Identifier
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| Citation Text
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list | Functions Label
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list | Functions Start End
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|---|---|---|---|---|---|---|
2022AandA...662A..34D__Hyde_&_Bernardi_2009_Instance_1
|
The marginalized probability for Mgal and Mhalo is shown in Fig. 6. The contours correspond to the 68% and 95% intervals and the black dot marks the position of the best model. We note that the best model is outside the 68% region of the marginalized probability. The N-dimensional likelihood near the best model forms a shallow valley that extends toward the 68% confidence region. The best models clearly prefer a relatively narrow region in the Mgal − MDM space. An interesting result from this figure, is that a baryon-only model (i.e., MDM = 0) does not perform much worse than a model with a DM halo, although the data prefers a model with a DM halo. The relatively weaker dependence with the DM component may be a consequence of the small size of the Einstein ring, since at this radius the baryon component dominates the projected mass (see Table 1 below). The inferred mass of the galaxy (baryons) in the best model (5.4 × 1010 M⊙) is marginally consistent with the estimated stellar mass derived from the velocity dispersion – stellar mass correlation (Hyde & Bernardi 2009; Zahid et al. 2016; Cannarozzo et al. 2020). Assuming the velocity dispersion in Mörtsell et al. (2020), and the models in Zahid et al. (2016), Cannarozzo et al. (2020), the predicted stellar mass is ≈2.5 × 1010 M⊙. This is about a factor two less than the stellar mass inferred from the lens model. However, a factor two uncertainty is typical in mass estimations from the velocity dispersion (see the references above). Since the galaxy and halo masses are partially degenerate (see Fig. 6), it is possible that the galaxy mass is lower than the one at the maximum of the likelihood. Another possibility is that the gas mass is similar to that of the stars, increasing the baryonic mass by a factor two compared with the stellar mass. This is in principle feasible based on the gas to stellar mass ratios, which is close to 1 for early type galaxies (similar to the lens considered in this work) and stellar masses ∼1010 M⊙ (Calette et al. 2018). However, in the central part of the galaxy the stellar component is still expected to dominate, especially if a bulge is present, as suggested by the data. An alternative way of estimating the contribution from the gas is by comparing with galaxies of similar morphological type. According to Casasola et al. (2020), and considering a morphological type for the lens galaxy between Sa and Sb, the gas fraction for this type of galaxy is ≈20%. In Schruba et al. (2011), gas and molecular surface densities as high as O(100) M⊙ pc−2 can be found in areas with large star formation rates. In the most extreme cases of star formation rates, gas surface densities of ≈1000 M⊙ pc−2 can be found. However, even at these extreme star formation rates, the contribution from the gas to the baryonic mass is still subdominant with respect to the stellar contribution from our fiducial model.
|
[
"Hyde & Bernardi 2009"
] |
[
"The inferred mass of the galaxy (baryons) in the best model (5.4 × 1010 M⊙) is marginally consistent with the estimated stellar mass derived from the velocity dispersion – stellar mass correlation"
] |
[
"Similarities"
] |
[
[
1063,
1083
]
] |
[
[
865,
1061
]
] |
2020ApJ...892...53A__Connolly_et_al._2018_Instance_1
|
These new limits, in conjunction with the inconsistency of isotropic flux interpretations, leave no room for an astrophysical interpretation of AAEs in the context of the standard model for time windows as short as 103 s. However, it has been shown that these events can be explained using physics beyond the standard model, as many models suggest that AAEs lend support for axionic dark matter, sterile neutrinos, supersymmetry, or heavy dark matter (Anchordoqui et al. 2018; Connolly et al. 2018; Dudas et al. 2018; Fox et al. 2018; Huang 2018; Abdullah et al. 2019; Anchordoqui & Antoniadis 2019; Borah et al. 2019; Chauhan & Mohanty 2019; Cherry & Shoemaker 2019; Chipman et al. 2019; Cline et al. 2019; Collins et al. 2019; Esmaili & Farzan 2019; Esteban et al. 2019; Heurtier et al. 2019a, 2019b; Hooper et al. 2019). Many of these models, excluding the axionic dark matter explanation (Esteban et al. 2019) or those heavy dark matter scenarios that are tuned to prevent signatures in IceCube (Hooper et al. 2019), can be constrained by this nonobservation at IceCube. Dedicated tests to quantify these constraints are beyond the scope of this work and may be the focus of a future study. In addition to explanations that incite new physics, it has recently been suggested that AAEs could be explained by downward-going CR-induced EASs that reflected off of subsurface features in the Antarctic ice (Shoemaker et al. 2019). Another possible explanation could be coherent transition radiation from the geomagnetically induced air shower current, which could mimic an upgoing air shower (Motloch et al. 2017; de Vries & Prohira 2019). Explaining these anomalous events with systematic effects or confirming the need for new physics requires a deeper understanding of ANITA’s detection volume. Efforts such as the HiCal radio frequency pulser, which has flown alongside ANITA on the last two flights (Prohira et al. 2018), are already underway to try to characterize the various properties of the Antarctic ice surface.
|
[
"Connolly et al. 2018"
] |
[
"These new limits, in conjunction with the inconsistency of isotropic flux interpretations, leave no room for an astrophysical interpretation of AAEs in the context of the standard model for time windows as short as 103 s. However, it has been shown that these events can be explained using physics beyond the standard model, as many models suggest that AAEs lend support for axionic dark matter, sterile neutrinos, supersymmetry, or heavy dark matter",
"Many of these models,",
"can be constrained by this nonobservation at IceCube."
] |
[
"Motivation",
"Motivation",
"Motivation"
] |
[
[
477,
497
]
] |
[
[
0,
450
],
[
824,
845
],
[
1021,
1074
]
] |
2020MNRAS.494.3413T__Shidatsu_&_Done_2019_Instance_2
|
The existence of winds is shown by blueshifted absorption lines from highly ionized ions. These are only seen in soft state but not in hard state (Ponti et al. 2012), anticorrelated with the radio jet which is seen in the hard state but not in the soft. This was thought to be evidence that the wind was magnetically driven by the same field as was responsible for the jet, but in a different geometric configuration (Miller et al. 2012). However, in Tomaru et al. (2019, hereafter Paper I) we show instead that thermally driven winds can explain this switch (see also Done, Tomaru & Takahashi 2018; Shidatsu & Done 2019). Thermal driving produces a wind by irradiation from the central source heating the surface of accretion disc up to the Compton temperature ($T_\text{IC} \sim 10^7 \!-\!10^8\, \text{K}$), which is hot enough for its thermal energy to overcome the gravity at large radii. The characteristic radius at which the wind can be launched is called the Compton radius, defined by RIC = μmpGM/kTIC ∼ 105 − 106Rg (Begelman, McKee & Shields 1983). Paper I show the first modern radiation hydrodynamic simulations of thermal (and thermal-radiative) winds designed to investigate the switch in wind properties between the hard and soft states changing illumination spectra. These simulations were tailored to the BHB system H1743−332, where there is Chandra high-resolution data in both states giving detailed spectral information on the wind or its absence (Miller et al. 2012; Shidatsu & Done 2019). They incorporate radiation force on the electrons, both bound and free, as they show that this is important factor driving the escape of the thermal wind in the fairly high Eddington fraction (L/LEdd ∼ 0.2–0.3), fairly low Compton temperature (TIC ∼ 0.1 × 108 K) characteristics of the soft state. The only other modern hydrodynamic simulation of thermal winds (e.g. Luketic et al. 2010; Higginbottom & Proga 2015; Higginbot et al. 2016) has not included radiation pressure, which is important in setting the velocity structure for L ≥ 0.3LEdd as required here (Paper I).
|
[
"Shidatsu & Done 2019"
] |
[
"These simulations were tailored to the BHB system H1743−332, where there is Chandra high-resolution data in both states giving detailed spectral information on the wind or its absence"
] |
[
"Uses"
] |
[
[
1488,
1508
]
] |
[
[
1283,
1466
]
] |
2022AandA...659A.124H__Liu_et_al._(2013b)_Instance_2
|
Combining different samples from various instruments at different redshifts therefore inevitably introduces ENLR size–luminosity relations with different slopes α depending on the details of target selection and analysis approaches. Slopes ranging from α = 0.22 ± 0.04 (Greene et al. 2012), α = 0.25 ± 0.02 (Liu et al. 2013b), α ∼ 0.3–0.4 (Hainline et al. 2013; Chen et al. 2019a), to α ∼ 0.5 (Bennert et al. 2002; Husemann et al. 2014) are reported in the literature. The slopes solely inferred from the CARS data are consistent with those reported by Greene et al. (2012) and Liu et al. (2013b) and are therefore on the shallower side of previous estimates. Nevertheless, the scatter in the observed relation is significant and measured slope variations might be entirely attributed to the observationally induced biases as discussed above. A slope of α = 0.5 is reminiscent of the BLR size-luminosity relation, but would require a constant ionization parameter U that demands a constant density with radius. This is not observed for the ENLR on kiloparsec scales (e.g., Bennert et al. 2006; Kakkad et al. 2018) and more detailed photoionization calculations are required to predict the shallower slopes inferred for most studies (Dempsey & Zakamska 2018). We cannot study the radial variations of U as our snapshot MUSE observations are not deep enough to map the electron density given the too low S/N of the [S II] doublet on kpc scales. However, the photoionization calculations do not take into account variable ionizing flux from AGN on 105 yr time scales (Schawinski et al. 2015) and the various geometrical intersections of the ionizing radiation field with the gas distribution of the galaxies. The CARS survey is least biased with regard to RENLR given the narrow redshift range and large dynamic range offered by MUSE (see Fig. 13). Therefore, the CARS survey is one of the best data set to explore the origin of the significant scatter in ENLR size–luminosity relation and search for additional factors or more fundamental parameters controlling the ENLR size.
|
[
"Liu et al. (2013b)"
] |
[
"The slopes solely inferred from the CARS data are consistent with those reported by Greene et al. (2012) and",
"and are therefore on the shallower side of previous estimates."
] |
[
"Similarities",
"Similarities"
] |
[
[
578,
596
]
] |
[
[
469,
577
],
[
597,
659
]
] |
2019AandA...630A.151P__Gallerani_et_al._(2011)_Instance_1
|
Capturing the full information content of the cosmic large-scale structure requires a field-based approach to infer the entire three-dimensional cosmic large-scale structure from observations. This poses a particular challenge for the analyses of Ly-α forest observations, which provide sparse inherently one-dimensional information along the lines of sight. Various approaches to perform three-dimensional density reconstructions from one-dimensional Ly-α forests have been proposed in the literature (e.g. Kitaura et al. 2012; Cisewski et al. 2014; Stark et al. 2015a; Ozbek et al. 2016; Horowitz et al. 2019). Gallerani et al. (2011) and Kitaura et al. (2012) proposed a Gibbs sampling scheme to jointly infer density and velocity fields and corresponding power-spectra. However, these approaches assume matter density amplitudes to be log-normally distributed. The log-normal distribution reproduces one- and two-point statistics but fails to reproduce higher-order statistics associated with the filamentary dark matter distribution. In an attempt to extrapolate information from one-dimensional quasar spectra into the three-dimensional volume, Cisewski et al. (2014) applied a local polynomial smoothing method. Ozbek et al. (2016) and Stark et al. (2015a) employed a Wiener filtering approach to reconstruct the three-dimensional density field between lines of sight of Ly-α forest data. In order to reproduce higher-order statistics, Horowitz et al. (2019) recently used a large-scale optimization approach to fit a gravitational structure growth model to Ly-α data, showing that this approach allows recovering the more filamentary structure of the cosmic web. Although the approach improves over linear and isotropic Wiener filtering approaches, it shows systematic deviations of reconstructed matter power-spectra and underestimates density amplitudes at scales corresponding to the mean separation between lines of sight (Horowitz et al. 2019).
|
[
"Gallerani et al. (2011)"
] |
[
"and Kitaura et al. (2012) proposed a Gibbs sampling scheme to jointly infer density and velocity fields and corresponding power-spectra.",
"However, these approaches assume matter density amplitudes to be log-normally distributed. The log-normal distribution reproduces one- and two-point statistics but fails to reproduce higher-order statistics associated with the filamentary dark matter distribution."
] |
[
"Background",
"Compare/Contrast"
] |
[
[
613,
636
]
] |
[
[
637,
773
],
[
774,
1038
]
] |
2016AandA...586A..16B__Mol_&_Defrise_(2004)_Instance_1
|
The problem we are facing consists in the reconstruction of two maps of a given region of an astrophysical object: the former consists in point sources, such as stars, of very high intensity, the latter of smooth structures surrounding these sources. The standard approach, i.e. simply solving Eq. (2), fails since the presence of the point sources destroys all the information about the smooth structures. Hence, the main idea is to consider the object x as the sum of two components, namely x = xP + xE, where xP represents the point sources and xE the extended source. This approach was proposed for the first time in De Mol & Defrise (2004), assuming that the positions of the point sources are known. Denoting by pj the position of the jth source, we can write (4)\begin{equation} \label{2:point} x_{\rm P}=\sum_{j=1}^qc_j\delta(p_j) , \end{equation}xP=∑j=1qcjδ(pj),where q is the total number of sources, cj is the jth unknown intensity, and δ(pj) is the delta function centred in pj. Thus, xP is a vector with zero entries except in the q positions corresponding to the known locations of the sources. Then, instead of computing f1 on the whole object x in Eq. (2), we only regularize the extended source, since the structure induced on xP already works as a regularization. This requires a slightly modification in the computation of f0. We introduce the vector c = (c1,c2,...,cq)t containing the intensities of the sources, and we define \hbox{$\overline{x} = (c^t,x_{\rm E}^t )^t$}x=(ct,xEt)t and the matrix \hbox{$\mathcal{H}=[\overline{H}, H]$}ℋ=[H,H], where \hbox{$\overline{H}=[h_{p_1},h_{p_2},\dots,h_{p_q}]$}H=[hp1,hp2,...,hpq], with hj denoting the jth column of H. We are hence led to solve(5)\begin{equation} \label{2:solve} \tilde{x} =\arg\min_{x\in\mathcal{C}} KL(\mathcal{H} \overline x+b;g)+\beta f_1\left(x_{\rm E}\right) . \end{equation}˜x=argminx∈𝒞KL(ℋx+b;g)+βf1(xE).It may happen that \hbox{$\tilde x$}˜x has a loss in contrast, even using the optimal value for β. To overcome this difficulty, we then propose the use of the inexact Bregman procedure, which permits the use of an overestimation of the regularization parameter, and at the same time allows us to obtain a contrast enhancement.
|
[
"De Mol & Defrise (2004)"
] |
[
"Hence, the main idea is to consider the object x as the sum of two components, namely x = xP + xE, where xP represents the point sources and xE the extended source. This approach was proposed for the first time in",
"assuming that the positions of the point sources are known."
] |
[
"Uses",
"Uses"
] |
[
[
621,
644
]
] |
[
[
407,
620
],
[
646,
705
]
] |
2022ApJ...926..155S__Umehata_et_al._2017_Instance_1
|
However, the sensitivity limits of most current submillimeter surveys only allow for the study of the most extremely star-bursting systems (e.g., Asboth et al. 2016; Geach et al. 2017; Simpson et al. 2019), which may represent merely a tip of an iceberg of dust-obscured star formation in the early Universe. A potential population of more typical dusty star-forming galaxies at high-z is still to be found (Wang et al. 2017). The Atacama Large Millimetre/submillimetre Array (ALMA) has now opened a new avenue to refine our understanding of dusty galaxies at high redshifts, enabling to uncover faint SMGs down to a flux level of 0.1–1 mJy. Several ALMA blind surveys have been performed and allowed to detect and characterize the faint SMGs across cosmic times (e.g., Aravena et al. 2016; Wang et al. 2016b; Dunlop et al. 2017; Umehata et al. 2017; Franco et al. 2018; Williams et al. 2019; Yamaguchi et al. 2019; González-López et al. 2020). Based on the ALMA survey of GOODS-South field over an area of 69 arcmin2, Franco et al. (2018) found that 20% of the 1.1 mm sources are not detected with HST down to a depth of H ∼ 28 mag, and suggested that they are massive main-sequence star-forming galaxies at z > 4 (see also, Yamaguchi et al. 2019; Umehata et al. 2020). A similar fraction of HST-dark galaxies has also been found in the ALMA [C ii] survey of main-sequence galaxies at 4.4 z 5.9 (∼14%, Gruppioni et al. 2020). Conversely, the existence of such HST-dark galaxies can also be uncovered by focusing the reddest galaxies in the IRAC and H bands (H-[4.5 μm] > 4.0), namely H-dropouts (Huang et al. 2011; Caputi et al. 2012; Wang et al. 2016a). Follow-up continuum observations with ALMA of a sample of 63 H-dropouts have yielded detections of 39 sources down to an 870 μm flux density of 0.6 mJy (Wang et al. 2019). They further suggested that the ALMA-detected H-dropouts are the bulk populations of massive (M
⋆ ≳ 1010.3
M
⊙) star-forming galaxies at z > 3 with the contribution to star formation rate density an order of magnitude higher than that of equivalently massive LBGs. To uncover the nature of H-dropouts and the critical role they play in the cosmic evolution of massive star-forming galaxies, we need to explore the fainter population that might have even fainter (sub)millimeter fluxes.
