adsabs/WIESP2022-NER huggingface dataset:
Paper: 2021MNRAS.505.1053W (fulltext)
reduce the in equation ( 60 ). Thus, it can achieve arbitrary precision by increasing the bit length. However, the round off error will also propagate and accumulate through the extrapolation table in Table 1 . As the order of the GBS extrapolation increases, the round off will accumulate faster and faster. Thus the additional mantissa bits becomes less and less efficient in increasing the integration accuracy. Therefore, reducing the round off error in GBS extrapolation is also important in order to achieve the same precision with less mantissa bits and less CPU time. 3.1.1 Improvement on GBS extrapolation In the BS extrapolation, higher order methods can be obtained from lower order methods by extrapolation, as described in equation ( 52 ). However, the round off error will also propagate and accumulate through the extrapolation table. If the round off error from the base integration method is ε, −ε, ε, ... in the first column, then the round off error propagation with the extensively used step sequence n i = 2 i suggested by Deuflhard ( 1983 ) will behave as in Table 2 . As shown in the table, T 8, 8 with 16th order precision magnified the round off error to become 100 times larger. Thus, we will lose at least 2 significant digits from the extrapolation. This becomes more problematic for higher order extrapolations, as required by arbitrary precision integration methods with extremely low tolerance. Table 2. Round off error propagation in the Bulirsch–Stoer extrapolation with step sequence n i = 2 i . To reduce the round off error in the extrapolation process, there are two fine-tunings that can be performed on the extrapolations. The first is a step sequence choice. The step sequences n i = 1, 2, 3, 5, 8, 12, 17, 25, 36, 51, 73, ...(Fukushima 1996 ) is better than the extensively used sequence n i = 2 i , where the extrapolation coefficient reduces the round-off error propagation. Table 3 shows the error propagation with the new step sequence. We see that it performs much better than the original sequence suggested by Deuflhard ( 1983 ). Table 3. Round off error propagation with optimal step sequence in the BS extrapolation. The second method to reduce the round off error in the extrapolation process is to extrapolate the raw increment dξ of the integration instead of the integrated result ξ. Since the GBS extrapolation is linear, it is possible to do the following Because the raw increment dξ is usually tiny compared to the integrated results, the round of error in the first column of Table 1 is ∼eps × |dξ| instead of ∼eps × |ξ|. Then, the round off error accumulated from the extrapolation table will be smaller than the error accumulated from the direct extrapolation. We note, however, that the initial Δ T i , 1 should be the raw increment evaluated from the basic integrator, not from the subtraction of T i , 1 ( H ) − T i , 1 (0). The close value subtraction will introduce additional round off errors and will significantly degrade the
Paper: 2019MNRAS.482.4253T (acknowledgments)
ACKNOWLEDGEMENTS We thank Elisabeth Krause and Fabian Schmidt for their useful discussions. This work was in part supported by Grant-in-Aid for Scientific Research from the Japan Society for the Promotion of Science (JSPS) Promotion of Science (No. 23340061, 26610058, 15H03654, and 17H01131), Ministry of Education, Culture, Sports, Science and Technology (MEXT) Grant-in-Aid for Scientific Research on Innovative Areas (No. 15H05887, 15H05893, 15K21733, and 15H05892), JSPS Program for Advancing Strategic International Networks to Accelerate the Circulation of Talented Researchers, and Japan Science and Technology Agency (JST) CREST Grant Number JPMJCR1414. Numerical computations were in part carried out on Cray XC30 at Centre for Computational Astrophysics, National Astronomical Observatory of Japan.
