adsabs/WIESP2022-NER huggingface dataset:
Paper: 2021A&A...647A..85D (fulltext)
outlines the global expansion of the FR. Later on, this prominence eruption was accompanied by a CME observed by the Large Angle and Spectrometric Coronagraph (LASCO, Brueckner et al. 1995 ) onboard the SOHO satellite and it is given in the SOHO/LASCO CME catalog 1 ( Yashiro et al. 2004 ). The CME first appears in the LASCO C2 field-of-view at ≈15:45 UT. The speed of the CME leading edge is 452 km s −1 so it is less than twice its velocity in the low corona (Fig. 3 b). At the distances observed by C2, the CME was already decelerating with a mean value of deceleration of 12 m s −2 . 2.2. Photospheric magnetic field We analyze the AR magnetic configuration with the HMI magnetic field data. Since the studied eruption was at the limb on 02 March 2015, the photospheric magnetic field is not available. Therefore, we follow its associated AR evolution when it was on the solar disk. AR NOAA 12290 was on the eastern limb on 18 February 2015. One day later, it started to show on the solar disk as a small bipolar AR with a negative leading and a positive following polarity. During its disk crossing, it progressively grew in size as its magnetic field got dispersed by convective motions. Then, AR 12290 was well in its decay phase when it reached the western limb. Figure 4 shows four magnetograms of the longitudinal field component with a time cadence of one day. The AR was located at ≈15° of longitude after the central meridian on 25 February 2015 in panel a, and at ≈55° of longitude on 28 February 2015 in panel d. The magnetograms are spatially de-rotated to the central meridian so that the spatial extension of the polarities could be compared between the different panels. However, we kept the longitudinal component so projection effects appeared as the AR approached the limb (mostly on the limb side of the leading polarity where fake positive polarities appear). The magnetic field mostly disperses with time in each polarity so that the polarities become more extended. Then, the global bipolar magnetic configuration is progressively larger in size, which is expected to induce a global expansion of the coronal loops. This dispersion also implies that the polarities of opposite signs become into contact at the PIL. This induces cancellations of the magnetic flux, for example, at the locations surrounded by a green ellipse in Figs. 4 a and c. These cancellations of flux are best seen in the movie associated with Fig. 4 . Such cancellations imply a progressive buildup of an FR above the PIL ( van Ballegooijen Martens 1989 ; Amari et al. 2003 ; Green et al. 2018 ). Finally, the FR becomes unstable leading to an eruption (Sect. 2.1 ). Fig. 4. Evolution of the AR magnetic field from 25 to 28 February 2015. The locations of the flux cancellation are shown by green ellipses. The location of the H α filament on 28 February 2015 is displayed in panel d with a red contour. A movie of magnetic field during 25−28 is available online . The above analysis shows that
Paper: 2019MNRAS.483.5223B (acknowledgments)
ACKNOWLEDGEMENTS This research was conducted by the Australian Research Council Centre of Excellence for All Sky Astrophysics in 3 Dimensions (ASTRO 3D), through project number CE170100013. The authors wish to thank E. Zackrisson for useful discussion in the early stages of this work, D. Nicholls for assistance in constructing photoionization models using Modeling And Prediction in PhotoIonized Nebulae and Gasdynamical Shocks (MAPPINGS) V, and the anonymous referee for their comments that have helped to improve the focus, clarity, and relevance of this work. JC acknowledges research support from Australian Research Council Future Fellowship grant FT130101219. TG is grateful to the LABEX Lyon Institute of Origins (ANR-10-LABX-0066) of the Université de Lyon for its financial support within the programme ‘Investissements d’Avenir’ (ANR-11-IDEX-0007) of the French government operated by the National Research Agency (ANR). (Some of) The data presented herein were obtained at the W. M. Keck Observatory (programs 2015A_W012LA, 2015A_W033M, and 2016A_W034LA), which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W. M. Keck Foundation. The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Maunakea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain.
Paper: 2016ApJ...820...83S (acknowledgments)
We thank Zara Scoville for proof reading the manuscript and Sue Madden for suggesting use of the SPIRE fluxes for calibration. This paper makes use of the following ALMA data: ADS/JAO.ALMA# 2013.1.00034.S. This paper makes use of the following ALMA data: ADS/JAO.ALMA#2013.1.00111.S. ALMA is a partnership of ESO (representing its member states), NSF (USA), and NINS (Japan), together with NRC (Canada), NSC and ASIAA (Taiwan), and KASI (Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO, and NAOJ. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. R.J.I. acknowledges support from ERC in the form of the Advanced Investigator Programme, 321302, COSMICISM.