|
[
"Umehata et al. 2017"
] |
[
"Several ALMA blind surveys have been performed and allowed to detect and characterize the faint SMGs across cosmic times (e.g.,"
] |
[
"Background"
] |
[
[
830,
849
]
] |
[
[
642,
769
]
] |
2018ApJ...860...24P__Warmuth_2015_Instance_2
|
Figure 13 shows the temporal evolution of the density, ρ, plasma flow velocity, vx, position of the wave crest, PosA, phase speed, vw, and magnetic field component in the z-direction, Bz, for the primary waves in every different case of initial amplitude, ρIA. In Figure 13(a), we observe that the amplitude of the density remains approximately constant at their initial values until the time when the shock is formed and the density amplitude of the primary wave starts decreasing (see Vršnak & Lulić 2000), i.e., at t ≈ 0.03 (blue), t ≈ 0.04 (red), and t ≈ 0.055 (green). For the case of ρIA = 1.3 (magenta), a decrease of the amplitude of the primary can hardly be observed, as expected for low-amplitude wave (Warmuth 2015). One can see that the larger the initial amplitude, ρIA, the stronger the decrease of the primary wave’s amplitude, which is consistent with observations (Warmuth & Mann 2011; Muhr et al. 2014; Warmuth 2015). The amplitudes decrease to values of ρ ≈ 1.6 (blue), ρ ≈ 1.5 (red), and ρ ≈ 1.4 (green) until the primary wave starts entering the CH. Due to the fact that the waves with larger initial amplitude enter the CH earlier than those with small initial amplitude, we can see in Figure 13(a) that the tracking of the parameters of the faster waves stops at an earlier time than the one for the slower waves. A similar behavior to the one of the density, ρ, can be observed for the plasma flow velocity, vx, in Figure 13(b) and the magnetic field component, Bz, in Figure 13(e). Here, the amplitudes decrease from vx = 0.75, Bz = 1.9 (for ρIA = 1.9, blue), vx = 0.6, Bz = 1.7 (for ρIA = 1.7, red), vx = 0.45, Bz = 1.5 (for ρIA = 1.5, green), and vx = 0.27, Bz = 1.3 (for ρIA = 1.3, magenta) to vx = 0.55, Bz = 1.6 (for ρIA = 1.9, blue), vx = 0.46, Bz = 1.5 (for ρIA = 1.7, red), vx = 0.36, Bz = 1.4 (for ρIA = 1.5, green), and vx = 0.25, Bz = 1.25 (for ρIA = 1.3, magenta). Figure 13(c) shows how the primary waves propagate in the positive x-direction. In all four cases of different initial amplitude, ρIA, the phase speed decreases slighty (consistent with observations; see Warmuth et al. 2004 and Warmuth 2015) until the waves enter the CH at different times, i.e., the values for the phase speed start at vw ≈ 2.2 (for ρIA = 1.9, blue), vw ≈ 1.9 (for ρIA = 1.7, red), vw ≈ 1.7 (for ρIA = 1.5, green), and vw ≈ 1.4 (for ρIA = 1.3, magenta) and decrease to vw ≈ 1.5 (for ρIA = 1.9, blue), vw ≈ 1.39 (for ρIA = 1.7, red), vw ≈ 1.2 (for ρIA = 1.5, green), and vw ≈ 1.13 (for ρIA = 1.3, magenta).
|
[
"Warmuth 2015"
] |
[
"One can see that the larger the initial amplitude, ρIA, the stronger the decrease of the primary wave’s amplitude, which is consistent with observations"
] |
[
"Similarities"
] |
[
[
922,
934
]
] |
[
[
729,
881
]
] |
2018ApJ...866...93L___2012c_Instance_1
|
Interestingly, from ∼11:33:40 UT, the PAD of the suprathermal electrons changes to cigar type, in association with a dramatic drop of electron flux (Figure 2(a)). This change of electron PAD and flux is related to a magnetic dip structure that is manifested by a conspicuous decrease of the magnetic field strength (Figure 2(b)). In contrast to the magnetic field enhancement observed at ∼11:33:08 UT, the magnetic dip structure should arise from the local expansion of flux tubes, which are driven by the two opposite flows (see the shaded region in Figure 2(c)). The observed electron cigar distribution and the associated magnetic dip are a strong indication of the betatron cooling effect. This betatron-mediated cigar distribution has been suggested in previous studies (Fu et al. 2011b, 2012c, 2013b; Liu et al. 2017c) but never clearly observed. These observations, for the first time, show a direct link between the cigar distribution and betatron cooling. Note that in the trailing edge of the magnetic dip region, a magnetic hump structure, indicated by the sharp enhancement of magnetic field, is observed at ∼11:33:55 UT (Figure 2(b)). This magnetic hump, reminiscent of the Earthward-propagating dipolarization-front structure typically generated by magnetic reconnection in the midtail (Fu et al. 2012d, 2012e, 2013a, 2015, 2016, 2017; Liu et al. 2013, 2018a, 2018b, 2018c; Cao et al. 2017; Peng et al. 2017; Yao et al. 2017; Chen et al. 2018), was possibly formed due to the local contraction of flux tubes driven by the ion flow with an increasing velocity in the Earthward direction (Figure 2(c)). Associated with this magnetic hump, a very weak pancake distribution of suprathermal electrons is observed (Figure 2(a)), and strong waves near the electron gyrofrequency are also observed (Figures 2(e) and (f)). These waves are whistler-mode; they indicate that the betatron acceleration inside the magnetic hump was accompanied by non-adiabatic effects (Fu et al. 2009, 2010a, 2010b, 2012a, 2012b) and is still ongoing (Fu et al. 2011a; Khotyaintsev et al. 2011; Wang et al. 2017; Yang et al. 2017).
|
[
"Fu et al.",
"2012c"
] |
[
"The observed electron cigar distribution and the associated magnetic dip are a strong indication of the betatron cooling effect. This betatron-mediated cigar distribution has been suggested in previous studies",
"but never clearly observed. These observations, for the first time, show a direct link between the cigar distribution and betatron cooling."
] |
[
"Similarities",
"Similarities"
] |
[
[
776,
785
],
[
793,
798
]
] |
[
[
565,
774
],
[
825,
964
]
] |
2018AandA...616L...2K__Frew_et_al._(2016)_Instance_1
|
The distances to planetary nebulae (PNe) have always faced the difficulty that nearby targets were lacking that could be reached well by direct methods. Trigonometric parallaxes have been obtained in a homogeneous long time-line campaign by the US Naval Observatory (USNO; Harris et al. 2007) and from the Hubble Space Telescope (HST; Benedict et al. 2009). Other studies (Acker et al. 1998; Smith 2015) showed that Hipparcos spacecraft parallaxes do not seem to be reliable. It was assumed that contamination by the emission of the surrounding nebulae caused these problems. Another model-independent method for distances to PNe are a cluster membership, as studied extensively by Majaess et al. (2007, 2014), and as discussed in Frew et al. (2016). In addition to these model-independent methods, a wide variety of statistical, model-dependent individual distance scales have been derived. The most frequently used of these are certainly those that are based on surface brightness versus angular sizes. They sometimes include optical depth corrections. All these methods have to be calibrated against a data set of nebulae with known distances. The older, widely used method is based on the 6 cm radio continuum flux, either using the ionized mass concept of Daub (1982) in the calibrations of Cahn et al. (1992) and Stanghellini et al. (2008), or by means of the radio continuum brightness temperature as used by van de Steene & Zijlstra (1994) and calibrated with a Galactic bulge sample. The newest model developed by Frew et al. (2016) is based on similar ideas, but makes use of the optical Hα surface brightness and a wide set of various calibrators. Moreover, they use a completely homogeneous data set for the brightness data derived earlier by themselves (Frew et al. 2013). Smith (2015) and Frew et al. (2016) described the underlying physics and assumptions for all these methods in detail. With the upcoming Gaia project (Gaia Collaboration 2016), a new era was expected to start for many classes of objects. The first step into this was described by Stanghellini et al. (2017) based on the combined Tycho + Gaia DR1 solution called TGAS (Michalik et al. 2015). With the second Data Release of Gaia (hereafter GDR2; Gaia Collaboration 2018), a complete homogeneous data set based only on Gaia measurements is available now for the first time. We present here the comparison of this new data set with common previous calibrations of PNe distances. Moreover, we compare it to the preliminary TGAS results in Stanghellini et al. (2017). Finally, we discuss possible caveats using the current GDR2.
|
[
"Frew et al. (2016)"
] |
[
"Another model-independent method for distances to PNe are a cluster membership, as studied extensively by Majaess et al. (2007, 2014), and as discussed in"
] |
[
"Background"
] |
[
[
731,
749
]
] |
[
[
576,
730
]
] |
2019MNRAS.489.4669S__Bigiel_et_al._2010_Instance_2
|
In Fig. 7 we compare UGC 1378’s SFR density versus gas surface density (the Schmidt–Kennicutt relation) to data in the literature. The gas surface density corresponds to H i calculated from the 0th moment map from Mishra et al. (2017) in the same areas as SFR density. Points for the HSB and LSB discs are plotted as black and grey circles, respectively. We plot the mean SFR and H i surface density for the entire galaxy with a large open circle. The black line corresponds to the relation with an exponent of 1.4 found by Kennicutt (1998). Triangles give results for LSB galaxies published by Wyder et al. (2009), and bright and faint crosses show normal spiral galaxies from Kennicutt (1998) – total and H isurface densities. A blue line shows the best-fitting relation for the H isurface density of Bluedisk galaxies from Roychowdhury et al. (2015). We also plot the SFR in the outer regions of spiral galaxies (Bigiel et al. 2010, square symbols). In Fig. 7 the UGC 1378 measurements lie between normal spirals and LSB galaxies. The HSB disc data lie above the relation plotted for normal spirals, possibly indicating that the SFR is boosted by the bar driving gas to the star-forming rings. We cannot account for molecular gas since there are no available measurements for UGC 1378. The contribution of molecular gas would likely move the HSB disc of UGC 1378 towards the locus of normal galaxies. Because the HSB SFR of UGC 1378 is close to the predicted SFR from the Kennicutt (1998) relation obtained from H i densities (faint crosses in Fig. 7). The LSB disc of UGC 1378 lies below the correlation and accounting for molecular gas would only increase the deviation from the normal Schmidt–Kennicutt relation. Similar deviations are observed in ‘classical’ LSB galaxies, Bluedisk galaxies (Roychowdhury et al. 2015), outer parts of HSB spiral galaxies (Bigiel et al. 2010), and H idiscs in early-type galaxies (Yıldız et al. 2017). These deviations for LSB galaxies are at least partially explained by their lower gas densities leading to lower SFRs (Abramova & Zasov 2011). A recent episode of gas accretion on to the disc of UGC 1378 may also contribute to a lower SFR if the gas is not yet fully participating in the star formation. Lutz et al. (2017) studied a sample of very H i rich galaxies and proposed that very high specific angular momentum in H irich galaxies prevents the accreted gas from being transported to the mid-plane of the disc and being converted into stars. This mechanism may act to preserve giant gaseous discs.
|
[
"Bigiel et al. 2010"
] |
[
"Similar deviations are observed in",
"outer parts of HSB spiral galaxies"
] |
[
"Similarities",
"Similarities"
] |
[
[
1862,
1880
]
] |
[
[
1719,
1753
],
[
1826,
1860
]
] |
2022ApJ...928...51B__Yadav_et_al._2016_Instance_1
|
Relatively few such global convective dynamo studies have been conducted in the domain of M-dwarf stars. Early work by Browning (2008) considering lower mass FC M dwarfs found that the deep CZ could support very strong nonaxisymmetric fields, which strongly quenched the star’s differential rotation. Later, more turbulent simulations of FC M-dwarf stars led to a number of interesting results. Yadav et al. (2015a) found strong, axisymmetric fields which were statistically steady in time and recovered many of the observed characteristics of M-dwarf surface fields. A somewhat slower rotating model (Yadav et al. 2015b) revealed that flux concentration by merging downflow lanes could lead to the formation of large, persistent high-latitude starspots in these stars. A still slower rotating model (Yadav et al. 2016) built large-scale, axisymmetric, cycling magnetic fields of somewhat lower amplitude which did not eliminate the star’s differential rotation, reminiscent of the distributed αΩ-type dynamos prevalent in solar-like contexts. Brown et al. (2020) presented the first simulation of a stratified, rotating FC star whose computational domain extended to r = 0, finding a preference for hemispheric dynamo states. Similar models were recently explored by Käpylä (2021) for a broad range of rotation rates, which yielded a variety of results. At slow rotation, differential rotation was antisolar, and dynamo action dipole dominant. Moderate rotation rates yielded solar-like profiles, which were then magnetically quenched in the fastest rotators, yielding nonaxisymmetric dynamos. In Bice & Toomre (2020; hereafter BT20), we presented an exploration of the influence exerted by a tachocline in more massive, shell-convecting M-dwarf stars as a contributing factor to the break in observed magnetic activity across the tachocline divide. Our models produced a wide variety of field configurations, nearly all of which led to quenching of the differential rotation to a significant degree. We found that including a tachocline in models of early M-dwarf stars led to their surface fields being more favorable for rapid stellar spin-down, which may contribute to the formation of the tachocline divide.
|
[
"Yadav et al. 2016"
] |
[
"A still slower rotating model",
"built large-scale, axisymmetric, cycling magnetic fields of somewhat lower amplitude which did not eliminate the star’s differential rotation, reminiscent of the distributed αΩ-type dynamos prevalent in solar-like contexts."
] |
[
"Background",
"Background"
] |
[
[
801,
818
]
] |
[
[
770,
799
],
[
820,
1043
]
] |
2017ApJ...836L...4S__Scholer_&_Burgess_2007_Instance_2
|
At quasi-perpendicular shocks, the average shock structure is dominated by a foot of reflected ions, which is upstream of the shock ramp where the major thermalization and deceleration occurs. Non-stationarity in the form of rippling of the surface or steepened whistler waves (Moullard et al. 2006; Lobzin et al. 2007) is an intrinsic feature of the shock, but this is generally manifest as minor perturbations on top of an otherwise stationary shock ramp. Simulations have predicted that if the fraction of ions reflected by the shock front becomes sufficiently high, the quasi-perpendicular shock can become periodically reforming on timescales of the ion gyroperiod. Various theories have been suggested for such non-stationarity, including self-reformation where a new shock ramp grows at the edge of the foot (Biskamp & Welter 1972a; Lembège & Dawson 1987), whistler-induced reformation (Biskamp & Welter 1972b; Scholer & Burgess 2007), kinetic instabilities such as the Buneman and modified two-stream instability (e.g., Cargill & Papadopoulos 1988; Matsukiyo & Scholer 2003, 2006b; Scholer et al. 2003; Scholer & Burgess 2007; Matsumoto et al. 2013), and gradient catastrophe of nonlinear whistler waves due to steepening (Krasnoselskikh et al. 2002). However, it has not been until recently that such non-stationarity has been confirmed with in situ spacecraft observations. In a survey of Cassini shock crossings at Saturn, Sulaiman et al. (2015) found evidence of a periodically reforming shock, pulsating at a period near 0.3 of the ion gyroperiod in the unperturbed upstream medium. This period agrees with the time taken for a specularly reflected proton to gyrate across the foot and return to the main shock ramp. Sulaiman et al. (2015) also report that these periodic non-stationary shocks are primarily found in the very high Mach number regime, which gives evidence for a relation between Mach number and reformation. The main processes behind the non-stationary behavior of these very high Mach number shocks, such as the details of the ion- and electron-scale processes acting within the shock transition, remain elusive.
|
[
"Scholer & Burgess 2007"
] |
[
"kinetic instabilities such as the Buneman and modified two-stream instability (e.g.,"
] |
[
"Background"
] |
[
[
1111,
1133
]
] |
[
[
943,
1027
]
] |
2016AandA...588A..25M__Weiss_&_Ferguson_2009_Instance_1
|
There are some indications that available models of the post-AGB and CSPN phases are not accurate enough. First, the two available grids of post-AGB models (Vassiliadis & Wood 1994; Blöcker 1995a) do not agree with each other on the predicted timescales (Zijlstra et al. 2008). Second, consistency between the masses of white dwarfs and those of CSPNe seems to require faster evolutionary speeds than predicted by both sets of models (Gesicki et al. 2014). Third, present models of the CSPNe phase are unable to explain why the cut-off of the PNe luminosity function is constant in most galaxies (Marigo et al. 2001, 2004). Lastly, post-AGB stellar evolution models, computed with updated physics in a reduced mass range (Kitsikis 2008; Weiss & Ferguson 2009), show a strong disagreement with the previous grids. This is not a surprise since many improvements have been carried out in the field of stellar physics in recent decades. Most importantly, available grids have been computed with opacities, which are now 45 years old (Cox & Stewart 1970b,a) before the big changes introduced by the OPAL (Iglesias & Rogers 1996), and Opacity Project (Seaton 2007) redeterminations. Similarly, nuclear reaction rates, equation of states, conductive opacities, and neutrino emission rates adopted in the models date from the early eighties and even earlier. In addition, Herwig et al. (1997) showed that the existence of carbon stars at low luminosities can be explained by the addition of mixing beyond the formal convective boundaries during the thermal pulses (TP) on the AGB. Finally, Marigo (2002) showed that C-rich molecular opacities are essential to predict the correct effective temperatures once the AGB models become carbon rich (NC/NO> 1, by number fractions). This is particularly important because of the impact of effective temperatures on the mass loss rates. While all these improvements in stellar modeling have been implemented in AGB stellar models, and very detailed and exhaustive grids and models are available (Weiss & Ferguson 2009; Cristallo et al. 2009, 2011; Ventura & Marigo 2010; Karakas 2010; Lugaro et al. 2012; Constantino et al. 2014; Doherty et al. 2015), the inclusion of these improvements in post-AGB stellar models is still missing. It is time for a recomputation of the post-AGB models in the light of all these advances.
|
[
"Weiss & Ferguson 2009"
] |
[
"Lastly, post-AGB stellar evolution models, computed with updated physics in a reduced mass range",
"show a strong disagreement with the previous grids."
