adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019MNRAS.487.2824D (fulltext)
counting pairs of i and j objects with unitary weights; and the weighted case, accumulating the product of the weights w i * w j , selected randomly (in this case) from a Gaussian distribution with (σ, μ ) = (1.0,0.1). This is a mock-up of a common scenario for deriving 2PCFs from real surveys, where numerous bias and observational effects are accounted by means of properly chosen weights. From an implementation point of view, this means dealing with extra weight vectors and at least an extra floating-point multiplication deeply nested inside loops. For this reason, Gundam uses slightly different versions of the counting algorithm to accommodate these two cases. From the figure it can be seen that our code, which is based in a completely different algorithm, performs as well and even better than Corrfunc . In ( r p , π) space, it can be up to 1.4 × faster. In angular space Gundam can be slightly slower. However, if non-unitary weights are required, our code can run up to 30 per cent faster. The timings for very few particles, e.g. below few times 10 4 , should be ignored as the dominant component there is most likely the overhead of grid creation or array manipulation. Gundam , as well as Corrfunc , are intended primarily for processing much larger samples. Figure 5. Single-core run time to calculate DD pairs with Gundam and Corrfunc codes for samples up to 2 × 10 6 objects, using random Gaussian weights or a single constant weight equal to 1 (unweighted pair count case). Performance is evaluated in ( r p , π) space (left) and angular space (right). Bottom panels show the speed-up factor of Gundam respect to Corrfunc for the weighted and unweighted case. 3.2 Multithread We run a series of tests in the described set-up, running the same pair counts with 2, 3, and 4 threads that execute in each of the physical CPU cores. There is little benefit in using extra virtual threads (i.e. the Hyper-Threading technology of Intel CPUs) because the code is already computationally intensive, making extensive use of available CPU resources and floating-point pipelines. Also, since logical processors share the caches, they might actually compete for cache utilization, which can lead to a degradation of performance. Therefore, we only report up to four threads, noting that eight threads can potentially bump the performance by 25 per cent. In Fig. 6 , we can see that below 3–4 × 10 4 objects there is a bias towards slightly longer run times independent of the number of threads. As said before, this is due to the overhead of allocating memory and arranging data into convenient structures. In fact, a naive pair counting algorithm might probably run faster for very small samples. For larger samples, the overhead is less relevant and the advantage of using multiple threads in parallel starts to become significant. For around 10 6 particles in ( r p , π) space, we get a close to constant speed-up factor of 1.9, 2.8, and 3.6, when we duplicate, triplicate, or
Paper: 2020A&A...640A..51C (fulltext)
13 C ratio is larger. Fig. 7 12 CN/ 13 CN ratio as a function of n H , for different temperatures, at t 1 ( top right panel ), t 2 ( bottom left panel ), and t 3 ( bottom right panel ). The large pink circles in the bottom panels represent the result of the model that simulates the physical conditions in the centre of the pre-stellar core B1b observed by Daniel et al. (2013) . The 12 CN/ 13 CN ratio observed by Daniel et al. (2013) is represented by the green star. We note that the error bar is within the symbol. In all panels, the black horizontal dashed line represents the initial 12 C/ 13 C ratio of 68. We also performed a parameter-space exploration for a fixed n H of 2 × 10 4 cm −3 , varying the temperatures from 10 up to 50 K and the cosmic-ray ionisation rate between ζ low and ζ high . In particular, as already done in Sect. 3.2 , we focused the analysis on HNC, HCN, and CN ( Figs. 11 – 13 ). The main trend is that the 12 C/ 13 C ratios are highest for ζ high . This is related to the fact that secondary photon reactions maintain a high efficiency of low-temperature isotopic exchange reactions that involve C 3 , and the atomic 12 C/ 13 C ratio (and that of related molecules) remains higher than 68 (see also Fig. C.2 , right panel). 4 Conclusions We developed a new chemical network to study, in detail, how important isotopic exchange reactions are to the chemistry of carbon-containing species in the low temperature (10–50 K) gas in star-forming regions with n H = 2 × 10 3 –2 × 10 7 cm −3 . In particular, we suggest the occurrence of 13 C exchange involving C 3 with atomic carbon and study the possible consequences of this reaction. The main results and conclusions of this work are summarised below: 1. We found that reaction ( 10 ) is mainly important for T gas 27 K (owing to its exothermicity), and leads to 12 C/ 13 C 68 for molecules that form from atomic carbon (e.g. C 2 , C 3 , and nitrile-bearing species). This behaviour occurs in the time period between the conversion of atomic carbon into CO, andthe time when CO is almost completely frozen out onto dust grains. 