adsabs/AstroCLIMB huggingface dataset:
Paper: Self-ionizing Galactic Winds
Authors: Kartick C. Sarkar, Amiel Sternberg, Orly Gnat
Figure caption: Column densities for “warm” ions showing that the O IV (IP = 75 eV) and O V (IP = 110 eV) dominate the total oxygen outflowing mass. Different line styles are—dotted: PIE, thin-solid: NEI-SR, thick solid: NEI-SR with ISM absorption.
Paper: Simulated Bars May Be Shorter but Are Not Slower Than Those Observed: TNG50 versus MaNGA
Authors: Neige Frankel, Annalisa Pillepich, Hans-Walter Rix, Vicente Rodriguez-Gomez, Jason Sanders, Jo Bovy, Juna Kollmeier, Norm Murray, Ted Mackereth
Figure caption: Same as Figure 2 for two TNG50 galaxies whose bar properties differ substantially when inspected based on stellar mass surface density vs. stellar light. In the top, we show a galaxy with an unrealistically large ﹩{R}_{{\rm{b}},\mathrm{obs}}^{\mathrm{TNG}50}﹩/﹩{R}_{{\rm{b}}}^{\mathrm{TNG}50}﹩ ratio, from an originally short bar that hardly seems identifiable in the image. In the bottom, we show a TNG50 galaxy with a bar nonvisible in the mock image. Most of such galaxies have small bars and contribute to increase the ﹩{ \mathcal R }﹩ value of the simulated galaxies, whereas they would not show up in an observed sample.
Paper: Simulated Bars May Be Shorter but Are Not Slower Than Those Observed: TNG50 versus MaNGA
Authors: Neige Frankel, Annalisa Pillepich, Hans-Walter Rix, Vicente Rodriguez-Gomez, Jason Sanders, Jo Bovy, Juna Kollmeier, Norm Murray, Ted Mackereth
Figure caption: Average relative difference of bar (left) and galaxy (right) properties between the high- and lower-resolution runs TNG50 and TNG50-2 between galaxies that started with the same initial conditions. On average and across the halo mass range explored here, bars in the lower-resolution run are slower (blue curve below 0 roughly by 40%–50%), larger (orange), and slightly stronger (green). The right panel shows that on average, matched galaxies in the lower-resolution run are smaller (see text) and contain less stellar mass in their half mass–radius (thick blue). The black line shows that across high- and lower-resolution runs, the halo mass is identical. In all panels, shades show the scatter around the mean relative differences between high- and lower-resolution runs.
Paper: Simulated Bars May Be Shorter but Are Not Slower Than Those Observed: TNG50 versus MaNGA
Authors: Neige Frankel, Annalisa Pillepich, Hans-Walter Rix, Vicente Rodriguez-Gomez, Jason Sanders, Jo Bovy, Juna Kollmeier, Norm Murray, Ted Mackereth
Figure caption: Spherically and sample averaged mass density in the simulated galaxies used in the analysis (top) and the resulting circular velocity curve (bottom), for both the high-resolution run (solid orange), the lower-resolution run (gray), and the high-resolution run that is smoothed in post-processing with a kernel of the lower-resolution softening length (dashed orange). This brings the central densities of the TNG50-1 and TNG50-2 galaxies to the same scale, especially near R  =  0.5 kpc (a small fraction of the average bar size). The same conclusions are reached when averaging in bins of stellar mass instead of the whole sample, which we do not show here for clarity.
Paper: Simulated Bars May Be Shorter but Are Not Slower Than Those Observed: TNG50 versus MaNGA
Authors: Neige Frankel, Annalisa Pillepich, Hans-Walter Rix, Vicente Rodriguez-Gomez, Jason Sanders, Jo Bovy, Juna Kollmeier, Norm Murray, Ted Mackereth
Figure caption: Distributions of observed and simulated bar properties using 3 different mass estimates to match the TNG50 galaxies to the observed ones in MaNGA. Using the mass in IFU definition tends to select galaxies of higher mass. As the pattern speed correlates with stellar mass, it shifts the pattern-speeds histogram slightly to higher pattern speeds. On the other hand, using the M elpetro definition, which makes different modeling assumptions in the stellar mass measurements, lead to selecting TNG50 galaxies with lower stellar masses. Overall, different stellar mass definitions do not significantly affect our conclusions.
