adsabs/WIESP2022-NER huggingface dataset:
Paper: 2018MNRAS.480.4754G (fulltext)
uranium, in the very metal-poor star BPS CS31082-001. This star, also called Cayrel’s star, was found to be more metal deficient than the globular clusters, and was likely born in the Galaxy during very early times. This measurement, carried out with the high-resolution UVES spectrograph of the ESO/VLT telescope, gave rise to the determination of the age of formation of U and Th in the early Galaxy, using, for the first time, the U/Th ratio. The derived uranium abundance yielded an age 12.5 ± 3 Gyr, which led to the best estimate of the age of the Galaxy and consequently provided a lower limit to the age of the Universe. However, as mentioned by Cayrel et al. ( 2001b ), the accuracy of this uranium dating is still affected by uncertainties in the measured abundance ratio, and in the calculated U/Th production ratio. The improvement of this situation depends not only on a better estimation of the production ratio U/Th, on advances in nuclear astrophysics models and on the investigation of other stars in which U and Th can be detected, but also on a comprehensive knowledge of the radiative properties for the potentially observable spectral lines, notably for the strongest U ii electric dipole transitions. To our knowledge, there have been very few available calculations of radiative parameters for singly ionized uranium so far. The main reason is that the complexity of the strongly mixed states involved in the lowest configurations, with five electrons distributed on the 5f, 6d, 7s, and 7p subshells, and the fragmentary knowledge of the experimental spectrum of this ion made theoretical computations extremely difficult. Most of the available oscillator strengths were deduced from quite old line intensities measured on arc spectra (Meggers, Corliss Schribner 1961 , Corliss Bozman 1962 , Voigt 1975 , Corliss 1976 , Kurucz 1995 ) if we except the results published by Nilsson et al. ( 2002 ) who obtained accurate gf values by combining branching fraction measurements with laboratory lifetimes determined using laser spectroscopy (Lundberg et al. 2001 ). These latter data concern however only a small number of strong U ii lines. This lack of reliable radiative parameters in singly ionized uranium justifies the work reported in this paper. An additional motivation was also found in the new analysis of the experimental spectrum recently published for this ion by Meftah et al. ( 2017 ). More precisely, we carried out detailed calculations of oscillator strengths for the strongest U ii transitions using the pseudo-relativistic Hartree-Fock (HFR) method (Cowan 1981 ) with core-polarization corrections (HFR+CPOL; Quinet et al. 1999 , 2002 , Quinet 2017 ). This allowed us to draw up a list of 38 visible spectral lines that could be used in cosmochronology. 2 AVAILABLE ATOMIC DATA IN U <sc>ii</sc> In their critical compilation related to actinides, Blaise Wyart ( 1992 ) listed some preliminary U ii energy levels, using the emission lines listed by Steinhaus et
Paper: 2021A&A...650A..46E (fulltext)
brought by these results on asymmetric jet launching models. Lovelace et al. (2010) , Dyda et al. (2015) and Lii et al. (2014) have all studied the launching of asymmetric outflows from the magnetospheres of rapidly rotating stars in the propeller regime. They find that the presence of a small dipole component leads to the formation of one-sided accretion funnel flows accompanied by episodic outflows in the opposite hemisphere. The dominant direction of ejection changes stochastically on timescales of ≃30–50 inner disk orbital periods. The inclusion of a disk magnetic field does not significantly affect these results ( Dyda et al. 2015 ; Lii et al. 2014 ). On longer timescales, the jets will therefore look essentially symmetric. Such models therefore cannot account for the velocity asymmetry sustained over typically an order of magnitude longer timescales observed in the DO Tau bipolar jet. A more promising scenario is provided by recent numerical works studying the launching of outflows from magnetised disks. If originating from an MHD disk-wind, a difference of a factor 4 in launching radius would account for a difference of a factor 2 in terminal velocities. Alternatively, different lambda values could be achieved between the two surfaces of the disk. The pioneering studies of Fendt Sheikhnezami (2013) and Dyda et al. (2015) , using a simplified α prescription for the disk viscosity, have shown that indeed asymmetric high velocity jets can develop and endure over long lasting timescales, ≃ a few hundred orbital timescales at the inner disk radius, depending on the disk magnetic properties. Bipolar jet velocity asymmetries of up to a factor 2–2.5 are observed in such models. Recent global non-ideal MHD simulations of magnetised disks, including more comprehensive disk microphysics and evolving both hemispheres, have confirmed these early results and shown that pronounced asymmetries between the lower and upper disk halves of the disk can develop and launch asymmetric outflows which persist over timescales of at least 1000 times the orbital period at the inner disk radius ( Béthune et al. 2017 ; Bai 2017 ; Gressel et al. 2020 ; Riols et al. 2020 ). However, most of these recent simulations do not currently include the very central regions of the disk ( r ≤ 1 au) from which the high-velocity jets are launched. 