adsabs/FOCAL huggingface dataset:
This n = 1, CT = 216 case corresponds to a new asteroseismic-calibrated version of the original TS dynamo. The parameter CT is introduced here to account for uncertainties on the adopted timescale for the damping of the azimuthal field (while the α parameter in Fuller et al. 2019 was introduced for the saturated value of ωA, hence the relation CT = α3). We thus find that the damping timescale adopted for the azimuthal field in the original TS dynamo has to be increased by a factor of about 200 to correctly reproduce the asteroseismic data of evolved stars. Ideally, one would expect a value of CT closer to unity for a well-defined physical process. We however recall here that this timescale is known to be quite approximated as mentioned by Spruit (2002) who indicated that the original estimates of the dynamo process were made by neglecting all multiplying factors of order unity and that these factors could sometimes compound to rather large numbers. A first uncertainty is related to the exact value of the growth rate of the Tayler instability which seems to be somewhat smaller than the adopted value of ωA (e.g., Goldstein et al. 2019). In the same way, the correction to this growth rate in the case of fast rotation (i.e., Ω ≫ ωA) introduces another uncertainty that also seems to overestimate its value (Ibáñez-Mejía & Braithwaite 2015). Another source of uncertainty is related to the fact that these estimates of the growth rate are only based on the fastest growing non-axisymmetric m = 1 Tayler mode, while the small-scale dynamo process could perfectly be dominated by modes with different values of m (see e.g., Ibáñez-Mejía & Braithwaite 2015). In this context, the damping timescale adopted in the original TS dynamo seems to correspond more to a minimal than an exact value; it is thus interesting to be able to constrain its value from asteroseismic data and it is then not surprising to deduce a longer timescale in this way. Of course, it is difficult to speculate whether such a large increase of about two orders of magnitude is really physically motivated in the framework of the original TS dynamo and we can only await numerical simulations performed under more realistic stellar conditions to obtain some answers on this point.
Citation: Spruit (2002)
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Among these latter complex molecules, formamide (HCONH2) is of crucial importance. Chemical reactions of molecules containing H, C, N, and O such as formamide are considered a plausible pathway for synthesis of biomolecules under prebiotic conditions (Oparin 1938). Formamide is the simplest molecule containing the peptide bond, which is known to be the basis for assembling proteins and polypeptides starting from amino acids, with a crucial role in the biotic processes of life on Earth. Formamide was first detected in the gaseous phase in two high-mass star forming regions: Orion-KL, an active star forming region, and the giant molecular cloud SgrB2 (e.g., Turner 1991; Nummelin et al. 1998; Halfen et al. 2011); later it was detected in the comet Hale Bopp (Bockelee-Morvan et al. 2000), whose chemical composition is suspected to be similar to the chemical composition of the primitive Solar Nebula. Furthermore, during landing of Philae aboard Rosetta mission, in situ mass spectrometer data inferred the presence of formamide in the comet nucleus with the highest abundance after water (Goesmann et al. 2015). In recent years, formamide was also observed in two types of low-mass star forming environment: shocked regions by protostellar jets (e.g., Codella et al. 2017) and hot corinos (Kahane et al. 2013; López-Sepulcre et al. 2015; Marcelino et al. 2018; Imai et al. 2016; Oya et al. 2017; Lee et al. 2017). In high-temperature regions (> 100 K) such as hot corinos, thermal desorption is responsible for sublimation of frozen mantles into the gas phase. Moreover, outbursting young stars like V883 Ori are good new targets for looking for organic complex molecules that thermally desorb from icy mantles. The interpretation of observations can benefit from information coming from the laboratory, where it is possible to simulate the thermal desorption process and UV irradiation of formamide under simulated space conditions. Here we report two laboratory analyses. We focus both on photostability and on thermal desorption of pure formamide, and in the presence of grains, before and after UV irradiation. Section 2 presents our choice of minerals used as analog samples and the experimental setup. In Sect. 3, in situ UV irradiation of pure formamide ice and its adsorption by minerals at 63 K are investigated by infrared spectroscopy (FTIR). Section 4 presents a temperature-programmed desorption (TPD) analysis of pure formamide ice and in the presence of TiO2 dust, before and after UV irradiation.
