adsabs/FOCAL huggingface dataset:
A number of sources of systematic error disproportionately affect the soft X-rays. One important example is absorption by partially ionized material around the AGN. Although the response of the absorbing gas to changes in the irradiating flux from the AGN does contribute its own time lag, this should only influence the lag-energy spectrum for Fourier frequencies lower than those of interest for our analysis (Silva, Uttley & Costantini 2016). The influence of absorption on the soft X-ray region of the spectrum (e.g. Miller, Turner & Reeves 2008; Miller et al. 2010), however, is potentially much more important and can influence the shape of the lag-energy spectrum (I19). It will therefore be prudent to select the AGN with the clearest view of the inner regions. The soft X-ray region of the reflection spectrum is also the most sensitive to modelling assumptions such as vertical disc structure (Nayakshin, Kazanas & Kallman 2000; Done & Nayakshin 2007; Różańska et al. 2011; Vincent et al. 2016). Moreover, X-ray reverberation models currently always assume that the time taken for photons to be reprocessed and reemitted in the disc atmosphere is very small compared with the light-crossing delays. This is a very good assumption for the fluorescence and scattering processes dominant for E ≳ 3 keV. However, the time taken for soft excess photons to be reprocessed will be longer, since these photons undergo enough interactions to approximately thermalize. This thermalization time is still expected to be small, but it is not yet clear if it is small enough to be neglected entirely. It is also important to note that the model we have explored here does not include a warm (kTe ∼ 0.1 keV), optically thick (τ ∼ 10–40) corona in addition to the hot corona, as is often invoked to explain the observed soft excess in AGNs (Czerny & Elvis 1987; Middleton et al. 2009; Done et al. 2012; Petrucci et al. 2018; Ursini et al. 2020). Models 1, 2, and 4 therefore effectively assume that the observed soft excess is dominated by reflection (e.g. Jiang et al. 2018; García et al. 2019). If in reality there is indeed a warm corona, its presence would make it more difficult to constrain the soft X-ray shape of the reflection spectrum (Xu et al. 2021). Since we believe our models to be most robust for E ≳ 3 keV, we ran an alternative fit ignoring E  3 keV (Model 3), and found that the 3σ error on H0 almost doubles (dH0 ≈ 60 km s−1 Mpc−1) compared to the other fits. It is therefore desirable to also include soft X-rays, but it is encouraging that constraints can even be achieved without. Ignoring soft X-rays entirely is an extreme measure. An alternative approach would be to keep the soft X-rays but trial a variety of different models, for instance with and without a warm corona. This approach would likely return uncertainties somewhere between the two extreme scenarios explored here.
Citation: Różańska et al. 2011
Citation function: Compare/Contrast
The compact blobs instead likely have a different origin. The fact that compact blobs are unresolved even at the HST resolution, that they are found at ∼1 kpc distance from the galaxy barycentre, that they have relatively small stellar masses (≲15 per cent of the underlying disc), but are actively forming stars suggests that they are star-forming regions likely originated due to disc instability and fragmentation of the galaxy disc (Bournaud et al. 2014; Mandelker et al. 2017). The in situ formation of the compact blobs is further supported by their metallicity. In fact, while the disc properties are consistent with the stellar mass–metallicity relation of z ∼ 2 star-forming galaxies (e.g. Maiolino et al. 2008; Zahid et al. 2014), compact blobs instead show metallicities inconsistent with the mass–metallicity relation (Fig. 7, bottom right panel). They have a comparable metallicity to the discs, but ∼1.5 dex lower stellar masses, so they are ∼1 dex above the mass–metallicity relation. Metallicity measurements for statistical sample of galaxies with M⋆ ∼ 108 M⊙ at z ∼ 2 are still lacking and therefore at the low-mass end we are showing an extrapolation of the mass–metallicity relation derived for galaxies with M⋆ ≳ 109 M⊙. In Fig. 7, we also show the average location of dwarf galaxies with M⋆ ∼ 108 M⊙ at z ≲ 1 (Kirby et al. 2013; Calabrò et al. 2017; Hidalgo 2017). Despite some of them may have gas-phase metallicities up to 12 + log (O/H) ∼ 8.5 (Sánchez Almeida et al. 2018), on average our compact blobs seem to be ∼0.5 dex more metal-rich than dwarf galaxies (Fig. 7). The high metallicity of compact clumps further suggests that they formed in situ, due to the gravitational collapse of pre-enriched gas in unstable regions of the galaxy disc. The young ages of the blobs reported in Fig. 7 support these conclusions. In fact, the metallicity of star-forming regions is altered in about one galactic dynamical time (≳100 Myr), increasing due to the active star formation and internal production of metals. The fact that our sample clumps with metallicity measurements have ages ≲ 50 Myr points towards the conclusion that they formed in situ from metal-rich gas, since this time-scale is too short for the gas to be self-enriched due to internal star formation (Bournaud 2016). Finally, the metallicity of our sample of clumps does not clearly correlate with their star formation rate, as indeed expected if they formed from pre-enriched material, although larger statistical samples are needed to confirm this finding.