] |
[
"Differences",
"Differences"
] |
[
[
737,
758
]
] |
[
[
624,
720
],
[
761,
812
]
] |
2019AandA...628A.110M__Kryukova_et_al._(2012)_Instance_2
|
Deriving the completeness limits of the WISE photometry is mandatory to assess the reliability of our catalogue of starless cores. We examined the histograms of the number of mid-infrared (MIR) sources versus magnitude; taking into account the effects of the cuts required to fulfil the criteria of Koenig et al. (2012), rough completeness limits are [3.6] ~ 14, [4.6] ~ 12, [12] ~ 9 and [22] ~ 7. These values are 1–3 mag brighter than the sensitivity limits quoted in the WISE Explanatory Supplement3 for the relevant sky region. Once converted into flux units and, for example, compared with the models of Class I and Class 0 sources of 0.5 M⊙ by Whitney et al. (2004), it can be seen that the completeness limits at 3.6 and 4.6μm are faint enough to detect such objects taking into account a distance of 700 pc and a further foreground reddening up to AV = 20. Even in the worst case of edge-on discs, these objects would be detectable at 3.6 and 22μm. Furthermore, the completeness limit at 22 μm is faint enough to allow detection of Class I and Class 0 sources of even-lower-mass central objects. Alternatively, one can compute the bolometric luminosity following Kryukova et al. (2012). Starting from our completeness limit at 22 μm, after conservatively dereddening it by AV = 20, we assumed a spectral index γ = −2 (see Table 3 for definition) to compute the MIR luminosity from Eq. (6) of Kryukova et al. (2012). Equation (7) of Kryukova et al. (2012) then yields Lbol ~ 1.7–2.8 L⊙, depending on whether the NIR flux is neglected (which may be the case) or extrapolated from γ = −2. A comparison with the birthline of Palla & Stahler (1993) indicates a mass of ~ 0.4–0.5 M⊙ for the central protostar. For the sake of comparison, we can roughly estimate the completeness limit in central masses of the Herschel protostellar cores in Giannini et al. (2012) using their quoted completeness limit at 70 μm of 0.21 Jy and following Dunham et al. (2008). By using Eq. (2) of Dunham et al. (2008), scaled to a distance of 700 pc, we found that the flux density at 70 μm translates into a bolometric luminosity of the central (proto)star Lbol ~ 0.28 L⊙ (we note that Dunham et al. 2008 indicate this luminosity as Lint). We highlight the fact that the 70 μm emission is in principle a more sensitive protostellar tracer than WISE. However, this contrasts with the much lower number of protostellar cores found by Giannini et al. (2012), which may be due to a poorer effective sensitivity because of their selection criteria.
|
[
"Kryukova et al. (2012)"
] |
[
"Starting from our completeness limit at 22 μm, after conservatively dereddening it by AV = 20, we assumed a spectral index γ = −2 (see Table 3 for definition) to compute the MIR luminosity from Eq. (6) of"
] |
[
"Uses"
] |
[
[
1400,
1422
]
] |
[
[
1195,
1399
]
] |
2020ApJ...903....6M__Bueno_2010_Instance_1
|
In this paper, we consider resonance scattering on a two-level atom with an infinitely sharp and unpolarized lower level. As for the frequency redistribution, we consider both CRD and angle-averaged PRD. In an unmagnetized one-dimensional spherically symmetric atmosphere, the polarized radiation field is axially symmetric and is described by
. We also take into account the effects of radial velocity fields. We solve the concerned transfer equation in the CMF and in the nonrelativistic regime of velocity fields. To efficiently handle the sphericity effects, we solve the spherical transfer equation in the (p, z) coordinate system (Hummer & Rybicki 1971), which is also called the tangent-ray method. Here z is the distance along the tangent rays and p is impact parameter (see Figure 1 in Megha et al. 2019). Following Frisch (2007), we express the Stokes vector components in terms of their irreducible components. From here on we present all the basic equations in the irreducible basis (see Sampoorna & Trujillo Bueno 2010 for details). The CMF polarized PRD transfer equation for a spherically symmetric medium in (p, z) representation under the nonrelativistic limit is given by (see also Equation (17) of Megha et al. 2019)
1
where x denotes the nondimensional frequency and
denotes the irreducible Stokes vector with “+” and “−” referring to the outgoing and incoming rays, respectively. The radial optical depth is defined as dτr = −χl(r)dr, where r is the radial distance and χl(r) is the line averaged absorption coefficient. In the CMF the monochromatic optical depth along the tangent ray is given by dτ = [φ(x) + βc]dτr/μ, where φ(x) is the line absorption profile function and βc = χc(r)/χl(r) with χc(r) denoting the continuum absorption coefficient. The direction cosine of the tangent ray about the radius vector of the intercepting radial shell is given by
. In Equation (1),
denotes the CMF term, which has the form
2
where χ(r, x) = χl(r)φ(x) + χc(r), and
3
The symbol V denotes the ratio of radial (vr) to the thermal (
) velocities. In Equation (1),
represents the irreducible CMF total source vector and is given by
4
where the line source vector is of the form
5
Here ϵ gives the probability of destruction of photons by inelastic collisions, Bν0 is the Planck function at the line center frequency ν0, and
. The continuum is assumed to be unpolarized. Therefore, the continuum source vector is given by
. The frequency-averaged PRD mean intensity vector is given by
6
where
is the 2 × 2 nonmagnetic angle-averaged PRD matrix (Domke & Hubeny 1988; Bommier 1997), and
7
The Rayleigh phase matrix
in the irreducible basis is given in Appendix A of Frisch (2007). It is useful to rewrite Equation (7) as
8
where
9
|
[
"Sampoorna & Trujillo Bueno 2010"
] |
[
"From here on we present all the basic equations in the irreducible basis (see",
"for details)."
] |
[
"Uses",
"Uses"
] |
[
[
1006,
1037
]
] |
[
[
928,
1005
],
[
1038,
1051
]
] |
2015MNRAS.450..763M__Rieke_et_al._2004_Instance_1
|
In this work, we focus on the 56 most massive (M* ≥ 1011 M⊙) galaxies at 1.4 ≤ z ≤ 2, 31 of which have spectroscopic redshifts. Two more objects entered the original sample, but they have been excluded from this study, because their WFC3/HST images are not available, or too noisy to perform SB fitting, due to the proximity of saturated stars. The sample was culled from the K-selected (K(Vega) 22) multiband catalogue of Daddi et al. (2007a, hereafter D07), including data from all the available filters in GOODS-S, i.e. HST/ACS optical, F435W (B), F606W (V), F775W (I), and F850LP(z), VLT/ISAAC near-IR, J, H, K, and Spitzer/Infrared Array Camera (IRAC), 3.6, 4.5, 5.8, and 8.0 μm (for more details on the data sets see Giavalisco et al. 2004). The optical/near-IR photometry was then complemented with the 24 μm catalog (Daddi et al., in preparation), built as summarized in Section 2.1 from the Multi-Band Imaging Photometer for Spitzer (MIPS) images (Rieke et al. 2004). Far-IR fluxes were extracted (Daddi et al., in preparation) from the publicly released PACS 70–160 μm data from Herschel GOODS (Elbaz et al. 2011), and SPIRE 250 μm data from Herschel Multi-tiered Extragalactic Survey (HerMES; Oliver et al. 2010). For objects without spectroscopic information, we used photometric redshifts from the public GOODS-MUSIC catalogue, which agree well with the spectroscopic ones (Δz/(1 + z) ≃ 0.03) for galaxies at z 2 (cf. Grazian et al. 2006, 2007). Although the uncertainties on photometric redshifts could result in the inclusion in the sample of a few lower redshift contaminants, we avoided applying colour criteria to pre-select high-z objects, in favour of completeness. The HST/WFC3/F160W H-band image mosaic (>5σ point source sensitivity for H160 27.7, AB system) drizzled to a pixel scale of 0.06 arcsec, was exploited to study galaxy morphology in the optical rest frame with a very high resolution (FWHM ∼ 0.18 arcsec ≃ 1.5 kpc). For more details on the observations and data reduction see Grogin et al. (2011) and Koekemoer et al. (2011).
|
[
"Rieke et al. 2004"
] |
[
"The optical/near-IR photometry was then complemented with the 24 μm catalog",
", built as summarized in Section 2.1 from the Multi-Band Imaging Photometer for Spitzer (MIPS) images"
] |
[
"Uses",
"Uses"
] |
[
[
958,
975
]
] |
[
[
749,
824
],
[
855,
956
]
] |
2018MNRAS.473.2000T__Noutsios_et_al._2011_Instance_2
|
The launch of the Fermi Gamma-ray Space Telescope has spurred on the search for pulsars in γ-rays (Grenier & Harding 2015), yielding over 2001 detections and triggering multiwavelength observations. While pulsars are common targets in the X-rays, they are very challenging targets in the optical and very few of them have been identified (see Mignani et al. 2016, and references therein). Here, we report on Large Binocular Telescope (LBT) observations of an isolated pulsar, PSR J2043+2740 (Taylor, Manchester & Lyne 1993), detected by both AGILE (Pellizzoni et al. 2009) and Fermi (Abdo et al. 2010; Noutsios et al. 2011). It was discovered as a radio pulsar (Ray et al. 1996) and later on as an X-ray source by XMM–Newton (Becker et al. 2004), although X-ray pulsations have not yet been found. PSR J2043+2740 is one of the very few non-recycled pulsars older than 1 Myr detected in γ-rays, with a characteristic age τc = 1.2 Myr, inferred from the values of its spin period Ps = 0.096 s and its derivative $\dot{P}_{\rm s} = 1.27 \times 10^{-15}$ s s−1 (Ray et al. 1996). This also yields a rotational energy loss rate $\dot{E}_{\rm rot} = 5.6 \times 10^{34}$ erg s−1 and a surface dipolar magnetic field Bs = 3.54 × 1011 G.2 Although PSR J2043+2740 does not have a very large spin-down power compared to young (∼1–10 kyr) pulsars (∼1036–1038 erg s−1), it is still a factor of 2 larger than that of middle aged γ-ray pulsars (∼0.1–0.5 Myr), such as Geminga, PSR B0656+14 and PSR B1055−52, all detected in the optical, thanks to their distances ≲ 500 pc (Abdo et al. 2013). The distance to PSR J2043+2740 is uncertain owing to the lack of a radio parallax measurement. The radio dispersion measure (DM = 21.0 ± 0.1 pc cm−3; Ray et al. 1996) gives a distance of 1.8 ± 0.3 kpc from the NE2001 model of the Galactic-free electron density (Cordes & Lazio 2002). A slightly smaller distance (1.48 kpc) is inferred from the model of Yao, Manchester & Wang (2017). The hydrogen column density towards the pulsar obtained from the X-ray spectral fits (NH ≲ 3.6 × 1020 cm−2; Abdo et al. 2013) suggests a distance of a few hundred pc (He, Ng & Kaspi 2013), although these estimates depend on the model X-ray spectrum. Such a distance would make PSR J2043+2740 a viable target for deep optical observations, never carried out until now, and might be compatible with the debated association (Noutsios et al. 2011) with the Cygnus Loop supernova remnant (SNR) at $540^{+100}_{-80}$ pc (Blair, Sankrit & Raymond 2005).
|
[
"Noutsios et al. 2011"
] |
[
"Such a distance would make PSR J2043+2740 a viable target for deep optical observations, never carried out until now, and might be compatible with the debated association"
] |
[
"Motivation"
] |
[
[
2383,
2403
]
] |
[
[
2211,
2381
]
] |
2017AandA...607A...9M__Weideman_1994_Instance_1
|
Having a line list that is as complete as possible is crucial for doing proper radiative transfer computations in high temperature atmospheric environments. However, computing the line opacities for a large number of lines can be computationally challenging. For example, the ExoMol line list of CH4 contains on the order of 1010 lines (Yurchenko & Tennyson 2014). Computing the exact pressure and temperature broadened Voigt profile for each of these lines is computationally extremely demanding. There are several approximate methods for computing Voigt profiles available in the literature (e.g. Humlícek 1982; Weideman 1994; Zaghloul & Ali 2012). Also, much effort is spend succesfully on making these approximate methods faster and more accurate (see e.g. Poppe & Wijers 1990; Letchworth & Benner 2007). These methods all focus on obtaining a given accuracy of the exact shape of the Voigt profile by applying mathematical approximations to decrease the computation time. While these methods still require significant computation time, they are now routinely applied to compute Voigt profiles in many applications. The fastest code able to compute Voigt profiles of large numbers of lines at this moment is the HELIOS-k code (Grimm & Heng 2015). This code is able to compute ~ 105 lines per second on a dedicated NVIDIA K20 GPU based machine. This implies that the computation of 1010 lines still requires on the order of one day for a single point in pressure temperature space. Usually a grid of pressure and temperature points is required. Thus, there is the need for an even faster method. In addition, we have to make sure that there are no systematic errors in the computations because a small systematic error in a single line can become large when computed for 1010 lines. Thus, we seek a method that is statistically exact, preserves the integrated opacities and the average shape of the lines, and computes the line profile accurately for the stronger lines.
|
[
"Weideman 1994"
] |
[
"Computing the exact pressure and temperature broadened Voigt profile for each of these lines is computationally extremely demanding. There are several approximate methods for computing Voigt profiles available in the literature (e.g."
] |
[
"Background"
] |
[
[
614,
627
]
] |
[
[
365,
598
]
] |
2019MNRAS.490.2155S__Blake_et_al._2016_Instance_1
|
Modern optical imaging surveys measure the positions and ellipticities of millions of galaxies; from them, the galaxy overdensity field as well as the gravitational lensing shear field can be derived. The two-point auto and cross-correlations of these two fields are the two-point correlation functions of cosmic shear, galaxy–galaxy lensing and galaxy clustering. A joint analysis of these correlation functions can break degeneracies between cosmological and nuisance parameters, leading to tighter cosmological constraints (Joachimi & Bridle 2010). Several earlier studies have indeed considered such joint analyses (Cacciato et al. 2013; Mandelbaum et al. 2013; More et al. 2015; Kwan et al. 2017; Nicola, Refregier & Amara 2017), albeit very few in a modified gravity context. Among the latter, Joudaki et al. (2018) recently performed a combined analysis of cosmic shear tomography, galaxy–galaxy lensing tomography, and redshift space multipole power spectra using data from KiDS-450 (∼450 deg2 of cosmic shear data from the KiDS survey) and two overlapping spectroscopic surveys, the 2-degree Field Lensing Survey4 (2dFLenS; Blake et al. 2016) and the Baryon Oscillation Spectroscopic Survey5 (BOSS; Dawson et al. 2013). They found that none of the extended cosmologies considered were simultaneously favoured in a model selection sense and able to resolve the discordance with Planck, except for an evolving dark energy component with a time-dependent w0 − wa equation of state. Amon et al. (2018) presented a measurement of EG, a statistic combining measurements of weak gravitational lensing, galaxy clustering, and redshift space distortions, proposed as a consistency test of General Relativity (Zhang et al. 2007). They determined the value of EG using data from the KiDS, 2dFLenS, BOSS, and Galaxy And Mass Assembly6 (GAMA; Driver et al. 2009, 2011; Liske et al. 2015) surveys; their results show that measurements of the EG statistic cannot be conducted as consistency checks of General Relativity until the aforementioned tension in cosmological parameters is resolved, and their EG measurements favour a lower matter density cosmology than the CMB. Recently, DES Collaboration (2019) presented a combined analysis of galaxy clustering and weak gravitational lensing from the first-year data of the Dark Energy Survey, targeting modifications of the metric potentials that would be a signal of modified gravity. They found that their constraints are compatible with a cosmological constant scenario.
|
[
"Blake et al. 2016"
] |
[
"Among the latter, Joudaki et al. (2018) recently performed a combined analysis of cosmic shear tomography, galaxy–galaxy lensing tomography, and redshift space multipole power spectra using data from KiDS-450 (∼450 deg2 of cosmic shear data from the KiDS survey) and two overlapping spectroscopic surveys, the 2-degree Field Lensing Survey4 (2dFLenS;"
] |
[
"Background"
] |
[
[
1133,
1150
]
] |
[
[
782,
1132
]
] |
2016AandA...596A..59F__Sobotka_et_al._1993_Instance_1
|
Sunspot light bridges can be formed by the fragmentation of the umbra in the decay phase or during the merging of different magnetized areas during the formation of the sunspot in complex active regions (Bray & Loughhead 1964; Garcia de La Rosa 1987). At the last stages of a sunspot, the photospheric-like conditions are recovered as a consequence of the splitting of the umbra, and a granulation pattern similar to that of the quiet Sun is found in light bridges (Vazquez 1973), although the light-bridge convection cells differ significantly from normal granules (Lagg et al. 2014). Understanding the decay of sunspots requires the study of the magnetic and dynamical structure of light bridges. Previous studies have classified light bridges based on their morphological properties. Light bridges separating two umbral regions are called strong light bridges (e.g., Sobotka et al. 1993, 1994; Jurčák et al. 2006; Rimmele 2008). Their brightness is similar to that of the penumbra, and they typically appear between two regions with the same polarity. Faint light bridges (e.g., Lites et al. 1991; Sobotka et al. 1993) are elongated bright structures that penetrate the umbra. They are composed of rows of bright grains with sizes comparable to umbral dots. Light bridges can exhibit different internal structures. Most of them show several segments along their length that resemble granules (Berger & Berdyugina 2003), while some light bridges are similar to the bright filaments seen in the penumbra (Lites et al. 2004). All types show a weakened magnetic field strength relative to the surrounding umbra and more strongly inclined field lines (e.g. Beckers & Schröter 1969; Lites et al. 1991; Leka 1997), which in some case even exhibit a polarity reversal (Lagg et al. 2014). Observations are consistent with a reduced field strength in the photosphere, with a magnetic canopy extending from either side of the light bridge and merging at the top (Jurčák et al. 2006; Lagg et al. 2014). The velocity field measured in light bridges shows evidence of their convective origin (Rimmele 1997; Hirzberger et al. 2002; Lagg et al. 2014). Recently, Schlichenmaier et al. (2016) have also found evidence that the umbral magnetic field is wrapped around light bridges. The authors defined a new type, the plateau light bridge, with Y-shaped dark canals that resemble penumbral grains and suggest the presence of inclined magnetic fields. Several dynamic phenomena have been detected at the chromosphere above light bridges, including recurrent plasma ejections visible in Hα (Roy 1973; Asai et al. 2001), Ca ii H (Shimizu et al. 2009), and other bands (AIA 1600 and 1700 Å, IRIS 1330 and 1400 Å, Toriumi et al. 2015), brightness enhancements in the TRACE 1600 Å band (Berger & Berdyugina 2003), or fan-shaped jets observed in Hα (Robustini et al. 2016). All these processes are considered to be caused by the interaction of the light bridge magnetic field with the surrounding atmosphere through magnetic reconnection events.
|
[
"Sobotka et al. 1993",
"Sobotka et al. 1993"
] |
[
"Previous studies have classified light bridges based on their morphological properties. Light bridges separating two umbral regions are called strong light bridges (e.g.,",
"Their brightness is similar to that of the penumbra, and they typically appear between two regions with the same polarity. Faint light bridges (e.g.,",
"are elongated bright structures that penetrate the umbra. They are composed of rows of bright grains with sizes comparable to umbral dots."