Citation: Halfen et al. 2011
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Figure 1a shows for each cluster a visualization of the resulting background-subtracted color-magnitude diagram, produced with an approach similar to that of van der Burg et al. (2016). Specifically, we subtracted the residual background contamination as follows: 1) We start from a candidate cluster member sample obtained as discussed in Sect. 3.1, and a control field sample from the GOODS-S field (see Sect. 2.3) that is selected in the same manner. 2) For each galaxy in the candidate member sample, we calculate a “weight” that corresponds to the statistical excess of the candidate member sample over the control field density at the magnitude and colors of the given galaxy. Weights are calculated as follows (see van der Burg et al. 2016, for a more detailed description): first, all candidate member weights are initially set to 1. Then, for each galaxy in the control field sample we subtract the corresponding “background contamination” from the candidate member sample by appropriately reducing the weights of all candidate members that lie within a distance in the color (m814–m140) – color (m140–[3.6]) – magnitude (m140) space given by their photometric uncertainties (1σ), with a minimum distance of 0.3 mag, effectively resulting in a smoothing of galaxy densities in the color-color-magnitude space. If no galaxies are found within this distance, we double the search distance, and then if necessary increase the search distance to 1.3 times the distance to the closest galaxy. This criterion allows the full subtraction of the background contamination estimated from all galaxies in the control field sample, while reducing the weights of those candidate members that are more similar to the galaxies in the control field. For each considered control-field galaxy, the weights of all selected candidate members identified with the above criterion are reduced so that the contribution of the considered field galaxy (normalized by the areas of the probed cluster region and of the control field) is removed. At the end of this procedure, the contributions of all galaxies from the field sample have been subtracted from the candidate member sample.
Citation: van der Burg et al. (2016)
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Citation: van der Burg et al. 2016
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The young starburst inferred by the detections of high ionization emission line of He II λ4686 and the blue WR bump in this and previous works (Guseva et al. 2000; Brinchmann, Kunth & Durret 2008) is confirmed by the age estimates made here for the bright and faint regions in Mrk 22 as ∼4 and ∼10 Myr, respectively. Unlike previous works, we carried out abundance analysis for both the regions separately. We found an appreciable metallicity difference of ∼0.5 dex between the bright region [12 + log (O/H) ∼ 8] and the faint region [12 + log(O/H) ∼ 7.5]. The separation between two regions is ∼0.6 kpc. Typical metallicity gradients in normal spiral galaxies have been found between −0.009 and −0.231 dex kpc−1, with an average gradient of −0.06 dex kpc−1 (Zaritsky, Kennicutt & Huchra 1994). The observed metallicity difference between the two regions in Mrk 22 is too large to be explained as a normal galactic metallicity gradient. The chemical composition as measured from the gas-phase metallicity [12 + log(O/H)] shows various degree of spatial variations in different types of dwarf galaxies. For instance, shallow gradient in metallicity is seen in SBS 0335−052 (Papaderos et al. 2006) while no significant variations were seen in Mrk 35 (Cairós et al. 2007). A study on a large sample indicates that normal BCD galaxies are chemically homogeneous (Kobulnicky & Skillman 1996; Papaderos et al. 2006; Kehrig et al. 2008; Cairós et al. 2009; Pérez-Montero & Contini 2009; Pérez-Montero et al. 2011; Hägele et al. 2011; García-Benito & Pérez-Montero 2012; Lagos & Papaderos 2013). On the other hand, the metallicity of extremely metal-poor galaxies is usually not homogeneous within the galaxy, with the low metallicity seen in regions of intense star formation (Papaderos et al. 2006; Izotov & Thuan 2009; Levesque et al. 2011; Sánchez Almeida et al. 2013, 2014, 2015). However, large metallicity gradients are not common in dwarf galaxies. The simplest explanation for large metallicity difference in a single system is a recent merger of two galaxies with different metallicity. In a few cases, significantly large metallicity differences between star-forming regions in dwarf galaxies were seen and understood in terms of recent tidal interactions or mergers (López-Sánchez, Esteban & Rodríguez 2004a,b; López-Sánchez, Esteban & García-Rojas 2006; López-Sánchez & Esteban 2009, 2010). The evolution in terms of metallicities in interacting dwarf galaxies is fairly complex as it can depend on various factors such as mixing of metals with the interstellar medium (ISM), possible outflows of metals, and inflow of metal-poor gas in tidally interacting systems.