Paper: 2016MNRAS.461.4227H (fulltext)
decrement, the flattest BLR decrements (lowest-ratio) are similar to the narrow-line decrement; i.e. the lower envelope is a 45-deg line parallel to the reddening vector. The null hypothesis, that the minimum BLR decrement is independent of the NLR decrement (see for example, Fig. 2 , panel VI), is excluded at a probability p = 0.002. There is a lack of AGNs with NLR decrements steeper than the BLR decrement. Points to the right of the 45-deg line are consistent with measuring errors. Only one point lies more than two standard deviations to the lower right of the red line. For the majority of AGNs, the BLR decrement is steeper than the NLR decrement. There is no significant difference in the median NLR decrement for the AGNs with very steep BLR Balmer decrements compared with those with very flat BLR Balmer decrements. These results are consistent with the earlier findings of De Zotti Gaskell ( 1985 ) and Glikman et al. ( 2007 ). De Zotti Gaskell ( 1985 ) found that Balmer decrements for nearby AGNs were steeper for the broad lines than the narrow lines, and Glikman et al. ( 2007 ) found that, for the FIRST–2MASS red quasar survey, broad lines seem to be reddened more than the sum of broad and narrow lines. Our results strongly imply that: the dust that reddens the NLR also reddens the BLR, but most of the dust reddening the BLR lies between the BLR and the NLR. The dust causing the common reddening of both the NLR and BLR could be associated with the NLR or it could be dust in the host galaxy. De Zotti Gaskell ( 1985 ) studied nuclear continuum colours and both broad- and narrow-line Balmer decrements of nearby AGNs as a function of the orientation of the host galaxy. They found that for more highly inclined disc galaxies, the nuclear colours were redder and the Balmer decrements steeper. Since the axis of symmetry of an AGN is not necessarily aligned with the rotation axis of the galaxy (see for example, Tohline Osterbrock 1982 ; Valtonen 1983 ; Battye Browne 2009 ), the greater nuclear reddening of inclined galaxies must be due in part simply to dust in the plane of the galaxy. 4 THE RELATIVE STRENGTH OF [O <sc>iii</sc>] VERSUS THE BLR REDDENING As noted above, intervening dust not only causes reddening, but also extinction. If the dust is located between the NLR and the BLR, we would expect higher extinction of the BLR and AD to make [O iii ] stronger relative to the BLR and the continuum. 4.1 Selection of objects [O iii ] λ5007 is the strongest narrow line in the optical and the Balmer lines are the strongest broad lines. Therefore, a very large number of measurements are potentially available in the Sloan Digital Sky Survey (SDSS). Dong et al. ( 2008 ) have carefully measured Hα/Hβ ratios after subtraction of narrow lines and Fe ii , etc. Comparison of their ratios with the simple Gaussian fits of the SDSS pipeline reveals quite a few very discrepant measurements in the SDSS data base and that the difference between the Dong et al. (
Paper: 2018A&A...611A..49Y (acknowledgments)
Acknowledgements IRIS is a NASA Small Explorer mission developed and operated by LMSAL, with mission operations executed at NASA Ames Research Center and major contributions to downlink communications funded by the Norwegian Space Centre through a European Space Agency PRODEX contract. The Solar Dynamics Observatory (SDO) data that we used are provided courtesy of NASA/SDO and the Atmospheric Imaging Assembly and Helioseismic and Magnetic Imager science teams. The authorsthank the anonymous referee for suggestions. We thank Y. Chen, Z.J. Ning, Y. Li, and L.P. Yang for valuable discussions. The work at Peking University is supported by NSFC under contracts 41574168, 41231069, 41274172, 41474148, 41421003. LM Yan is supported by National Postdoctoral Program for Innovative Talents (grant BX201600159). J.S. H.E. is also supported by “National Youth Talent Support Program” of China. L.D. Xia is supported by NSFC under contracts 41274178 and 41474150.
Paper: 2017A&A...607A.104B (fulltext)
model. This is partially caused by the inner region of the circumbinary disk ( ), which is thinned out by the gravitational forces of the two stars (e.g., Avramenko et al. 2017 ). However, this effect only occurs if the separation between the two stellar components is on the order of . A smaller separation would lead to a behavior that is similar to circumstellar disks around a single star. If the separation is significantly larger, the dust and gas around the binary can no longer be described with a single circumbinary disk. In circumstellar disks with a single star, the inner region has the highest density and usually the highest magnetic field strength (see Fig. 1 in Sect. 5.4 ; Flock et al. 2017 ). Therefore, as mentioned in Sect. 6.3 , circumbinary disks lack a significant fraction of circularly polarized emission if the inner high magnetic field region is missing. In addition, the different directions of circulation of the toroidal magnetic field decrease the circularly polarized emission from regions with similar velocities relative to the observer if not spatially resolved. Furthermore, the estimation of the LOS magnetic field strength from circumstellar disks with a single star is dominated by the inner region characterized by high magnetic field strength. In contrast, a Zeeman observation of a circumbinary disk is able to provide information about the magnetic field farther out in the disk if the requirements for such an observation are fulfilled. However, this requires that a significant part of its innermost region be depleted by the gravitational forces of both stellar components. 