adsabs/FOCAL huggingface dataset:
We wanted to discuss the observed periods at sunspot. A period of nearly five minutes can be detected in Doppler velocities of CO 3-2 R14 and 7-6 R67 lines, which is consistent with previous findings in white light images or continuum spectrum (Beckers & Schultz 1972; Lites 1988; Nagashima et al. 2007; Yuan et al. 2014; Su et al. 2016), and might be considered the solar p-mode waves in the photosphere (Thomas 1985; Bogdan 2000; Solanki 2003). While a period of roughly three minutes is found in the Doppler velocities of the CO 3-2 R14 and Mg II k lines, which agrees closely with the previous observational results in UV–infrared lines or images at the sunspot umbra (e.g., Solanki et al. 1996; Bogdan 2000; Fludra 2001; Maltby et al. 2001; Centeno et al. 2008; Tian et al. 2014; Khomenko & Collados 2015; Yang et al. 2017). They are explained as the resonant modes of sunspot oscillations (Uexkuell et al. 1983; Thomas 1984; Gurman 1987; Khomenko & Collados 2015). On the other hand, the CO 3-2 R14 line is believed to provide the information in the upper photosphere or the temperature minimum region (Uitenbroek 2000; Ayres 2002; Ayres et al. 2006). Moreover, a time delay of about two minutes is measured between the CO 3-2 R14 and Mg II k lines, as shown in Fig. 9. So, the three-minute oscillation at the sunspot umbra could come from the upper photosphere or the temperature minimum region and then propagate to the chromosphere, supporting the interpretation that propagating waves above the sunspots originate from the lower solar atmosphere (De Moortel et al. 2002; O’Shea et al. 2002; Brynildsen et al. 2004; Khomenko & Collados 2015; Krishna Prasad et al. 2015). In other words, the three-minute oscillation can be regarded as the upwardly propagating slow magnetoacoustic waves (e.g., Su et al. 2016; Chae et al. 2019; Cho & Chae 2020; Feng et al. 2020b). Finally, Chae et al. (2017) found that the three-minute oscillation in the light bridge or umbral dots of a sunspot originates from the photosphere (see aslo Cho et al. 2019), which is consistent with our results and further suggests that the CYRA data is reliable.
Citation: Maltby et al. 2001
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It is important to emphasize that any estimate of [Na/Fe] should be actually taken as an upper limit. Note that we can compute the theoretical response of stellar spectra to [Na/Fe] only for stars hotter than 3500 K (see Section 3.2). As discussed in Appendix A, a linear extrapolation of Na responses to cooler temperatures would likely lead to lower [Na/Fe] estimates, by about 0.1 dex. Moreover, as already noticed by CvD12b, atomic NaI transitions in the atmospheres of late-type stars are prone to substantial departures from LTE (Bruls, Rutten & Shchukina 1992; Gehren et al. 2006; Andrievsky al. 2007; Lind et al. 2011). In a Sun-like star, LTE calculations predict weaker lines (Allende Prieto et al. 2003) requiring corrections for the strongest lines, which can be as high as an effective change in the abundance of ∼0.5 dex. For a lower temperature star (with Teff ∼ 4000 K), more relevant for models having old ages (as those of our sample of ETGs), we may expect NLTE corrections in the range of 0.1–0.2 dex (based on fig. 4 of Lind et al. 2011). Hence, we may expect that NLTE models could enhance the predicted absorption in Na indices, resulting in lower inferred values of [Na/Fe]. However, note that our LTE-based methodology can match NaD, $\rm Na\,\small {I}8190$, $\rm Na\,\small {I}1.14$, and $\rm Na\,\small {I}2.21$, simultaneously, suggesting that NLTE corrections should be approximately the same for all four Na lines – an important aspect to test with future models. We emphasize that although we have introduced four free-fitting parameters to match the four Na lines (the $\alpha _{{\rm Na}_j}$ constants; see equation 1), in practice, the values of αNaD, $\rm \alpha _{{\rm Na}\,\small {I}1.14}$ and $\rm \alpha _{{\rm Na}\,\small {I}2.21}$ are fully consistent with those of α-MILES and CvD12a models. Effectively, we are able to fit the four Na-sensitive line strengths of the seven X-Shooter spectra – spanning a range of age, metallicity, and [α/Fe] – based on only one ‘extra’ free-fitting parameter (i.e. the $\rm \alpha _{{\rm Na}\,\small {I}8190}$).