Paper: 2019A&A...625A..10N (acknowledgments)
This work was supported by the Young Investigators Program of the Helmholtz Association, by the Deutsche Forschungsgemeinschaft (DFG) within the Collaborative Research Center SFB 876 “Providing Information by Resource-Constrained Analysis”, project C3 and by the European Commision through the ASTERICS Horizon2020 project (id 653477). We would like to thank the H.E.S.S., MAGIC, VERITAS, and FACT collaborations for releasing the data that were used. This work made use of astropy ( Astropy Collaboration 2013 ) and sherpa ( Freeman et al. 2001 ). The authors are indebted to Abelardo Moralejo and Hans Dembinski for their useful suggestions on the statistical and systematic uncertainty estimation. We are grateful to the anonymous referee for improving the paper with helpful comments.
Paper: 2016ApJ...821..112K (acknowledgments)
The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica. This paper makes use of the following ALMA data: ADS/JAO.ALMA#2011.0.00172.S. ALMA is a partnership of ESO (representing its member states), NSF (USA), and NINS (Japan), together with NRC (Canada) and NSC and ASIAA (Taiwan), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO, and NAOJ. K.S. was supported by the grant NSC 102-2119-M-001-001-MY3. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities Inc. We thank the anonymous referee for thorough and constructive comments. Facilities : SMA , ALMA .
Paper: 2018A&A...613A..24K (acknowledgments)
Acknowledgements We thank G. Ambrazevičiūtė, V. Dobrovolskas, and D. Prakapavičius for their input during the various stages of this project. We also thank V. Vansevičius for the useful comments and discussions. This research was funded by a grant (MIP-089/2015) from the Research Council of Lithuania. HGL and AK acknowledge financial support by the Sonderforschungsbereich SFB 881 “The Milky Way System” (subproject A4) of the German Research Foundation (DFG). HGL and MS acknowledge funding from the Research Council of Lithuania for the research visits to Vilnius.
Paper: 2018ApJ...867L..31C (fulltext)
at +1542d (Stritzinger et al. 2012 ), and SN 2010jl at +461 day and +573 day (Fransson et al. 2014 ). Other coronal lines detected in SN 2017ens are similar to those seen in SN 2010jl (Figure 3 , bottom panel): [Fe x ] λ 6374.5 is strong, as are [Fe xi ] λ 7891.8, [Ne v ] λλ 3345.8, 3425.9, [Ca v ] λ 6086.8, and [Ar x ] λ 5533.2. The presence of these lines is indicative of a highly ionized and dense CSM, although we do not detect the highest-ionization coronal lines such as [Fe xiv ] λ 5302.9 and [Ar xiv ] λ 4412.3, which were seen in SN 2005ip. The flux ratio of the [O iii ] λ 4363 to λ 5007 lines is a function of the CSM density and temperature. Following Fransson et al. ( 2014 , their Figure 26), we use our measured flux ratio, log( λ 4363/ λ 5007) = −0.22, to constrain the CSM electron density to lie between 10 6 and 10 8 cm −3 for T e = 50,000 to 10,000 K. This density range is consistent with that observed for SN 2010jl. From our mid-resolution X-Shooter data, we resolve narrow P-Cygni profiles on top of the ∼2000 km s −1 wide Balmer and Paschen lines. We measure the blueshifted wavelength from the absorption component of the H γ , H β , and H α P-Cygni profiles, which suggests that the unshocked CSM has a low velocity of ∼50 km s −1 . A similar velocity of ∼60 km s −1 is obtained from the P-Cygni profile of the He i λ 10,830 line. Moreover, we measure the FWHM intensity of the wide components, such as H α (2500 ± 700 km s −1 ), H β (2300 ± 400 km s −1 ), Pa γ (2000 ± 200 km s −1 ), and He i λ 10,830 (2200 ± 200 km s −1 ). We also detect narrow absorption lines from the Balmer series (no clear emission), spanning H to H33 (3659 Å). In addition, we see emission from the H ii region close to the host-galaxy center (see Figure 4 , marked B1), as part of a faint galaxy (SDSS J120409.47–015552.4) with g = 21.92 ± 0.24 mag ( M g ≈ −16.5 mag). These lines have a slightly different redshift ( z = 0.1084) than SN 2017ens. In particular, the (noisy) detection of the weak auroral [O iii ] λ 4363 line indicates a low host-galaxy metallicity of ∼0.04–0.4 Z ⊙ using the direct T e -based method. If we instead use the empirical N2 metallicity diagnostic (Pettini Pagel 2004 ), we measure Z = 0.3 ± 0.2 Z ⊙ 4. Bolometric Light Curve and Model Fitting