4.3 The origin of jet axis wiggling As seen in Sect. 4.1 , the DO Tau bipolar jet shows a quasi-sinusoidal small amplitude wiggling of its axis. Similar wiggling has been recently revealed at the base of other T Tauri jets (e.g. Dougados et al. 2000 ; Anglada et al. 2007 ; Garufi et al. 2019 ). This wiggling is often interpreted as tracing the dynamical perturbation due to an unseen companion. A companion can induce wiggling of the jet axis in two different ways. The first one is orbital motion of the jet source ( Masciadri Raga 2002 ; hereafter ‘orbital scenario’). Alternatively, a mis-aligned companion can cause retrograde precession
Paper: 2018A&A...619A..85Z (fulltext)
fact that we do not know the interstellar reddening for most of these stars, and that the assumed spectral types were only derived on the basis of their photometric indices (i.e. a kind of lower limit). However, the stars having early spectral types have mostly been studied during the last century, and it is only during the last decades with the use of large telescopes did the focus of stellar astronomers shift to the more late and fainter systems. We also have more systems with longer orbital periods (> 30 d) than appear in Bulut Demircan (2007) . Fig. 3. Distribution of omega angles with respect to the eccentricity. We can also address the issue of selection effects for our sample of stars. These “serendipitous” discoveries should be more likely detected as eccentric when the two minima are displaced from their positions 0.0 and 0.5 in phase diagram, but the binaries which have ω close to 90° or 270° should be only rarely detected. This can more easily be done when the photometry is sufficiently precise and also when the eccentricity is higher (i.e. the durations of both eclipses are significantly different). Hence, we tried to test this distribution of ω versus eccentricity in our Fig. 3 . One can see that the diagram is almost uniformly covered with data points, but the number of systems with higher eccentricity is still rather low (see Sect. 4 ). 4. Period eccentricity relation The diagram which should be discussed in more detail is the period-eccentricity distribution of our systems. We plotted our results together with those already published before by Bulut Demircan (2007) and also with the spectroscopic binaries from the SB9 catalogue by Pourbaix et al. (2004) in the Fig. 4 . As can be seen, orbital circularization plays a role in binaries with shorter periods ( Raghavan et al. 2010 , Duchêne Kraus 2013 ), while practically all systems that have periods below one day are already circular. Besides the plotted observed data we have also plotted several solid lines representing the limits of very close periastron approaches of both components (i.e. 1.5 × R ⋆ ) when these likely collide with each other. These periastron distances were computed for different spectral types (B to M) according to their typical radii and masses (see Zombeck 1990 ) with the assumption that both components are similar to each other (same masses and radii). As we can see, the sample of eclipsing apsidal motion systems has increased significantly with our new data set and should be very helpful in future when analysing the P-e diagram in detail. However, we still have to be very careful when interpreting these results with our new data. For some systems a significant change of eccentricity could appear for the case when the star still has poor data coverage of both light curve and the O–C diagram. The P-e diagram still lacks of such systems which have significant eccentricity close to its upper limit for a particular period. The system with the highest
Paper: 2018MNRAS.477.2760M (fulltext)