Citation: Lind et al. 2011
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Citation: Lind et al. 2011
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We performed an analysis on simulated XMM maps of point sources, presented in Brunner et al. (2008) for similarly large XMM exposures in the Lockman Hole, detecting no extended emission in the simulated maps containing the detected by XMM point sources. The higher sensitivity of Chandra towards the detection of point sources allows us to make a statistical assessment of the effect of sub-threshold (for XMM) AGNs toward the detection of extended emission. Performance of XMM observations was accompanied by deepening the Chandra data within one year from each other, which makes Chandra maps suitable for XMM point source contamination analysis, limiting the effect of AGN long-term variability (Salvato et al. 2011; Paolillo et al. 2004). We have computed the variation of unresolved point source flux on the detection scales for XMM, using the Chandra image, masking out the sources detected in the XMM analysis. The constructed Chandra flux map has been further smoothed with a Gaussian of 16′′ width, approximating the effects of the XMM PSF. In the map, the uniform distribution of the faintest point sources results in nearly constant emission, which we subtract following the procedure for local cosmic background estimates for XMM, while bright sources and clustered sources make an enhancement. We find the contamination by point sources unresolved by XMM to the flux of identified extended sources is below the 5% level of the extended source’s flux. The highest peaks in the contamination map are associated with stand-alone sources near the (XMM) detection threshold, which by chance happened not to coincide with any of the detected groups and would contribute 30% to the faintest group flux. The importance of these sources is even higher in shallow surveys (Mirkazemi et al. 2015), to a degree requiring matched detection thresholds between point-like and extended sources, effectively removing faint extended sources from consideration. The importance of point source removal in XMM data is mentioned also in other cluster publications (Hilton et al. 2010; Pierre et al. 2012).
Citation: Salvato et al. 2011
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One remarkable feature present in the central region of Orion is an explosive event that occurred 550 ± 25 years ago (J. Bally, priv. comm.) and was revealed by the three runaway stars BN, n, and I (Gómez et al. 2005; Rodríguez et al. 2005, 2017), and by the CO and H2 fingers (e.g., Allen & Burton 1993; Zapata et al. 2009; Nissen et al. 2012; Youngblood et al. 2016; Bally et al. 2017). However, Luhman et al. (2017) showed that the object n is no longer a runaway member because its real proper motion is much lower than previously estimated; but conversely another object, named x, is moving away at high speed from the same explosion center. J. Bally (priv. comm.) confirms the fast proper motion of x and the absence of movement of n from ground-based H2 images 14 years apart. X is further out having passed our 20% beam coupling mark (see Fig. 1 of Paper I), and therefore does not appear in the figures presented in this Letter. Zapata et al. (2011) and Orozco-Aguilera et al. (2017) in their follow-up work with ALMA proposed that the hot core (HC) is externally heated despite its high temperature, and that the heating source could be the nearby explosion. Similarly, Blake et al. (1987), Wang et al. (2011), and Favre et al. (2011) advocated that the Compact Ridge is also externally heated, although the heating source should not be the same since we presented evidence in Paper I that the Compact Ridge has not yet been affected by the impact of the explosion. A possibility could be the outflow from source I hitting the Compact Ridge (Liu et al. 2002). In Paper I, we also presented evidence for an interaction between the explosive event and the main components of the Orion KL region including the HC, several infrared (IR) components (Rieke et al. 1973), and methyl formate (CH3OCHO; hereafter MF) peaks (Favre et al. 2011). We showed that the IRc6/MF5 and IRc20/MF4 sources, west of the explosion center, display emission lines of various species having only red wings, while sources on the east and south sides display emission lines having only blue wings. We also confirmed that excited emission lines are found preferentially surrounding the explosion center and that complex organic molecules (COMs) rich in oxygen (O-COMs) do not occupy the same volumes as CN rich COMs (CN-COMs). We identified the ethylene glycol peak (EGP) to be coincident with a hollow sphere of material, which we interpreted to have originated from the impact of a “bullet” launched from the explosion center (Favre et al. 2017; Wright & Plambeck 2017). We also proposed that the Compact Ridge (MF1) is sufficiently far away from the rest of the KL region to have not yet been perturbed by the explosion, the evidence being the absence of asymmetric emission line wings and the narrowness of the lines themselves (∼1 km s−1). In this Letter, we study further the interaction of the explosion blowout with the surrounding gas and dense sources.