2. We performed a detailed study of the 12 C/ 13 C ratios for the nitrile-bearing species HCN, HNC, and CN, as a function of density and for different temperatures. Our model can partially reproduce the carbon isotopic ratios derived by Daniel et al. (2013) and Magalhães et al. (2018) from observations towards low-mass pre-stellar cores. 3. We also used our grid of models to evaluate how far we are from the correct 12 C/ 13 C ratio, assumed to be 68 in the local ISM, to derive the 14 N/ 15 N ratios in the sample of high-mass star-forming regions analysed by Colzi et al. (2018a ). Our chemical model predicts that we are overestimating or underestimating the assumed 12 C/ 13 C ratios by factors ~0.8–1.9, even up to ~3.5, depending on the evolutionary time that we choose for the analysis. Considering these factors, the average nitrogen fractionation value of ~330 derived by
Paper: 2020ApJ...902...92R (fulltext)
this solution to be unrealistic. Figure 7. T eff –log( g ) diagram of the primary star in EVR-CB-004 (red star). The EHB/BHB evolutionary tracks for different stellar masses (bottom to top: 0.471, 0.473, 0.475, 0.480, 0.490, 0.500, and 0.510 M ⊙ ), that is, increasing hydrogen envelope mass (0.000, 0.002, 0.004, 0.009, 0.019, 0.029, and 0.039 M ⊙ , respectively), and solar metallicity according to Dorman et al. ( 1993 ) are shown with solid lines. In addition, the post-AGB tracks according to 0.524 M ⊙ are displayed with dotted lines (Bloecker 1995 ), and the pre-He WD track calculated with MESA (Paxton et al. 2011 , 2013 , 2015 , 2018 , 2019 ) is shown with dashed lines. The hot subdwarfs are confirmed binaries with WD companions taken from Kupfer et al. ( 2015 ). The BHB stars are taken from Saffer et al. ( 1997 ), Ramspeck et al. ( 2001 ), Vennes et al. ( 2007 ), Naslim et al. ( 2010 ), Copperwheat et al. ( 2011 ), Østensen et al. ( 2012 ), Geier et al. ( 2014 ), and Schneider et al. ( 2018 ). The plotted error bars include 1 σ statistical and systematic uncertainties as presented in the text (see Section 3 for details). If, instead, the progenitor filled its Roche lobe after He burning started, the envelope would get stripped and form a He core–burning hot subdwarf star that is expected to burn He for ≈100–150 Myr. Depending on its hydrogen envelope, it is considered to be an EHB star (for hydrogen envelopes ≲0.01 M ⊙ ) or a BHB star (for hydrogen envelopes of a few hundredths M ⊙ ). Once burning exhausts the He in the core, the star evolves toward hotter temperatures. As the core contracts, residual hydrogen is predicted to burn in a shell, pushing the surface to a larger radius (Østensen 2009 ) and hence lower surface gravities. This is seen in the solid tracks shown for different masses in Figure 7 . This stage of the evolution is expected to last for only ≈10–20 million yr and is commonly referred to as post-EHB or post-BHB evolution. A helpful discussion of EHB/BHB stars and their evolution can be found in Moehler ( 2001 ), Østensen ( 2009 ), and Heber ( 2016 ). Figure 7 shows the position of the primary of EVR-CB-004 in the T eff –log( g ) diagram, which is consistent with a post-BHB sequence with a mass of ≈0.5 M ⊙ . We note here the limitation of using the EHB/BHB evolutionary models in their current form to describe the primary of EVR-CB-004, since their radii exceed the Roche radius of EVR-CB-004 between leaving the BHB and the present. We suggest more involved modeling, beyond the scope of this work, to explore the post-BHB space specific to EVR-CB-004. If the post-BHB interpretation is correct, the primary of EVR-CB-004 is even rarer, as we would have to have caught the object during this transitioning state. The only other reasonably similar system (compact binary with a WD companion, ellipsoidal deformation, Doppler boosting, gravitational limb darkening, similar mass, and an old evolved primary) we found in the
Paper: 2019ApJ...871...11Q (acknowledgments)
A.C.N.Q. and P.G. are grateful for support from the National Science Foundation grant AST1412648 and NASA/STScI grants GO12055, GO-14268, and GO-14769. The research of L.C. is supported by the Comité Mixto ESO-Chile and the DGI of Universidad de Antofagasta. We would also like to thank the anonymous referee for the helpful comments that added to the depth of this paper. We appreciate the significance of Maunakea to the indigenous Hawaiian community and acknowledge that we are fortunate to be able to use observations taken from its summit.