] |
[
"Background",
"Background",
"Background"
] |
[
[
870,
889
],
[
1101,
1120
]
] |
[
[
699,
869
],
[
932,
1081
],
[
1122,
1260
]
] |
2021AandA...655A..12T__Tang_et_al._2017b_Instance_5
|
Using the RADEX3 non local thermodynamic equilibrium (LTE) modeling program (van der Tak et al. 2007) with collisional rate coefficients from Wiesenfeld & Faure (2013), we modeled the relation between the gas kinetic temperature and the measured average of para-H2CO 0.5 × [(322–221 + 321–220)/303–202] ratios, adopting a 2.73 K background temperature, an average observational linewidth of 4.0 km s−1, and column densities N(para-H2CO) = 2.7 × 1012 and 3.7 × 1012 cm−2 for N113 and N159W, respectively. The results are shown in Fig. 5. The values of the para-H2CO column density were obtained with APEX data (beam size ~30″; Tang et al. 2017b), which cover similar regions. Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in Tang et al. 2017b or Fig. 4 in Tang et al. 2018a; this was also shown in Fig. 13 and discussed in Sect. 4.3.1 of Mangum & Wootten 1993) as long as all lines are optically thin. Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in Tang et al. 2017b), we modeled it at spatial densities 104, 105, and 106 cm−3 in Fig. 5. It appears that Tkin at n(H2) = 105 cm−3 is consistently lower than values at 104 and 106 cm−3 by ≲23% and ≲34%, respectively, for Tkin ≲ 100 K. Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions (Mangum & Wootten 1993; Tang et al. 2017a,b, 2018b). Following the method applied by Tang et al. (2017b) in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3–2) line ratio in Fig. 5. Apparently, TLTE agrees well with Tnon-LTE at volume densities n(H2) ~ 105 cm−3 as long as Tkin ≲ 100 K. Previous observations show that para-H2CO (3–2) is sensitive to gas temperature at density 105 cm−3 (Ginsburg et al. 2016; Immer et al. 2016; Tang et al. 2017b). The spatial density measured with para-H2CO (303–202) and C18O (2–1) in N113 and N159W is n(H2) ~ 105 cm−3 on a size of ~30″ (Tang et al. 2017b). Therefore, here we adopt 105 cm−3 as an averaged spatial gas density in the N113 and N159W regions.
|
[
"Tang et al. (2017b)"
] |
[
"Following the method applied by",
"in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3–2) line ratio in Fig. 5."
] |
[
"Uses",
"Uses"
] |
[
[
1558,
1577
]
] |
[
[
1526,
1557
],
[
1578,
1705
]
] |
2021ApJ...907...15H__Geiss_1982_Instance_1
|
Figure 2 shows the ratios of the relative fluences (element/Mg) measured in Genesis bulk SW (Figure 2(a)) and regime targets (Figure 2(b)) after normalization to the respective solar abundance ratios (Table 7 in Appendix A.5). The bulk SW data for K and Fe in silicon (marked by asterisks in Figure 2(a)) are the preliminary fluences from Rieck (2015), Rieck et al. (2016), and Burnett et al. (2017). If SW had the composition of the solar photosphere, the ordinate value in Figure 2 would be 1 for all elements. The bulk SW shows a fractionation of elemental abundances correlating with FIP, as previously observed by spacecraft (e.g., Geiss 1982; Bochsler 2009). However, the better accuracy and element coverage of Genesis data compared to in situ measurements allows further insights into this FIP-related fractionation. The low-FIP elements (Figure 2(a)) are consistent with a flat pattern, at least below 7 eV, as many pre-Genesis studies suggested. This consistency would make these SW data easier to apply to cosmochemistry, as they would suggest that low-FIP elements are unfractionated relative to each other and hence their SW abundance ratios equal photospheric ratios. However, Figure 2(a) is also consistent with a monotonic increase of the low-FIP elements with decreasing FIP, which could be viewed as a trend continuing that of the high-FIP elements between Ar and C discussed below. The Na abundance given here is about 20% higher than that of Burnett et al. (2017). Na fluences derived from backside depth profiling of diamond-like-C collectors are roughly a factor of 2 lower than ours, which are based on Si collectors (Rieck 2015; Rieck et al. 2016). Jurewicz et al. (2019) propose that diffusion of surface contamination Na might have enhanced fluences from Si collectors. Because the amounts of surface contamination is highly variable, the Na fluences from Si should show significant scatter; however, replicate analyses of bulk and regime samples agree to within 2%–11% (Table 4, Appendix A.1). More analyses are required to resolve this important discrepancy. A ∼20% lower fractionation of K compared to that of Na, if real, would be surprising although within 2σ, both values still overlap.
|
[
"Geiss 1982"
] |
[
"The bulk SW shows a fractionation of elemental abundances correlating with FIP, as previously observed by spacecraft (e.g.,",
"However, the better accuracy and element coverage of Genesis data compared to in situ measurements allows further insights into this FIP-related fractionation."
] |
[
"Similarities",
"Compare/Contrast"
] |
[
[
637,
647
]
] |
[
[
513,
636
],
[
665,
824
]
] |
2022MNRAS.517.4986G__Chambers_1999_Instance_1
|
Wang et al. (2017), motivated by the fact that previous studies on high-eccentricity migration focused only on the total efficiency of HJs formed, analyzed the efficiency of each high-eccentricity mechanism, trying to understand which mechanism is the dominant one in HJ formation. They considered multiplanetary systems containing from two to five planets with equal masses (1MJ), initially in circular and near-coplanar orbits with a host star of $1\, {\rm M}_\odot$ and $1\, {\rm R}_\odot$. Moreover, various initial semimajor axis conditions were considered, depending on the initial mutual separation between the planets. Their numerical simulations were performed using the classical version of the Mercury code (Chambers 1999), which does not include general relativity (GR) effects and tidal interaction with the central star. They studied how the initial number of planets, the spatial separation between them, and the location of the inner planet influence the efficiencies of high-eccentricity mechanisms. As a result, they found that the Kozai–Lidov mechanism plays the most important role in HJ production. The restriction made in the previous study, regarding not including the contribution of GR and realistic initial mass configurations (unequal masses), makes us wonder about the implications of these contributions on the efficiency of each high-eccentricity mechanism in the activation of the HJ production process. The main effect of GR is to cause the apsidal precession of planetary orbit, where the rate of precession is faster for planets close to the star and with eccentric orbits due to the term a5/2(1 − e2) in the denominator, being a and e the planet’s semimajor axis and eccentricity, respectively (Einstein 1916; Misner, Thorne & Wheeler 1973). Studies related to the influence of GR on the dynamic evolution of systems with more than one massive planet have been carried out, but only considering some particular systems (e.g. Adams & Laughlin 2006; Migaszewski & Goździewski 2009; Veras & Ford 2010; Zhang, Hamilton & Matsumura 2013; Marzari & Nagasawa 2019, 2020). Therefore, the aim of our study is to obtain a broader and more detailed description of the efficiency of each high-eccentricity mechanism in the activation of the HJ production process, including the contribution of GR and different initial planetary mass configurations. The analysis is made by solving the numerical simulation of the exact equations of motion, in the context of general N-body problem. Several initial conditions are considered, changing the initial mass and number of planets, the semimajor axis of the inner planet and the location of the other planets in the system.
|
[
"Chambers 1999"
] |
[
"Their numerical simulations were performed using the classical version of the Mercury code",
"which does not include general relativity (GR) effects and tidal interaction with the central star. They studied how the initial number of planets, the spatial separation between them, and the location of the inner planet influence the efficiencies of high-eccentricity mechanisms. As a result, they found that the Kozai–Lidov mechanism plays the most important role in HJ production. The restriction made in the previous study, regarding not including the contribution of GR and realistic initial mass configurations (unequal masses), makes us wonder about the implications of these contributions on the efficiency of each high-eccentricity mechanism in the activation of the HJ production process."
] |
[
"Motivation",
"Motivation"
] |
[
[
719,
732
]
] |
[
[
627,
717
],
[
735,
1434
]
] |
2018AandA...620A..46M__Bergin_et_al._2000_Instance_1
|
Astrochemical models have always dedicated special attention to molecular oxygen. With a cosmic abundance twice that of C, atomic O is the third most abundant element in space. In dense clouds, standard gas phase chemical models therefore suggest a comparable ratio of CO and O2 after times ≥ 3 × 105 yr (e.g. Woodall et al. 2007), where O2 is supposed to be formed especially via OH + O → O2 + H. The OH here can be formed by the dissociative recombination of H3O+, H3O+ +e− → OH + 2H. However, observations with the Submillimeter Wave Astronomy Satellite (SWAS) by Goldsmith et al. (2000) towards Orion and with Odin by Larsson et al. (2007) towards ρ Oph A showed a significant difference between model predictions and measurements. The O2 abundances found were more than 100 times smaller than those predicted by models (Goldsmith et al. 2000). Better agreement with observations can be obtained if freeze-out of O atoms onto dust grains is taken into account in gas-grain chemical models (Bergin et al. 2000; Viti et al. 2001), with consequence surface production of H2O and O2, which may trap a significant fraction of oxygen, leaving only some atomic O and CO in the gas phase. Observations conducted by Liseau et al. (2012) led to a O2 column density of N(O2) = 5.5 × 1015 cm−2 with an upper limit of abundance of N(O2)/N(H2) ∼ 5 × 10−8 in warm gas (T > 50 K) and to N(O2) = 6 × 1015 cm−2 with a little higher abundance in cold gas (T 30 K). Liseau et al. (2012) stated that detecting gas phase O2 might be so difficult because the O2 abundance is transient in ρ Oph A and O2 is no longer detectable after ∼2 × 105 yr. A relatively large amount of O2 has only been found with Herschel in Orion as reported by Goldsmith et al. (2011). This source is quite warm (≥180 K), leading to a grain temperature of ≥100 K. At this temperature the grains are warm enough to desorb H2O ice and keep a large amount of oxygen with a big fraction in the form of O2 in the gas phase. Another explanation for the high O2 abundance found by Goldsmith et al. (2000) is that low-velocity C-shocks might be responsible for the increase of molecular oxygen in the gas phase.
|
[
"Bergin et al. 2000"
] |
[
"Better agreement with observations can be obtained if freeze-out of O atoms onto dust grains is taken into account in gas-grain chemical models",
", with consequence surface production of H2O and O2, which may trap a significant fraction of oxygen, leaving only some atomic O and CO in the gas phase."
] |
[
"Motivation",
"Motivation"
] |
[
[
994,
1012
]
] |
[
[
849,
992
],
[
1031,
1184
]
] |
2015MNRAS.451.4328K__Coddington_1994_Instance_1
|
We estimate the DPL parameters in equation (9) via Monte Carlo simulations. We compare the real structure function of an AGN to simulations computed from mock light curves generated using the DPL model of equation (9). We generate ‘mock’ light curves using the Timmer & König (1995) method. To create a single mock light curve, pseudo-random numbers are generated using the Fast Mersenne Twister SFMT19937 generator seeded with hardware-generated random numbers (generated using Intel RDRAND instruction) to ensure that the random number sequences are free of artificial correlations (Coddington 1994) induced by poor random seed choices. At this intermediate stage, the mock light curve is oversampled by a factor of 10 i.e. we generate points at 10 × the required cadence in order to avoid sampling issues. To include low-frequency modes that are not adequately characterized by the length of the observed light curve, the intermediate mock light curve is much longer than is ultimately required to make the final mock; mock light curves generated in this manner are capable of exhibiting low-frequency modes longer than the length of the observed data. Fast Fourier Transforms (FFTs) are most efficient for data sequences that are a power of 2; for this reason, we pick the intermediate (including the oversampling) to be of length 223. This results in the intermediate mock light curve being between 15 × to 45 × the length required for the final mock light curve depending on the actual length of the observed light curve. We pick a uniformly distributed random segment of the intermediate overly long light curve that has the same length as the observed light curve and generate another stream of uncorrelated Gaussian random deviates to simulate the white-noise properties of Kepler instrumentation noise. After adding this ‘measurement noise’, we set data points corresponding to the unobserved cadences in the real light curve to 0. This procedure creates a final mock light curve with identical sampling and noise properties to the real light curve. Fig. 8 shows the true light curve (orange) along with an example mock light curve (light green) for the Sy 1 AGN kplr006932990 illustrating what the mock light curves look like for the best-fitting DPL parameters for this object.
|
[
"Coddington 1994"
] |
[
"To create a single mock light curve, pseudo-random numbers are generated using the Fast Mersenne Twister SFMT19937 generator seeded with hardware-generated random numbers (generated using Intel RDRAND instruction) to ensure that the random number sequences are free of artificial correlations",
"induced by poor random seed choices."
] |
[
"Uses",
"Uses"
] |
[
[
585,
600
]
] |
[
[
291,
583
],
[
602,
638
]
] |
2016MNRAS.461.1719C__Fu_et_al._2012_Instance_3
|
HATLAS12-00 had already been identified as a candidate gravitationally lensed galaxy as a result of its high submm flux (i.e. F500 > 100 mJy), red Herschel colours and the lack of a bright optical or radio counterpart (see e.g. Negrello et al. 2010 for a discussion of the selection of lens candidates in H-ATLAS and other Herschel surveys). This source was therefore observed spectroscopically in the submm. A CO spectroscopic redshift of 3.26 was first suggested by Z-spec (Bradford et al. 2004) observations, then subsequently confirmed by observations by the CARMA interferometer (Leeuw et al., in preparation) and the Zpectrometer instrument (Harris et al. 2007) on the Greenbank Telescope (Harris et al. 2012; see also Fu et al. 2012). Additional followup observations in the optical, near-IR, submm and other wavelengths were targeted at the lensed z = 3.26 source and the foreground objects responsible for the lensing, resulting in detailed analyses of this lensing system by Fu et al. (2012) and Bussmann et al. (2013). Their conclusions are that the z = 3.26 source HATLAS12-00 is subject to gravitational lensing, with a magnification of 9.6 ± 0.5 in both the submm continuum and CO, and 16.7 ± 0.8 in the K′ band, by two foreground galaxies, one at a spectroscopically determined redshift of 1.22, and another with photometry suggesting that it lies at a similar redshift. The submm photometry of HATLAS12-00 at 890 μm acquired with the Submillimeter Array (SMA) as part of this programme (Fu et al. 2012; Bussmann et al. 2013) is fully consistent with the 870 μm and 850 μm fluxes derived for this source from the LABOCA and SCUBA2 data to be presented here. The spectral energy distribution (SED) of the lensed source, after correcting for the lensing amplification, is well matched by the optically thick SED model for Arp220 from Rangwala et al. (2011), with a lensing-corrected far-IR luminosity of 1.2 ± 0.2 × 1013 L⊙, and an implied star formation rate of 1400 ± 300 M⊙ yr−1. In many ways the unlensed properties of this object match those of the broader population of bright submm selected galaxies first discovered by the SCUBA submm imager (see e.g. Chapman et al. 2005; Clements et al. 2008; Michałowski, Hjorth & Watson 2010). The unlensed 870 μm flux of this object would be ∼7.7 mJy.
|
[
"Fu et al. 2012"
] |
[
"The submm photometry of HATLAS12-00 at 890 μm acquired with the Submillimeter Array (SMA) as part of this programme",
"is fully consistent with the 870 μm and 850 μm fluxes derived for this source from the LABOCA and SCUBA2 data to be presented here."
] |
[
"Similarities",
"Similarities"
] |
[
[
1503,
1517
]
] |
[
[
1386,
1501
],
[
1541,
1672
]
] |
2020AandA...635A..47H___2005_Instance_1
|
Finally, the multiphase nature of galactic outflows implies that measurements of the outflow properties based on a single gas phase can lead to misleading conclusions (for a discussion, see e.g., Cicone et al. 2018b). Historically, systematic studies of galactic outflows in nearby and high-z galaxies have focused on the ionized gas – for example, as observed as broad wing emission in the spectra of the Hα, [O III] or Paα lines – (e.g., Heckman et al. 1990; Rupke & Veilleux 2013a; Woo et al. 2016; Harrison et al. 2016; Förster Schreiber et al. 2019; Ramos Almeida et al. 2019) and the atomic phase – based on the Na D or Mg II lines in absorption (e.g., Heckman et al. 2000; Rupke et al. 2002, 2005; Weiner et al. 2009; Roberts-Borsani & Saintonge 2019). The molecular component of outflows, on the other hand, has been much more difficult to study. Great progress was made with the Herschel Space Observatory using the OH 119 μm line in absorption to study molecular outflows in Seyfert and luminous infrared galaxies (Fischer et al. 2010; Sturm et al. 2011; Veilleux et al. 2013; Bolatto et al. 2013; Spoon et al. 2013; George et al. 2014; Stone et al. 2016; González-Alfonso et al. 2017; Zhang et al. 2018). More recently, the advent of powerful millimeter-wave interferometers such as the Atacama Large Millimeter/submillimeter Array (ALMA) and the NOrthern Extended Millimeter Array (NOEMA) are rapidly increasing the number of molecular outflows detected based on observations of the CO line (e.g., Combes et al. 2013; Sakamoto et al. 2014; García-Burillo et al. 2014; Leroy et al. 2015; Feruglio et al. 2015; Morganti et al. 2015; Dasyra et al. 2016; Pereira-Santaella et al. 2016, 2018; Veilleux et al. 2017; Fluetsch et al. 2019; Lutz et al. 2020). At high-z, so far only a handful of large-scale, molecular outflows have been studied in QSOs (e.g., Cicone et al. 2015; Vayner et al. 2017; Feruglio et al. 2017; Carniani et al. 2017; Fan et al. 2018; Brusa et al. 2018), sub-millimeter galaxies (e.g., Spilker et al. 2018), and main-sequence, star-forming galaxies (e.g., Herrera-Camus et al. 2019).
|
[
"Rupke et al.",
"2005"
] |
[
"Historically, systematic studies of galactic outflows in nearby and high-z galaxies have focused on the ionized gas",
"and the atomic phase – based on the Na D or Mg II lines in absorption (e.g."