Paper: Simulated Bars May Be Shorter but Are Not Slower Than Those Observed: TNG50 versus MaNGA
Authors: Neige Frankel, Annalisa Pillepich, Hans-Walter Rix, Vicente Rodriguez-Gomez, Jason Sanders, Jo Bovy, Juna Kollmeier, Norm Murray, Ted Mackereth
Figure caption: A rescaling of the particles mass and galactocentric radii to homogenize the central stellar mass and bar size induces an alignment of pattern speeds between the lower- and high-resolution runs of the TNG50 simulation.
Paper: Simulated Bars May Be Shorter but Are Not Slower Than Those Observed: TNG50 versus MaNGA
Authors: Neige Frankel, Annalisa Pillepich, Hans-Walter Rix, Vicente Rodriguez-Gomez, Jason Sanders, Jo Bovy, Juna Kollmeier, Norm Murray, Ted Mackereth
Figure caption: Average baryon dominance across the sample of simulated galaxies in bins of ﹩{ \mathcal R }﹩ for the TNG50-1 and TNG50-2 runs. On average, the lower-resolution run (TNG50-2) is less baryon dominated than the high-resolution one, but the differences across runs are much smaller than the scatter within runs highlighted by the shaded area for TNG50-1. We find no clear trend between the baryon dominance and the ﹩{ \mathcal R }﹩ within a simulation run.
Paper: Simulated Bars May Be Shorter but Are Not Slower Than Those Observed: TNG50 versus MaNGA
Authors: Neige Frankel, Annalisa Pillepich, Hans-Walter Rix, Vicente Rodriguez-Gomez, Jason Sanders, Jo Bovy, Juna Kollmeier, Norm Murray, Ted Mackereth
Figure caption: Distribution contours of simulated galaxies (green) and observed galaxies (orange) in the bar size–strength plane. For both simulated and observed galaxies, strong bars are also longer. But simulated bars are here systematically weaker and shorter than observed ones.
Paper: Simulated Bars May Be Shorter but Are Not Slower Than Those Observed: TNG50 versus MaNGA
Authors: Neige Frankel, Annalisa Pillepich, Hans-Walter Rix, Vicente Rodriguez-Gomez, Jason Sanders, Jo Bovy, Juna Kollmeier, Norm Murray, Ted Mackereth
Figure caption: Comparing the circular rotation speed Ωcirc at a characteristic radius (a fourth of the bar size measured in the lower-resolution galaxies) obtained in (1) the high-resolution run (gray circles), (2) the high-resolution run smoothed with the softening length of the lower-resolution run (orange triangles), to (3) the lower-resolution run (x-axis). The dashed line shows y = x for comparison.
Paper: Orbital Decay of Double White Dwarfs: Beyond Gravitational-wave Radiation Effects
Authors: G. A. Carvalho, R. C. dos Anjos, J. G. Coelho, R. V. Lobato, M. Malheiro, R. M. Marinho, J. F. Rodriguez, J. A. Rueda, R. Ruffini
Figure caption: Value of ﹩{\eta }_{\max }﹩ as a function of ω 0, given by Equation (13), for selected values of α.
Paper: The Effect of Shock-wave Duration on Star Formation and the Initial Condition of Massive Cluster Formation
Authors: Daisei Abe, Tsuyoshi Inoue, Rei Enokiya, Yasuo Fukui
Figure caption: Panel (a): scatter plot of the peak column density vs. the dense gas mass M dense(t sat) for all runs summarized in Table 1. Before computing the peak column density of the system, we take a smoothing of the column density using the Gaussian kernel function of a width of 0.5 pc, corresponding to the typical beamwidth for massive star-forming regions. The color of the points represents t dur/t ff indicating the influence of self-gravity. Panel (b): scatter plot of the peak column density vs. the estimated number of OB-type stars N OB. N OB is estimated using Equation (5). The black line is the fitting line in Figure 9(b) of Enokiya et al. (2021).
Paper: The Effect of Shock-wave Duration on Star Formation and the Initial Condition of Massive Cluster Formation
Authors: Daisei Abe, Tsuyoshi Inoue, Rei Enokiya, Yasuo Fukui
Figure caption: Column-density maps of the result of model v24t0.23 at time t = 0.3, 0.8, 1.6, and 1.8 Myr (from top to bottom). Left row (panels (a), (b), (c), and (d)): y–z plane column densities. Right row (panels (a′), (b′), (c′), and (d′)): x–y plane column densities.