Paper: 2018MNRAS.481.5350S (fulltext)
fraction, χ i , of the gas that inflows from the i + 1th annular ring to the i th ring during the simulation time-step of 1 Myr. If we assume that the entire interstellar gas fraction within an annular ring participates in the radial inflow, the percentage area of the annular rings involved in the radial mass inflow across a ring during 1 Myr with the chosen radial inflow velocities of 0, 0, 0.5, 1, 2, 3, and 3.7 km s −1 would be 0, 0, 0.02, 0.05, 0.09, 0.14, and 0.17 per cent, respectively, from the first to the seventh ring. This corresponds to a substantial inflow of gas causing a strong depletion of gas as well as the total surface matter density in the outer annular rings at the cost of matter accumulation in the inner rings. In order to circumvent this excessive inflow of gas towards the centre, we assumed 0.1 mass fraction of the total available gas fraction to be actively involved in the inflow. Thus, we used the values of 0, 0, 0.002, 0.005, 0.009, 0.014, and 0.017 per cent for χ i , from the first to the seventh ring. The assumed radial profiles of the halo surface mass density in Models I and J are presented in Table 1 . The radial stellar formation efficiency parameter, , in these simulations is assumed to be unity over the entire Galaxy. As discussed later in Section 4 , the radial gas inflow provides a natural mechanism to obtain the observed elemental abundance gradients without adopting a preferred set of stellar efficiency parameters. Finally, the choice of the simulation parameters made in this work is certainly not restricted (Mott et al. 2013 ). We have also simulated a Model J in order to incorporate the radial gas inflow superimposed over the radial gas mixing. The mixing was performed in a manner identical to that of Model F with a mixing velocity of 1 km s −1 . We have not specifically performed the simulation with radial gas outflows that are mainly caused by the energetic winds from supernovae and other massive stars. The radial gas mixing models take care of this aspect to an extent. 3 RESULTS The major aim of this work was to develop a Monte Carlo based GCE model for the Milky Way Galaxy. We ran a large number of simulations to optimize the various simulation parameters based upon the observational constraints. A selective set of simulation results for our three best models A, H, and J are presented in Figs 1 – 10 and Table 1 . We have avoided a detailed discussion and presentation of the simulation results for Models B–G and I in the main text. The results of these models can be found in figs 11–31 in the supplementary data . However, we summarize the major findings of these models in the main text. We present a comprehensive account of the deduced evolution of the , the [Fe/H], the metallicity , and the stellar birth rate for the various annular rings of 2 kpc width each, starting from 2 to 18 kpc from the Galactic centre. The deduced temporal evolution of the total surface mass density (stars +
Paper: 2015A&A...575A.118O (acknowledgments)
The authors would like to thank the anonymous referee for constructive criticism, as well as Shane P. O’Sullivan for providing his unpublished rotation measure synthesis data, and S. Ann Mao for providing her newest data set in digital form. The results in this publication have been derived using the NIFTy 4 package ( Selig et al. 2013 ), as well as the HEALPix 5 package ( Górski et al. 2005 ). This research has made use of NASA’s Astrophysics Data System. It was supported by the DFG Forschergruppe 1254 “Magnetisation of Interstellar and Intergalactic Media: The Prospects of Low-Frequency Radio Observations”. T.A. is supported by the Japan Society for the Promotion of Science (JSPS). B.M.G. acknowledges the support of the Australian Research Council through grant FL100100114. Some of the calculations were carried out on the Canada Foundation for Innovation funded CITA Sunnyvale cluster.