] |
[
"Background",
"Background"
] |
[
[
680,
692
],
[
699,
703
]
] |
[
[
218,
333
],
[
582,
657
]
] |
2022AandA...663A.105P__Finoguenov_et_al._2010_Instance_1
|
Cluster radio relics are usually found in the outskirts of merging galaxy clusters. They exhibit elongated morphologies and high degrees of polarisation above 1 GHz (up to 70%, Ensslin et al. 1998; Bonafede et al. 2014; Loi et al. 2019; de Gasperin et al. 2022). The resolved spectral index in radio relics shows a gradient: it steepens towards the cluster centre and flattens towards the outskirts. Their size can reach up to ∼2 Mpc, and high-resolution observations have revealed filamentary structures within relics themselves (Di Gennaro et al. 2018; Rajpurohit et al. 2020, 2022a,b; de Gasperin et al. 2022). The Largest Linear Sizes (LLS) and radio powers of relics are correlated, as well as the integrated spectral index and the radio power (van Weeren et al. 2009b; Bonafede et al. 2012; de Gasperin et al. 2014). Relics trace ICM shock waves with relatively low (M 3) Mach numbers (Finoguenov et al. 2010; Akamatsu et al. 2013; Shimwell et al. 2015; Botteon et al. 2016). The acceleration of electrons is believed to proceed via diffusive shock acceleration (DSA) in the ICM (Ensslin et al. 1998; Roettiger et al. 1999), in which particles scatter back and forth across the shock front gaining energy at every crossing. Nevertheless, this mechanism has been shown to be rather inefficient in accelerating electrons from the thermal pool (Vazza & Brüggen 2014; Vazza et al. 2016; Botteon et al. 2020a; Brüggen & Vazza 2020; see Brunetti & Jones 2014 for a review). Recently, it has been suggested that seed electrons could originate from the tails and lobes (driven by AGN outflows) of cluster radio galaxies (Bonafede et al. 2014; van Weeren et al. 2017; Stuardi et al. 2019), which alleviates the requirements of high acceleration efficiencies at cluster shocks (e.g., Markevitch et al. 2005; Kang et al. 2012, 2017; Botteon et al. 2016; Eckert et al. 2016). In some cases, double relics have been detected on opposite sides of the cluster centre (e.g., Rottgering et al. 1997; van Weeren et al. 2010, 2012b; Bonafede et al. 2012; de Gasperin et al. 2015a). In these clusters it is possible to constrain the merger history, providing important information about the formation processes of relics.
|
[
"Finoguenov et al. 2010"
] |
[
"Relics trace ICM shock waves with relatively low (M 3) Mach numbers"
] |
[
"Background"
] |
[
[
894,
916
]
] |
[
[
823,
892
]
] |
2019MNRAS.490.3061V__Myers_2009_Instance_1
|
On the other hand, the line-shift-absence conundrum of ZP74 is easily explained through geometrical considerations. Essentially, the arguments leading to this conundrum assume that the collapse is roughly spherically symmetric and monolithic, so that the infall motions are coherent, and directed towards a single collapse centre, at the geometrical centre of the cloud. This assumption is inconsistent with our current understanding of the structure of MCs, which are known to be far from spherically symmetric, and instead consist of an intricate and inhomogeneous network of filaments and clumps within them (e.g. Bally et al. 1987; Feitzinger et al. 1987; Gutermuth et al. 2008; Juvela, Pelkonen & Porceddu 2009; Myers 2009; André et al. 2010; Henning et al. 2010; Men’shchikov et al. 2010; Molinari et al. 2010; Arzoumanian et al. 2011; Busquet et al. 2013). The central clumps (‘hubs’) appear to accrete from the filaments, while in turn the filaments seem to accrete radially from their surroundings (Schneider et al. 2010; Kirk et al. 2013; Peretto et al. 2014; Gong et al. 2018; Lu et al. 2018; Williams et al. 2018; Shimajiri et al. 2019). Thus, the geometry is far from being spherically symmetric, and therefore the accreting gas is not isotropically distributed around the collapse centres (the hubs). In addition, the velocity field is highly complex and chaotic (e.g. Gómez & Vázquez-Semadeni 2014; Zamora-Avilés, Ballesteros-Paredes & Hartmann 2017; Gómez, Vázquez-Semadeni & Zamora-Avilés 2018), so there is no reason to expect a systematic redshift of the absorption lines produced in the gas surrounding the hubs. Instead, the accretion flow is most directly observed as velocity-centroid gradients along the filaments, directed towards the hubs. Indeed, synthetic CO observations of simulations of the regime often show only marginal or no evidence for infall profiles, due to the chaotic motions and perhaps velocity crowding effects, although the line profiles do look similar to observed ones (e.g. Heitsch et al. 2009; Heiner, Vázquez-Semadeni & Ballesteros-Paredes 2015; Clarke et al. 2018). Nevertheless, recent dedicated searches for evidence of infall signatures in CO lines from GMCs have met with success. For example, Schneider et al. (2015) have found the classical combination of self-absorbed and blue-skewed optically thick lines (12CO (3–2)) together with centrally peaked optically thin (13CO (1–0)) lines, indicating collapse in the molecular gas surrounding IRDC G28.37+0.07, while Barnes et al. (2018) have measured shifts between the lines of 12CO (tracing gas in the outer parts of the cloud) and 13CO (tracing gas deeper into the cloud) in the pc-scale, massive clumps of the CHaMP survey, finding systematic velocity differentials between the two lines that imply an average mass accretion time-scale of ∼16 Myr, consistent with the time-scales we discuss in this paper (cf. Section 7.1 and Fig. 13).
|
[
"Myers 2009"
] |
[
"This assumption is inconsistent with our current understanding of the structure of MCs, which are known to be far from spherically symmetric, and instead consist of an intricate and inhomogeneous network of filaments and clumps within them (e.g."
] |
[
"Differences"
] |
[
[
717,
727
]
] |
[
[
371,
616
]
] |
2019AandA...630A..98S__Cenko_et_al._2012a_Instance_1
|
The close approach of a star to a supermassive black hole (SMBH) can lead to the destruction of the stellar body in a process known as a tidal disruption event (TDE; Hills 1975). Gravitationally bound material returns to the black hole and is accreted, giving rise to a flare whose electromagnetic signature peaks in the extreme ultraviolet (EUV) band (Rees 1988; Ulmer 1999). These flares were first detected in the soft X-ray band by ROSAT (Komossa et al. 2004; Bade et al. 1996; Komossa & Greiner 1999; Komossa & Bade 1999), later by XMM-Newton and Chandra (Esquej et al. 2007; Saxton et al. 2012a, 2017; Maksym et al. 2010; Lin et al. 2015) (see review by Komossa 2017), and also in the UV band by GALEX (Gezari et al. 2006, 2008, 2009). In recent years, large-area optical surveys have detected candidate TDEs emitting at temperatures of a few ×104 K (van Velzen et al. 2011; Cenko et al. 2012a; Gezari et al. 2012; Arcavi et al. 2014; Holoien et al. 2016), ostensibly too cool to be coming from an accretion disc (e.g. Bonning et al. 2007). This optical radiation has been interpreted as being due to reprocessing of the accretion radiation by an optically thick screen (Metzger & Stone 2016; Roth & Kasen 2018; Dai et al. 2018) or to emission from shocks (Piran et al. 2015). Super-Eddington accretion in the initial phase of the disruption causes a large-scale, radiation-driven outflow of material from the central engine (Strubbe & Quataert 2009), which Metzger & Stone (2016) showed would initially completely absorb the radiation from the central engine and convert it into optical/UV photons with an effective temperature similar to that observed. In this model, the screen density is expected to drop after a few months to the point where the inner thermal radiation would become visible, with the delay time and ratio of X-ray to optical/UV flux depending on the line of sight (Metzger & Stone 2016; Dai et al. 2018). Observationally, the evidence for differences in the X-ray and UV/optical timescales is mixed. The X-rays may have lagged the UV by ∼32 days in ASASSN-14li (Pasham et al. 2017), but broadly fell on the same timescale (Brown et al. 2017), as they did in 2MASX 0740-85 (Saxton et al. 2017). In SDSS J1201+30 the UV flux did not change, while the X-rays dropped by a factor of 100 (Saxton et al. 2012a), whereas in ASASSN-15oi the X-ray luminosity was quite low (LX ∼ 1041 ergs s−1) at the peak of the optical flare, but 200–300 days later had increased to LX ∼ 1042ergs s−1 (Holoien et al. 2018; Gezari et al. 2017).
|
[
"Cenko et al. 2012a"
] |
[
"In recent years, large-area optical surveys have detected candidate TDEs emitting at temperatures of a few ×104 K"
] |
[
"Background"
] |
[
[
881,
899
]
] |
[
[
742,
855
]
] |
2017ApJ...835..154H__Lee_et_al._2014_Instance_1
|
The simplest explanation for these metal abundances, i.e., that they reflect the yields of normal core-collapse SNe (averaged over the stellar initial mass function (IMF)), fails to predict anything like the observed stellar abundances of the extremely metal-poor stars ([Fe/H] −3). Of course, at these low metallicities, the number of progenitor SNe enriching the ISM may be small, so models typically allow for metal-poor or metal-free progenitor stars, and an arbitrary mix of progenitor stellar masses (i.e., assuming the abundances might come from just one or at most a few SNe with individual explosion and progenitor parameters fitted to the observations). However, even with these degrees of freedom, the models still often fail to explain the abundances of certain individual species at the order-of-magnitude level (see, e.g., Nomoto et al. 2006; Heger & Woosley 2010; Lee et al. 2014; Placco et al. 2015), although they undoubtedly explain many of the observed abundance ratios. For the lowest metallicity stars observed ([Fe/H] −4), and in particular for the CEMP stars, these remaining discrepancies have led to more “exotic” models with a number of free parameters, invoking a mix of normal/faint SNe and hypernovae (with variable explosion energies of ∼1051–1054 erg); jets, prior “failed explosions,” and fallback episodes; rotation and adjustable mixing layers allowing for a tunable stellar abundance profile in the progenitor stars; and pollution of the stars via companions (e.g., Tominaga et al. 2007; Ishigaki et al. 2014; Takahashi et al. 2014; Abate et al. 2015). These additions can improve the agreement with observations; however, there is still no consistent theoretical scenario that simultaneously explains most of the observed stars, and even the best-fit models for many individual stars can still have order-of-magnitude discrepancies with certain outlier elements (see Tominaga et al. 2014; Frebel et al. 2015; Placco et al. 2015, and references therein).
|
[
"Lee et al. 2014"
] |
[
"However, even with these degrees of freedom, the models still often fail to explain the abundances of certain individual species at the order-of-magnitude level (see, e.g.,",
", although they undoubtedly explain many of the observed abundance ratios."
] |
[
"Differences",
"Similarities"
] |
[
[
880,
895
]
] |
[
[
665,
837
],
[
916,
990
]
] |
2015MNRAS.446.1140T__Murray_et_al._2010_Instance_2
|
Recently, forms of feedback that are fundamentally different from SNe have been shown to be essential to galaxy formation. Murray, Quataert & Thompson (2010) analysed the dynamical effects of several forms of stellar feedback on parent molecular clouds. In their models they include momentum input from ionized gas in H ii regions, shocked stellar winds, hot gas pressure, protostellar jets and cosmic rays. Murray et al. (2010) conclude that radiation pressure (RP) on dust grains is likely to be the dominant form of feedback in star-forming galaxies. A variety of other studies have reached the same conclusions, placing the combination of RP and photoionization of gas by massive stars as the dominant mechanism for disruption of molecular clouds and internal regulation of the SF process (Indebetouw et al. 2009; Krumholz & Matzner 2009; Murray et al. 2010; Andrews & Thompson 2011; Hopkins, Quataert & Murray 2011; Lopez et al. 2011; Pellegrini, Baldwin & Ferland 2011). RP alone might also be the only mechanism that explains galactic fountains and the warm gas outflows observed in absorption in high-redshift galaxies (Murray, Ménard & Thompson 2011). In addition, recent numerical work by Krumholz & Thompson (2012) shows that radiation feedback fully accounts for the large gas velocity dispersions measured in young star clusters in the MW. There are at least three reasons why radiative feedback is an essential ingredient of the galaxy formation process. First, observations show that molecular clouds begin to disperse shortly after the O stars form and before the first SNe explode and deposit their energy into the gas (Kawamura et al. 2009). Secondly, the total energy output of a stellar cluster is dominated by radiation. The rate of radiative energy output by O and B stars is ∼200 times larger than the average power injected by SNe and stellar winds during the lifetime of the most massive stars. Thirdly, it is difficult to explain the large gas turbulence values observed in star-forming regions without including the momentum input by radiation (Murray et al. 2010).
|
[
"Murray et al. 2010"
] |
[
"A variety of other studies have reached the same conclusions, placing the combination of RP and photoionization of gas by massive stars as the dominant mechanism for disruption of molecular clouds and internal regulation of the SF process"
] |
[
"Background"
] |
[
[
843,
861
]
] |
[
[
554,
792
]
] |
2015ApJ...799..149J___2014_Instance_3
|
With our joint analysis of stellar mass fraction and source size, we find a larger stellar mass fraction than earlier statistical studies. In Figure 2, we compare our determination of the stellar surface density fraction to a simple theoretical model and to the best fit of a sample of lens galaxies by Oguri et al. (2014). The simple theoretical model is the early-type galaxy equivalent of a maximal disk model for spirals. We follow the rotation curve of a de Vaucouleurs component for the stars outward in radius until it reaches its maximum and then simply extend it as a flat rotation curve to become a singular isothermal sphere (SIS) at large radius (see details in the Appendix). The ratio of the surface mass density of the de Vaucouleurs component to the total surface mass density is shown as a dashed curve in Figure 2. We also show as a gray band the best fit for the stellar fraction in the form of stars determined by Oguri et al (2014) in a study of a large sample of lens galaxies using strong lensing and photometry, as well as the best model using a Hernquist component for the stars and an NFW halo for the dark matter with and without adiabatic contraction, also from Oguri et al. (2014). We have used the average and dispersion estimates for the Einstein and effective radii available for 13 of the objects in our sample from Oguri et al. (2014), Sluse et al. (2012), Fadely et al. (2010), and Lehár et al. (2000; see Table 1) as an estimate of RE/Reff in Figure 2. The average value and dispersion of the sample is RE/Reff = 1.8 ± 0.8. This also averages over the different radii of the lensed images. The agreement of our estimates with the expectations of the simple theoretical model and with estimates from other studies (Oguri et al. 2014) is quite good. For comparison, the estimate of Pooley et al. (2012; using the Einstein and effective radii estimates for 10 out of 14 of their objects from Schechter et al. 2014) seems somewhat lower than expected at those radii. The range of stellar mass fractions from MED09 for source sizes in the range 0.3â15.6 light days is also shown in Figure 2. In this case, the discrepancy between our estimate and their reported value of α = 0.05 is completely due to the effect of the source size. Although accretion disk sizes are known to be smaller in X-rays, recent estimates are in the range of 0.1â1 light-days, depending on the mass of the black hole (see Mosquera et al. 2013), and these finite sizes will increase the stellar surface densities implied by the X-ray data. Another possible origin for this discrepancy is that Pooley et al. (2012) use the macro model as an unmicrolensed baseline for their analysis. It is well known that simple macro models are good at reproducing the positions of images, but have difficulty reproducing the flux ratios of images due to a range of effects beyond microlensing. Recently, Schechter et al. (2014) found that the fundamental plane stellar mass densities have to be scaled up by a factor 1.23 in order to be compatible with microlensing in X-rays in a sample of lenses with a large overlap with that analyzed by Pooley et al. (2012). It is unclear how this need for more mass in stars at the position of the images found by Schechter et al. (2014) can be reconciled with the apparently low estimate of mass in stars at those radii by Pooley et al. (2012). Our estimate of the stellar mass fraction agrees better with the results of microlensing studies of individual lenses (Keeton et al. 2006; Kochanek et al. 2006; Morgan et al. 2008, 2012; Chartas et al. 2009; Pooley et al. 2009; Dai et al. 2010) that reported values in the range 8%â25%, and with the estimates from strong lensing studies (see for example Jiang & Kochanek 2007; Gavazzi et al. 2007; Treu 2010; Auger et al. 2010; Treu et al. 2010; Leier et al. 2011; Oguri et al. 2014) which produced stellar mass fractions in the range 30%â70% integrated inside the Einstein radius of the lenses.
|
[
"Oguri et al. (2014)"
] |
[
"as well as the best model using a Hernquist component for the stars and an NFW halo for the dark matter with and without adiabatic contraction, also from"
] |
[
"Uses"
] |
[
[
1193,
1213
]
] |
[
[
1039,
1192
]
] |
2021MNRAS.502.1312C__Recchia,_Blasi_&_Morlino_2016_Instance_1
|
The dynamical importance of CRs is even more uncertain. This is in part because most early work on this question focused only on galactic conditions similar to those found locally (Jokipii 1976; Badhwar & Stephens 1977; Ghosh & Ptuskin 1983; Chevalier & Fransson 1984; Boulares & Cox 1990; Ko, Dougherty & McKenzie 1991; Ptuskin 2001), and/or focused largely on the question of how and whether CRs can drive galactic winds originating in the ionized, low-density medium found several scale heights above galactic planes (Ipavich 1975; Breitschwerdt, McKenzie & Voelk 1991; Zirakashvili et al. 1996; Ptuskin et al. 1997; Zirakashvili & Völk 2006; however, for an exception see Breitschwerdt, McKenzie & Voelk 1993). More recent numerical and analytic models have continued in this vein (e.g. Everett et al. 2008; Jubelgas et al. 2008; Samui, Subramanian & Srianand 2010; Wadepuhl & Springel 2011; Uhlig et al. 2012; Booth et al. 2013; Pakmor et al. 2016; Simpson et al. 2016; Recchia, Blasi & Morlino 2016, 2017; Ruszkowski, Yang & Zweibel 2017; Pfrommer et al. 2017; Buck et al. 2019), rather than address the question of whether CRs represent a significant contribution to the support of the neutral material that dominates the total mass budget and occupies at least $\sim 50{{\ \rm per\ cent}}$ of the volume (e.g. Dekel et al. 2019) near the mid-plane. Indeed, the vast majority of published simulations that include CR transport do not resolve the neutral phase or galactic scale heights (∼100 pc), and those that do (e.g. Hanasz et al. 2013; Salem & Bryan 2014; Salem, Bryan & Corlies 2016; Chan et al. 2019) generally assume that CR transport in the neutral ISM is identical to that in the ionized ISM (though see Farber et al. 2018), an assumption that is almost certainly incorrect (e.g. Zweibel 2017; Xu & Lazarian 2017; Krumholz et al. 2020). Only a few published models attempt to address the question of CR pressure support in the neutral ISM for non-Solar neighbourhood (mostly starburst or Galactic Centre) conditions (e.g. Thompson et al. 2006; Socrates, Davis & Ramirez-Ruiz 2008; Lacki, Thompson & Quataert 2010; Lacki et al. 2011; Crocker et al. 2011; Crocker 2012; Lacki 2013; Yoast-Hull, Gallagher & Zweibel 2016; Yoast-Hull & Murray 2019; Krumholz et al. 2020).
|
[
"Recchia, Blasi & Morlino 2016"
] |
[
"More recent numerical and analytic models have continued in this vein (e.g.",
"rather than address the question of whether CRs represent a significant contribution to the support of the neutral material that dominates the total mass budget and occupies at least $\\sim 50{{\\ \\rm per\\ cent}}$ of the volume",
"near the mid-plane."