Citation: Papaderos et al. 2006
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Citation: Papaderos et al. 2006
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In the past, as well as in the more recent literature, there have been many attempts to explain the different chemical evolutionary paths of different MW components, particular those of the thin and thick discs. The outcome of these studies is that the observed different chemical evolutionary paths are related to differences in the main physical processes that drive galaxy evolution, the most significant of which are the gas accretion time-scale and the star formation efficiency and, possibly, radial migration (Larson 1972; Lynden-Bell 1975; Pagel & Edmunds 1981; Matteucci & Greggio 1986; Matteucci & Brocato 1990; Ferrini et al. 1994; Prantzos & Aubert 1995; Chiappini et al. 1997, 2001; Portinari & Chiosi 1999; Bekki & Tsujimoto 2011; Micali et al. 2013; Sahijpal 2014; Snaith et al. 2014; Grisoni et al. 2017; Grand et al. 2018; Spitoni et al. 2021). A good agreement between observations and theoretical predictions for the Galaxy is obtained by models that are based on the assumption that the disc formed via the infalling of gas (Chiosi 1980; Matteucci & Francois 1989; Chiappini et al. 1997). The formation of the different components is associated with distinct sequential main episodes of gas accretion (infall phases) that, at first, rapidly accumulates in the central regions and then, more slowly, in the more external ones, according to the so-called ‘inside-out scenario’ (Chiappini et al. 2001). In particular, the three-infall model, devised by Micali et al. (2013), is capable of reproducing the abundance patterns of the MW halo, thick and thin disc at once. In this model, the halo forms in a first gas infall episode of short timescale (0.2 Gyr) and mild star formation efficiency, ν = 2 Gyr−1, lasting for about 0.4 Gyr. It is immediately followed by the thick disc formation, characterised by a somewhat longer infall timescale (1.2 Gyr), a longer duration (about 2 Gyr) and a higher star formation efficiency (ν = 10 Gyr−1). Finally, star formation continues in the thin disc with a longer infall timescale (6 Gyr in the solar vicinity) and is still continuing to this day, with a star-formation efficiency of ν = 1 Gyr−1. The [O/Fe] vs. [Fe/H] path is thus continuous across the regions populated by halo, thick, and thin disc stars. While Micali et al. (2013) described the chemical enrichment as continuous across the three different infall stages, Grisoni et al. (2017) used also an alternative scheme where the thin and thick disc components evolve separately, in a parallel approach (see also Chiappini 2009). In such a parallel approach, the disc populations are assumed to form in parallel but to proceed at different rates. The gas infall exponentially decreases with a timescale that is 0.1 Gyr and 7 Gyr, for the thick and thin disc, respectively. This alternative approach better reproduces the presence of the metal-rich α-enhanced stars in the [Mg/Fe] vs. [Fe/H] diagram obtained with the recent AMBRE data (Mikolaitis et al. 2017).
