adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016MNRAS.457.3988B (fulltext)
continuum of depths and durations from the many different scales of planetesimal impacts that would occur. The clustering of dips at D1500 also points to these events being correlated which is hard to reconcile with this scenario, though the planetesimals in the belts could be shepherded by planets into confined azimuthal regions (e.g. Wyatt 2003 ; Nesvorný, Vokrouhlický Morbidelli 2013 ). 5.2 Aftermath of giant impact in planetary system A possible way around the issues in Section 5.1 is to invoke dust thrown off in a single collision, perhaps analogous to the Earth–Moon system forming event (Jackson Wyatt 2012 ). In this case there need not be an underlying asteroid belt, as the collision could be between planets whose orbits recently became unstable, or between growing planetary embryos. Such events are expected to result in strong IR excesses (e.g. Jackson Wyatt 2012 ; Genda, Kobayashi Kokubo 2015 ), and are indeed seen in systems such as HD 172555 where giant impacts are the favoured explanation (Lisse et al. 2009 ). In this scenario, the putative collision would need to have occurred between the WISE observation taken in Kepler Q5 and the first large dip at D800. The dip at D1500 is then interpreted as the same material seen one orbit later, with the ∼750 d period implying an orbit at ∼1.6 au. The difference in the dip structure from D800 to D1500 could arise because the clump(s) created in the original impact are expanding and shearing out. This scenario therefore predicts that KIC 8462852 may now have a large mid-IR excess, but the most recent IR observations taken in 2015 January with Spitzer IRAC show no significant excess for KIC 8462852 (Marengo et al. 2015 ). However, non-detection of an excess would not necessarily rule this scenario out, as the dust levels derived in Section 4 (which account for the dust seen passing in front of the star) were shown to be consistent with a non-detection. A more robust prediction is that future dimming events should occur roughly every 750 d, with one in 2015 April and another in 2017 May. Two new issues arise with this scenario however. First, if the period of the orbiting material is a few years, what is the origin of the two small 0.5 per cent dips seen in the first few hundred days (D140 and D260; Table 1 ), and why did they not repeat 750 d later? It is a concern that these could require the existence of an outer planetesimal belt, which may contradict the lack of infrared emission to this star. Perhaps more problematic is the probability that this star (of unknown age) should suffer such an event that occurs within a few-year window between the WISE observation and the end of the prime Kepler mission, and that the geometry of the system is such that material orbiting at ∼1.6 au lies almost exactly between us and the star. Taking this few year window, the main-sequence lifetime, and an optimistic estimate for the scale height of giant impact debris, and the number of Kepler stars observed,
Paper: 2019MNRAS.489.5900D (fulltext)
[α/Fe] = +0.4 Dartmouth isochrone. The dashed outline is the final adopted selection window as discussed in Section 2.2 . The agreement between the observations and the theoretical isochrone is good: for 0.6 ≤ ( g − i ) 0 ≤ 0.8 where the isochrone is approximately constant at m i = 0.029 mag, the mean of the globular cluster stars is m i = 0.033 with a standard deviation σ = 0.040 mag. The observed scatter in the globular cluster m i values is large but there is no obvious indication of a systematic offset between the observations and the theoretical isochrone. The scatter is presumably driven by uncertainties in the v magnitudes as the corresponding i 0 , ( g − i ) 0 combined colour–magnitude diagram (generated using an appropriate relative distance modulus offset) shows only a small scatter in ( g − i ) 0 : σ( g − i ) 0 = 0.021 mag about the mean RGB locus. 