] |
[
"Background",
"Background",
"Background"
] |
[
[
975,
1004
]
] |
[
[
715,
790
],
[
1086,
1311
],
[
1337,
1356
]
] |
2021ApJ...919..140S__Bartos_et_al._2017_Instance_1
|
Resonant dynamical friction may have applications beyond the relaxation of IMBHs examined in this paper. It may affect all objects in stellar clusters much more massive than the individual constituents of the disk, if present, including massive stars, stellar mass black holes (BHs), or the center of mass of massive binaries. Furthermore, it is also expected to operate in any type of disk with a high number of particles, including active galactic nucleus (AGN) accretion disks. Previously, it has been argued that stars and BHs crossing the disk on low-inclination orbits get captured by Chandrasekhar dynamical friction into the disk (Bartos et al. 2017; Panamarev et al. 2018; Tagawa et al. 2020). An interesting implication is that, if BHs settle into the disk, they interact dynamically and form BH–BH binaries efficiently, and frequent dynamical interactions and gas effects drive the BHs to merger, producing gravitational waves (GWs) detectable by LIGO, VIRGO, and KAGRA (McKernan et al. 2014, 2018; Bartos et al. 2017; Leigh et al. 2018; Yang et al. 2019; Tagawa et al. 2020, 2021; Samsing et al. 2020). Mergers are also facilitated by Lidov–Kozai oscillations in anisotropic systems (Heisler & Tremaine 1986; Petrovich & Antonini 2017; Hamilton & Rafikov 2019). The results in this paper show that resonant dynamical friction may accelerate the capture of objects in the accretion disks by a factor proportional to the SMBH mass over the local disk mass for large orbital inclinations. Pressure and viscosity in a gaseous disk do not inhibit the orbit-averaged torque from the IMBH, which leads to realignment and the warping of the disk (Bregman & Alexander 2012). Thus, RDF may efficiently catalyze the alignment of the orbital planes of BHs even in low-luminosity AGN or Seyfert galaxies with relatively small disk masses, which may not be possible for Chandrasekhar dynamical friction. In fact, this mechanism extends the scope of the “AGN merger channel” for GW source populations even beyond low-luminosity AGN and Seyfert galaxies, as it may organize BHs into disks also in nonactive galaxies with nuclear stellar disks.
|
[
"Bartos et al. 2017"
] |
[
"Previously, it has been argued that stars and BHs crossing the disk on low-inclination orbits get captured by Chandrasekhar dynamical friction into the disk"
] |
[
"Background"
] |
[
[
639,
657
]
] |
[
[
481,
637
]
] |
2015AandA...573A.138B__Deubner_(1975)_Instance_1
|
Solar-like oscillations are excited stochastically by motions in the convective envelope of stars. For low-mass stars, they are found in all evolutionary states, between the main sequence and horizontal branch of helium-core burning stars (e.g. Leighton et al. 1962; Frandsen et al. 2002; Carrier et al. 2003; Hekker et al. 2009; Chaplin et al. 2011; Huber et al. 2011; Kallinger et al. 2012; Mosser et al. 2013) and were even detected in the M5 super giant α Her (Moravveji et al. 2013). These very characteristic oscillations lead to a nearly regular spaced comb-like pattern in the power spectrum. It was shown by Deubner (1975) that this ridge structure is governed by the degree of oscillation modes and resembles the predictions made by Ando & Osaki (1975). Empirically, the frequency patterns of solar-like oscillations were described through scaling relations by Kjeldsen & Bedding (1995). Since then, these relations have been tested and revised from large sample studies based on high-precision space photometry of red giants in clusters and in eclipsing binaries besides single stars (e.g. Corsaro et al. 2012b; Frandsen et al. 2013; Kallinger et al. 2010). Indications of non-radial oscillation modes were found in the variations of the absorption lines of bright red giants (Hekker et al. 2006; Hekker & Aerts 2010), but were firmly established in a large set of red giants observed with the CoRoT satellite (De Ridder et al. 2009). The identification of dipole mixed modes extended the sensitivity of the seismic analyses towards the core of evolved stars also (Beck et al. 2011; Bedding et al. 2011; Mosser et al. 2011). The analysis of solar-like oscillations enabled us to unravel many open questions on stellar structure and evolution, such as constraining the internal rotational gradient (Elsworth et al. 1995; Beck et al. 2012; Deheuvels et al. 2012) or determining the evolutionary status in terms of nuclear burning of a given red giant star (Bedding et al. 2011; Mosser et al. 2011).
|
[
"Deubner (1975)"
] |
[
"It was shown by",
"that this ridge structure is governed by the degree of oscillation modes and resembles the predictions made by Ando & Osaki (1975)."
] |
[
"Background",
"Background"
] |
[
[
617,
631
]
] |
[
[
601,
616
],
[
632,
763
]
] |
2017AandA...606A.113G__Carollo_et_al._(2013)_Instance_2
|
In this context, massive (ℳ>1011 M⊙) PGs (MPGs) deserve particular attention. These systems are expected to evolve mainly through (dry) mergers (e.g. Hopkins et al. 2009; De Lucia & Blaizot 2007). If this is the case, in this mass range we should detect a stronger signal of the size-growth with respect to a lower mass range. So far, because MPGs are extremely rare, there have been very few studies that have investigated the combined evolution of the number density and of the age of MPGs as a function of their compactness (Carollo et al. 2013; Fagioli et al. 2016). Carollo et al. (2013) found that the number density of massive quiescent and elliptical galaxies with Re 2.5 kpc decreases by about 30% from z~1 to z~0.2 and that their U−V colours are consistent with passive evolution. They concluded that the driving mechanism for the average size-growth of the whole population is the appearance at later epochs of larger quiescent galaxies. More recently, Fagioli et al. (2016, hereafter F16) analysed the spectroscopic properties of ~500 MPGs (defined as galaxies with absent or very weak emission lines and no MIPS detections) at 0.2 z 0.8 in the zCOSMOS-bright 20 K catalogue (Lilly et al. 2007). From the analysis of stacked spectra of small and large MPGs, they dated the stellar content of these groups and found that the two sub-populations have similar ages. The authors concluded that, in this mass regime, the size growth of individual galaxies through dry mergers is the most likely explanation for the increase in the mean effective radius of the whole population. A recent analysis by Zahid et al. (2016) on the physical properties of compact post starburst galaxies at 0.2z0.8 with ℳ>1011 M⊙ provides new insights. On the basis of both their number density and of their ages, which have been found to be 1 Gyr, the authors suggest that this class of objects are the progenitors of compact quiescent galaxies. They conclude that a substantial fraction of dense quiescent galaxies at z0.8 are newly formed.
|
[
"Carollo et al. (2013)"
] |
[
"found that the number density of massive quiescent and elliptical galaxies with Re 2.5 kpc decreases by about 30% from z~1 to z~0.2 and that their U−V colours are consistent with passive evolution. They concluded that the driving mechanism for the average size-growth of the whole population is the appearance at later epochs of larger quiescent galaxies."
] |
[
"Background"
] |
[
[
571,
592
]
] |
[
[
593,
948
]
] |
2022MNRAS.515.2914E__Lada_&_Lada_2003_Instance_1
|
Also addressing the evolution of a stellar cluster is work by Parker & Goodwin (2009) and Parker et al. (2009) in which N-body simulations are used to analyse the prevalence of planets susceptible to the Kozai effect from a binary companion, and the stability of binaries in a dense cluster environment respectively. Parker & Goodwin (2009) found that around $20{{\ \rm per\ cent}}$ of all exoplanets should at one point in their lives be in the presence of a binary companion that has been sufficiently inclined by its cluster environment such that Kozai cycles can occur. This is an intriguing finding that bolsters the idea that a binary companion perturbed by the cluster can have an appreciable effect on the planets, with the caveat that the authors only examined the evolution of binary orbits, as their simulations lacked planetary bodies. Additionally, they particularly focused on very dense clusters similar to Orion, with a half-mass radius of only 0.1 pc. This is not typical for embedded clusters, which have typical half mass radii of ∼0.8 pc (Lada & Lada 2003), and therefore would have less frequent interactions between cluster stars and a particular binary. Nevertheless, they showed that a dense cluster environment can significantly alter the architecture of a stellar binary which in turn can affect a protoplanetary disc or mature planet system. Parker et al. (2009) focused on similarly dense cluster environments, but instead explored the longevity of moderately wide (${\sim }10^3 \, \text{au}$) to ultrawide (${\gt}10^4 \, \text{au}$) binaries. They found that cluster environments strip away all ultra-wide binary companions, and that the denser clusters, with half mass radii of 0.1–0.2 pc do not retain any binaries with separations ${\gt}10^3 \, \text{au}$. The less dense clusters, with half mass radii 0.4–0.8 pc, do retain some of these moderately wide binaries. The authors noted that as ultrawide binaries are often stripped in only a few cluster crossing times, these very separated binaries may form in isolation An alternate channel is that ultrawide binaries are formed during cluster dissolution (Kouwenhoven et al. 2010).
|
[
"Lada & Lada 2003"
] |
[
"This is not typical for embedded clusters, which have typical half mass radii of ∼0.8 pc",
", and therefore would have less frequent interactions between cluster stars and a particular binary."
] |
[
"Differences",
"Differences"
] |
[
[
1059,
1075
]
] |
[
[
969,
1057
],
[
1076,
1176
]
] |
2020MNRAS.495.4508E__Heinke_et_al._2014_Instance_1
|
Several qLMXBs have been identified in GCs and in the Galactic field (for some examples, see table 4 in Guillot et al. 2009 and references therein). While LMXBs in the field were detected following the onset of a bright accretion outburst, most qLMXBs in GCs, including all those with the highest flux at Earth, have not shown accretion activity.3 Most of these sources have only been spectrally identified based on their similarities to field LMXBs, observed during quiescence (e.g. Cen X-4 or Aql X-1). Previous works have confirmed that H-atmosphere models accurately describe the spectra of qLMXBs, with radii in the range 10–15 km, as expected for NSs, either from single sources (e.g. Heinke et al. 2006a; Webb & Barret 2007; Guillot, Rutledge & Brown 2011; Heinke et al. 2014; Bogdanov et al. 2016), or from statistical analyses of multiple qLMXBs (e.g. Guillot et al. 2013; Guillot & Rutledge 2014; Lattimer & Steiner 2014; Guillot 2016; Steiner et al. 2018). However, in some cases the accreted material may not be hydrogen, but helium (e.g. Servillat et al. 2012; Catuneanu et al. 2013; Heinke et al. 2014). One way to circumvent this is to identify the nature of the donor star, i.e. to determine the nature of the material transferred on to the NS (e.g. with the detection of an H α emission line, presumably originating in a faint accretion disc, Haggard et al. 2004). The possibility of helium (or heavier element) atmospheres is well-founded on the existence of ultracompact X-ray binaries (UCXB), with white dwarfs or helium-dominated donors4 (e.g. Zurek et al. 2009; Altamirano et al. 2010; Sanna et al. 2017; Cadelano et al. 2019). In fact, around 1/3 of the LMXBs in GCs with constraints on the companion nature, possess a white dwarf donor (Bahramian et al. 2014). Since NS He-atmosphere models have harder spectra than H-atmosphere models, using the incorrect composition for the observed thermal emission can result in biases of the inferred radii (Servillat et al. 2012; Heinke et al. 2014).
|
[
"Heinke et al. 2014"
] |
[
"Previous works have confirmed that H-atmosphere models accurately describe the spectra of qLMXBs, with radii in the range 10–15 km, as expected for NSs, either from single sources (e.g."
] |
[
"Background"
] |
[
[
764,
782
]
] |
[
[
505,
690
]
] |
2022MNRAS.511.5797M__Liu_&_Lai_2019_Instance_1
|
A variety of formation channels have been proposed for binary black holes (BBHs; see e.g. Mapelli 2021 for a recent review): BBH mergers can be the outcome of isolated binary evolution via common envelope (Tutukov & Yungelson 1973; Bethe & Brown 1998; Portegies Zwart & Yungelson 1998; Belczynski, Kalogera & Bulik 2002; Belczynski et al. 2008, 2016a; Dvorkin et al. 2016, 2018; Eldridge & Stanway 2016; Mapelli et al. 2017, 2019; Stevenson, Berry & Mandel 2017; Kruckow et al. 2018; Spera et al. 2019; Belczynski et al. 2020; Klencki et al. 2021; Olejak, Belczynski & Ivanova 2021; Tanikawa et al. 2021a), stable mass transfer (Giacobbo, Mapelli & Spera 2018; Neijssel et al. 2019; Bavera et al. 2021; Gallegos-Garcia et al. 2021; Shao & Li 2021), or chemically homogeneous evolution (de Mink & Mandel 2016; Mandel & de Mink 2016; Marchant et al. 2016; du Buisson et al. 2020; Riley et al. 2021). Alternatively, BBHs can form dynamically in triples (e.g. Antonini, Toonen & Hamers 2017; Silsbee & Tremaine 2017; Fragione & Silk 2020; Arca Sedda, Li & Kocsis 2021a; Vigna-Gómez et al. 2021), multiples (e.g. Fragione & Kocsis 2019; Liu & Lai 2019, 2021; Hamers & Safarzadeh 2020), young star clusters (YSCs; Banerjee, Baumgardt & Kroupa 2010; Mapelli 2016; Banerjee 2017, 2021; Di Carlo et al. 2019, 2020a; Kumamoto, Fujii & Tanikawa 2019, 2020), globular clusters (GCs; Portegies Zwart & McMillan 2000; Tanikawa 2013; Samsing, MacLeod & Ramirez-Ruiz 2014; Rodriguez, Chatterjee & Rasio 2016; Askar et al. 2017; Fragione & Kocsis 2018; Hong et al. 2018; Choksi et al. 2019; Kamlah et al. 2022), and nuclear star clusters (NSCs; Antonini & Rasio 2016; Petrovich & Antonini 2017; Antonini, Gieles & Gualandris 2019; Arca Sedda 2020; Arca Sedda et al. 2020; Fragione, Loeb & Rasio 2020). Furthermore, gas torques in active galactic nucleus (AGN) discs trigger the formation of BBHs and speed up their mergers (e.g. Bartos et al. 2017; Stone, Metzger & Haiman 2017; McKernan et al. 2018; Yang et al. 2019; Ishibashi & Gröbner 2020; Tagawa, Haiman & Kocsis 2020). Finally, primordial black holes (BHs), born from gravitational collapses in the early Universe, might also pair up and merge via gravitational wave (GW) emission (e.g. Carr & Hawking 1974; Carr, Kühnel & Sandstad 2016; Sasaki et al. 2016; Ali-Haïmoud, Kovetz & Kamionkowski 2017; Clesse & García-Bellido 2017; De Luca et al. 2021).
|
[
"Liu & Lai 2019"
] |
[
"Alternatively, BBHs can form dynamically in",
"multiples (e.g."
