray light curves ﹩{L}_{{\rm{X}}}={10}^{-2}{\dot{M}}_{\bullet }{c}^{2}﹩ for a series of models that adopt fiducial parameters (M ⋆ = M ⊙, M • = 2 × 106  M ⊙, β = 1, α = 10−2, H/r = 0.3, ζ = 2) but varying the SMBH feedback parameter η and envelope power-law density index ξ as marked.
Paper: Cooling Envelope Model for Tidal Disruption Events
Authors: Brian D. Metzger
Figure caption: Models calculated for the same parameters as in Figure 3, but now including the effects of SMBH feedback on the envelope structure for two different values of η = 1.5 × 10−2 (solid lines) and η = 10−3 (dotted lines). Powerful feedback (large η) acts to slow the rate of envelope cooling and accretion, flattening the late-time ν L  ν  optical light-curve decay. A dashed black line shows ∝ t −3/2 decay, as roughly expected absent efficient feedback.
Paper: Cooling Envelope Model for Tidal Disruption Events
Authors: Brian D. Metzger
Figure caption: Optical-band light curves ν L  ν  at ν = 6 × 1014 Hz and proxy X-ray light curves ﹩{L}_{{\rm{X}}}={10}^{-2}{\dot{M}}_{\bullet }{c}^{2}﹩ for a series of models that adopt fiducial parameters (β = 1, α = 10−2, H/r = 0.3, ζ = 2) but varying the mass of the star and SMBH as marked. We assume an SMBH feedback parameter ﹩\eta ={\eta }_{\min }({M}_{\star },{M}_{\bullet })﹩. (Equation (24)).
Paper: A Low-inclination Neutral Trans-Neptunian Object in an Extreme Orbit
Authors: Ying-Tung Chen, 英同 陳, Marielle R. Eduardo, Marco A. Muñoz-Gutiérrez, Shiang-Yu Wang, 祥宇 王, Matthew J. Lehner, Chan-Kao Chang, 展誥 章
Figure caption: Heatmaps of the orbital evolution of 2016 SD106 clones and Haumea Family members in simulations with and without the HP. The white cross and green dot show the initial orbits of 2016 SD106 and HP, respectively. The initial orbits in each simulation are overplotted as black dots, while light blue dots show the final distribution of surviving particles after 1 Gyr. The observational q-gap suggested by Oldroyd & Trujillo (2021) is outlined with a red square. The black curve shows the stability limit for distant SDOs as derived by Batygin et al. (2021); above this line in the a–q plane, particles are quickly destabilized and ejected from the solar system.
Paper: A Low-inclination Neutral Trans-Neptunian Object in an Extreme Orbit
Authors: Ying-Tung Chen, 英同 陳, Marielle R. Eduardo, Marco A. Muñoz-Gutiérrez, Shiang-Yu Wang, 祥宇 王, Matthew J. Lehner, Chan-Kao Chang, 展誥 章
Figure caption: Same as Figure 3, but for plutinos and twotinos.
Paper: A Low-inclination Neutral Trans-Neptunian Object in an Extreme Orbit
Authors: Ying-Tung Chen, 英同 陳, Marielle R. Eduardo, Marco A. Muñoz-Gutiérrez, Shiang-Yu Wang, 祥宇 王, Matthew J. Lehner, Chan-Kao Chang, 展誥 章
Figure caption: The orbital distribution of 1000 clones of 2016 SD106 after simulations lasting 0.1 Gyr (orange dots) and 1 Gyr (blue dots). The red dots indicate the initial conditions. The clone orbits are much more stable than typical Centaurs/SDOs, with Δq < 1 au over time. Only a few percent of clones evolve to a > 1000 au.
Paper: Formation of Moons and Equatorial Ridge around Top-shaped Asteroids after Surface Landslide
Authors: Ryuki Hyodo, Keisuke Sugiura
Figure caption: Schematic summary of our paper. Panel (a): a rubble-pile asteroid spins up due to, for example, the YORP effect, small impacts, a close encounter with a planet, or re-accumulation after a catastrophic impact. Panel (b): a surface landslide occurs when a critical spin state is realized and a top-shaped figure is formed (Sugiura et al. 2021). Panel (c): a particulate disk spreads due to inelastic collisions and gravitational interactions among particles. Panel (d): a moon is gravitationally accreted outside the Roche limit of the central top-shaped body, and an axisymmetric equatorial ridge is formed due to the re-accretion of disk particles. Panels (e) and (f): the formed moon is lost or remains, depending on the long-term orbital evolution between the moon and the primary.