2.1.2 Known EMP stars An alternative approach is to investigate the location of known EMP stars in the metallicity-sensitive diagram. To perform this task we have compiled from the literature a list of stars with [Fe/H] ≤ −2.5, as determined from high-dispersion spectroscopy. The list was then cross-matched with the SkyMapper DR1.1 data base. In carrying out the match-up we have used tighter limits on the SkyMapper photometry in order to reduce the effects of photometric uncertainty on the location of the stars in the metallicity-sensitive diagram. Specifically, we required e_g_psf AND e_i_psf ≤ 0.02, e_v_psf ≤ 0.03, and E( B − V ) SFD 0.1 mag, and this resulted in a sample of 294 stars. The majority of the stars are drawn from Yong et al. ( 2013 ), Jacobson et al. ( 2015 ), Barklem et al. ( 2005 ), Roederer et al. ( 2014 ), and Cohen et al. ( 2013 ); the effective temperatures, gravities, and metallicities from the latter two sources have been adjusted to ensure consistency with the other sources. The location of the known abundance stars in the metallicity-sensitive diagram is shown in Fig. 3 . We note that the most iron-poor star known, star SMSS J031300.36−670839.3, is the dark coloured five-point star symbol at (( g − i ) 0 , m i ) = (0.62, −0.09) in the figure, while the star HE 0107−5240, which has [Fe/H] ≈ −5.4, is the dark coloured five-point star symbol at (0.59, −0.135). The subgiant star HE 1327−2326 ([Fe/H] ≈ −5.7) is located at (0.29, 0.19), and is also plotted as a dark coloured five-point star symbol in Fig. 3 . In general, the majority of the stars plotted are consistent with expectations in that they lie at notably bluer m i values at a given ( g − i ) 0 than the bulk of the population. However, this is clearly not the case for a small number of known EMP stars that fall well below their ‘expected’ location in the metallicity-sensitive diagram. As will be discussed in more detail in Section 4.1 , generally these stars are strongly enhanced in carbon. Figure 3. SkyMapper metallicity-sensitive diagram for all objects satisfying the basic selection criteria. Overplotted
Paper: 2020MNRAS.491.4232S (fulltext)
Additional details of our numerical set-up are described in Appendix A . Mass-loss is a key ingredient in novae, which plays an important role in setting the time-scale of the eruption by removing mass from the white dwarf envelope (Livio et al. 1990 ; Kato Hachisu 1994 ). However, the exact mechanisms giving rise to the unbound outflows, as needed to produce variable internal shocks, are uncertain and not addressed in this work. Relatively slow outflows, of up to several hundreds of km s −1 , may be launched from near the binary L 2 Lagrange point. These could be driven for instance by frictional drag (e.g. Livio et al. 1990 ; Shankar, Livio Truran 1991 ; Lloyd, O’Brien Bode 1997 ; however, see Kato Hachisu 1991 ), non-axisymmetric gravitational torques from the binary (Pejcha, Metzger Tomida 2016a ), or in the outflow of a circumbinary disc formed from matter ejected from L 2 that initially remains gravitationally bound (Pejcha, Metzger Tomida 2016b ). These slow outflows, with velocities of order the binary escape speed, would plausibly be concentrated in the binary equatorial plane. Outflows can also occur from the white dwarf directly, with little or no influence from the binary. These could be driven by radiation pressure acting on the iron opacity bump (e.g. Kato Hachisu 1994 ), or by any large (near- or super-Eddington) localized energy deposition in the white dwarf envelope, e.g. from radioactive heating or damping of convectively excited waves (e.g. Quataert et al. 2016 ). Outflows launched from deeper locations in the potential well of the white dwarf would be expected to possess higher velocities ∼2000–5000 km s −1 and a more spherical geometry than the slower outflows from L 2 . Lacking a first principles model, we are motivated on phenomenological grounds to consider two outflow modes characterized by distinct mass-loss rates, , and asymptotic velocities, v w . Somewhat arbitrarily (as the precise values will vary between novae), we assume a ‘fast’ mode outflow with and a ‘slow’ outflow mode with These are spherically equivalent mass-loss rates. To obtain the true value, must be multiplied by the fraction of the total solid angle covered by each outflow component, f Ω . For instance, f Ω ≈ 1 for the fast spherical outflow, while f Ω ≈ 0.3 for the slow equatorially focused outflow. 1 The top panel of Fig. 2 shows the chosen time evolution of and v w for our two fiducial models (‘single transition’ and ‘multiple transition’). In the single-transition model (red lines in Fig. 2 ), the outflow velocity is increased from slow ( v w = v s ) to fast ( v w = v f ) only once at t ≈ 0, while in the multiple-transition model (blue lines in Fig. 2 ) the outflow velocity fluctuates between slow and fast modes several times. Local maxima (‘flares’) are observed in classical nova light curves on time-scales ranging from hours to weeks (e.g. Strope et al. 2010 ; Walter et al. 2012 ; Henze et al. 2018 ; Aydi et al. 2019 ). To generate