Citation: Grand et al. 2018
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Extensive studies have attempted to probe the near-infrared EBL through different methods. Galaxy counts measures emission from resolved galaxies and the observations are extrapolated to estimate contributions from faint galaxies below the detection limit (Keenan et al. 2010; Domínguez et al. 2011; Helgason et al. 2012; Driver et al. 2016; Saldana-Lopez et al. 2021; Koushan et al. 2021). This sets a lower bound to the integrated galaxy light (IGL) component of EBL, as EBL from diffuse emission or sources not associated with galaxies is not included. Direct measurements using absolute photometry capture all emission in the EBL (Bernstein 2007; Levenson et al. 2007; Tsumura et al. 2013; Matsumoto et al. 2015; Sano et al. 2015; Matsuura et al. 2017; Zemcov et al. 2017; Lauer et al. 2021; Sano et al. 2020), however, absolute photometry is challenging since systematic errors and foreground contamination have to be tightly controlled. Fluctuation analysis is an alternative approach that is less susceptible to foreground contamination (e.g., zodiacal light, Galactic cirrus), since foregrounds and the EBL signal have distinct spatial and spectral correlation features (Kashlinsky et al. 2005; Thompson et al. 2007; Matsumoto et al. 2011; Cooray et al. 2012; Kashlinsky et al. 2012; Zemcov et al. 2014; Mitchell-Wynne et al. 2015; Seo et al. 2015; Kim et al. 2019; Matsumoto & Tsumura 2019); nevertheless inferring the absolute intensity of the EBL from fluctuation measurements depends on model assumptions. Another limitation of absolute photometry and fluctuation measurements is the inferred EBL redshift resolution. As the measured intensity is a projection of all emission along the line of sight, one way to infer the redshift dependency of the EBL is to measure the opacity of gamma-ray photons from individual blazars. Near-infrared photons along the line of sight will interact with gamma-ray photons by pair production, and thus the redshift composition of EBL can be constrained by the absorption features in blazar spectra (e.g., Aharonian et al. 2006, 2007; MAGIC Collaboration et al. 2008; Abdo et al. 2010; Ackermann et al. 2012; H. E. S. S. Collaboration et al. 2017; Ackermann et al. 2018; Abeysekara et al. 2019; Acciari et al. 2019; Abdalla et al. 2020). This method enables a redshift tomography of the EBL by observing blazars from different distances. However, these estimates have low spectral resolution and depend on the assumption of intrinsic blazar spectra.
Citation: Matsumoto et al. 2015
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We have used 3.0–75.0 keV NuSTAR data to probe any cyclotron line feature. To describe the continuum of 4U 1700–37, we have applied the NPEX model [cons*TBpcf*(powerlaw*npex+gaus+gaus)], following the previous work of Jaisawal & Naik (2015a). The NPEX model has been created by adding two cutoffpl models with their cutoff energies tied to each other and keeping the photon index of one to be frozen at –2.0. For the best fit, the χ2/d.o.f is found to be 577.64/475. The fit shows some residuals in the overall spectrum. We added a Gaussian absorption model around 39 keV (following the previous work of Jaisawal & Naik 2015a), but the best fit gives the line energy as $15.44^{+0.56}_{-0.53}$ keV with χ2/d.o.f = 487.32/472. The width and the depth of the line are found to be $5.47^{+0.90}_{-0.78}$ keV and $1.29^{+0.51}_{-0.39}$, respectively. The chance probability of the line has been computed using the ftest task in xspec. The F-test with this absorption line gives an F value  = 29.2 and a chance probability of 2.61 × 10−17(Table 3). If we add another Gaussian absorption line at 38.9 keV (Energy value frozen) the best-fitting χ2/d.o.f is found to be 486.8/470. This indicates that the second absorption line is not required for the fit. If we use only one Gaussian absorption line and freeze the line energy at 38.9 keV then the width and the depth of the line are found to be $4.54_{-0.87}^{+0.90}$ keV and $2.87_{-1.05}^{+1.24}$, respectively with a χ2/d.o.f  = 549.2/473. The ftest gives a chance probability of the 38.9 keV line to be 6.4 × 10−6. So, with NPEX model we find two valid model combinations of the data. One, the presence of a Gaussian absorption line at ∼15 keV, two, the presence of a Gaussian absorption line at 38.9 keV. But, the presence of both lines together is not supported by the data. The 10.0–70.0 keV flux of the source is found to be (2.26 ± 0.01) × 10−9 erg cm-2s-1, much lower than the value (5.6 ± 0.3) × 10−9 erg cm-2s-1, previously reported from SUZAKU data (Jaisawal & Naik 2015a).