Paper: The Effect of Shock-wave Duration on Star Formation and the Initial Condition of Massive Cluster Formation
Authors: Daisei Abe, Tsuyoshi Inoue, Rei Enokiya, Yasuo Fukui
Figure caption: Panel (a): the evolution of the total dense gas mass in the model v24t0.23. The blue, red, and green points represent the total masses with densities above 103, 104, and 105 cm−3, respectively. Panel (b): the same as panel (a), but for model v24t1.6.
Paper: The Effect of Shock-wave Duration on Star Formation and the Initial Condition of Massive Cluster Formation
Authors: Daisei Abe, Tsuyoshi Inoue, Rei Enokiya, Yasuo Fukui
Figure caption: Evolution of M dense(t) = M sink,tot(t) + M Av8(t) for the results of models v24t0.23 (left) and v24t1.6 (right), and the fitting curve obtained using a trial function ﹩\alpha \tanh [2\pi (t-{t}_{0})/{t}_{{\rm{w}}}]+\beta ﹩ (red line), where α, t 0, t w, and β are fitting parameters. Time t sat ≡ t 0 + t w/4 is defined as the time at which the total gas mass associated with the star formation is measured (dashed line).
Paper: The Effect of Shock-wave Duration on Star Formation and the Initial Condition of Massive Cluster Formation
Authors: Daisei Abe, Tsuyoshi Inoue, Rei Enokiya, Yasuo Fukui
Figure caption: Column-density maps in the result of model v24t1.6 at time t = 0.3, 0.8, 1.6, and 1.8 Myr (from top to bottom). Left row (panels (a), (b), (c), and (d)): y–z plane column densities. Right row (panels (a′), (b′), (c′), and (d′)): x–y plane column densities.
Paper: The Effect of Shock-wave Duration on Star Formation and the Initial Condition of Massive Cluster Formation
Authors: Daisei Abe, Tsuyoshi Inoue, Rei Enokiya, Yasuo Fukui
Figure caption: Schematic of the initial condition. The color bar represents the density magnitude, the black lines represent the initial magnetic field lines, the θ B represents the angle between the y-axis and the initial magnetic field lines, and the gray arrows represent the orientations of the converging flows.
Paper: The Effect of Shock-wave Duration on Star Formation and the Initial Condition of Massive Cluster Formation
Authors: Daisei Abe, Tsuyoshi Inoue, Rei Enokiya, Yasuo Fukui
Figure caption: Sink mass histogram at t = t sat in the results of models v14t0.38 (orange), v14t0.84 (red), and v141.9 (brown).
Paper: The Effect of Shock-wave Duration on Star Formation and the Initial Condition of Massive Cluster Formation
Authors: Daisei Abe, Tsuyoshi Inoue, Rei Enokiya, Yasuo Fukui
Figure caption: Sink mass histogram at t = t sat in the models v24t1.6 (red) and v24t1.6AMR (filled-in region).
Paper: A Data-constrained Magnetohydrodynamic Simulation of the X1.0 Solar Flare of 2021 October 28
Authors: Daiki Yamasaki, Satoshi Inoue, Yumi Bamba, Jeongwoo Lee, Haimin Wang
Figure caption: Comparison of RUN A and RUN B in terms of the vertical velocity, V  z  (a)–(c) and current density normalized by magnetic field, ∣ J ∣/∣ B ∣ (d–f). (a) and (b) show 2D distribution of V  z  in a vertical x − z plane at y = 0.34, obtained from RUN A and RUN B at t = 0.85, respectively. (c) shows their 1D distributions at the location of (x, y) = (0.50, 0.34) indicated by the white line marked in other panels. Likewise, (d, e) show the 2D distribution, and (f), the 1D distribution of ∣ J ∣/∣ B ∣ obtained from RUN A at t = 5.11 and RUN B at t = 5.06, respectively. Regions of enhanced ∣ J ∣/∣ B ∣ are regarded as edges of the MFR.