7. Conclusions We studied the requirements to observe the Zeeman splitting of the CN lines in circumstellar disks and their dependence on selected disk parameters. Such observations can be used to resolve degeneracies and obtain complementary information about the magnetic field if combined with polarimetric continuum observations. However, this is only applicable, if the polarimetric continuum observations are mainly caused by elongated aligned dust grains. In particular, we evaluated whether Zeeman observations of circumstellar disks should be performed spatially resolved or unresolved. With spatially resolved Zeeman observations, the strength and structure of the magnetic field in the LOS direction can be reconstructed for most regions in circumstellar disks. Only the region close to the star cannot be traced. However, the required sensitivity (, if ) is far beyond the capabilities of recent and upcoming instruments/observatories. For spatially unresolved Zeeman observations, we obtained different results. Under the assumption of Keplerian rotation and a toroidal structure of the magnetic field, two independent average magnetic field strengths can be obtained from the circularly polarized fraction. These average values are sensitive to changes in the strength and structure of the magnetic field inside the disk. In addition, they can be used to
Paper: 2018ApJ...853..186Z (fulltext)
leads to the inclusion of a number of RGB stars that are located close to the selection boundary, while retaining the color selection in ( m F435W − m F555W ) will still lead to rejection of RC stars. The RGB sample resulting from this updated selection run includes 207 stars. We repeated all of our subsequent analysis and found a new C index value of 0.063 mag, which is almost the same as our initial result, C index = 0.062 mag. We therefore conclude that the potential loss of a small number of genuine RGB stars by our initial selection method does not significantly influence our measurement of the cluster’s RGB broadening, nor our final results. 4. Discussion Our results show that the artificial stellar population is characterized by a similar dispersion in C F336W,F435W,F814W as the observed RGB sequence of NGC 1783. This demonstrates that the photometric uncertainties fully account for the RGB’s broadening in NGC 1783. The color index used, C F336W,F435W,F814W , is sensitive to the stellar N and He abundances, which are two of the most representative elements that exhibit large abundance variations in star clusters (Gratton et al. 2004 ). The consistency of the observed RGB’s width with that defined by our artificial SSP therefore indicates that our RGB sample is best represented by single-valued N and He abundances. There are a small number of outliers in the observed RGB distribution, in particular two stars at the far left (see the top panel of Figure 6 ). These objects are unlikely statistical fluctuations, given that they deviate by >5 σ from the Gaussian peak. Alternatively, they might be field stars that have not been properly removed by our field star decontamination algorithm. If they were cluster members, they would need to have unrealistically extreme metallicities. Note that these two data points appear both in the scatter diagram in the top panel of Figure 6 and in the observed histogram in the figure’s bottom panel. We suspect that these two stars might result from stochasticity in the field’s RGB sample (see Figure 5 ), rather than from random effects associated with our field star decontamination. We explored whether these objects might be binary systems instead. If we assume that the primary components ( M 1 ) of such binary systems are genuine RGB stars, the secondary components ( M 2 ) must have lower masses. We tested the appearance of binary systems with different mass ratios, M 2 / M 1 . For very low-mass companions, the change in color is almost negligible, while for equal-mass binaries no color difference is expected, compared with the single-star case. The top panel of Figure 6 includes binary sequences for M 2 / M 1 = 0.7, 0.8, and 0.9. It is clear that the outliers at the far left cannot be binary systems. An intermediate-age cluster with an age of 1.8 Gyr and a mass of 1.8 × 10 5 M ⊙ , NGC 1783 is comparably as massive as other star clusters that show clear evidence of MSPs (e.g., Martocchia et al. 2018 ). If
Paper: 2017MNRAS.472.1392S (acknowledgments)
Acknowledgements We would like to thank John Helly for technical assistance during the course of this work and Spoorthy Raman for reading the manuscript and providing useful suggestions for improving the text. This work was supported by the Science and Technology Facilities Council [ST/L00075X/1]. This work used the DiRAC Data Centric system at Durham University, operated by the Institute for Computational Cosmology on behalf of the STFC DiRAC HPC Facility ( www.dirac.ac.uk ). This equipment was funded by BIS National E-infrastructure capital grant ST/K00042X/1, STFC capital grant ST/H008519/1, and STFC DiRAC Operations grant ST/K003267/1 and Durham University. DiRAC is part of the National E-Infrastructure.