need of high angular resolution observations to robustly discriminate different flavours of magnetically regulated collapse models, since an hourglass at large scales can arise from different initial conditions, and outcomes of the protostellar accretion process. The B335 case could be more similar to the L1157 one if L1157 would be observed at similar physical scales (50 au), although our B335 ALMA map does not show a clear complete hourglass morphology at the envelope scale as the one observed in L1157. This is due to a combination of distance (L1157 is 2–3 times farther than B335), and filtering of the extended emission from the lower column density outer envelope, in our ALMA observations. L1527 in Taurus is another case of protostar seen almost perfectly edge-on, although it is probably more evolved than B335 (Chen et al. 1995 ; Ohashi et al. 1996 ). In L1527, the 1.3 mm polarized continuum at scales of 500–1000 au shows a magnetic field mostly aligned along the equatorial plane (Hull et al. 2014 ; Segura-Cox et al. 2015 ), while at larger scales it is roughly aligned with cavity walls (Davidson et al. 2014 ). This source has a larger disc radius, ≃60 au, and rotation is detected (Ohashi et al. 2014 ). Segura-Cox et al. ( 2015 ) use this example to suggest the topology of the B-field lines could be a method to help identify candidate Class 0 disc sources with R > 10 au because a transition from poloidal to radial field in the equatorial plane should be detected only in the sources where large rotational energy led to the formation of large discs: our observations of an equatorial component in B335, which harbours no disc at R > 10 au, challenge this simple view, and we suggest both strong pinching and strong coiling of the field lines can produce such a pattern. We also stress that inclination effects (B335 is an isolated edge-on source) and a simple environment might (i) help to enhance the detection of the strongly pinched magnetic field in the equatorial plane, and (ii) identify whether the magnetic field is regulating the evolution of the collapsing matter. For example, in more dynamic and/or clustered environments, polarization studies do not always reveal hourglass shapes, nor well-organized fields (Zhang et al. 2014 ). Among the few other cases where the magnetic field has been resolved well enough, the Serpens protostar SMM1 observed with ALMA (Hull et al. 2017 ) shows a complex magnetic field morphology, although mostly aligned with the outflow's cavity. 5 CONCLUSIONS Our ALMA 233 GHz polarized dust emission map shows a well-organized pattern at scales 50–500 au, mostly distributed along two structures: the pronounced east–west outflow cavity walls where the field is aligned with the outflow direction and reminiscent of the large-scale east–west native field, and the high-density equatorial plane where the field line segments appear to be mostly aligned with the equatorial plane. Because observations suggest there is only room for
Paper: 2017ApJ...838...37S (acknowledgments)
This research was supported by NASA/ HST Guest Observing grant #14277 (PI: JTS). B.D.O.’s contributions were supported by a NASA/ HST Archive/Theory grant #14308. Facilities : HST /COS, WIYN, APO 3.5m.
Paper: 2016ApJ...833...36F (acknowledgments)
The research has been sponsored by an Odysseus grant of the FWO Vlaanderen. The results were obtained in the KU Leuven GOA project GOA/2015-014 and by the Interuniversity Attraction Poles Programme initiated by the Belgian Science Policy Office (IAP P7/08 CHARM). Part of the simulations used the infrastructure of the VSC—Flemish Supercomputer Center, funded by the Hercules Foundation and the Flemish Government—Department EWI. We acknowledge fruitful discussion with Hugh Hudson, Yang Guo and Patrick Antolin.
Paper: 2019ApJ...885..158W (fulltext)
the text. A close examination of Figure 1 discloses that all UVBs evolve with time and most UVBs last for several minutes. In particular, we note that UVB 8, when viewed in Figure 1 (e), appears to be one larger UVB, while when it is viewed in Figure 1 (d) 68 s earlier, it appears to be two closely related but seemingly separate UVBs. UVB 7 appear to last for the shortest time: slightly more than one minute. Figure 1. Panels (a)–(f): time series of IRIS SJI 1400 A logarithmic intensity (DN/s). All panels are normalized between 0.1 and 1000 DN/s, with black being the lowest intensity and white being the highest intensity. The preceding/next three slit positions for each plot are marked with dotted lines. The contours mark the identified UVBs where the intensity is a factor of 24 above the image median. We enumerate the UVBs that have been measured by the slit. Each UVB’s number is mark on top of itself. Table 1 UVBs with Spectrogram (SG) Measurements Notes. If more than one UVBs are observed at a time on the slit, we enumerate them from the top to the bottom along the slit. a We refer to Figure 1 SJI image panels (a)–(f) to denote slit position in this table. If there is no corresponding SJI image at a SG data time, the preceding/next SJI image is referred to. The slit position is subsequently noted as the ±(1–3) slit position of that SJI image. For example, “a+2” means that the corresponding slit position is along the second dotted line to the right of the slit position in Figure 1 (a), and “b−2” means that the