Paper: A Data-constrained Magnetohydrodynamic Simulation of the X1.0 Solar Flare of 2021 October 28
Authors: Daiki Yamasaki, Satoshi Inoue, Yumi Bamba, Jeongwoo Lee, Haimin Wang
Figure caption: Decay index distribution in 3D. (a) The index is plotted in vertical cross section of x–z plane at y = 0.34. The initial magnetic field lines (t = 0) in the simulation are plotted with the same color convention as used in Figure 3(b). (b) Azure lines are the potential field lines surrounding the MFR. (c) The field lines of the potential field are projected on the x–z plane at y = 0.34. The green line is the vertical line of (x, y) = (0.50, 0.34) selected for plotting 1D distribution of the decay index and the magnetic field strength in (d).
Paper: A Data-constrained Magnetohydrodynamic Simulation of the X1.0 Solar Flare of 2021 October 28
Authors: Daiki Yamasaki, Satoshi Inoue, Yumi Bamba, Jeongwoo Lee, Haimin Wang
Figure caption: Comparison of magnetic fields in RUN A and RUN B. The magnetic fields at t = 4.26 in RUN A (a) and those at t = 4.22 in RUN B (b). The same color convention as in Figure 3 is used to identify the three groups of field lines. Orange box indicated by orange arrow in (b) shows the region where we set the velocity to zero in RUN B. It lies in 0.45 ≤ x ≤ 0.54, 0.31 ≤ y ≤ 0.40, and 0.02 ≤ z ≤ 0.06, satisfying ∣ J ∣ ≥ 10.
Paper: A Data-constrained Magnetohydrodynamic Simulation of the X1.0 Solar Flare of 2021 October 28
Authors: Daiki Yamasaki, Satoshi Inoue, Yumi Bamba, Jeongwoo Lee, Haimin Wang
Figure caption: Observed and computed flare ribbons. Green lines in all the panels show the polarity inversion lines. (a) Flare ribbons in the GONG Hα line image at 15:37:10 UT. (b) The radial component of the photospheric magnetic field overlaid with the enhanced Hα intensity as red contours at 15:37:10 UT. (c) Flare ribbons in the AIA 1600 Å channel at 15:37:02 UT. (d) The computed flare ribbons from the MHD simulation of RUN A (red) plotted over the magnetogram (grayscale).
Paper: Stellar Obliquity from Spot Transit Mapping of Kepler-210
Authors: Adriana Valio, Alexandre Araújo
Figure caption: Short cadence light curve of Kepler-210 observed by the Kepler satellite.
Paper: Intermediate-mass Black Holes on the Run from Young Star Clusters
Authors: Elena González Prieto, Kyle Kremer, Giacomo Fragione, Miguel A. S. Martinez, Newlin C. Weatherford, Michael Zevin, Frederic A. Rasio
Figure caption: Mass of IMBHs as a function of the ejection time from the cluster. The IMBHs escaping in a binary are indicated in yellow, while those escaping as single objects are marked in blue. The different shapes indicate the dynamical interactions responsible for the ejection of the IMBHs.
Paper: Intermediate-mass Black Holes on the Run from Young Star Clusters
Authors: Elena González Prieto, Kyle Kremer, Giacomo Fragione, Miguel A. S. Martinez, Newlin C. Weatherford, Michael Zevin, Frederic A. Rasio
Figure caption: The normalized BH mass spectrum for our models listed in Table 1. The solid lines presented BHs formed through the collapse of a massive star while the dashed lines are products of binary BH mergers. The colors indicate the value for the high-mass binary fraction. The blue shaded region indicates the upper mass gap and the dashed/dotted line the beginning of the IMBH region. The shaded histogram is taken from González et al. (2021).
Paper: Intermediate-mass Black Holes on the Run from Young Star Clusters
Authors: Elena González Prieto, Kyle Kremer, Giacomo Fragione, Miguel A. S. Martinez, Newlin C. Weatherford, Michael Zevin, Frederic A. Rasio
Figure caption: Formation history for one of the IMBHs formed in our cluster models. The color indicates the stellar type and binaries are marked with a loop. The times of each collision are indicated, as well as the masses at the time of the next collision. Lastly, the core mass of the progenitor is indicated in parenthesis. Circles in gray indicate stars that are present in the dynamical interaction but do not contribute to the merger product.
