adsabs/WIESP2022-NER huggingface dataset:
Paper: 2021ApJ...907...33Y (fulltext)
Right panels present the distributions of the position angles of the magnetic field orientations as a function of distance from the protostars. If multiple protostars are present in one core, the distances between the magnetic field orientations and the protostar closest to the Stokes I intensity peak are calculated. In the middle and right panels, black dashed lines denote the mean magnetic field orientations computed from the averaged Stokes Q and U emission in the area of the dense cores, and red solid and dashed–dotted (if there are multiple outflows) lines show the position angles of the outflows. The names of the protostars are labeled in the order from western to eastern sources above the middle panel in each row. Figure 12.16. L1157 Same as Figures 3 and 4, but left panels only present the magnetic orientations (small orange segments) detected in the 40′′ × 40′′ regions (dashed open squares) around the sample protostars as well as the mean orientations of the magnetic field (long thick orange segments) in the individual dense cores/clumps. Left panels show the Stokes I maps (gray scale) of the dense cores, where the sample protostars are embedded, and green contours delineate the area of the dense cores identified by Clumpf ind . Yellow dots show the locations of the protostars, and blue and red arrows present the orientations of their blue- and redshifted outflows. Black horizontal segments denote the spatial scale of 0.02 pc (or ∼4000 au). Middle panels present the distributions of the position angles of the individual magnetic field orientations detected with the JCMT POL-2 observations. Right panels present the distributions of the position angles of the magnetic field orientations as a function of distance from the protostars. If multiple protostars are present in one core, the distances between the magnetic field orientations and the protostar closest to the Stokes I intensity peak are calculated. In the middle and right panels, black dashed lines denote the mean magnetic field orientations computed from the averaged Stokes Q and U emission in the area of the dense cores, and red solid and dashed–dotted (if there are multiple outflows) lines show the position angles of the outflows. The names of the protostars are labeled in the order from western to eastern sources above the middle panel in each row. Appendix D Dependences of Misalignment on Source Properties Figure 13 compares the measured misalignment angles with the sizes, total fluxes, mean intensities, and mean polarization percentages of the dense cores and the distances to the sources. The total flux and mean intensity are expected to be proportional to the total mass and mean column density of a dense core, respectively. These comparisons show that the measured misalignments do not depend on the mass and density of the dense cores. We note that the magnetic field structures in the nearby sources can be better resolved than in the distant sources. The comparison also
Paper: 2020MNRAS.492.3364W (fulltext)
task of solving equation ( 69 ). The problem then assumes the form of an inhomogeneous Helmholtz equation and the solution becomes particularly simple when expanding the radial dependence in modified spherical Bessel functions that fulfil the boundary conditions. Like in the classical gravity problem, Legendre polynomials remain the representation of choice for the latitudinal dependence. These basis functions allow a very efficient semi-analytical solution to the problem. Each of the calculations presented here required only a few seconds of run time on a standard 4-core notebook. There has been a discussion whether the DSG term could be neglected when inverting high precision gravity observations of Jupiter and Saturn for zonal flow properties. Our new formulation allows us to quantify the relative impact of the DSG for each gravity harmonic, practically independent of the considered zonal flow or background state. A special case arises for degree l = 1. For the background density with polytropic index unity, the l = 1 solution comprises the case where the two buoyancy contributions in the TGWE cancel. This corresponds to the homogeneous solution of the Helmholtz equation. Zonal flow and gravity perturbation then decouple, and it becomes impossible to draw any conclusion on the zonal flows from the respective gravity contribution. Kong et al. ( 2017 ) seem to have noticed the related problems without realizing their origin. However, this is of little practical interest since the origin is generally chosen to coincide with the centre of gravity so that l = 1 contributions vanish. Table 5 compares the relative DSG impact with the precision of most recent gravity measurement for Jupiter and Saturn. The even harmonics J 2 – J 6 are not listed since they are dominated by the rotational deformation of the planet. For Jupiter’s J 3 , J 5 , and J 7 coefficients, the relative impact of DSG is comparable to the error and should thus be taken into account when inverting gravity harmonics for zonal flow properties. This agrees with the results and conclusion by Kong et al. ( 2017 ). The error in the higher order observed moments may decrease as the Juno mission progresses. For Saturn, J 3 , J 5 , and J 10 seem precise enough to warrant including DSG effects. The estimates of Kong et al. ( 2017 ) and Galanti et al. ( 2019 ) for the relative impact of the DSG are compatible with our results. Including the DSG term generally increases the amplitude of the gravity coefficients. Table 5. Relative error of gravity harmonics for Jupiter (Iess et al. 2018 ) (second column) and Saturn (Iess et al. 2019 ) (third column). The fourth column shows S l , the relative impact of the DSG for radial profile f 3 also listed in Table 4. In addition to the TWE, Cao Stevenson ( 2017 ) also solve the Euler equation for a polytrope of index unity and a geostrophic equatorial zonal jet. The respective results not only provide the perturbed density structure but also the
Paper: 2021ApJ...909..171D (fulltext)
analysis of photochemical models of Titan (Dobrijevic et al. 2008 ). It was found that the chemistry of C 2 H 2 and C 2 H 4 could be bistable although the underlying chemical pathways, involving neutral–neutral processes, differ from those studied here. The physical environment of Titan is significantly different from that of dense interstellar clouds; nevertheless, as H 2 is much less abundant on Titan, this work suggests that Titan’s ion–molecule chemistry could also lead to bistable solutions. Finally, the interstellar chemistry of carbon is much more complex than studied here, where we have only considered hydrocarbons up to C 2 . Autocatalytic processes in addition to those described above are likely present. For example, ethene, when viewed as a dimer (CH 2 ) 2 takes part in autocatalysis for CH 2 in reaction sequences similar to those for acetylene. Although present in our network, we found this played no role in producing bistable solutions. Analysis of larger, more extensive, carbon chemistry networks may lead to the discovery of more deeply embedded, and so less apparent, autocatalytic cycles. If one identifies the higher hydrocarbons as dimers then diacetylene (C 4 H 2 ) and the carbon tetramer (C 4 ) could act as respective intermediates for C 2 H and C 2 autocatalysis. The possibility for autocatalysis involving larger organic molecules remains to be explored; anion chemistry (e.g., Cordiner Charnley 2012 ; McElroy et al. 2013 ) may become as important as the cationic reactions considered here. 5. Conclusions We have found new bistable solutions in reduced chemical sets comparable to simplistic gas-phase chemical models of dense molecular clouds. We have identified several autocatalytic cycles in the ion–molecule chemistry of the major volatile elements: O, N, and C. In nitrogen chemistry there are two possible autocatalytic processes. An autocatalytic cycle driven by the reaction of He + with N 2 produces bistability. We find that this model has bistable solutions at high densities ( n H > 10 5 cm −3 ), which are present at all values of the elemental N abundance considered. The bistable solutions in the reduced model of the oxygen chemistry occur at low densities but are similarly present at all values of elemental O depletion, including none. Thus, bistable solutions are a fundamental property of interstellar chemistry; introduction of other elements, carbon in particular, acts to suppress these solutions by inhibiting autocatalysis through reactions with O 2 and OH. The carbon chemistry exhibits bistable solutions that appear around 10 2 cm −3 , or less, but only when an elemental depletion equivalent to that of the diffuse interstellar medium is adopted. We have identified three possible autocatalytic processes. We find that two autocatalyses occur in the acetylene chemistry, with CH, rather than atomic C, being the autocatalyst. These involve and C 2 H 2 in processes analogous to those found in the oxygen autocatalysis; their
Paper: 2016A&A...587A..50A (fulltext)