Citation: Zahid et al. 2014
Citation function: Similarities; Differences
Single-photon VUV excitation from the N2(X1





) ground state can only populate the ungerade states. Many of these ungerade states can be coupled by spin–orbit or vibronic interactions. Hence, only the ungerade states of N2 are calculated and presented. Because of the potential diffuse nature of the excited states, their wavefunctions were calculated using a large basis set composed by the aug-cc-pVQZ quality, which is augmented by 3s and 2p diffuse Gaussian-type orbitals (GTOs; Dunning 1989; Kendall et al. 1992; Woon & Dunning 1995). In addition, the expected high density of the electronic states in this energy range should lead to the mixing of their electronic wavefunctions. Therefore, a multiconfigurational approach is adopted where the electronic computations were carried out using the Complete Active Space Self Consistent Field (CASSCF; Knowles & Werner 1985; Werner & Knowles 1985) approach followed by the internally contracted Multi Reference Configuration Interaction (MRCI; Knowles & Werner 1988; Werner & Knowles 1988) technique as implemented in MOLPRO (Werner et al. 2012). We followed the procedure established in Spelsberg & Meyer (2001), Ndome et al. (2008), and Hochlaf et al. (2010a, 2010b). Briefly, the CASSCF active space is constructed by the valence molecular orbitals (MOs) of N2 increased by one sg and one pg MOs. This allows better relaxation of the wavefunctions of the N2 electronic states whose configurations differ in their s and p orbital occupations as those of interest at present. In MRCI, all CASSCF configurations were taken into account as reference. Finally, the CASSCF wavefunctions were used to evaluate the spin–orbit matrix elements. These spin–orbit computations were performed in Cartesian coordinates, where the CASSCF wavefunctions were used as the multielectron basis for the two-step spin–orbit coupling calculation (Llusar et al. 1996; Zeng et al. 2011) through the Breit–Pauli Hamiltonian (Berning et al. 2000) as implemented in MOLPRO.
Citation: Kendall et al. 1992
Citation function: Uses
The aforementioned form of power-law primordial power spectrum is a prediction of inflation where the scalar field (inflaton) slowly rolls down to the bottom of the flat inflationary potential. With the constraints on the tilt and an upper bound on the amplitude of tensor perturbation with respect to scalar perturbation, various surveys have ruled out a wide class of models. However, fundamental questions such as the energy scale of inflation and the detailed shape of the potential remain unanswered. Note that any changes in the nearly flat potential will eventually lead to certain features in the spectrum. Local glitches in the potential including rapid change of its amplitude, or the break in its first or second derivatives (Starobinsky 1992; Starobinsky 1998; Adams, Cresswell & Easther 2001; Covi et al. 2006; Joy, Sahni & Starobinsky 2008; Joy et al. 2009; Hazra et al. 2010; Miranda, Hu & Adshead 2012; Benetti 2013; Cadavid & Romano 2015; Chluba, Hamann & Patil 2015), false vacuum decay (leading to open inflation, in particular) (Linde 1999; Linde, Sasaki & Tanaka 1999; Bousso, Harlow & Senatore 2014), or an inflection point in the potential (Allahverdi & Mazumdar 2006; Jain et al. 2009), or oscillations in the potential (Ashoorioon & Krause 2006; Pahud, Kamionkowski & Liddle 2009; Biswas, Mazumdar & Shafieloo 2010; Flauger et al. 2010; Aich et al. 2013; Hazra 2013; Peiris, Easther & Flauger 2013; Easther & Flauger 2014; Meerburg, Spergel & Wandelt 2014; Motohashi & Hu 2015; Miranda et al. 2016) all lead to local and non-local oscillations in the spectrum. Direct reconstruction of the primordial spectrum from the Planck data (Hazra et al. 2014c) hints at large-scale oscillations, an intermediate-scale burst of oscillations, and persistent high-frequency oscillations within intermediate to small scales. While these types of features can be obtained by different potentials, in this work we will be using the WWI, which is known the provide these local and non-local features in a unified framework.