Citation: Jaisawal & Naik (2015a)
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Citation: Jaisawal & Naik 2015a
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Sudden outbursts of NH3 simultaneously with H2S detected with the ROSINA-DFMS instrument on the Rosetta S/C point to the presence of abundant ammonium hydrosulphide in or on carbonaceous grains from comet 67P/Churyumov-Gerasimenko. There seems to be a clear distinction between the nucleus ice, where H2S and NH3 exist independently and grains, where they desorb together. S2 is much more abundant on grains compared to water than in the ice of the comet, while S3 is found only in grain impacts. This higher abundance points to radiolysis in these grains, which means they must have been exposed to energetic particles over an extended time. While for operational reasons, S4 could not be measured close to the dust impacts, S4 was clearly identified in periods where the coma was very dusty (Calmonte et al. 2016). Longer sulphur chains very likely are refractory, not sublimating at temperatures reached in the instrument or on grains in the coma. While Sn can also be formed from pure H2S ice by photo processing (Cazaux et al. 2022), the fact that S3 is clearly related to dust and is not found in the normal nucleus ice, where H2S is quite abundant, indicates that S3 is a product of radiolysis of the ammonium salt. In addition, photo processing of H2S results not only in Sn, but also in H2S2 (Cazaux et al. 2022), a species not detected in the DFMS m/z 66 and m/z 65 (HS2) spectra. This exposure rules out a contemporary formation of the salt on the surface or in the interior of the comet or a formation of the salt in the mid-plane of the protoplanetary disc, while the comet accreted. A pre-stellar formation is therefore likely. The salt is semivolatile, less volatile than water and could probably have survived quite high temperatures. It seems that on these grains, acids and ammonia are all locked in salts, be it sulphur, halogens, or carbon bearing acids like HOCN. If indeed, a relatively large part of sulphur and nitrogen is therefore in a semivolatile state in these grains, then the depletion of nitrogen in comets and of sulphur in star-forming regions could probably be explained, primarily because salts escape detection unless they experience temperatures above water sublimation. With the JWST S/C in orbit, there is hopefully the possibility to detect salts, or at least several of the acids in ices, which are supposed to be part of ammonium salt, like HOCN, H2CO, and formamide while looking for ammonium salts in star-forming regions and possibly comets.
Citation: Cazaux et al. 2022
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Solar energetic particle (SEP) events observed in interplanetary space present a large variability in terms of intensity, composition, and spatial and temporal extent (e.g., Kahler et al. 1999; von Rosenvinge & Cane 2006; Gopalswamy 2012). In particular, the observed SEP intensity time profiles result from both (i) the injection history of SEPs as they are released from their acceleration sites, and (ii) the transport processes undergone by the particles as they travel through the interplanetary medium from their sources to the particle detectors on board spacecraft. The transport effects pose the main difficulty in reconstructing the processes of particle release occurring at the Sun from in situ spacecraft observations. In order to correctly extract interplanetary transport effects, an effective treatment of the SEP transport processes in the heliosphere is required. Several approaches have been used over the years to characterize the SEP transport conditions (e.g., Reid 1964; Earl 1974, 1976; Roelof 1975, 1979, 2008; Hamilton 1977; Beeck et al. 1987; Kallenrode et al. 1992; Ruffolo 1995; Kocharov et al. 1998; Ruffolo et al. 1998; Laitinen et al. 2000, 2013; Dröge 2003; Qin et al. 2005; Wang et al. 2006; Agueda et al. 2008, and references therein). Such approaches allow us to reproduce spacecraft SEP observations and thereby infer SEP injection histories. The particle angular distributions relative to the local direction of the magnetic field (i.e., pitch-angle distributions (PADs)) have proven to be relevant in these analyses. A given intensity time profile can be fitted using a large variety of combinations of the injection profile at the Sun and interplanetary transport conditions, although this non-uniqueness can be constrained by a simultaneous analysis of the omni-directional intensity and the computed anisotropy time profiles (e.g., Schulze et al. 1977). More recently this problem has been further constrained by including the most direct form of data in the analysis, that is, the measured directional distributions together with the angular response of the detector (Agueda et al. 2008, 2009a, 2014).
Citation: Reid 1964
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