Using all of our available UV-through-NIR photometry, we built a pseudobolometric light curve for SN 2017ens using the prescription from Inserra et al. ( 2018b ). The results are very similar to those derived when using a blackbody fit, as expected as our photometry covers a large wavelength range. From a polynomial fit to the bolometric data we obtain L bol = (5.86 ± 0.20) × 10 43 erg s −1 at peak and an integrated energy of (3.53 ± 1.42) × 10 50 erg. To fit our bolometric light curve, we used a two-component model consisting of a central heating and an interaction component. First, the centrally heated component uses the standard Arnett method (Arnett 1982 ; Inserra et al. 2013 ). We tested three possible
Paper: 2020MNRAS.498..517F (fulltext)
but relatively insensitive to temperature. From fig. 1 of Keenan et al. ( 1995 ), the observed λ1483.3/λ1486.5 ratio corresponds to an electron density of (2.0–3.6) × 10 5 cm −3 in the temperature range (1–2) × 10 4 K. We conclude that the CSM density should be safely lower than the critical densities of the semiforbidden lines. 3.2.4 Line profiles Fig. 7 compares the velocity profile of several UV and optical lines at ∼3100 d. The Hα line extends to ∼2600 km s −1 on the blue side and only ∼1000 km s −1 on the red. Hβ, He i λ5876, and Mg ii λλ2796, 2803 have similar line profiles. While the blue side can be well fitted with an electron-scattering wing, the line centroid is shifted to the blue, with a ‘shoulder’ at ∼−800 km s −1 . As shown by Taddia et al. ( 2020 ) and Dessart et al. ( 2015 ), this type of profile can be explained by a combination of emission from the shock wave and emission from the pre-shock CSM. The emission from the shock is responsible for the blueshifted shoulder, while the central component is coming from pre-ionized gas in the CSM. Both components are affected by electron scattering in the low-velocity CSM. This gives rise to the smooth, blue wing shortward of the ‘shoulder.’ Because of the smoothing by the electron scattering, the velocity of the ‘shoulder’ is lower than the shock velocity. Figure 7. Line profiles of Hα (black), Hβ (green), He i λ5876 (red), Lyα (magenta), and N iv ] λλ1483.3, 1486.5 (blue). Note the narrower line profile of the N iv ] lines, as well as the contribution of the semiforbidden λ1483.3 component on the blue side. Both Hβ and He i λ5876 have line profiles similar to that of Hα, as expected if these arise from similar regions (Fig. 7 ). The Lyα line has a profile broadly consistent with Hα, although strongly distorted by a central absorption feature from both the Galactic and host-galaxy ISM. The C iii ], N iii ], N iv ], and N v UV lines are considerably narrower than Hα, and consistent with being unresolved at the ∼300 km s −1 instrumental resolution of STIS. These lines do not show the velocity shift of the Balmer and He i lines, and most likely come from the same optically thin, highly ionized region of the CSM as the high-ionization optical lines. 3.2.5 Spectral evolution The optical spectra in this paper span one of the longest periods of any SN. Fig. 8 highlights only a very slow evolution of the spectrum over three epochs: 736, 3099, and 5105 d. Unlike most other SNe, there is no obvious transition to a nebular spectrum even at this extremely late stage. However, the line ratios do exhibit some interesting quantitative changes that are consistent with the shock wave propagating into a medium of decreasing density. Figure 8. Evolution of the optical spectrum from day 736 to 5105 with identification of the high-ionization lines at the top and colour coding and identification of the low-ionization lines for the day 3099 spectrum. We start with a look at some with the most
Paper: 2019MNRAS.484.1838B (fulltext)