Paper: 2016A&A...586A.135P (fulltext)
of freedom is 1. The errors on the elements of the Hessian matrix are propagated to . Since Eq. ( 9 ) is non-linear, we compute the error map on from Gaussian realizations of the errors on the coefficients of the matrix. The uncertainty in , computed within bins of values of the map, is , independent of the intensity at 353 GHz (Fig. A.1 ). Fig. A.1 Two independent estimates of the uncertainties in the determination of angles from the Hessian analysis, computed within bins of values of the map. The red squares represent the noise obtained by propagation of the errors on the coefficients of the Hessian matrix. The noise computed with the toy model (Sect. A.2 ) is plotted with black squares. Fig. A.2 Distribution functions of orientation angles for the selected ridges derived from the Hessian analysis on . They are computed within three bins of Galactic latitudes with equal numbers of data points. The mean absolute latitudes are , , and for the DFs plotted in black, orange and red, respectively. The image shows that the method does not introduce any bias as a function of Galactic latitude in estimating . Appendix A.2: Test on a toy model of the sky We have tested the Hessian algorithm on a toy model of the full sky built from a realization of a Gaussian map, , with an angular power spectrum, equal to a power law of index , computed with the procedure SYNFAST of HEALPix . The spectral index used is within the range of values found for power spectra of dust maps ( Miville-Deschênes et al. 2007 ). The mean value of is 0. We form, where , , , and are four factors chosen to fit the latitude profile, and the longitude profile at Galactic latitude , measured on the map. The values of these factors, including their sign, depend on the realization of . The toy model matches the large-scale structure of the Galactic dust emission, but it does not have its filamentary structure because it is computed from an isotropic Gaussian map. It assumes that the amplitude of the brightness fluctuations at a given scale is proportional to the brightness. This is in agreement with what has been reported for the emission at 100 m of the diffuse ISM ( Miville-Deschênes et al. 2007 ). We use the toy model to quantify the impact of the structure of the diffuse Galactic emission and data noise on the Hessian analysis. We run the Hessian analysis on the map with and without data noise computed as a Gaussian realization of the noise map on . The DFs of orientation angles, computed for the model with added noise, are flat at all Galactic latitudes, as shown in Fig. A.2 . We find no bias on the angle introduced by the large-scale gradient of the emission with Galactic latitude. We use the difference between the two angle maps computed on the map, with and without noise, to estimate the contribution of the data noise to the uncertainty in the Hessian angle (black dots in Fig. A.1 ). This uncertainty decreases for increasing values of (i.e. increasing signal-to-noise ratio).
Paper: 2018ApJ...857...24W (fulltext)
the cascade emission is shown. The magnetic field in the comoving frame is then . Basically, a higher B leads to a more important synchrotron radiation, while a smaller B leads to a more important IC radiation. Since electrons usually emit higher-energy photons via IC radiation than via synchrotron radiation, the flux of the cascade emission at high energies becomes higher for a smaller B . On the other hand, more high-energy photons in turn produce more e ± pairs at high energies, making a harder electron spectrum than that in the case of a large B , as is shown in Figure 7 . This explains why the spectrum is also harder in the case of B = 0.01 than in the case of B = 1 at keV energies, where the emission is dominated by synchrotron radiation of these (relatively) low-energy electrons in both cases. One can see that the gamma-ray flux around 100 MeV is not very sensitive to the value of B . This is because the injection rate of the secondary electrons that emit 100 MeV photons is dominated by the γγ -annihilation process (i.e., in the right-hand side of Equation ( 3 )), mainly relies on the photon field, and does not change with the magnetic field. On the other hand, the jet is already an electron calorimeter even if only considering IC radiation under the benchmark parameters. Therefore, we expect the total photon production rate at a certain energy is roughly equal to the energy injection rate of the electrons that emit these photons (i.e., L