Paper: 2016ApJ...817..115C (acknowledgments)
In this research, we have used and acknowledge with gratitude, the data from the AAVSO International Database, based on observations submitted to the AAVSO by variable star observers worldwide. A part of the data in this study was obtained from the NRO Data Archive System. The authors wish to express their gratitude to S. Oizumi, H. Yamamoto, S. Tanada, H. Toujima, T. Yasuda, Y. Arao, K. Ando, K. Ueda, D. Namikawa, N. Matsumoto, M. Kuroki, Y. Kodama, and M. Suzuki for their help with data analysis. We also acknowledge the contribution of Dr. S. Deguchi who provided us with scientific advice. Thanks to the referee for constructive comments.
Paper: 2017ApJ...846..120B (fulltext)
mostly the and lines, and Mg for XMM2235. We note that the S/N significantly decreases in the blue part of the spectrum because of the lower throughput of KMOS in the bluer part of the YJ band, which affects the uncertainties of our measurements. Moreover, the blue part of the spectrum we fit covers a shorter wavelength range than the red part. This could potentially affect the set of templates applied by the fitting procedure. To circumvent this issue, we adopt the same combination of templates derived in the fit of the full spectrum in both the fit of the blue and red part. As discussed in Section 3.3 , before performing the kinematic fit, we smooth the KMOS spectrum with a variable kernel to a common FWHM to match the maximum FWHM. This should prevent us from adding systematics related to the KMOS resolution in the blue and red part of the spectrum. Figure 9 shows a comparison of the results obtained by fitting the full spectrum and the blue part (left panels), the full spectrum and the red part (central panels), and red and blue part (right panels). The actual value is determined as the mode of the distribution over the 100 bootstrap realizations and the error as the standard deviation. The actual value of the stellar velocity dispersion from our spectra is consistent within the errors with the mode of the distribution of derived from the bootstrap realizations. Figure 9. Comparison between the stellar velocity dispersion measurements in the blue and red part of the spectrum for XMM2235 (upper panel) and XMM2215 (lower panel). From left to right, the panels show the comparison between the fit in the full spectrum and the blue part, the full spectrum and the red part, and between blue and red part, respectively. For XMM2235, we find that the kinematics derived from the blue part and from the full spectrum have a median difference () and a scatter of 70.11 , while by fitting the red region of the spectrum, we find an offset () and a scatter of 56.20 from the fit of the full spectrum. The median difference between the fit of the blue and red part of the spectrum is () with a scatter of 72.54 . In XMM2235 the quality of some data in the “blue” part is particularly poor, resulting in a few objects with systematically higher than in the full sample. For XMM2215, we find that the kinematics derived from the blue part and from the full spectrum have a median difference () and a scatter of 60.80 , while by fitting the red region of the spectrum, we find an offset of () and a scatter of 60.00 from the fit of the full spectrum. The median difference between the fit of the blue and red part of the spectrum is () with a scatter of 121.41 . This test aims at assessing the systematic effect we can have from deriving stellar velocity dispersions from different absorption lines at different redshift. It did not result in a rejection of galaxies, which were mostly discarded at the stage of fitting the full spectrum through all the bootstrap
Paper: 2020A&A...637A..92G (fulltext)
the scales of r > 1600 au ( j app ∝ r 1.6 on average, see Table F.2 ), consistent with the prestellar stage ( j ∝ r 1.6 ). We discuss the properties and physical origin of these two regimes in more details in the next sections. 5.3 Conservation of angular momentum in Class 0 inner envelopes In this section, we focus on the relatively constant values of specific angular momentum observed in the inner envelopes at scales of r ≤ 1600 au in the j ( r ) profiles due to rotation motions (see Fig. 17 ). From these flat profiles, we find that the matter directly involved in the formation of the stellar embryo has a specific angular momentum ~3 orders of magnitude higher than the onein T-Tauri stars ( j ~ 2 × 10 −7 km s −1 pc, Bouvier et al. 1993 ). We discuss constant values of specific angular momentum as conservation of angular momentum to test disk formation as a possible solution to the angular momentum problem. It is difficult to constrain the time evolution of specific angular momentum for a given particle from angular momentum distributions which are snapshots of the angular momentum distribution of all particles at a given