The centres of the two groups projected on the sky are not very different. Given the low density and extended nature of the foreground population, it is possible that it extends beyond our field of view. The statistical parameters obtained for each group distribution ( d N and σ N of Model 2, see Table 1 ) suggest that the two groups are spatially separate. We consider the larger population to be the main Cygnus OB2 association (Group 2) and consider the foreground population to be a separate group approximately ∼400 pc in the foreground (Group 1). The distance of the foreground group of ∼1350 pc puts it at a similar distance to Cygnus-X as a whole (see Rygl et al. 2012 ) suggesting that the main part of Cygnus OB2 is actually behind Cygnus-X by several hundred parsecs (though the line-of-sight depth of Cygnus-X is not well constrained). Consequently, the main group is more distant than previously thought, and therefore its stellar content will both be more luminous (approximately 1.5 times more luminous compared to the estimates in Wright et al. 2015 ) and more massive. Interestingly, this puts the distance to the main part of Cygnus OB2 closer to that originally derived by Massey Thompson ( 1991 ). 5.2 The foreground group We have identified 19 stars in the foreground group ( of the sample), seven of them classified as O-type stars. The bright BD + 40 4212 double system ( G = 9.39 mag) is included in this group, as well as the star HD 195213 ( G = 8.38 mag). This group includes approximately 10 of the total population of O-type stars in Cygnus OB2, and thus its total mass can be estimated as a similar fraction of the total mass of 16 500 M ⊙ estimated by Wright et al. ( 2015 ), i.e. 1650 M ⊙ , similar to that of the Orion Nebula Cluster (although according to our results, the estimation by Wright et al. 2015 will have to be corrected upwards). We note that the foreground group is appreciably more dispersed on the sky than the main group. The proper motions also suggest it to be more diffuse and less likely to be a bound group. This could suggest that it is a part of older foreground population that extends further outside our field of view. However, a detailed study of the physical properties of its stellar content is needed to establish the most probable scenario. 5.3 Potential contaminants Here, we discuss the sources identified as probable foreground or background contaminants (Group 3) and not part of either the main Cyg OB2 population or the foreground group. Foreground contaminants : HD 196305 is a very luminous star and has a parallax that places it at a distance of pc, in agreement with previous studies that suggest it to be a foreground contaminant (Chentsov et al. 2013 ). CCDMJ20323 + 4152AB has been reported as a visual double star by Gili Bonneau ( 2001 ) and therefore its binary nature could be affecting the parallax. MT91-426 and MT-170 also appear as foreground sources, despite the fact that Wright et al. ( 2015 ) proposed them
Paper: 2020MNRAS.493..952A (fulltext)
and a global alignment to an absolute reference frame, here defined by the Gaia DR2 catalogue (Gaia Collaboration 2018 ). As there are no telescope offsets between the paired F 814 W direct and G800L grism exposures, the alignment of the former is applied directly to the latter. Finally, we combine the F 814 W direct exposures into a rectified mosaic with astrodrizzle at a final sampling of arcsec pix −1 and generate a source catalogue and segmentation map from this mosaic using the sep software package. 9 The global contamination model for a given field is constructed in two passes. In the first, we assume that each object has a spectrum flat in units of f λ flux density with a normalization set by its observed flux in the F 814 W direct image (integrated over the object’s segmentation polygon). In the second pass, we step through objects in the catalogue sorted by increasing magnitude and compute a third-order polynomial fit to each spectrum after subtracting the contamination model from any neighbours (i.e. either the flat spectra of fainter sources or the polynomial models of brighter sources). This refinement is iterated three times. In both passes, the spatial template of the dispersed 2D model spectra is taken from the observed F 814 W cut-out within an object’s segmentation polygon. Note that this approach neglects morphological variation across the G800L bandpass; e.g. the PSF, continuum colour gradients, and morphological differences between line- and continuum-emitting regions. However, it provides a spatial model with much higher fidelity than any simple parametrized representation (e.g. Gaussian or Sérsic approximations). The above process automatically yields a source and contamination model for each galaxy detected in the reference image. These can be used to ‘decontaminate’ source spectra – i.e. remove overlapping light from neighbouring sources – to provide more accurate line and continuum