Paper: 2017ApJ...840...15S (fulltext)
which both β and IRX have been measured (Capak et al. 2015 ), but the galaxies with from Capak et al. ( 2015 ) and the galaxy presented by Watson et al. ( 2015 ) lie outside the region spanned by the simulated galaxies. Because our simulations and theirs cover disjoint regions of parameter space in terms of redshift and mass, it is not possible to make detailed comparisons between our results and theirs. However, their results highlight the potential importance of dust destruction for accurately predicting the UV properties of simulated galaxies and provide motivation for future studies that include more complex treatments of dust production, growth, and destruction. 9. Implications for Observations Our simulation results suggest that there is not a tight relation between IRX and β that applies to all galaxies at all redshifts. The simulated galaxies tend to lie near the M99 relation, which was determined based on observations of galaxies in the local universe. However, during merger-induced starbursts, even the galaxies can depart significantly from the relation: because IRX increases significantly but β is almost unaffected, the merger-induced starbursts, such as that shown in Figure 3 , tend to lie above the M99 relation. The simulations—of both disk galaxies and mergers—deviate more significantly from the M99 relation: even when they are not undergoing starbursts, they tend to lie above the relation. These results are consistent with observational works that have demonstrated that some classes of galaxies deviate significantly from the M99 relation. As already discussed above, observed DSFGs tend to lie above the M99 relation and do not exhibit any relationship between IRX and β (e.g., Bell et al. 2002 ; Goldader et al. 2002 ; Howell et al. 2010 ; Casey et al. 2014 ). Conversely, some galaxies lie significantly below the M99 relation (i.e., have lower IRX than expected from the relation given their β ). This was demonstrated for the lensed galaxies cB58 at (Pettini et al. 2000 ) and the Cosmic Eye (Smail et al. 2007 ) at . Similarly, Reddy et al. ( 2012 ) found that young systems (age Myr) at lie below the M99 relation. Moreover, we have demonstrated that the dust composition—specifically the strength of the 2175 Å feature—can have a very dramatic effect on the UV slope. Consequently, extrapolation of the locally calibrated relation to regimes for which it is not observationally constrained may be problematic if, e.g., the dust composition of high-redshift galaxies is significantly different from that of galaxies. How the strength of the UV bump varies with galaxy properties and redshift is still poorly understood, but there are some useful observational constraints. The commonly employed Calzetti et al. ( 1994 ) attenuation law, which was defined based on observations of local-universe starburst galaxies, does not exhibit a UV bump. Based on a study of ∼10,000 star-forming galaxies, Battisti et al. ( 2016 ) derived an average attenuation
Paper: 2018ApJ...863..103S (fulltext)
configuration. 6. Formation of Molecular Gas Due to Jet–ISM Interactions? Figures 10 – 12 display that a wall structure is clearly seen toward the locus of the extension of the precession axis of the SS 433 jets, in which the column density of H i increased precipitously. Enhanced CO emissions are also found to be nicely associated with the wall structure seen in H i emission. MC G40.331−4.302 has the H i counterpart, which displays more extended arc-like structure nearly perpendicular to the precession axis of the SS 433 jets (Figure 12 ). The positional coincidence between G40.331−4.302 and the precession axis of SS 433, as well as their coincident kinematic features, strongly suggests the physical connection between them (see Sections 4 and 5 ). We obtain that the column density, N H = N (H i ) + 2 × N (H 2 ), is about 2.1 × 10 21 cm −2 for the cloud G40.331−4.302, which corresponds to A V ≈ 1.1 mag assuming N H ≈ 1.9 × 10 21 cm −2 mag −1 × A V . At the column density, the molecules are shielded from UV radiation, and self-shielding allows significant H 2 molecules to exist. The total mass of the main concentration of the cloud is ∼3 × 10 3 M ⊙ within a radius of ∼66 (Table 1 ). The molecular fraction / N H is about 0.87, suggesting that the dominant of the gas in the cloud is in molecular form. Table 1 Properties of Gas toward G39.315−1.155 and G40.331−4.302 Notes. a The area of the emission with I (CO)>5 K km s −1 for G39.315−1.155 and I (H i )>130 K km s −1 and I (CO)>2.5 K km s −1 for G40.331−4.302. b The effective radius is calculated from (Area/3.14) 0.5 . c The mean intensity of the H i (70–90 km s −1 ) and CO (73–88 km s −1 ) emission. d The column density of the atomic and