Citation: Joy, Sahni & Starobinsky 2008
Citation function: Motivation; Motivation
While there are increasing efforts to try to explain the SFR dependence on the environment, by conducting surveys at high redshift (e.g. Hayashi et al. 2010; Matsuda et al. 2011; Sobral et al. 2011; Muzzin et al. 2012; Koyama et al. 2013; Darvish et al. 2014; Tal et al. 2014), so far such studies have not been able to fully reveal the physical processes leading to the ultimate quenching of (satellite) star-forming galaxies (e.g. Peng et al. 2010; Muzzin et al. 2012, 2014). Several strong processes have been proposed and observed, such as harassment (e.g. Moore et al. 1996), strangulation (e.g. Larson, Tinsley & Caldwell 1980) and ram pressure stripping (e.g. Gunn & Gott 1972; Fumagalli et al. 2014). Observations are also showing a variety of blueshifted rest-frame ultraviolet (UV) absorption lines which indicate that most star-forming galaxies at least at z ∼ 1–2 are able to drive powerful gas outflows (e.g. Shapley et al. 2003; Weiner et al. 2009; Kornei et al. 2012) which may play a significant role in quenching, particularly if those happen in high-density environments. Evidence of such galactic winds have also been seen in e.g. Förster Schreiber et al. (2009) through broad components in the rest-frame optical Hα and [N ii] emission line profiles (e.g. Genzel et al. 2011). Spatially resolved observations allow for constraints on the origin of the winds within galaxies, and on the spatial extent of the outflowing gas, which are essential to derive mass outflow rates. In field environments, it is expected that such outflows will not be able to escape the halo (as long as it is massive enough and it is not a satellite), and in many conditions would likely come back and further fuel star formation (e.g. Hopkins et al. 2014). However, in the most massive clusters, such strong outflows will likely result in significant amounts of gas being driven out of the subhaloes that host star-forming galaxies, enriching the intracluster medium (ICM) and quickly quenching star-forming galaxies with the highest SFRs/highest outflow rates.
Citation: Moore et al. 1996
Citation function: Background
The close proximity in the sky of M 16 and M 17, two of the nearest giant HII regions of our galactic neighborhood lying at a similar distance from the Sun, naturally leads to the question of whether they are physically related, and whether they may share a common origin (Moriguchi et al. 2002; Oliveira 2008; Nishimura et al. 2017). Both giant HII regions are projected on the contour of a giant bubble-shaped structure, outlined in the distribution of HI and CO emission as first noted by Moriguchi et al. (2002). Evidence for triggered star formation in M 16 has been examined in detail by Guarcello et al. (2010), who concluded it had been induced externally, and not by the activity of its associated cluster NGC 6611. This suggests that the formation of M 16 and M 17 could have been triggered by the expansion of the bubble, powered by a previous generation of massive stars near its center, thus representing an example of triggered star formation on the scale of several tens of parsecs (Elmegreen 1998). Given the ages of the giant HII regions, the timescale of expansion of a wind-blown bubble, and the short lifetimes of massive stars, it is to be expected that the most massive members of that previous generation may have exploded as supernovae several Myr ago. The spatial dispersion of the members of the association that must have taken place progressively during its existence, combined with the distance of 2 kpc to the M 16–M 17 complex and the large number of unrelated foreground and background stars in that general direction, would make it extremely difficult to identify even its currently hottest members still remaining on the main sequence. Moriguchi et al. (2002) noted the presence of O and early B stars in the area and proposed that they were part of a massive star population responsible for having caused the bubble, but a review of their properties shows them to be generally too bright to be at the distance of the bubble and the giant HII regions, and instead are more likely members of a foreground population.