and carbon-normal metal-poor stars (red crosses and blue dots, respectively). Because our model is computed for , we restrict our observed sample to stars with iron abundance within a factor of five from the model value, . This selection leaves us with CEMP, CEMP-, and CEMP- stars, that is more than of our entire carbon-rich sample. The theoretical and observed abundances differ in many aspects. 1. Observed carbon-rich and carbon-normal stars appear as two distinct and clearly separated populations. We compute the linear regression of the observed europium and barium abundances weighted with the observational errors, and we determine the following equations for the ratio of barium to europium in carbon-normal and CEMP- stars, respectively, represented in Fig. 2 with blue-dashed and magenta-solid lines. If we compute the linear fit for CEMP- stars with and CEMP- stars separately we find almost identical results as in Eq. ( 2 ). Equations ( 1 ) and ( 2 ) show that the empirical relation between [Eu/Fe] and [Ba/Fe] has almost the same slope in the two classes of stars, and CEMP- stars have [Ba/Fe] on average dex higher than carbon-normal stars at the same [Eu/Fe]. The relations described by Eqs. ( 1 ) and ( 2 ) are likely consequences of the barium-to-europium ratios produced in the - and in the -process, respectively. In contrast, in our simulation there is not a clear separation between the two groups (colour distributions in Fig. 2 ). Because - and -enrichments are independent, the initial [Eu/Fe] and [Ba/Fe] follow approximately the relation described by Eq. ( 1 ), whereas the final abundances depend on the amount of material transferred from the AGB primary stars onto their secondary companions. Higher accreted masses cause stronger abundance enhancements of carbon and barium. Because AGB stars generally do not produce high europium abundances, [Eu/Fe] is essentially determined by its initial value, except for stars with initial [Eu/Fe] less than about . Consequently, for each [Eu/Fe], which is given by the initial -enrichment, we find stars with all barium abundances up to two dex higher than predicted by a pure -process, and the correlation between [Eu/Fe] and [Ba/Fe] described by Eq. ( 2 ) is completely washed out. Lugaro et al. (2009 , 2012) and Abate et al. (2015c) criticise this scenario of independent - and -process contributions with analogous arguments. Fig. 2 [Ba/Fe] vs. [Eu/Fe] in carbon-normal and carbon-rich metal-poor stars (blue dots and red crosses, respectively). Stars enriched in -elements occupy the region above the horizontal and diagonal solid black lines, corresponding to zones A and B of the plot. According to our adopted definitions, CEMP- stars are in zone B. In zone C there are stars with , , and . CEMP- stars and carbon-normal stars with and are in zone D. The dashed-blue and solid-magenta lines are linear fits of the carbon-normal stars and CEMP- stars, respectively. The red and blue distributions represent the
Paper: 2020MNRAS.493.1907M (fulltext)
Zanazzi Lai 2018 ). If the initial inclination angle of the disc is small, the result of the evolution is that the disc tends to be coplanar with the binary star. If the initial inclination angle of the disc is large (i.e. >40°; Martin Lubow 2017 ) and the eccentricity of the binary star is non-negligible, then the disc will evolve into a polar orbit. The presence of high-inclination discs indicates that high-inclination CBPs might exist. However, the above conclusions are based on the evolution of the isolated CBD systems themselves. Besides, they can only evolve to high-inclination orbits in specific initial configurations, such as a high initial inclination and a large eccentricity for the binary. Other mechanisms are still needed to explain why the CBD has a large inclination initially. It is worth noting that the semimajor axes of binaries with high-inclination discs found already are relatively large (several tens of au), while the discs around close binaries are nearly coplanar (Kennedy et al. 2012 ; Czekala et al. 2019 ). For the Kepler CBPs, the semimajor axes of their host binaries are small, about 0.1–0.2 au. One possibility is that the discs around well-spaced binaries are susceptible to the surrounding environment, such as th close encounters with fly-by stars. It is generally believed that most stars, and therefore most planetary systems, were born in clusters or associations (Clarke, Bonnell Hillenbrand 2000 ; Lada Lada 2003 ; Pfalzner 2013 ; Hao, Kouwenhoven Spurzem 2013 ; Cai et al. 2017 ). Some open clusters later dissolved, forming the current field stars. A notable example is that our own Solar system may have formed in an open cluster. The chemical composition of objects in the Solar system, along with the orbital elements of Sedna, suggest that the Sun formed in a cluster with roughly 10 3 –10 4 stars (Adams 2010 ). Other work even suggests that the Sun was born in a massive cluster (10 4 –10 5 stars) because a cluster of several thousand stars is much too small to produce a nearby massive supernova (Dukes Krumholz 2012 ). Malmberg, Davies Heggie ( 2011 ) considered the effect of a close encounter between stars on the planetary systems in a single star system. They found that the fly-by of a star is one of the possible mechanisms for explaining the generally high eccentricity of exoplanes (mainly discovered by radial velocity methods). It is interesting to ask how the cluster environment affects CBPs. In fact, the first CBP discovered, PSR B1620–26, is located in a globular cluster (Sigurdsson et al. 2003 ). According to the theory of star formation, it is impossible to form a close binary directly (Moe Kratter 2018 ). One of the explanations for this is that they are formed by interaction with other stars (Martin, Mazeh Fabrycky 2015 ; Hamers, Perets Portegies Zwart 2016 ; Moe Kratter 2018 ), which means that they were in a star-dense environment previously. A cluster provides such a possibility. If a CBP was born in an