] |
[
"Background",
"Background"
] |
[
[
1132,
1146
]
] |
[
[
898,
941
],
[
1092,
1107
]
] |
2015AandA...584A..76S__Bernstein_et_al._(1995)_Instance_2
|
The only gas-phase process, which is predicted to efficiently lead to products, is the process involving ionized methanimine. According to the model by Vuitton et al. (2007), the amount in the upper atmosphere of Titan of ionized methanimine is small, but not negligible. The products of the reaction CH2NH + CH2NH+ all have a mass-to-charge ratio of 57, where an important contribution is given by the abundant carbocation C4H\hbox{$_{9}^{+}$}+9. In this condition, it is difficult to see if any of these species is present in small amounts in the ionosphere of Titan. Interestingly, we have also investigated the further reaction of 1+ with a third molecule of methanimine. The formation of the species (CH2NH)\hbox{$_{3}^{+}$}+3, starting from two molecules of methanimine and the ionic species CH2NH+, is a strongly exothermic reaction, being \hbox{$\Delta H_{0}^{\circ} = -305$}ΔH0◦=−305 kJ/mol at CCSD(T) level. The optimized structure of (CH2NH)\hbox{$_{3}^{+}$}+3 is shown in Fig. 10. We can conclude, therefore, that polymerization of methanimine in the gas-phase at low temperatures may well be initiated by the presence of an ionized molecule. As for the experiment on ice by Bernstein et al. (1995), it is well known that reaction barriers possibly present in gas-phase reactions are not significantly reduced when moving to ice-mediated reactions (see, for instance, Rimola et al. 2014). Normally, the tunneling effect is invoked to explain the observed reactivity, but in this case the reaction barriers are so high that it is difficult to think that a reaction sequence starting with dimerization of neutral methanimine molecules can account for the observed formation of hexamethylenetetramine or polymethylenimine. The reaction must start by involving a radical or an ionized species and not two neutral closed shell molecules. This was already noted by Vinogradoff et al. (2012), who suggested that the reaction between two neutral methanimine molecules is mediated by formic acid, which acts as a proton donor, and by Cottin et al. (2001), who irradiated mixed ice with protons. Notably, Vinogradoff et al. (2013) failed to see hexamethylenetetramine and polymethylenimine formation starting from pure CH2=NH ice. Since in the experiment by Bernstein et al. (1995) the ice was irradiated by VUV photons at the Lyman alpha wavelength, and we now know that methanimine can be efficiently ionized by those photons, we can also argue that ionization of several methanimine molecules can instead trigger the process in cold interstellar ices. This statement is in line with what is known for gas-phase polymerization of olefin (El-Shall 2008; Cottin et al. 2001). Alternatively, external strong energy sources that can induce local nonequilibrium conditions, could promote neutral-neutral dimerization. For instance, in a study by Zhou et al. (2010), in which acetylene ices were irradiated with energetic electrons, electronically excited acetylene molecules were invoked to account for the experimental observation of benzene formation.
|
[
"Bernstein et al. (1995)"
] |
[
"Since in the experiment by",
"the ice was irradiated by VUV photons at the Lyman alpha wavelength, and we now know that methanimine can be efficiently ionized by those photons, we can also argue that ionization of several methanimine molecules can instead trigger the process in cold interstellar ices."
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
2260,
2283
]
] |
[
[
2233,
2259
],
[
2284,
2556
]
] |
2020MNRAS.498.1319M__MacGregor_et_al._2017_Instance_1
|
While most systems with exo-Kuiper belts do not have known planetary mass companions, in a few of these it has been possible to directly image one, thus enabling the study of planet–disc interactions in more detail. There are well-known examples such as β Pic with a massive planet possibly warping the disc (Mouillet et al. 1997; Lagrange et al. 2012, 2019; Matrà et al. 2019); HR 8799 with four giant planets creating a scattered disc and possibly replenishing its warm dust closer in (e.g. Marois et al. 2010; Booth et al. 2016; Zurlo et al. 2016; Read et al. 2018; Wilner et al. 2018; Geiler et al. 2019; Faramaz et al., in preparation); HD 95086’s axisymmetric disc implying a low eccentricity of its 4MJup planet (Rameau et al. 2016; Su et al. 2017); and Fomalhaut having a narrow and eccentric planetesimal belt (Kalas, Graham & Clampin 2005; Acke et al. 2012; Boley et al. 2012; MacGregor et al. 2017), implying that its candidate companion on an eccentric orbit has a low mass (∼Earth or super-Earth) and is not sculpting the belt (Quillen 2006; Kalas et al. 2008; Chiang et al. 2009; Beust et al. 2014; Faramaz et al. 2015), or is not a compact object but rather the dusty aftermath of a recent planetesimal collision (Gaspar & Rieke 2020). Some exo-Kuiper belt host systems even have companions in the brown dwarf (BD) or low stellar mass regime, suggesting that their likely formation through gravitational instability (Boss 1997, 2003, 2011; Vorobyov 2013) is compatible with the formation of massive Kuiper belt analogues, e.g. HR 2562 (Konopacky et al. 2016), HD 193571 (Musso Barcucci et al. 2019), HD 92536 (Launhardt et al. 2020), and HD 206893 (Milli et al. 2017). The last one is the subject of this paper. For even more massive companions, Yelverton et al. (2019) found a significant lower detection rate of debris discs around binaries, with no discs detected in binaries with separations between 25 and 135 au (comparable to typical debris disc radii; Matrà et al. 2018b). This is likely due to dynamical perturbation inhibiting planetesimal formation or clearing any debris disc formed near those separations.
|
[
"MacGregor et al. 2017"
] |
[
"While most systems with exo-Kuiper belts do not have known planetary mass companions, in a few of these it has been possible to directly image one, thus enabling the study of planet–disc interactions in more detail. There are well-known examples such as",
"and Fomalhaut having a narrow and eccentric planetesimal belt"
] |
[
"Background",
"Background"
] |
[
[
887,
908
]
] |
[
[
0,
253
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[
757,
818
]
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2017AandA...599A..97H__Gratton_et_al._2012_Instance_1
|
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.
|
[
"Gratton et al. 2012"
] |
[
"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.,"
] |
[
"Background"
] |
[
[
1342,
1361
]
] |
[
[
996,
1318
]
] |
2022MNRAS.516.6194C__Tran_et_al._2001_Instance_1
|
In this paper, we focus on the bright elliptical galaxy NGC 5813. This is the central dominant member of a subgroup, hereafter referred to as the NGC 5813 group with an extensive diffuse X-ray emission. While part of the well-isolated NGC 5846 group ($z$ = 0.006578; e.g. Mahdavi, Trentham & Tully 2005a; Machacek et al. 2011) with a projected separation of ∼740 kpc, the two show no signs of an interaction between them. NGC 5813 itself has been assumed to be dynamically old (Emsellem et al. 2007) with no evidence of a recent major merger in its history, as indicated by the lack of any significant disturbances to its dusty circumnuclear disc (Tran et al. 2001). Regarding its atmosphere, the NGC 5813 group has a rather well defined regular morphology, consisting of three collinear pairs of cavities and associated shock fronts that are products of three distinct outburst events in the central AGN’s history (Randall et al. 2011). Such cavities often appear as a result of AGN-ICM or IGrM interaction with the radio lobes produced by the AGN, displacing the surrounding X-ray gas (e.g. McNamara & Nulsen 2012; Barai et al. 2016; Yang, Gaspari & Marlow 2019; Gastaldello et al. 2021). It has been estimated that an energy of 1.5 × 1056 and $4\times 10^{57}\,{\rm erg}$ has been released by the most and second most recent outbursts, respectively, suggesting that the most recent outburst might still be ongoing. These features are correlated with the group’s observed radio emission (Giacintucci et al. 2011), and can serve as clear indicators of AGN feedback being present. Previous studies of the group have also found moderate turbulent velocities of ∼175 km s−1 and a 3D Mach number of the order of 0.4, using resonant scattering (e.g. de Plaa et al. 2012; Ogorzalek et al. 2017), further indicating the presence of such a mechanism. These characteristics make NGC 5813, a very promising candidate for the study of the effects AGN feedback can have on the distribution of elements in a low-mass system. Additionally, this target has some of the deepest Chandra data of any galaxy group and, similar to M49 has been found to have an explicit anticorrelation between its metal abundance and the location of its radio lobes (Randall et al. 2015). That work already suggests that this result is dependent on the Fe-bias; employing a two-temperature fit brings the abundance in the region of the extended radio lobes in better agreement to their surrounding medium, although the central region appears to remain under-enriched.
|
[
"Tran et al. 2001"
] |
[
"NGC 5813 itself has been assumed to be dynamically old",
"with no evidence of a recent major merger in its history, as indicated by the lack of any significant disturbances to its dusty circumnuclear disc"
] |
[
"Background",
"Background"
] |
[
[
648,
664
]
] |
[
[
422,
476
],
[
500,
646
]
] |
2018AandA...620A..80M__Coutens_et_al._2018_Instance_2
|
The comparison with other hot corinos is not easy since we have a limited number of lines, in contrast to well-known sources that have been more extensively observed with ALMA and NOEMA, such as IRAS 16293-2422 (Jørgensen et al. 2016), NGC1333 IRAS 2A and 4A (Taquet et al. 2015; López-Sepulcre et al. 2017). In general, we find a similar inventory ofCOMs but with lower abundances in B1b-S. We are going to discuss only some trends and ratios between them. For example, one similarity with hot corinos are the higher abundances of O-bearing COMs with respect to N-bearing COMs. We observe a similar trend in B1b-S, where the highest abundances are obtained for CH3OCOH, CH3OCH3, and CH3CHO, while NH2CN and NH2CHO show the lowest abundances. Cyanamide has been recently detected towards the low-mass protostars IRAS 16293-2422 and NGC 1333 IRAS2A (Coutens et al. 2018), with observed NH2CN to NH2CHO ratios of 0.2and 0.02, respectively. These values are in the range of those observed towards the molecular clouds in the Galactic centre but lower than in Orion KL (Coutens et al. 2018). We obtain in B1b-S anabundance ratio of 0.25, similar to IRAS 16293-2422. Of the three possible isomers of C2 H4O2, we have detected CH3OCOH and CH2OHCHO. The observed CH3OCOH to CH2OHCHO ratio in B1b-S, ~20, is similar to that observed in the low-mass protostars in NGC 1333 and IRAS 16293-2422 (Taquet et al. 2015; Jørgensen et al. 2012). Acetic acid (CH3COOH) is the most stable but the least abundant of the three isomers (Lattelais et al. 2010). It has been observed in IRAS 16293-2422 with a ratio with respect to glycolaldehyde of ~11 (Jørgensen et al. 2016), consistent with the 5–15 upper limits in B1b-S. Glycolaldehyde and its corresponding alcohol, ethylene glycol, show similar abundances in B1b-S, slightly higher for CH2OHCHO. This is in contrast to other hot cores and hot corinos (Fuente et al. 2014; Jørgensen et al. 2016; Favre et al. 2017). However, since detections of these species are based on just a few lines, in particular for glycolaldehyde, column densities may be not well constrained and it is difficult to make definite conclusions. We have detected CH3CH2OCOH for the first time in a low-mass protostar, with a relatively high abundance of 10−11. This detection shows that the molecular complexity is high in young hot corinos and other species with a larger number of atoms could be present. It is noteworthy that this species has not been previously observed in the well-known hot corinos IRAS 16293-2422 and NGC 133IRAS4A. We have checked in the ALMA spectral line survey of IRAS 16293-2422 (PILS, Jørgensen et al. 2016), and no clear features are seen, although the observed frequencies are different. It is possible that the high level of line confusion and wide linewidths prevent the detection of weak lines in these more evolved sources.
|
[
"Coutens et al. 2018"
] |
[
"These values are in the range of those observed towards the molecular clouds in the Galactic centre but lower than in Orion KL"
] |
[
"Compare/Contrast"
] |
[
[
1066,
1085
]
] |
[
[
938,
1064
]
] |
2020MNRAS.497L..56Y__Verde,_Treu_&_Riess_2019_Instance_1
|
Over the past decade, direct measurements of the Hubble constant have achieved few percent in precision (Freedman 2017). Among the conducted measurements, the Supernova H0 for the Equation of State (SH0ES) team challenged the well-believed Hubble constant (H0) value inferred from the Planck cosmic microwave background (CMB) measurements assuming a flat Lambda cold dark matter (ΛCDM) model. In detail, the latest result from SH0ES is H0 = 74.03 ± 1.42 km s−1 Mpc−1 (Riess et al. 2019), which differs from thePlanck result H0 = 67.4 ± 0.5 km s−1 Mpc−1 by 4.4σ (Planck Collaboration VI 2018). Recently, the Carnegie-Chicago Hubble Program (CCHP) also presented a new and independent determination of H0 parameter based on a calibration of the Tip of the Red Giant Branch (TRGB) applied to Type Ia supernovae (SNIa) (Freedman et al. 2019, 2020). They find a value of H0 = 69.6 ± 2.5 km s−1Mpc−1, which is in the middle and consistent with the SH0ES and Planck values. A different analysis and calibration of the TRGB method was performed by the SH0ES team (Reid, Pesce & Riess 2019; Yuan et al. 2019) resulting in a slightly higher H0 value. To come to a robust conclusion, independent H0 probes with accuracy better than $2{{\ \rm per\ cent}}$ are crucial (Verde, Treu & Riess 2019). Among the possible independent probes, the time-delay strong lensing (TDSL) measurements, such as from the H0 Lenses in COSMOGRAIL’s Wellspring (H0LiCOW) collaboration (Bonvin et al. 2017; Suyu et al. 2017; Birrer et al. 2019; Wong et al. 2019), are the most precise to date. The latest constraint from a joint analysis of six gravitationally lensed quasars with measured time delays (Wong et al. 2019) indicates for a flat ΛCDM, $H_0=73.3^{+1.7}_{-1.8}$ km s−1 Mpc−1, a $2.4{{\ \rm per\ cent}}$ precision measurement, which are in agreement with local measurements from SNIa, but in 3.1σ tension with CMB. The forecasts of future 40 TDSL measurements suggest that the H0 would be constrained at $\mathcal {O}(1){{\ \rm per\ cent}}$ level (Shajib, Treu & Agnello 2018; Yildirim, Suyu & Halkola 2019). A more optimistic forecast for dark energy studies can be found in Shiralilou et al. (2019). On the other hand, one of the main obstacles for the lensing mass modelling, or to determine precise H0 value, is the mass-sheet degeneracy (Schneider & Sluse 2014; Xu et al. 2016). These issues in the H0LiCOW analysis have been frequently discussed in the literature (Sonnenfeld 2018; Kochanek 2019; Pandey, Raveri & Jain 2019). No direct evidence of bias or errors is found from a comparison of self-consistency among the individual lenses (Millon et al. 2019; Liao et al. 2020). Considering the fact that both the Planck and H0LiCOW’s H0 values are based on general relativity (GR) plus ΛCDM model, it inspires us to question the concordance cosmology model and investigate the modified gravity (MG).
|
[
"Verde, Treu & Riess 2019"
] |
[
"To come to a robust conclusion, independent H0 probes with accuracy better than $2{{\\ \\rm per\\ cent}}$ are crucial"
] |
[
"Motivation"
] |
[
[
1257,
1281
]
] |
[
[
1141,
1255
]
] |
2022MNRAS.517.2801W__Gallo_et_al._2014_Instance_2
|
During the X-ray ‘hard’ state, the radio and X-ray behaviour of BHXBs is correlated and has been studied in depth for many sources using quasi-simultaneous observations (e.g. Corbel et al. 2003; Gallo, Fender & Pooley 2003; Coriat et al. 2011; Corbel et al. 2013), and is known as the radio:X-ray plane. It was thought that all XBs followed a relation in the form of $L_{\rm Radio} \propto L_{\rm X-ray}^{0.6}$, based upon early observations of GX 339−4 (Hannikainen et al. 1998; Gallo et al. 2003; Corbel et al. 2003, 2013). This relation extends down to very low luminosities, i.e. into quiescence (Corbel et al. 2003, 2013; Plotkin et al. 2017; Tremou et al. 2020), and has been observed in other sources such as V404 Cygni (e.g. Corbel, Koerding & Kaaret 2008) and XTE J1118+480 (Gallo et al. 2014). However, further observations revealed the presence of another population of BHXBs which are less radio luminous than this relation, a so-called ‘radio-quiet’ branch, which followed $L_{\rm R} \propto L_{\rm X}^{1.4}$, such as H 1743−322 (e.g. Coriat et al. 2011; Williams et al. 2020). In some of these ‘radio-quiet’ objects, they are then found to re-join the ‘radio-loud’ branch when they go back into quiescence (Coriat et al. 2011; Carotenuto et al. 2021). The underlying cause of the split tracks (see Gallo, Miller & Fender 2012; Gallo et al. 2014; Gallo, Degenaar & van den Eijnden 2018, for a clustering analysis into the statistical robustness of this split) for BHXBs is not known, but it may be due to differences in the radiative efficiency of the accretion flow (Coriat et al. 2011; Koljonen & Russell 2019), an inclination effect of the source (Motta, Casella & Fender 2018), differences in the accretion disc contribution (Meyer-Hofmeister & Meyer 2014) or changes in the magnetic field (Casella & Pe’er 2009). For the purposes of this paper, we will refer to the original $L_{\rm R} \propto L_{\rm X}^{0.6}$ correlation sources as ‘radio-loud’ objects, and those that diverge on to the $L_{\rm R} \propto L_{\rm X}^{1.4}$ track as ‘radio-quiet’ sources.
|
[
"Gallo et al. 2014"
] |
[
"The underlying cause of the split tracks (see",
"for a clustering analysis into the statistical robustness of this split) for BHXBs is not known"
] |
[
"Background",
"Background"
] |
[
[
1341,
1358
]
] |
[
[
1266,
1311
],
[
1400,
1495
]
] |
2016ApJ...818..141V__Sellwood_2014b_Instance_1
|
Historically, the study of orbits in potentials has focused on periodic orbits. In systems like disk galaxies small perturbations to closed periodic orbits (e.g., the epicyclic and vertical perturbations of circular orbits) provided a good analytic description of most orbits. Self-consistent distribution functions are thought to be “parented” by stable periodic orbits (Arnold 1978). Early works (e.g., Contopoulos & Papayannopoulos 1980) identified and characterized the stability properties of the periodic orbit families in rapidly rotating bars. The most important periodic families in two-dimensional bars were identified as the prograde x1 family, which is elongated along the major axis of the bar, and the prograde stable x2 and unstable x3 families, which are elongated perpendicular to the bar (primarily found at small radii). The retrograde x4 (stable) orbit family is also elongated perpendicular to the bar at small radii, but becomes rounder as it extends to larger radii (for detailed description of orbit families and how they are identified see Contopoulos & Grosbol 1989; Sellwood & Wilkinson 1993; Binney & Tremaine 2008; Sellwood 2014b). In the frame of reference rotating with the bar, all of these families are characterized by a 1:2 resonance between the tangential oscillation frequency (Ωϕ) and the radial or epicyclic frequency (ΩR). Indeed, studies of orbits in 2D N-body bars largely confirmed the picture arising from the study of periodic orbits and showed that many regular orbits elongated along the bar were parented by x1 orbits, a small fraction were parented by retrograde x4 orbits (Sparke & Sellwood 1987), and none were parented by prograde x2 orbits. The realization that bars can also undergo buckling instabilities (Combes & Sanders 1981; Raha et al. 1991), which makes them develop substantial vertical thickness and peanut-shaped morphologies, led to the study of periodic orbits in three-dimensional bars (Pfenniger 1984; Martinet & de Zeeuw 1988; Pfenniger & Friedli 1991; Skokos et al. 2002a, 2002b). It was shown that the appearance of specific morphological features in images of bars, such as the X-shape and peanut features seen in edge-on bars and the boxy/rectangular isophotes and “ansae” of face-on bars, could be explained by orbits trapped around specific periodic orbit families (Patsis et al. 2002, 2003, 2010). The introduction of a third dimension did not drastically change the picture of the nature of periodic orbits, and it was found that 3D bars are composed primarily of vertical bifurcations (resonances) of the x1 family and a few additional families (e.g., Pfenniger & Friedli 1991; Skokos et al. 2002a, 2002b). Most studies of periodic orbits in analytic potentials consider prograde x2 orbits (but not retrograde x4 orbits) to be another fundamental building block of bars (Skokos et al. 2002a; Binney & Tremaine 2008). In our study we do not find any orbits parented by the periodic prograde x2 orbit in our initial bar model, but we do find orbits that are parented by the periodic retrograde x4 orbit—a result that is consistent with previous studies (Sparke & Sellwood 1987; Pfenniger & Friedli 1991; Voglis et al. 2007).
|
[
"Sellwood 2014b"
] |
[
"The retrograde x4 (stable) orbit family is also elongated perpendicular to the bar at small radii, but becomes rounder as it extends to larger radii (for detailed description of orbit families and how they are identified see",
"In the frame of reference rotating with the bar, all of these families are characterized by a 1:2 resonance between the tangential oscillation frequency (Ωϕ) and the radial or epicyclic frequency (ΩR)."