Citation: Wang et al. (2011)
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Although it is a simple concept, obtaining galaxy sizes is not an easy task and is subject to a number of assumptions. The most common way to derive galaxy sizes is by performing light-profile fitting assuming a given shape of the surface brightness profile using a χ2 minimization (e.g. Simard et al. 1999; Peng et al. 2002; Ravindranath et al. 2004; Daddi et al. 2005; Ravindranath et al. 2006; Trujillo et al. 2006; Akiyama et al. 2008; Franx et al. 2008; Tasca et al. 2009; Cassata et al. 2010, 2013; Williams et al. 2010; Mosleh et al. 2011; Huang et al. 2013; Ono et al. 2013; Stott et al. 2013; Morishita et al. 2014; van der Wel et al. 2014; Straatman et al. 2015; Shibuya et al. 2015). Another method assumes circular or elliptical apertures around a predefined galactic center and computes the size enclosing a given percentage of the total galaxy flux (e.g. Ferguson et al. 2004; Bouwens et al. 2004; Hathi et al. 2008; Oesch et al. 2010; Ichikawa et al. 2012; Curtis-Lake et al. 2016). A third approach, involving counting the number of pixels belonging to the galaxy to derive its size, was also explored in Law et al. (2007). Studies of galaxy sizes at z> 2 became possible with the deep imaging obtained with HST. The first reports on size evolution found that galaxy sizes as observed in the UV rest-frame were becoming smaller at the highest redshifts (Bouwens et al. 2003, 2004; Ferguson et al. 2004). We have now access to the size evolution up to z ~ 10 from the deepest HST imaging data (e.g., Hathi et al. 2008; Jiang et al. 2013; Ono et al. 2013; Kawamata et al. 2015; Holwerda et al. 2015; Shibuya et al. 2015). With the multiwavelength and near-infrared coverage of CANDELS (Grogin et al. 2011; Koekemoer et al. 2011) optical rest-frame measurements are reported up to z ~ 3 for a large collection of galaxies in diverse populations (e.g. Bruce et al. 2012; van der Wel et al. 2014; Morishita et al. 2014). At z ~ 2 the size of star-forming galaxies (SFGs) is, to first order, independent of the observed rest-frame bands (Shibuya et al. 2015). It is generally accepted that galaxy sizes tend to decrease with increasing redshift (e.g. Bouwens et al. 2003, 2004; Ferguson et al. 2004; Mosleh et al. 2012) and that galaxy sizes depend on stellar mass (e.g. Franx et al. 2008; van der Wel et al. 2014; Morishita et al. 2014) and luminosity (e.g. Grazian et al. 2012; Huang et al. 2013). However, some results point to a scenario consistent with no size evolution as seen in UV rest-frame from HST data (Law et al. 2007; Curtis-Lake et al. 2016) and, at a fixed stellar mass, from optical rest-frame ground-based data (Ichikawa et al. 2012; Stott et al. 2013).
Citation: Shibuya et al. 2015
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In recent years, the quest to observe embedded accreting planets as they form has come a long way as a result of combined observations that are covering the ultraviolet to the millimeter wavelength ranges. One instrument in particular has been playing a major role in this, due to its diffraction-limited and high-contrast observations: the Spectro-Polarimetric High contrast imager for Exoplanets REsearch (SPHERE; Beuzit et al. 2019), which covers the visible and near-infrared (VIS/NIR) wavelength ranges. Together with other instruments, it has revealed a large abundance of substructures that are now known to be commonly present in protoplanetary disks (PPDs) ranging from gaps and rings to spiral arms, cavities, and various asymmetric features (e.g., Garufi et al. 2016, 2018; Avenhaus et al. 2018; Andrews et al. 2018, DSHARP). These features are often interpreted as the result of embedded planets that interact with their natal circumstellar disk (CSD; e.g., Wolf & D’Angelo 2005; Fouchet et al. 2010; Ruge et al. 2013; Ober et al. 2015; Dong et al. 2016). However, other origins have been proposed as well (e.g., Flock et al. 2015; Zhang et al. 2015; Ruge et al. 2016; Gonzalez et al. 2017; Suriano et al. 2017, 2019; Dullemond & Penzlin 2018). Despite the ambiguity of their particular origins, the properties of the potential planets that may produce these features have been studied extensively, especially using hydrodynamics simulations (e.g., Bae et al. 2019; Toci et al. 2020; Calcino et al. 2020). Dong & Fung (2017), for instance, analyzed observed gap features in order to determine the masses of planets that may have caused them by using 2D and 3D hydrodynamics simulations with 3D radiative transfer simulations of five scattered light observations of CSDs around Herbig Ae/Be and T Tauri stars: HD 97048 (Ginski et al. 2016), TW Hya (van Boekel et al. 2017), HD 169142 (Momose et al. 2015), LkCa15 (Thalmann et al. 2016), and RX J1615 (de Boer et al. 2016). By assuming an α-viscosity model (Shakura & Sunyaev 1973) with αvisc = 10−3 and single gap-opening planets as origins, they deduced that the corresponding planetary masses are all of typical low-mass giant planets between about 0.1 and 1 MJ.