characterizations. This was critical in GLASS’ crowded cluster G102 and G141 pointings, but the ACS parallels are sparse enough that source contamination is often negligible (Section 3 ). Note that no spectra are modelled or extracted for objects not detected in the reference frame. As such, the catalogue presented here is truncated at a F 814 W flux of m 814 = 26 with various subsamples defined by brighter flux limits (Section 3 ). There are two main differences between grizli and the previous GLASS pipeline. First, grizli automatically handles HST grism spectra taken at one or more telescope roll angles that translate into different spectral dispersion PAs. This is implemented in that the 2D spectral models are computed in the pixel space of each separate grism exposure. This is the space in which the field-dependent grism dispersion configuration is measured and defined (e.g. Kümmel et al. 2009 ). This approach allows for evaluating model fits in the space of the original detector pixels with a well-calibrated noise model. Recent works
Paper: 2018ApJ...856L...9A (fulltext)
) assigned both peaks to transpolyacetylene at grain boundaries and surfaces of nanodiamonds. The weak peak at ∼1020 cm −1 along with the broad feature centered at ∼1270 cm −1 are also associated with nanodiamonds (Veres et al. 2007 ; May et al. 2008 ), and the broad doublet at ∼1800 and 1900 cm −1 may be due to overtones and/or combinations. In the 2000–4000 cm −1 region (not shown), the Raman spectra show broad CH-stretching peaks centered at ∼2900 cm −1 in addition to the second-order peaks of nanodiamonds and graphitic/amorphous carbon. Figure 2. Reflected-light image from the KBr pellet of the Kapoeta carbonaceous-chondrite material showing a cluster of diamond particles. Figure 3. Raman spectra of the Kapoeta nanodiamonds (see the text for details). Thus, in the cluster of particles shown in Figure 2 , each particle may be considered as an aggregate of nanodiamonds. To estimate the sizes of these nanodiamonds from their Raman spectra, we consider the particle-size dependence of the position and FWHM of the 1333 cm −1 diamond peak (Yoshikawa et al. 1993 ). Based on the analyses of 26 Raman spectra, the average diamond peak position and FWHM are 1332.6 cm −1 ( σ = 0.3 cm −1 ) and 4.0 cm −1 ( σ = 0.6 cm −1 ), respectively, suggesting a particle size of ≥50 nm for the Kapoeta diamonds that show the Raman diamond peak in their spectra. However, some spectra do not show the Raman diamond peak at 1333 cm −1 , but they display all the other peaks shown in Figure 3 (a), including the peaks and features associated with nanodiamonds at ∼1020, 1150, 1450 cm −1 , and ∼1270 cm −1 (Figure 3 (b)). The spectrum shown in Figure 3 (b) is very similar to the Raman spectra of ultrananocrystalline diamonds having an average particle size of ∼3–5 nm, i.e., single-digit nanodiamonds, where the 1333 cm −1 diamond peak is absent and they show the D- and G-bands along with the 1150 and 1450 cm −1 peaks (Popov et al. 2004 ). One difference is that the G-band in their spectra occur at a lower wavenumber (1560 cm −1 ), possibly due to differences in sp 2 /sp 3 ratios in the amorphous carbon phase. Thus, our Raman spectroscopic data suggest the existence of two populations of nanodiamonds in the Kapoeta CC material: nanodiamonds that are ≥50 nm and single-digit nanodiamonds. 3.3. Micro-FTIR Spectroscopy Figure 4 (a) shows a typical micro-FTIR spectrum taken from the diamond-enriched region (spot size is ∼70 μ m). It is dramatically different from the KBr-pellet spectrum (Figure 1 ) both in the high- and low-frequency sides. The position of the absorption bands and their possible assignments are listed in Table 1 . Figure 4. (a) Micro-FTIR spectrum collected from the diamond-enriched region in the KBr pellet of the Kapoeta carbonaceous-chondrite material (see the text for details). (b) Expanded view of the C–H stretch region. Table 1 Assignment of IR Bands in FTIR Spectrum of the Diamond-enriched Region of the Kapoeta CC Material 3.3.1. C–H Stretching Bands The C–H
Paper: 2016ApJ...830...81F (acknowledgments)
We thank the anonymous referee for a critical reading of the paper and valuable suggestions that helped improve the quality and clarity of this work. We also thank Charles Dermer, Tyce DeYoung, Markus Bötcher, and Bin Zhang for useful discussions.