Paper: Intermediate-mass Black Holes on the Run from Young Star Clusters
Authors: Elena González Prieto, Kyle Kremer, Giacomo Fragione, Miguel A. S. Martinez, Newlin C. Weatherford, Michael Zevin, Frederic A. Rasio
Figure caption: Mass ratio as a function of the IMBH mass in escaping binaries. The color gradient indicates the inspiral times. The crosses label binaries that do not merge within a Hubble time (t inspiral ≥ 12 Gyr). The inspiral times are calculated using the Peters equation (Peters 1964), taking the eccentricity and separation of the binary into account. The dashed blue and black curves indicate companions with masses larger than 40 and 120 M ⊙, respectively.
Paper: Intermediate-mass Black Holes on the Run from Young Star Clusters
Authors: Elena González Prieto, Kyle Kremer, Giacomo Fragione, Miguel A. S. Martinez, Newlin C. Weatherford, Michael Zevin, Frederic A. Rasio
Figure caption: Mass ratios of an interaction involving an IMBH–BH binary and an intruder. The horizontal axis shows the mass ratio between the IMBH companion and the IMBH while the vertical axis shows the mass ratio with the intruder. The color indicates if the intruder is a single object (magenta) or a binary (blue). The dashed line indicates an intruder with the same mass as the IMBH companion for cases where the intruder is a single object.
Paper: Intermediate-mass Black Holes on the Run from Young Star Clusters
Authors: Elena González Prieto, Kyle Kremer, Giacomo Fragione, Miguel A. S. Martinez, Newlin C. Weatherford, Michael Zevin, Frederic A. Rasio
Figure caption: Escape velocities [initial speed+kick velocity] of single IMBHs ejected from the clusters in km s−1. The subpopulations indicate IMBHs ejected due to strong GW recoil kicks during isolated BH mergers (blue) and dynamical mergers (magenta), as well as IMBHs receiving velocity kicks during strong dynamical interactions (yellow).
Paper: Intermediate-mass Black Holes on the Run from Young Star Clusters
Authors: Elena González Prieto, Kyle Kremer, Giacomo Fragione, Miguel A. S. Martinez, Newlin C. Weatherford, Michael Zevin, Frederic A. Rasio
Figure caption: The top panel shows the BH companion masses over time for three separate IMBHs. The middle panel shows the semimajor axes of the different binaries that the IMBH inhabits, and the bottom panel plots the eccentricity.
Paper: Intermediate-mass Black Holes on the Run from Young Star Clusters
Authors: Elena González Prieto, Kyle Kremer, Giacomo Fragione, Miguel A. S. Martinez, Newlin C. Weatherford, Michael Zevin, Frederic A. Rasio
Figure caption: In magenta we show the raw distribution of the masses of the BH companions of IMBHs. The blue outline is the companion mass distribution weighted by the lifetime of the binary.
Paper: Intermediate-mass Black Holes on the Run from Young Star Clusters
Authors: Elena González Prieto, Kyle Kremer, Giacomo Fragione, Miguel A. S. Martinez, Newlin C. Weatherford, Michael Zevin, Frederic A. Rasio
Figure caption: Example of an escaping BH binary that merges within a Hubble time.
Paper: Multiple Transitions in Vacuum Dark Energy and H 0 Tension
Authors: Hossein Moshafi, Hassan Firouzjahi, Alireza Talebian
Figure caption: One-dimensional likelihoods for Two-step models for H 0 based on CMB+BAO+SN data, comparing the effects of the sign of f 2 in Models 10 and 11. This plot is parallel to Figure 8 performed for Models 7 and 9, but now the value of f 1 is increased to f 1 = 0.25 and the values of f 2 for the dS and AdS cases are changed to f 2 = +0.15 and f 2 = −0.15, respectively.
Paper: Multiple Transitions in Vacuum Dark Energy and H 0 Tension
Authors: Hossein Moshafi, Hassan Firouzjahi, Alireza Talebian
Figure caption: Contour plots for Three-step models for H 0 vs. Ω m  based on “CMB+BAO+SN” data, comparing the effects of the signs of the fraction of dark energy in the second and third phases (f 2, f 3) with a fixed f 1 > 0. Note that the shaded area shows the measurement of H 0 done by the SH0ES team and its 1σ error (Riess et al. 2022).
Paper: Multiple Transitions in Vacuum Dark Energy and H 0 Tension
Authors: Hossein Moshafi, Hassan Firouzjahi, Alireza Talebian
Figure caption: Contour plots for Two-step models for H 0 vs. Ω m  based on CMB+BAO+SN data, comparing the effect of the sign of f 2, dS (f 2 > 0) or AdS (f 2 < 0) in Models 7 and 9. Note that the shaded area shows the measurement of H 0 done by the SH0ES team and its 1σ error (Riess et al. 2022).