Citation: Moriguchi et al. (2002); Moriguchi et al. (2002)
Citation function: Background; Compare/Contrast
While new data will refine our understanding of the interplanetary medium, it is necessary for future models to be consistent with both new and archival data. It is precisely this need that motivates the COSMOGLOBE project1, which aims to create a framework that will allow for the refinement of astrophysical models jointly with the raw data from complementary experiments. This form of joint analysis is already being explored within the framework of WMAP (Watts et al. 2022) and LiteBIRD (Aurlien et al., in prep.), in combination with Planck Low Frequency Instrument data (BeyondPlanck 2022). ZE is an especially promising direction for joint analysis, in part because HFI made observations of ZE at complementary wavelengths to DIRBE. While the DIRBE model was modified during the HFI analysis, no attempt was made to improve upon the geometrical representation of the model components using the larger effective dataset. Our understanding of the interplanetary medium has improved since the development of the DIRBE IPD model (see, e.g., Reach et al. 1997; Reach 2010a,b). The AKARI satellite, which observed in the infrared at wavelengths between 6 and 180 µm, detected small-scale structures in the ZE which were not well-characterized by the DIRBE model (Ootsubo et al. 2016). While the DIRBE model has been successful in describing the large-scale diffuse ZE, modern high-resolution high-frequency and infrared experiments will require IPD models that can more effectively resolve the small-scale structures in the ZE. As such, an update of the community state-of-art ZE model is long overdue. In this COSMOGLOBE framework, such a model refinement is a natural byproduct of a joint analysis of DIRBE, HFI, and other data. To improve upon this model, it is essential to have the model agree with the data at all wavelengths. The modeling of ZE is in the process of being implemented in the Commander framework, which will ultimately be used for the joint processing of HFI and DIRBE data. However, a stand-alone code is useful for agile model building and data analysis. As such, ZodiPy will function as the first step toward this goal in the COSMOGLOBE framework, in addition to making ZE corrections on arbitrary data more accessible to the community.
Citation: BeyondPlanck 2022
Citation function: Background; Background; Background
Amongst the oldest stellar systems known to exist in the Milky Way (MW) are metal-poor globular clusters (GCs). These accumulations of stars do not seem to have undergone substantial star formation for extended periods. Given the limited quality of the available data, for a long time color-magnitude diagrams (CMDs) of GCs appeared to be narrow and could be readily described by a single isochrone. These observations have justified the establishment of the long-lasting paradigm that considers CGs as prime examples of simple stellar populations (SSPs), that is, the results of very short bursts of star formation in their natal clouds. However, improved photometric precision indicates the presence of sub-populations in the cluster CMDs that are inconsistent with the SSP assumption, for a number of luminous GCs in a variety of bandpasses. Thus, early detections of chemical abundance variations (e.g., Cohen 1978) could be more easily explained in a scenario involving several populations. Moreover, in recent years evidence has grown supporting the statement that GCs are generally composed of two or three chemically distinct populations. These subpopulations are separated by a few tens to hundreds of Myr in age and show vastly varying abundances of light elements such as C, N, O, Na, Mg, and Al (see, e.g., Carretta et al. 2009b; Gratton et al. 2012, and references therein). Theoretical considerations (see, e.g., D’Ercole et al. 2008, 2011) imply that GCs could have lost the majority of the initial stellar content of the first population, which consequently should have ended up in the Galactic halo. In fact, numerous studies found metal-poor GCs to be consistent with the abundance trends of the MW halo at equally low metal content (e.g., Pritzl et al. 2005; Koch et al. 2009; Koch & McWilliam 2014; Villanova et al. 2016). We address this scenario by adding NGC 6426 to the short list of metal-poor clusters with available information on detailed chemical abundances. There are only two GCs in the Harris catalog (Harris 1996, 2010 edition) more metal poor than NGC 6426. At 12.9 ± 1.0 Gyr, the cluster is the oldest in the age compilation by Salaris & Weiss (2002). At a galactocentric distance of Rgc = 14.4 kpc and a galactic latitude of 16.23° it is located in the transition region between inner and outer halo. Previous studies found consistent [Fe/H]1 values: −2.20 ± 0.17 dex (Zinn & West 1984), −2.33 ± 0.15 (Hatzidimitriou et al. 1999), and −2.39 ± 0.04 dex (Dias et al. 2015). The latter value originates from the very first spectroscopic analysis of NGC 6426 at low resolution, which also stated [Mg/Fe] = 0.38 ± 0.06 dex. To date, there has been no study further addressing the detailed metal content of this cluster.