Paper: Formation of Moons and Equatorial Ridge around Top-shaped Asteroids after Surface Landslide
Authors: Ryuki Hyodo, Keisuke Sugiura
Figure caption: Evolutions of the disk mass (left panel; not including the mass of the largest moons) and the masses of the largest three moons (right panel). At around t ∼ 1 × 105 s with a time interval of Δt ∼ 104 s, the landslide occurs and a particle disk is formed (left panel). The disk spreads and starts to form gravitational clumps with masses larger than ∼ 0.001M c at t ∼ 5.5 × 105 s, where M c is the mass of the central primary. Two large moons are formed by the end of our simulations (t  =  25 × 105 s; ∼700 hr).
Paper: Formation of Moons and Equatorial Ridge around Top-shaped Asteroids after Surface Landslide
Authors: Ryuki Hyodo, Keisuke Sugiura
Figure caption: Overall disk evolution after the landslide. Gray particles represent member particles of the top-shaped primary. Dark-gray particles represent disk particles and member particles of gravitational clumps (including the formed moon at later epochs). The black circle indicates the Roche limit of the primary (r R ∼ 2.5R c, where R c  =  500 m is used).
Paper: Formation of Moons and Equatorial Ridge around Top-shaped Asteroids after Surface Landslide
Authors: Ryuki Hyodo, Keisuke Sugiura
Figure caption: Same as Figure 2, but edge on views.
Paper: Magnetic Switchbacks Heat the Solar Corona
Authors: M. Akhavan-Tafti, J. Kasper, J. Huang, L. Thomas
Figure caption: (a) Scatter plot of the relative magnitudes of the normal magnetic field and proton velocity across the leading switchback transition regions. The color bar represents the heliocentric distance of the observed switchback in solar radii [R S]. (b) Scatter plot of RD-to-TD event count ratio as a function of heliocentric distance at the leading switchback transition regions. The color bar represents the average proton velocity in the leading quiet region. The blue histogram indicates the event count inside each distance bin (bin width = 1 R S). The solid red curve represents the exponential fit, y = a e −bx , where a and b are the y-intercept (at x = 0) and the decay rate in the unit [R S −1], respectively. The light red shaded region corresponds to the 95% confidence interval within which the observed RD-to-TD ratios within each radial bin may fall. The dashed black line denotes a two-term Gaussian model fit to the RD-to-TD data: y = a 1 exp(-((x-b 1)/c 1)2) + a 2 exp (−((x-b 2)/c 2)2), where a 1 = 12.88, b 1 = 37.11, c 1 = 2.454, a 2 = 8.0e+14, b 2 = 554.9, c 2 = 84.8. The correlation factor is R 2 = 0.72. (c) The superposed epoch analysis of ∣B∣ (red), proton density N (blue), proton velocity ∣V∣ (green), and proton radial temperature T (black) within the five distinct regions, relative to the leading quiet solar wind (QL). The error bars represent the standard deviation in each region.
Paper: Magnetic Switchbacks Heat the Solar Corona
Authors: M. Akhavan-Tafti, J. Kasper, J. Huang, L. Thomas
Figure caption: Concept illustration of the generation mechanism and temporal evolution of a magnetic switchback (dark blue). Magnetic reconnection between closed coronal loops (orange line) and open field lines (gray line) is shown by a red cross. The reconnection inflow (V in) and outflow (V out) regions are marked by black and green arrows, respectively. The magenta arrows indicate the magnitude of B Norm relative to B QL. The shade of blue indicates magnetic pressure (P M) at the switchback spike region.
Paper: Magnetic Switchbacks Heat the Solar Corona
Authors: M. Akhavan-Tafti, J. Kasper, J. Huang, L. Thomas
Figure caption: Magnetic field profile of a magnetic switchback event in rotational (R), tangential (T), and normal (N) coordinates during the first PSP encounter. A switchback comprises a magnetic spike that is separated from the pristine solar wind by a transition region. The radial component of the magnetic field reverses inside the transition region. The interval is 120 s in duration.