γ ∼ L e ). A large B only changes the way the electrons radiate their energy away but does not increase the photon production rate. Although the energy of the gamma-ray-emitting electrons is different in the IC-dominated case with a small B from that in a synchrotron-dominated case with a large B , the resultant gamma-ray flux is at the the same level, since the energy injection rate of electrons in the relevant energy range is roughly constant with respect to electron energy (i.e., ). Figure 6. Same as Figure 3 , but for the equipartition coefficient for the magnetic field in the GRB jet B . The dashed and dotted lines correspond to the IC and synchrotron radiation, respectively. Figure 7. Quasi-steady-state spectrum of the cascaded electrons in the comoving frame of the GRB jet with fiducial parameters (except for the equipartition coefficient of the magnetic field). Purple, blue, and red curves are for B = 1, 0.1, and 0.01, respectively, while solid and dashed curves are the contribution by the BH and photomeson processes, respectively. 3. Application to LAT GRBs In this section, we apply our calculations to explain the high-energy spectrum of GRB 090926A, GRB 090902B, and GRB 080916C, which represent three different types of spectral feature. The origin of prompt keV/MeV emission is beyond the scope of this work, and we just regard it as a preexisting target photon field for the BH and photomeson processes. In our calculations below, we employ a flat spectrum for the injected proton spectrum ( s =
Paper: 2021A&A...651A..57N (fulltext)
a revised model for a He shell detonation Type Ia SN. In any case, we think it is likely that the Mn/Fe, Ni/Fe, Cu/Fe, and Zn/Fe enhancements in G 112-43/44 and the three α -poor stars are caused by element contribution from a special kind of SN to pockets of gas from which these stars were formed. Such inhomogeneous mixing of SN products is more likely to affect the chemical evolution of dwarf galaxies than higher-mass systems because in lower-mass galaxies the gas has a relatively long cooling time, resulting in episodic star-formation bursts ( Revaz et al. 2009 ; Venn et al. 2012 ). In view of this, it is interesting that G 112-43/44, as a member of the Helmi streams, was probably born in a dwarf galaxy and that one of the three α -poor halo stars with enhancements of Mn/Fe, Ni/Fe, and Zn/Fe also has Helmi stream kinematics 6 , as seen from Fig. 1 . G 4-36 and BPS BS 16920-17 have kinematics similar to stars belonging to the Gaia -Enceladus population, but BPS CS 22966-43 lies in the Helmi box with positive V Z . Not all stars belonging to the Helmi streams have unusual abundances. Roederer et al. (2010) made a detailed abundance analysis of 12 subgiant and giant stars with Helmi stream kinematics, and Gull et al. (2021) recently added seven stars to this work. Within the typical errors of the abundance determinations (0.10–0.15 dex), [Mn/Fe], [Ni/Fe], [Cu/Fe], and [Zn/Fe] agree with the [X/Fe]-[Fe/H] trends of Galactic halo stars. The stars investigated, however, have [Fe/H] ≲ −1.5 and high [ α /Fe] ratios (e.g., [Mg/Fe] ∼ 0.3), so it remains to be seen if the Helmi streams contain a population of low- α stars at higher metallicities and if some of these stars have over-abundances of Mn, Ni, Cu, and Zn relative to Fe. Low- α stars are well known in dwarf galaxies (see the review by Tolstoy et al. 2009 ). In the Sculptor dSph galaxy, for example, [Mg/Fe] shows a declining metallicity trend from [Mg/Fe] ≃ 0.4 at [Fe/H] = −1.8 to [Mg/Fe] ≃ −0.2 at [Fe/H] = −1.0. Interestingly, Skúladóttir et al. (2017) find indications of a cosmic scatter in [Zn/Fe] of the Sculptor stars (i.e., from about −0.6 dex to +0.4 dex) and a positive correlation between [Ni/Fe] and [Zn/Fe]. At a distance of ∼85 kpc, the giants observed in Sculptor are, however, faint (17.0 V 18.5), and the errors of [Zn/Fe] are of the order of ±0.3 dex. It will require more precise [Ni/Fe] and [Zn/Fe] values as well as an improved precision of [Mn/Fe] values ( North et al. 2012 ) to verify if stars with enhanced Mn/Fe, Ni/Fe, and Zn/Fe ratios are present in Sculptor. Another interesting case is the ultra-faint dwarf galaxy Horologium I, in which Nagasawa et al. (2018) found three very metal-poor ([Fe/H] ∼ −2.5) giant stars that are α -poor ([Mg/Fe] and [Ca/Fe] ≃ 0.0) and enhanced in [Mn/Fe] by about 0.4 dex relative to Galactic halo stars. However, their [Ni/Fe] ratios seem to be normal, and their Cu and Zn abundances were not determined. Again, we need more precise abundances,