Citation: Gratton et al. 2012
Citation function: Background
Citation: Carretta et al. 2009b
Citation function: Background
Various surveys of SNRs in our Galaxy and nearby galaxies have been carried out at radio, X-ray, Infrared (IR), and optical wavelengths. The first extragalactic SNR candidates were identified in the LMC by Mathewson & Healey (1964) and later confirmed with a combination of radio and optical techniques by Westerlund & Mathewson (1966). To date, a total of 60 SNRs have been confirmed in the LMC with an additional 14 suggested candidates (Maggi et al. 2016; Bozzetto et al. 2017; Maitra et al. 2019). However, sensitivity and resolution limitations severely reduce the effectiveness of the past and present generations of radio and X-ray searches for SNRs in galaxies beyond the Small and Large Magellanic Clouds (MCs) (Goss et al. 1980; Long et al. 1981; Cowan & Branch 1985; Matonick et al. 1997; Matonick & Fesen 1997; Millar, White & Filipovic 2012; Galvin & Filipovic 2014; Sasaki et al. 2018; Lin et al. 2020; Sasaki 2020). As a result, optical studies have produced the largest number (∼1200) of new extra-Magellanic SNR candidates. Optical extragalactic searches for SNRs are mainly done by using an emission line ratio criterion of the form [S ii]/H α > 0.4–0.5 (Mathewson & Clarke 1973; Dodorico, Dopita & Benvenuti 1980; Fesen 1984; Blair & Long 1997; Matonick & Fesen 1997; Dopita et al. 2010b; Lee & Lee 2014; Vučetić et al. 2019b, a, 2018; Lin et al. 2020). This criterion separates shock-ionization from photoionization in SNRs from H ii regions and Planetary Nebulae (PNe) (Frew & Parker 2010). SNR radiative shocks collisionally excite sulphur ions in the extended recombination region resulting in S+, hence the larger contribution of [S ii] accounting for an increase of the [S ii] to H α ratio. In typical H ii regions, sulphur exists predominantly in the form of S++, yielding low [S ii] to H α emission ratios. Ratios from narrow-band imaging are usually verified spectroscopically, since [N ii] lines at 6548 and 6584 Å can contaminate the H α images at an unknown and variable level. Spectroscopic observations of such emission nebulae also can provide other evidence of shock heating, such as strong [O i] λ6300 emission, elevated [N ii] to H α with respect to H ii regions, or high [O iii] electron temperatures, verifying the candidate as being an SNR (Blair, Kirshner & Chevalier 1981, 1982; Long et al. 1990; Smith et al. 1993; Blair & Long 1997). Although somewhat biased as an isolated criterion, this method is proven and a good way of identifying ordinary radiatively cooling SNRs in nearby galaxies. We note that young, Balmer-dominated SNRs (Chevalier, Kirshner & Raymond 1980) would be missed by this criterion.
Citation: Lin et al. 2020
Citation function: Background
Citation: Matonick & Fesen 1997
Citation function: Background
Citation: Lin et al. 2020
Citation function: Background
Citation: Vučetić et al.; 2018
Citation function: Background
We need to build data sets to train and evaluate the model. Then the data sets are prepared in the following way: (1) if the AR does not flare within 24 hr after the observation time, the No-flare (weaker than C1.0) label is assigned to the magnetogram sample in the same AR. (2) If the C/M/X-level flare occurs within 24 hr after the observation time, the corresponding flare label (i.e., C, M, or X) is assigned to the magnetogram sample. Note that there are a number of ARs producing recurring flares with different flare levels within 24 hr. For the first flare of one AR, the corresponding flare label is assigned to the magnetogram sample within 24 hr after the observation time. Then for the following flares of the same AR, the corresponding labels are assigned to the magnetogram sample in the period from the end of its prior flare to the end of this flare. (3) We adopt a four-level AR classification scheme based on the maximum GOES-level flare an AR ever yields (Song et al. 2009; Yuan et al. 2010; Liu et al. 2017). In other words, ARs are further categorized into four levels (i.e., No-flare, C, M, and X) if they yield at least one flare with such GOES-level criterion but no flares with higher GOES-level criterion: “Level = X” indicates that an AR yields at least one X-level flare; “Level = M” indicates that an AR yields at least one M-level flare but no X-level flares; “Level = C” indicates that an AR yields at least one C-level flare but no M/X-level flares; “Level = No-flare” indicates that an AR only yields microflares (weaker than C1.0 flares). Finally, we gather 870 ARs and 136134 magnetogram samples in total, including 443 X-level, 6534 M-level, 72412 C-level, and 56745 No-flare level samples. Note that the magnetogram samples with multiple ARs (Bobra et al. 2014) are not included in our work. For the ≥M class, magnetogram samples of M/X-level flare in an AR are defined as positive class, while all the others are defined as negative class. For the ≥C class, magnetogram samples of C/M/X-level flare in an AR are defined as positive class, while all the others are defined as negative class.
Citation: Song et al. 2009
Citation function: Uses