corresponding slit position is along the second dotted line to the left of the slit in Figure 1 (b). Note that “a+2” and “b−2” mark the same slit position, likewise “b+1” and “c−3” also mark the same slit position. b Each UVB’s position along Solar-Y is taken from the center (highest intensity) of the UVB as measured by the slit. Figure 2 is an overview of UVB 1's optically thin Si iv 1393.78/1402.77 A lines. Note that the Si iv 1394 A/1402 A ratio can be used to test whether the ion is formed under optically thin or thick conditions. A ratio close to 2, which is the case for the UVBs studied in this paper, is compatible with the lines being optically thin (e.g., Peter et al. 2014 ; Yan et al. 2015 ). However, optically thick lines can still have a ratio of about 2, so one cannot use the line ratio to conclusively state that the lines are optically thin. Panels (a) and (b) show that the UVB has remarkably intensive brightening and broadening of the Si iv lines, being orders of magnitude brighter in spectral radiance and significantly wider in wavelengths than the background Si iv emissions. To form the Si iv line, one needs a temperature of at least ∼80,000 K (Peter et al. 2014 ) or higher. The strong brightening of this line indicate strong heating of the UVB region plasmas to ∼10 5 K. Note that there is a significant absorption line (Ni ii 1393.33 A) on the left wing of the 1393.78 A line, as seen before (e.g., Young et al. 2018 ).
Paper: 2020MNRAS.491.5555L (fulltext)
9 ) and ( 10 ), and then substituting into equations( 4 ) or ( 7 ) allows us to calculate the discrete terms. We now need to calculate the continuous terms. 2.2 Calculating the transition coefficients for the continuous terms The radial wavefunction of the hydrogen for the discrete R nl and the continuum state can be found in Landau Lifshitz ( 1981 ) and Bethe Salpeter ( 1957 ), where is the confluent hypergeometric function. Then, we can derive the transition coefficients required in equations( 4 ) and ( 7 ) from equations( 13 ) and ( 14 ) using with where m j 1 and j 3 in equations( 16 ) and ( 17 ), respectively, and denotes the m th derivative of the term in square brackets. The evaluation of equation( 18 ) for m 1 can be found in the textbook (Bethe Salpeter 1957 ). Evaluating the continuous term in equation( 7 ) evidently requires evaluating equations( 16 ) and ( 17 ). The latter involve summations over j up to ( n f l 1). Since in principle n f , practical computation forces us to assess the number of summation terms required in order to reach a sufficiently accurate convergence. The transition amplitudes are given by equations( 4 ) and ( 7 ) and the overall continuum term is the sum of the Rayleigh and Raman continuum states, and , which are obtained from where as shown in equations( 4 ) and ( 7 ). Fig. 3 illustrates the rapid cross-section convergence by the fourth term in and with f 2, 3, and 4 s i.e. the green dotted curve is the one adopted. Figure 3. Illustration of the KHT absorption cross-section KHT (relative to Thomson scattering) for the first five lines in the Lyman series. The relevant KHT cross-sections are equations ( 57 ) and ( 58 ) and the Thomson cross-section is given in equation( 5 ). The purposes are (i) to illustrate that the continuous terms are of decreasing importance as more continuous terms are added, and (ii) that the same is true down the Lyman series. Hence, we are justified in simplifying (i.e. truncating) the series calculation for continuous terms, as described in Section 2 . To more clearly distinguish between the various curves, the vertical scale cuts off the line centre (shown in later figures). Five curves are illustrated in each panel, in which all discrete terms are included, but the number of continuous terms is increased from zero to four: grey solid line: no continuous terms are included. ( ) is evaluated from equation( 8 ) with . Blue solid line: the most important continuous transition (Rayleigh scattering) is included (i.e. from initial state i 1 s to all continuous states and back to the final state, f 1 s ) such that ( ) is evaluated from equation( 8 ) with and keeping . Red dashed line: the first two continuous transitions are included (from initial state i 1 s to all continuous states, back to f 1 s and 2 s ) such that ( ) is evaluated with and , where is calculated from equation( 19 ) and the remaining for all possible states above 2 s . Yellow dotteddashed line: the first three
Paper: 2019MNRAS.490.5210R (acknowledgments)
ACKNOWLEDGEMENTS We thank Michael Zevin, Christopher Berry, and Moritz Hübner for helpful comments on the manuscript. This work is supported through Australian Research Council (ARC) Future Fellowships FT150100281, FT160100112, Centre of Excellence CE170100004, and Discovery Project DP180103155. This is LIGO Document P1900228.