molecular gas is calculated from 1.823 × 10 18 × I mean (H i ) and 2 × 10 20 × I mean (CO), respectively. See the text. e Parameter d 4.9 is the distance to the cloud in units of 4.9 kpc (see Section 4 ). f The volume density of the atomic and molecular gas is calculated from the mass and the effective radius of the emission. The virial mass exceeds the total mass of the cloud by more than one order of magnitude, indicating that the cloud is in a gravitationally unbound state. The fragmented subclouds (e.g., with a size of ∼1 pc) in MC G40.331−4.302 may dissipate on a timescale of several Myr (sound crossing time) without other support. Alternatively, the MC may be confined by an external pressure of ≳4.4 × 10 4 K cm −3 (e.g., ρ (H i ) σ 2 (H i )) when the magnetic pressure and the self-gravity are negligible. We emphasize that G40.331−4.302 is the only detected MC within the range of l = 349–451, −51 ≲ b ≲ −30, and V LSR ≳ 70 km s −1 . It indicates that such an MC is located about 370 pc below the Galactic plane, where a kinematic distance of 4.9 kpc is adopted. However, the MCs within the solar circle are concentrated in the Galactic plane with a Gaussian FWHM of ∼90–100 pc (e.g., see Section 4.5 and Figure 6 in Heyer Dame 2015 ). This suggests that
Paper: 2015ApJ...806..264S (acknowledgments)
This work was partially supported by NASA grants ATP12-0139 and ATP13-0077. We thank Alessandra Buonanno, Francesc Ferrer, Ted Jacobson, Henric Krawczynski, Tzvi Piran, Laleh Sadeghian, and Joe Silk for helpful comments and discussion.
Paper: 2016ApJ...828...94C (acknowledgments)
The research of E.C. and J.C.W. is supported by NASA grant HST-AR-13276.002-A. E.C. thanks the Enrico Fermi Institute for its support via the Enrico Fermi Fellowship.
Paper: 2018ApJ...854...34W (fulltext)
a stronger field provides more magnetic energy to damp. We fit the slopes of this “big trend” for all the curves by appropriately removing a few points far away from this “big trend” and taking the average of these slopes. Interestingly, both viscous and magnetic dissipations obey an identical scaling law:The dashed lines in Figures 5 (b) and (c) denote this scaling law. This scaling law for the relation of the two dissipations and magnetic fields reveals that there exists self-similarity for the magnetic effect on dynamical tide. Figure 5. (a) Kinetic energy, (b) viscous dissipation, and (c) magnetic dissipation vs. S 2 at various ω. Pm = 1.0. The two dashed lines denote the scaling law ∝ S 2 . 3.2. Investigation of Pm In this subsection we investigate Pm . Similar to Figure 2 , Figure 6 shows viscous dissipation, magnetic dissipation, total dissipation, and the ratio of magnetic to viscous dissipations versus tidal frequency at various Pm with S 2 fixed to be 0.1, namely in a strong field regime. The first three subfigures show that for all the tidal frequencies, a larger Pm leads to higher viscous dissipation and lower magnetic dissipation but total dissipation is almost independent of Pm . This is reasonable because Pm is the ratio of viscosity to magnetic diffusivity (see Equation ( 9 )) but it cannot influence the total diffusivity. More interestingly, Figure 6 (d) shows that magnetic dissipation wins out over viscous dissipation for and this ratio is approximately equal to Pm −1 . In the astrophysical situation Pm is very small, and it can be inferred that magnetic dissipation dominates over viscous dissipation with a moderate or stronger magnetic field. Note again that with a very weak field, e.g., S ≤ 10 −4 , viscous dissipation wins out over magnetic dissipation even at a small Pm , e.g., (Lin Ogilvie 2018 ). Figure 6. (a) Viscous dissipation, (b) magnetic dissipation, (c) total dissipation, and (d) ratio of magnetic to viscous dissipations vs. tidal frequency at various Pm . . To find the scaling laws for Pm , we plot the two dissipations versus Pm as shown in Figure 7 . The dependence on Pm looks more regular than the dependence on S in Figure 5 because the field strength influences the resonant frequencies, whereas diffusivities cannot. By taking the average of the slopes for different tidal frequencies, we obtain the scaling lawsThe dashed lines in Figures 7 (a) and (b) denote these two scaling laws. The ratio of magnetic to viscous dissipations is indeed close to Pm −1 , as shown in Figure 6 (d). Note that our parameter regime is not the real regime, e.g., E and Pm are not too small, so these scaling laws might not be realistic. Lacking knowledge of the physical properties in stars and planets, we do not know the real parameters, and even estimating the order of magnitude is not convincing. These scaling laws obtained in the moderate parameter regime at least give us some qualitative evidence that a moderate or stronger