Paper: 2016MNRAS.462L..36G (fulltext)
luminosity an excellent proxy of the star formation rate. In Arp 220, the L ′ CS(7–6) -to- L ′ CS(4–3) ratio is equal to 0.76. Using the calibration of Zhang et al. ( 2014 ), this leads to an L IR estimate of 1.6 × 10 12 L ⊙ , consistent with the L IR determined by Rangwala et al. ( 2011 ) using Herschel /Spectral and Photometric Imaging Receiver (SPIRE) data (i.e. 1.8 × 10 12 L ⊙ ). The relation linking L CS(4–3) and L IR can now be calibrated with SEPIA Band 5. H 2 O – submillimetre water line emission is considered to be an important gas cooling mechanism in warm molecular clouds and the H 2 O line (ν rest = 183.3 GHz) only requires moderate densities ( n (H 2 ) = 10 4 –10 5 cm −3 ) and low temperatures ( T ∼ 40K) to be excited (Cernicharo et al. 1994 ). We confirm the detection of H 2 O(3 3,1 –2 2,0 ) emission in Arp 220 (Fig. 1 , middle), with a peak flux of 0.14 Jy and a width at half-power of 332±31 km s −1 . These values are consistent with the peak flux (0.17 Jy) and width at half-power (310 km s −1 ) obtained with IRAM data by C06 (see section 4.1 for further analysis). CH 3 OH – extragalactic methanol maser emission has been detected towards many objects (Henkel et al. 1987 ; Sjouwerman et al. 2010 ). In particular, millimetre methanol transitions are known to be sensitive (but sometimes degenerate) tracers of temperature and density. Millimetre CH 3 OH emission has been detected in Arp 220 by Martín et al. ( 2011 ), particularly the 5–4 group (241.8 GHz). Salter et al. ( 2008 ) also report an absorption feature from the 6.7 GHz methanol line. Our SEPIA observations give us an access to the 193.5 GHz methanol (4–3) group (Fig. 1 , bottom). The peak flux is 68 mJy. The water and methanol lines have similar widths (equal to 333 km s −1 for CH 3 OH). We note that this is at the lower end of the line width range (280–590 km s −1 ) estimated from thermal molecules in Arp 220 by G09 . HC 3 N – molecules from the cyanopolyyne family like HC 3 N are standard tracers of the dense gas in Galactic star-forming regions but are also detected in extragalactic environments (Aladro et al. 2011 ; Lindberg et al. 2011 ) in many transitions (from 9–8 to 28–27). Arp 220 is one of the very few ‘HC 3 N-luminous’ galaxies (i.e. that the HC 3 N(10–9) line intensity is at least 15 per cent of the HCN(1–0) line intensity; Aalto et al. 2002 ; Martín et al. 2011 ). The HC 3 N(20–19) line is narrow and detected at 16 per cent of the HNC(2–1) luminosity. Its profile looks symmetric but its peak is shifted. This is consistent with the asymmetry observed in the 10–9 transition by Aalto et al. ( 2002 ) and could suggest that most of the emission emerges from the eastern nucleus. C 34 S – we observe a strong and broad C 34 S(4–3) line at 192.8 GHz. The line is unfortunately at the edge of the velocity range observed. An HC 3 N vibrationally excited line resides in a similar frequency range (192.6 GHz; see Fig. 1 , bottom)). Vibrationally excited emission of this
Paper: 2020MNRAS.498.3503S (fulltext)
of gravitational wave after supernova explosion. In this study, since we consider only one progenitor model and one EOS, it is difficult to say how this relation is independent of the models, i.e. equation ( 7 ) cannot be used to analyse data from future gravitational wave observations. Even so, this relation seems to be independent of the models at least in the early phase, as shown in Fig. A2 in Appendix A . Anyway, additional models should be considered in the future. Figure 6. In the left-hand panel, frequencies of the f - and g 1 -modes are shown as a function of the square root of the PNS average density, where the thick solid and the thick dotted lines are the fitting formula given by equations ( 7 ) and ( 8 ). In the right-hand panel, frequencies of the f - and g 1 -modes are shown as a function of the surface gravity of the PNS, where the thick-solid line denote the fitting formula given by equation ( 9 ). The relation similar to equation ( 7 ) has already been proposed, as a function of x in Sotani Sumiyoshi ( 2019 ): and as a function of in the unit of M ⊙ km −2 in Torres-Forné et al. ( 2019b ): although in Torres-Forné et al. ( 2019b ) the