time during the collapse phase. During the collapse of a core initially in either solid-body rotation or differential rotation, particles conserve their specific angular momentum during the accretion on the stellar embryo ( Cassen Moosman 1981 ; Terebey et al. 1984 ; Goodwin et al. 2004 ). In the case of a protostellar envelope with a density profile ρ ∝ r −2 , an observed flat profile j ( r ) = constant requires, since each particle at different radii has the same specific angular momentum, an initially uniform distribution of angular momentum. This does not agree with the steep increase in specific angular momentum we observe at scales of r > 1600 au in the j ( r ) profiles. The break in the specific angular momentum profile could be due either to a faster collapse of the inner envelope caused by an initial inner density plateau ( Takahashi et al. 2016 ) or to a change of dominant mechanisms responsible for the observed velocity gradients from inner to outer scales of the envelope. In our sample, we distinguish eight sources with a relatively flat j ( r ) profile in the inner envelope ( β 0.5, see Table F.1 ): L1448-2A, L1448-NB, L1448-C, IRAS2A, SVS13-B, IRAS4B, L1527, and GF9-2. We estimate a centrifugal radius that would be obtained when the mass currently observed at ~100 au collapses and based on the mean value of specific angular momentum observed today ⟨ j 100 au ⟩ as follows: Table 4 Centrifugal radius R cent assuming angular momentum conservation. Notes. ( a ) Weighted mean of specific angular momentum in the inner envelopes (50 au r ≤ 1600 au). ( b ) Range of the object mass at 100 au, the minimum and maximum values are defined in Sect. 5.3 . ( c ) Centrifugal radii estimated from ⟨ j 100 au ⟩ and M 100 au using Eq. ( 2 ), assuming conservation of angular momentum. ( d ) Candidate disk radius determined from the
Paper: 2018MNRAS.480L..54D (fulltext)
peak value and rms noise. Both weighting methods provide comparable (12 per cent difference) integrated flux densities. We followed the same strategy to obtain the image of the line emission, by cleaning with different velocity resolutions up to 6.3 km s −1 . The H30 α recombination line was not detected in the data cube, using either Briggs or natural weighting. We were able to clean each velocity channel to 0.1 mJy beam −1 (rms), which may be taken as an upper limit to the line flux density. The recombination line flux estimated in local thermodynamic equilibrium using the emission measure from Section 3 falls below our detection limit. 3 RESULTS The continuum emission of the shell at 230 GHz is shown in Fig. 1a . This is the best resolution image calculated from our data. Fig. 1b displays the best resolution image contours superposed to the natural weighting, lower resolution map, which confirms the structures seen in the best resolution image. An even lower resolution image, aimed to match the Hubble Space Telescope ( HST ) point spread function at H α, simulates the expected emission of an optically thin recombination line ( Fig. 1c ). This image, derived directly from the high resolution image ( Fig. 1a ) using a Gaussian kernel convolution, is remarkably similar to the ring/bipolar shell morphology proposed by Harvey et al. ( 2018 ) for this nova. Most, if not all the large-scale structures seen in Fig. 1c is due to the blending of a distribution of smaller unresolved sources. Large contiguous emission regions that could be identified as ring or bipolar caps, are absent above the 2.5σ level in ALMA high resolution image, although they may be present at fainter flux densities. The total measured shell flux density is 9.4 ± 1.5 mJy. Figure 1. ALMA image of the V5668 Sgr shell at 230 GHz. The top panel (a) shows the full resolution continuum image above 3.0σ. Major tick marks are spaced by 100 mas. The middle panel (b) shows the natural weighted image with five superposed contour levels from the image in panel (a), starting at 2.5σ. The bottom panel (c) displays the image given in panel (a) after a Gaussian convolution to match the HST spatial resolution at H α. The flux unit is Jy beam −1 for panels (a) and (b), and arbitrary units for panel (c). North is at the top and east to the left. The power spectrum of the continuum image ( Fig. 2 ) shows that a significant fraction of the emission is formed in spatially resolved structures. It also indicates that features above the noise level are present at scales as small as 55 mas or 1 × 10 15 cm, i.e. near the resolution limit of the observations. A distance of 1540 pc is assumed in this work (Banerjee et al. 2016 ), a value that is in agreement with both the observed shell expansion parallax for V exp = 650 km s −1 and maximum magnitude and rate of decay relations. No parallax to this object could be found in Gaia DR2. Measurement of individual clump sizes in the image are consistent with