Paper: 2016MNRAS.462.3993B (acknowledgments)
We thank the referee for a constructive report. BJS is supported by NASA through Hubble Fellowship grant HF-51348.001 awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc. , for NASA, under contract NAS 5-26555 . TW-SH is supported by the DOE Computational Science Graduate Fellowship, grant number DE-FG02-97ER25308 . CSK is supported by NSF grants AST-1515876 and AST-1515927 . This paper used data obtained with the MODS spectrographs built with funding from NSF grant AST-9987045 and the NSF Telescope System Instrumentation Program (TSIP), with additional funds from the Ohio Board of Regents and the Ohio State University Office of Research. This paper is based on data acquired using the LBT. The LBT is an international collaboration among institutions in the United States, Italy, and Germany. LBT Corporation partners are: The University of Arizona on behalf of the Arizona university system; Istituto Nazionale di Astrofisica, Italy; LBT Beteiligungsgesellschaft, Germany, representing the Max-Planck Society, the Astrophysical Institute Potsdam, and Heidelberg University; The Ohio State University, and The Research Corporation, on behalf of The University of Notre Dame, University of Minnesota and University of Virginia.
Paper: 2019MNRAS.483..480V (fulltext)
rotational periods cannot be confidently measured because double-peaked and evolving light curves have been observed in many variable brown dwarfs to date (e.g. Apai et al. 2017 ; Vos et al. 2018 ). Longer duration follow-up observations are necessary to robustly measure the rotational periods of the variable objects detected in this survey. 5 SENSITIVITY TO VARIABILITY SIGNALS We construct a sensitivity plot for each observation to determine our sensitivity to variability signals of given amplitudes and periods. We inject simulated sinusoidal curves into random permutations of each target light curve. For targets found to be variable in the survey we divide the light curve by a polynomial fit to the light curve before injecting the simulated sinusoidal signals. The 1000 simulated sine curves have peak-to-peak amplitudes of and periods of 1.5−20 h, with randomly assigned phase shifts. Each simulated light curve is put through our periodogram analysis, which allows us to produce a sensitivity plot, showing thepercentage of recovered signals as a function of amplitude and period. Sensitivity plots for all light curves are shown in the bottom panel of Fig. 5 for variable objects and Fig. A3 (available online) for non-variables in the survey. Figure 5. Light curves, periodograms, and sensitivity plots for variable objects. Top panel: Relative photometry of target. Middle panel: Periodogram of target light curve (black) and periodograms of reference stars (grey). Blue dashed lines show the and significance thresholds. Bottom panel: Sensitivity plot showing the percentage of recovered signals for injected sinusoidal signals of various variability amplitude and periods. 6 SIGNIFICANT DETECTIONS OF VARIABILITY We detect significant variability in seven objects in the survey. Variable objects and their estimated variability amplitudes are presented in Table 2 . We present the light curves of these variable objects along with their reference stars in Fig. 4 . We show their periodograms and sensitivity plots in Fig. 5 . We discuss the first epoch variability detection of each object below. Table 2. Measured first epoch variability amplitudes for variability detections. 2MASS J00452143+1634446 – Gagné et al. ( 2014a ) classify the L2 object 2M0045+16 as a very low gravity brown dwarf, with H α emission and unusually red colours. It has been identified as a bona fide member of the Argus association (30−50 Myr), giving it an estimated mass of 14.7 ± 0.3 M Jup (Gagné et al. 2015a ), however recent work by Bell et al. ( 2015 ) has called into question the validity of the Argus association. We observed 2M0045+16 on 2014 November 11 using NTT SofI. We detect highly significant variability in this object at this epoch. The light curve of 2M0045+16 and the reference stars used for detrending are shown in Fig. 4 and the periodogram and sensitivity plots are shown in Fig. 5 . A lack of stars in the field resulted in only three references stars suitable for