Citation: Avenhaus et al. 2018
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Both Greco etÂ al. (2009) and Servidio etÂ al. (2011) show that a remarkable similarity exists between the PDF of the field magnitude from 2D MHD simulations with a guide field and the PDF from ACE observations in the supersonic solar wind near Earth. This suggests that the common occurrence of strong magnetic field discontinuities in the solar wind (e.g., Li 2008) might be attributed to the ubiquitous presence of inertial-scale quasi-2D coherent magnetic field structures or flux ropes interspersed with current sheets and reconnection x-points. Additional arguments in support of this view come from theoretical considerations (e.g., Zank & Matthaeus 1992, 1993) and observations (Matthaeus etÂ al. 1990; Bieber etÂ al. 1996) that MHD turbulence in the supersonic solar wind can be thought of comprising to lowest order a dominant 2D component and a minor parallel propagating AlfvÃ©n wave component. Furthermore, recent observations in the supersonic solar wind suggest the common occurrence of magnetic islands with different sizes in the inertial range together with evidence for island merging near the heliospheric current sheet (Khabarova etÂ al. 2014). A statistical solar wind data analysis by Osman etÂ al. (2014) also supports the idea that strong magnetic field discontinuities have a larger likelihood of being associated with current sheets and x-point reconnection outflows and have a higher possibility that a connection exists between current sheets and such outflows. With today's higher resolution measurement capability, identification of x-point reconnection outflows in the solar wind has become quite common (e.g., Gosling 2010). Gosling (2010) found that most reconnection events are observed in the slow supersonic solar wind (â¼2â3 events per day near Earth), while in the fast supersonic solar wind much less are found (â¼0.6 events per day). Reconnection events are also regularly observed in association with coronal mass ejections (Gosling 2010) where inertial-scale flux-rope merging was reported at the current sheet behind a coronal mass ejection close to the Sun by Song etÂ al. (2012), and current sheet reconnection was observed at the turbulent leading edge of an interplanetary coronal mass ejection near Earth by Chain & Munoz (2011), for example.
Citation: Matthaeus etÂ al. 1990
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Fig. 1 clearly shows that in increasing fluence the processes of dissociation are greater, and thus new molecular species come to light. We identify the new species by comparing the peak positions with similar works on acetone ices. New observed peaks were: a flat peak in around 3 μm that could be associated to amorphous water, but since several other alcohols have a similar peak and water is not a highly expected acetone radiation product (see Section 3), we could not make any strong assignment to this peak. Some new peaks in the range of 2000–3500 cm−1, which we assign to isopropanol, ketene, and carbon monoxide according to the work of Hudson (2018). Other two peaks in the region 2000–600 cm−1 we assign to CH4 and -CH=CH- bending (Hudson 2018; Matrajt et al. 2013). The most sharp and visible peak is the CO2 one. CO2 could be a contaminant, however the CO2 peak is not present in our control experiments, such as monitoring the ice sample before irradiation to confirm that without irradiation sample remains unaltered, or also irradiation of the bare substrate. Also, the CO2 peak increases with fluence showing this is a molecule being formed due to irradiation. Due the vacuum chamber pressure to be 10−9 mbar, some residual gas could contaminate the sample, but this is accounted for our error bars estimates. The H2CCO peak can be overlapped with a CO one since both are expected radiation products and have a peak near 2128 cm−1. However, as stated in Hudson (2018), the shift of this peak to higher wavenumbers with increasing fluence indicates this peak is most due to formation of ketene. In comparison with the work of Hudson (2018), he observed peaks for CH4, CO, H2CCO, and C3H8O. Here, we observed and confirm a peak for CO2 as well other similar peaks presented in the work of Hudson (2018) for isopropanol and ketene. Photo-irradiation of pure acetone ice with X-rays was also performed in Almeida et al. (2014), where authors employed a photo-stimulated desorption (PSID) and took the positive PSID spectra to analyse acetone fragmentation. They have thus detected lots of cations as fragments, such as HCO+, CO+, H+, C2CH$_3\, ^+$, CO$_2\, ^+$, H2CO+, etc. Tables 1 and 2 present a compilation of peak positions (in wavenumbers and wavelength), our corresponding assignments, band strengths and references used to determine molecular assignments. Question mark in Table 2 is a peak probably related to several different molecules being formed in the ice as explained before. Note that ketene band strength was taken from Berg & Ewing (1991) and it was also used as reference value in other recent works (Hudson & Loeffler 2013; Bergner, Öberg & Rajappan 2019). No reference band strength for C3H8O were found.