Paper: Magnetic Switchbacks Heat the Solar Corona
Authors: M. Akhavan-Tafti, J. Kasper, J. Huang, L. Thomas
Figure caption: Boxplot of the normalized magnetic pressure, categorized based on switchback magnetic discontinuities. On each box, the red bar indicates the median, and the bottom and top edges of the blue box indicate the 25th and 75th percentiles, respectively. The whiskers extend to the most extreme data points not considered outliers, and the outliers are plotted individually using the “+” marker symbol.
Paper: Magnetic Switchbacks Heat the Solar Corona
Authors: M. Akhavan-Tafti, J. Kasper, J. Huang, L. Thomas
Figure caption: Heat map of the RD-to-TD boundary value, determined based on the relative magnitudes of the normal magnetic field and proton velocity across the leading switchback transition regions. The bold rectangle marks the RD-to-TD boundary value 0.4, below which both ∣B Norm∣/∣B∣ and ∣V Norm∣/∣V∣ < 0.4.
Paper: The Relation between the Mass Accretion Rate and the Disk Mass in Class I Protostars
Authors: Eleonora Fiorellino, Łukasz Tychoniec, Carlo F. Manara, Giovanni Rosotti, Simone Antoniucci, Fernando Cruz-Sáenz de Miera, Ágnes Kóspál, Brunella Nisini
Figure caption: Mass accretion rate vs. disk dust mass. Red filled circles are the Class I sources where ﹩{\dot{M}}_{\mathrm{acc}}﹩ is the mean value between the results assuming the birthline and 1 Myr old evolutive track. The purple filled circles are other Class I sources from the literature. Empty circles are Class II objects from the Lupus sample. The black dashed lines correspond to several disk depletion times ﹩{t}_{\nu }={\dot{M}}_{\mathrm{acc}}/{M}_{\mathrm{disk}}﹩, from 103 yr (top) to 107 yr. A representative error bar for the mass accretion rate is shown in gray on the top left corner of the plot.
Paper: The Relation between the Mass Accretion Rate and the Disk Mass in Class I Protostars
Authors: Eleonora Fiorellino, Łukasz Tychoniec, Carlo F. Manara, Giovanni Rosotti, Simone Antoniucci, Fernando Cruz-Sáenz de Miera, Ágnes Kóspál, Brunella Nisini
Figure caption: Mass accretion rate vs. disk dust mass. The symbols are the same as in Figure 1. The blue line corresponds to our best fit for the overall sample of Class I YSOs, while the light-blue lines show a subsample of the results of some chains. The dashed gray shows the best fit for the Class II YSOs (Manara et al. 2016).
Paper: The Relation between the Mass Accretion Rate and the Disk Mass in Class I Protostars
Authors: Eleonora Fiorellino, Łukasz Tychoniec, Carlo F. Manara, Giovanni Rosotti, Simone Antoniucci, Fernando Cruz-Sáenz de Miera, Ágnes Kóspál, Brunella Nisini
Figure caption: Mass accretion rate vs. disk dust mass. The symbols are the same as in Figure 1. The brown, light-blue, and blue dashed lines are isochrones from Lodato et al. (2017) at t  ν  = 104, 105, and 106 yr, respectively. For each t  ν  we show in solid, dashed, and dashed–dotted lines isochrones for three different initial disk masses (0.01, 0.1, and 1 M ⊙, respectively).
Paper: First Detection of Transverse Vertical Oscillation during the Expansion of Coronal Loops
Authors: Qingmin Zhang, Chuan Li, Dong Li, Ye Qiu, Yanjie Zhang, Yiwei Ni
Figure caption: Time–distance plots of S3 in various EUV wavelengths of SDO/AIA. The speeds of contraction (black lines) are labeled. The expansion and vertical oscillation of the ACLs in 304 Å and 171 Å are denoted by “+” symbols in the top two panels.
Paper: First Detection of Transverse Vertical Oscillation during the Expansion of Coronal Loops
Authors: Qingmin Zhang, Chuan Li, Dong Li, Ye Qiu, Yanjie Zhang, Yiwei Ni
Figure caption: The coronal mass ejection (CME) observed by LASCO/C2 at 06:12:08 UT and STA/COR2 during 06:23:30−06:53:30 UT. The CME-driven shock with lower intensity is pointed by arrows.