Paper: Multiple Transitions in Vacuum Dark Energy and H 0 Tension
Authors: Hossein Moshafi, Hassan Firouzjahi, Alireza Talebian
Figure caption: One-dimensional likelihoods for Three-step models for H 0 based on “CMB+BAO+SN” data, comparing the effects of the signs of the fraction of dark energy in the second and third phases (f 2, f 3) with a fixed f 1 > 0.
Paper: Multiple Transitions in Vacuum Dark Energy and H 0 Tension
Authors: Hossein Moshafi, Hassan Firouzjahi, Alireza Talebian
Figure caption: One-dimensional likelihoods for Two-step models for H 0 based on CMB+BAO+SN data, comparing the effect of the sign of f 2, dS (f 2 > 0) or AdS (f 2 < 0) in Models 7 and 9.
Paper: Multiple Transitions in Vacuum Dark Energy and H 0 Tension
Authors: Hossein Moshafi, Hassan Firouzjahi, Alireza Talebian
Figure caption: Contour plots for Two-step models for H 0 vs. Ω m  based on CMB+BAO+SN data, comparing the effects of sign of f 2. The rest of the description is as in Figure 10.
Paper: MeV–GeV Gamma-Ray Emission from SNR G327.1–1.1 Discovered by the Fermi-LAT
Authors: Jordan Eagle, Daniel Castro, Tea Temim, Jean Ballet, Patrick Slane, Joseph Gelfand, Matthew Kerr, Marco Ajello
Figure caption: Best-fit broadband model fits for emission from the evolved PWN in G327.1–1.1 along with radio observations of the PWN (Ma et al. 2016), X-ray fluxes, and spectral indices (this work), new Fermi-LAT data (this work), and H.E.S.S. γ-ray emission (Acero et al. 2011). For comparison, we plot the broadband representation derived from the evolved broken power-law injection spectrum in Temim et al. (2015) in solid blue.
Paper: MeV–GeV Gamma-Ray Emission from SNR G327.1–1.1 Discovered by the Fermi-LAT
Authors: Jordan Eagle, Daniel Castro, Tea Temim, Jean Ballet, Patrick Slane, Joseph Gelfand, Matthew Kerr, Marco Ajello
Figure caption: Best-fit γ-ray spectral energy distribution between 300 MeV and 1 TeV over seven equally spaced logarithmic energy bins with 1σ statistical uncertainties (red) and the quadratic sum of systematic errors (black; see Section 3.3.4). TS values for each spectral bin are plotted as the green histogram. The best-fit power-law spectrum is fit with a photon index Γ γ  = 2.45 ± 0.13.
Paper: MeV–GeV Gamma-Ray Emission from SNR G327.1–1.1 Discovered by the Fermi-LAT
Authors: Jordan Eagle, Daniel Castro, Tea Temim, Jean Ballet, Patrick Slane, Joseph Gelfand, Matthew Kerr, Marco Ajello
Figure caption: Left: data and best-fit model for the “body” PWN component illustrated in Figure 2. Right: data and best-fit model for the “head” PWN component illustrated in Figure 2. In both panels, the nonthermal spectrum originating from the PWN is plotted in green and the thermal spectrum originating from the SNR is plotted in gray. The total X-ray spectral model is plotted in black.
Paper: MeV–GeV Gamma-Ray Emission from SNR G327.1–1.1 Discovered by the Fermi-LAT
Authors: Jordan Eagle, Daniel Castro, Tea Temim, Jean Ballet, Patrick Slane, Joseph Gelfand, Matthew Kerr, Marco Ajello
Figure caption: Left: 2 ° × 2° test statistic (TS) map of all events from 300 MeV to 2 TeV. Right: 2 ° × 2° 1–10 GeV TS map of all events. The maximum TS at the PWN/SNR position is ∼18 and ∼23 in the 300 MeV–2TeV and 1–10 GeV energy ranges, respectively. The source model used only excludes the point-source modeling emission at the location of G327.1–1.1. The 843 MHz (white) radio contours and H.E.S.S. (blue) and Fermi-LAT (black) 95% uncertainty regions are plotted in both panels.