Citation: Matrajt et al. 2013
Citation function: Uses
The beaming functions of the polar and fan beams are given by f sinmϕ and p coskϕ, respectively. Here ϕ is the angle between the magnetic field axis and the photon propagation and is thus a function of the angle between the magnetic and rotation axis (θ), the angle of the NS rotation axis, the observing angle (i), and the NS spin phase. The exponent values m and k can be as low as 1, but in the general case of the accretion column emission, they can have significantly higher values (Trümper et al. 2013). We constructed various pulse profiles for different combinations of fan- and polar-beam emission patterns with different relative intensities of the fan and polar beams (i.e., FPolar∕FFan ranging from ~0.01 to ~30). We note that these values exceed any realistic configuration between the fan beam and the reflected polar beam. More specifically, since the polar beam is a result of reflection of the fan beam, it is only for a very limited range of observing angles that its contribution will (seemingly) exceed that of the fan beam, and thus a value of FPolar∕FFan ≳ 2, is unlikely. On the other hand, the increasing height of the emitting region, which would result in a smaller fraction of the fan beam being reflected by the NS surface, is limited to ≲ 10 km (Mushtukov et al. 2015), thus limiting the FPolar∕FFan ratio to values greater than 0.1 (see, e.g., Fig. 2 of Poutanen et al. 2013). Despite these physical limitations, we have extended our estimates to these exaggerated values in order to better illustrate the contribution of each component (fan and polar) to the PF. Moreover, a value of FPolar∕FFan exceeding 10 qualitatively describes the pencil-beam regime, which is expected at lower accretion rates. We estimated the PF for a range of observer angles and for different combinations of angles between the magnetic and rotation axis. An indicative result for a 45° angle between the NS magnetic field and rotation axis and for an rG∕rNS = 0.25 is plotted inFig. 12.
Citation: Trümper et al. 2013
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For the clustering measurements, we use the sample of galaxies compiled in Data Release 11 (DR11) of the SDSS-III project. The SDSS-III is a spectroscopic investigation of galaxies and quasars selected from the imaging data obtained by the SDSS (York et al. 2000) I/II covering about 11,000 deg2 (Abazajian et al. 2009) using the dedicated 2.5 m SDSS Telescope (Gunn et al. 2006). The imaging employed a drift-scan mosaic CCD camera (Gunn et al. 1998) with five photometric bands (





 and z; Fukugita et al. 1996; Smith et al. 2002; Doi et al. 2010). The SDSS-III (Eisenstein et al. 2011) BOSS project (Ahn et al. 2012; Dawson et al. 2013) obtained additional imaging data of about 3000 deg2 (Aihara et al. 2011). The imaging data was processed by a series of pipelines (Lupton et al. 2001; Pier et al. 2003; Padmanabhan et al. 2008) and corrected for Galactic extinction (Schlegel et al. 1998) to obtain a reliable photometric catalog. This catalog was used as an input to select targets for spectroscopy (Dawson et al. 2013) for conducting the BOSS survey (Ahn et al. 2012) with the SDSS spectrographs (Smee et al. 2013). Targets are assigned to tiles of diameter 3° using an adaptive tiling algorithm designed to maximize the number of targets that can be successfully observed (Blanton et al. 2003). The resulting data were processed by an automated pipeline which performs spectral classification, redshift determination, and various parameter measurements, e.g., the stellar-mass measurements from a number of different stellar population synthesis codes which utilize the photometry and redshifts of the individual galaxies (Bolton et al. 2012). In addition to the galaxies targeted by the BOSS project, we also use galaxies that pass the target selection but have already been observed as part of the SDSS-I/II project (legacy galaxies). These legacy galaxies are subsampled in each sector so that they obey the same completeness as that of the CMASS sample (Anderson et al. 2014).
Citation: Ahn et al. 2012
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