ramp up signal is identified as g 2 -mode in their classification. Equation ( 8 ) is derived for the f -mode frequency after the avoided crossing with the g 1 -mode with the PNS models provided by the 1D numerical simulations, which are eventually collapsed into black hole. In the left-hand panel of Fig. 6 , we also plot the thick dotted line given by equation ( 8 ). From this figure, we find that the f -mode frequency is well identified by equation ( 8 ) up to f ≃ 1.4 kHz, which corresponds to T pb ≲ 0.6 s from Fig. 5 (or Fig. 2 ), but it deviates for T pb ≳ 0.6 sec. This may come from that the PNS models for deriving equation ( 8 ) become more compact in the later phase due to the massive progenitor model. On the other hand, in order to see the correspondence between the eigenmodes calculated in this study and equation ( 9 ), we show the f - and g 1 -mode frequencies as a function of in the right-hand panel of Fig. 6 . This fitting formula seems to correspond to the f -mode rather than the g 1 -mode in the very early phase and also in the late phase, at least for comparing to our results. Finally, we make a comment with respect to the previous results about the gravitational wave signals, especially obtained via the general relativistic 3D simulation (Kuroda et al. 2016 ), which is done with a 15 M ⊙ progenitor model (Woosley Weaver 1995 ) and SFHx EOS (Steiner et al. 2013 ). Unlike most of the other simulation results, several modes of gravitational wave signals have been found in this simulation. The characteristic gravitational wave frequencies extracted by the time-frequency analysis (Kawahara et al. 2018 ) are shown in the left-hand panel of Fig. 7 , where the signal A corresponds to the ramp up signals. Nevertheless, considering the results with the PNS asteroseismology shown in Fig. 2
Paper: 2016A&A...591L...2V (fulltext)
HO (under laboratory conditions), which is consistent with the astronomical observation of solid CO in water-rich environments. Neill et al. ( 2014 ) performed a sensitive line survey of Sgr B2(N), and based on the HOCO species, they estimated a CO/CO ratio in the colder external envelope of between 0.010.1, implying a CO gas-phase abundance of 10–10, consistent with the finding of Minh et al. ( 1988 , 1991 ). Their lower limit to the CO gas-phase abundance of 10 relative to H indicates that the CO abundance in the gas is clearly enhanced, while most of the water is still frozen out in ices in the Sgr B2 envelope. Table 1 Properties of the observed HOCO and DOCO lines ( K). Notes. The upper limits have been computed using the rms around the line. The errors in the fluxes include the calibration and statistical uncertainties from the Gaussian fits of the lines. As part of the IRAM-30 m Large Program ASAI 1 (Lefloch et al., in prep.), we carried out a highly sensitive, unbiased spectral survey of the molecular emission of the L1544 prestellar core with high spectral resolution. This source is a prototypical starless core in the Taurus molecular cloud complex ( pc) on the verge of the gravitational collapse (Caselli et al. 2002 , and references within). It is characterised by a central high density ( cm) and a low temperature (7 K). We reported the detection of many oxygen bearing complex organic molecules produced through the release in the gas phase of methanol and ethene through non-thermal desorption processes (Vastel et al. 2014 , hereafter Paper I). The sensitivity of these IRAM observations also led to the detection of the hyperfine structure of CHCN (Vastel et al. 2015a ). In the present study we report on the detection of the HOCO ion in the L1544 prestellar core, as well as a tentative detection of its deuterated form, DOCO. 2. Observations and results The observations were performed at the IRAM-30 m towards L1544 () using the broad-band receiver EMIR at 3 mm, connected to an FTS spectrometer in its 50 kHz resolution mode. The beam of the observations is 29 and 23 at 85 and 106 GHz, respectively. Line intensities are expressed in units of main-beam brightness temperature (see Paper I for more details). The frequencies and other spectroscopic parameters of the HOCO transitions have been retrieved from the JPL database 2 from laboratory measurements by Bogey et al. ( 1988 ). Two lines of HOCO (the 4–3 and 5–4 at 85.531 and 106.913 GHz, respectively) lie in the frequency range covered by the ASAI survey for upper level energies lower than 30 K and were detected. Their spectra are shown in Fig. 1 . Table 1 reports the spectroscopic parameters as well as the properties of the two detected lines, obtained by Gaussian fitting. For the deuterated form of HOCO, DOCO, we used the JPL database with the spectroscopic parameters from Bogey et al. ( 1986 , 1988 ). Three transitions lie in the frequency range for upper level energies lower than 30 K.