] |
[
"Background",
"Background"
] |
[
[
1144,
1158
]
] |
[
[
840,
1064
],
[
1161,
1362
]
] |
2020MNRAS.498.2575W__Tampo_et_al._2020_Instance_1
|
The sample of DES RETs shows a preference for low-metallicity, strongly star-forming host environments. The PDF of their metallicities displays a strong similarity to the hosts of SESNe, as well as LGRBs. There is a clear difference to the PDF of SNe II, which follow SDSS field galaxies. The preference for low-metallicity systems is not as strong as for LGRBs or SLSNe, but the highest metallicities found in all three samples are very similar at around solar metallicity. This result is suggestive of a stripped-envelope, massive-star origin for RETs. The population of RET hosts lies, on average, between CCSNe and LGRBs/SLSNe in terms of both star formation and metallicity. A loose correlation exists between the luminosity and rarity of events, and the host galaxy conditions required for their formation − on average, rarer events occur in more extreme environments. The approximate rate of RETs (≥10−6 Mpc−3 yr−1; Drout et al. 2014; P18; Coppejans et al. 2020; Ho et al. 2020; Tampo et al. 2020, although the definition of RET varies in the above calculations) is ${\sim}1{{\ \rm per\ cent}}$ of the CCSN rate (Horiuchi et al. 2011; Li et al. 2011; Strolger et al. 2015), which itself is divided into the more common SNe II and sub-dominant SESNe (Kelly & Kirshner 2012; Frohmaier et al., submitted). At ${\sim}1{{\ \rm per\ cent}}$ of the CCSN rate, RETs are more common than SLSNe (${\sim}0.01{-}0.05{{\ \rm per\ cent}}$ of CCSNe; McCrum et al. 2015; Prajs et al. 2017; Frohmaier et al., in preparation) and LGRBs (intrinsically ${\sim}0.08{{\ \rm per\ cent}}$ when accounting for beaming; Graham & Schady 2016). These figures place the rate of DES RETs between extreme objects (SLSNe, LGRBs) and more common SNe (SNe II, SESNe) in terms of rate, matching the location of RET hosts in the various host galaxy parameter spaces presented in Section 5. While stressing rates are uncertain and host galaxy parameters span wide ranges for all transients, they are both linked to the respective transients’ progenitor channels. While it is likely that RETs are a heterogeneous population comprising several progenitor scenarios, it is reasonable to infer from the rates and the host properties that RETs are linked to very massive stars, potentially stripped of their envelopes, and possibly sharing some of the extreme properties of SLSN or LGRB progenitors such as rapid rotation and low metallicity. This hypothesis can be extended to posit that some RETs represent an intermediate and/or precursory step in the late stages of evolution of a massive star that is close to forming a SLSN or LGRB, whereby the initial collapse of the star occurs leading to shock breakout and subsequent cooling driving the RET light curve (P18), but the progenitor is sufficiently different to the progenitors of LGRBs and SLSNe such that the central engine either does not form or has properties that differ from the central engines of LGRBs or SLSNe, hence the lack of longer-term light curves.
|
[
"Tampo et al. 2020"
] |
[
"The approximate rate of RETs (≥10−6 Mpc−3 yr−1"
] |
[
"Uses"
] |
[
[
986,
1003
]
] |
[
[
875,
921
]
] |
2019MNRAS.485..189O__Ogiya_&_Mori_2014_Instance_1
|
Finally, we emphasize that although the parameter space covered by the DASH library is vast, it is by no means exhaustive. One obvious shortcoming, as discussed above, is that the DASH simulations are inadequate to describe major mergers with ${\cal M}\lesssim 100$. In those cases, dynamical friction due to the host, and self-friction due to tidally stripped material, cause the orbit of the subhalo to decay, exposing it to stronger tides. Another degree of freedom not covered here is the inner density slope of dark matter haloes. It is well known that observations of dwarf galaxies often suggest that their haloes have constant density cores, rather than the steep r−1-cusps predicted by dark matter-only simulations (e.g. Burkert 1995; Gentile et al. 2004; Oh et al. 2011; Hayashi & Chiba 2015). Such cores can be created within the CDM paradigm by a variety of baryonic processes (e.g. El-Zant, Shlosman & Hoffman 2001; Inoue & Saitoh 2011; Pontzen & Governato 2012; Ogiya & Mori 2014), and have a dramatic impact on the tidal evolution of subhaloes (Peñarrubia et al. 2010; Errani, Peñarrubia & Tormen 2015; Ogiya 2018). In addition, baryons modify the potentials of host- and subhaloes through the bulges and discs that they form at the halo centres, and these also strongly impact the tidal fields (Errani et al. 2017; Garrison-Kimmel et al. 2017). Finally, in the DASH simulations presented here, the host halo is assumed to be spherically symmetric, which allows us to completely specify each orbit with only two parameters (energy and angular momentum). Cosmological simulations, though, indicate that dark matter haloes are expected to be triaxial systems (e.g. Jing & Suto 2002; Allgood et al. 2006; Hayashi, Navarro & Springel 2007), consistent with the shapes of the gravitational potentials of galaxies and clusters as inferred from a variety of observations (e.g. Oguri et al. 2005; Corless & King 2007; Law & Majewski 2010). Triaxial systems have a much richer variety of orbits, which is likely to impact the tidal evolution of subhaloes.
|
[
"Ogiya & Mori 2014"
] |
[
"Such cores can be created within the CDM paradigm by a variety of baryonic processes (e.g."
] |
[
"Background"
] |
[
[
976,
993
]
] |
[
[
804,
894
]
] |
2021MNRAS.507.4316B__Tan_et_al._2014_Instance_1
|
One of the main debates in the star formation community is whether the massive young stellar objects are a scaled-up version of low-mass young stellar objects (YSOs) where disc-accretion plays the dominating role for gaining the stellar mass. For the formation of low-mass stars, bipolar outflows driven by the accretion discs are proposed to be the basic formation mechanism theoretically (Shu, Adams & Lizano 1987), and are also verified observationally (e.g. Bontemps et al. 1996; Richer et al. 2000; Arce et al. 2007, and references therein). However, on the other hand, understanding of the formation mechanism of massive stars is still elusive (Tan et al. 2014). Two major competing models for massive star formation are (i) core accretion via disc (McKee & Tan 2003) and (ii) competitive accretion (Bonnell et al. 2001). The most obvious way to distinguish between these two models might be the detection of the accretion disc around massive protostars. But a direct detection of accreting disc is difficult because the accretion disc is small and short-lived, and also because of complicated gas dynamics at that scale (Kim & Kurtz 2006). Here, the study of the properties of molecular outflows that are the manifestation of disc-accretion in young sources, could help us to improve our understanding of the underlying formation process (Shepherd & Churchwell 1996; Beuther et al. 2002; Molinari et al. 2002; Arce et al. 2007). If it is assumed that massive stars do form via an accretion disc similar to the low-mass stars, they should generate massive and powerful outflows (see de Villiers et al. 2014, and references therein). A few recent studies (e.g. de Villiers et al. 2014; Li et al. 2018) indeed found the applicability of the same scaling between outflow activity and clump masses for both low-mass and massive objects, suggesting a similar formation mechanism. However, an extensive study of molecular outflows toward massive star-forming regions is still lacking owing to their large distances and high level of clustering. Recent interferometric observations with the ALMA enable us to target such regions thanks to its high spatial resolution and sensitivity. Studies of outflows associated with Galactic massive star-forming regions may provide us clues to better understand the launching mechanism of molecular outflows, and hence, the underlying star formation mechanism.
|
[
"Tan et al. 2014"
] |
[
"However, on the other hand, understanding of the formation mechanism of massive stars is still elusive"
] |
[
"Background"
] |
[
[
651,
666
]
] |
[
[
547,
649
]
] |
2019AandA...630A..26M__Wozniakiewicz_et_al._(2012)_Instance_1
|
The majority of the particles collected by Stardust are olivine and pyroxene silicates with solar isotopic compositions, which suggests an origin in our solar system rather than an interstellar provenance. These polymineralic particles dominate those made of a single mineral even down to sizes smaller than 100 nm, indicating that the dust composition is surprisingly consistent at different scales and that the smallest subunits of the dust may be as small as tens of nanometers (Hörz et al. 2006; Zolensky et al. 2006). The sizes of these smallest single mineral impactors are similar to those of the nanocrystals determined by Rietmeijer (1993). As discussed above, they might also be existing in MIDAS dust particles and might be fused into the 100 nm features. Price et al. (2010) and Wozniakiewicz et al. (2012) investigated the sizes of particles smaller than 10 μm that impacted the aluminum foils of the Stardust probe. The distribution peaks at about 175 nm, but if we assume that the particles areagglomerates of smaller subunits, as indicated by their common polymineralic nature, then the subunit size distribution would peak at sizes below 100 nm (Price et al. 2010). A study of over 450 particles that do not seem to be agglomerates, that is, those that show single mineral impactors of silicate or sulfide, found geometric mean sizes of
$532^{741}_{-310}$
532
−310
+741
nm for the silicate particles and
$406^{491}_{-222}$
406
−222
+491
nm for the sulfides (Wozniakiewicz et al. 2013). These sizes are notably larger than the 175 nm (or less) found for the whole dataset. This large spread of subunit sizes could indicate a size distribution with a large width. No fits of these sizedistributions are available, but the figures in Wozniakiewicz et al. (2012) and Price et al. (2010) indicate that the differential sizes may follow a log-normal distribution. When we assume that the smallest subunit sizes are possibly between tens and hundreds of nanometers, the subunit size range found for MIDAS smallest features would be encompassed. The determination of the size distributions for the small Stardust particles and a detailed comparison to the distributions obtained for comet 67P could be the work of an interesting future project.
|
[
"Wozniakiewicz et al. (2012)"
] |
[
"Price et al. (2010) and",
"investigated the sizes of particles smaller than 10 μm that impacted the aluminum foils of the Stardust probe. The distribution peaks at about 175 nm, but if we assume that the particles areagglomerates of smaller subunits, as indicated by their common polymineralic nature, then the subunit size distribution would peak at sizes below 100 nm"
] |
[
"Uses",
"Uses"
] |
[
[
791,
818
]
] |
[
[
767,
790
],
[
819,
1161
]
] |
2019MNRAS.484.3307M__Hobbs_et_al._2005_Instance_1
|
For most of our models, we can already determine, or at least extrapolate, the final neutron star properties quite well, barring the possibility of late-time fallback. Except for the $12.5 \, \mathrm{M}_\odot$ model, mass outflow already dominates over mass accretion on to the PNS, and the PNS mass has practically stabilized at its final value. Correcting for the binding energy of the neutron stars, we obtain gravitational masses between $1.22 $ and $1.44 \, \mathrm{M}_\odot$, which is compatible with the distribution of observed neutron star masses (Özel & Freire 2016; Antoniadis et al. 2016; Tauris et al. 2017). While the neutron star kicks are still growing at the end of the simulations due to the long-range gravitational tug by the asymmetric ejecta, the subsequent acceleration of the neutron star can be smoothly extrapolated to obtain tentative final values in all but one case. The extrapolated kicks range from $11$ to $695 \, \mathrm{km}\, \mathrm{s}^{-1}$. Thus, the most extreme, ECSN-like models with the smallest helium cores can reproduce the very low kicks required to explain some double neutron star systems and pulsars in globular clusters (Tauris et al. 2017), while the models with higher He core masses are compatible with the typical kicks of young pulsars (Arzoumanian, Chernoff & Cordes 2002; Hobbs et al. 2005; Ng & Romani 2007). If the extrapolated kick of $1236 \, \mathrm{km}\, \mathrm{s}^{-1}$ for the $18 \, \mathrm{M}_\odot$ model of Müller et al. (2017a) is included, the 3D coconut-fmt models roughly span the full range of observed kick velocities. We see tentative evidence for a correlation of the kick velocity with the explosion energy as proposed by Janka (2017) and Vigna-Gómez et al. (2018, as a refinement of earlier ideas for progenitor-dependent kicks by Bray & Eldridge 2016). Our models suggest that this correlation may not be a tight one, however, and that the kicks may scatter between zero and an upper limit that scales with the explosion energy. Low kicks can be achieved in more energetic explosions if the explosion geometry is bipolar rather than unipolar, as has already been noted in 2D by Scheck et al. (2006). Such a bipolar explosion occurs in one of our seven simulations (the $12 \, \mathrm{M}_\odot$ model). Although there is some concern that the bipolarity may be connected to the grid geometry, we find unipolar models even in cases where we do not include strong aspherical seed perturbations in the convective O shell that break grid alignment; this suggests that the possibility of bipolar neutrino-driven explosions with low kicks is real in 3D. We also find a loose correlation between the neutron star mass and the kick velocity, which is in line with current observations, and partly theoretical expectations, of double neutron stars (Tauris et al. 2017), but cannot make as strong a case for this correlation based on our simulations. An investigation of a larger suite of supernova simulations of ultra-stripped stars is needed to confirm this hypothesis.
|
[
"Hobbs et al. 2005"
] |
[
"while the models with higher He core masses are compatible with the typical kicks of young pulsars"
] |
[
"Compare/Contrast"
] |
[
[
1328,
1345
]
] |
[
[
1191,
1289
]
] |
2016ApJ...827...31F__Lu_&_Yuan_1998_Instance_1
|
While treated as fully relativistic under strong gravity, we note that the current model is time-independent based on axisymmetric plasma. This assumption makes it impossible for us to predict any temporal nature of the soft excess considered in this work, e.g., spectral time variabilities associated with shock compression and cooling effects. The downstream plasma properties are numerically solved by considering adiabatic (nonradiative) Rankine–Hugoniot jump conditions as a pure mathematical discontinuity with no energy/mass loss. Hence, most of the heat generated at the shock front is advected with the downstream plasma. A more realistic shock process, on the other hand, is most likely accompanied by radiative cooling to some degree in which the post-shock plasma temperature may stay comparatively as cool as that of the upstream one, as in the isothermal shocks (e.g., Lu & Yuan 1998; Das et al. 2003; Fukumura & Tsuruta 2004; Fukumura & Kazanas 2007b). Radiative dissipation at the shock front could therefore drastically change the subsequent downstream plasma condition, which in turn alters the Comptonization process. In reality, furthermore, accreting plasma may be characterized by a two-temperature gas between electrons and ions (e.g., Shapiro et al. 1976; Mahadevan 1998; Manmoto 2000) unless the Coulomb coupling between the two is very efficient, whereas in this work we prescribed a single-fluid approximation for simplicity. Becker et al. (2011) have considered a particle transport process (e.g., bulk advection, spatial diffusion, and particle escape) via the effects of the first-order Fermi acceleration across a standing shock. In a more self-consistent scenario such a calculation of diffusive shock acceleration should be incorporated to reflect the energetic outflows/jets from the shock front. Although all these micro-physics should be addressed and incorporated into more sophisticated calculations of GRMHD simulations for completeness, this is beyond the scope of this paper.
|
[
"Lu & Yuan 1998"
] |
[
"A more realistic shock process, on the other hand, is most likely accompanied by radiative cooling to some degree in which the post-shock plasma temperature may stay comparatively as cool as that of the upstream one, as in the isothermal shocks (e.g.,"
] |
[
"Differences"
] |
[
[
883,
897
]
] |
[
[
631,
882
]
] |
2019ApJ...885...93F__McClintock_et_al._2006_Instance_1
|
The vertical structure of an accretion disk has been investigated in detail by Abramowicz et al. (1997). A general accurate expression for the vertical hydrodynamical equilibrium that is valid both for a thin and a slim disk has been derived in their work. The hydrostatic balance in the vertical direction of the accretion disk is described by
19
1
ρ
∂
p
∂
z
+
∂
ψ
∂
z
+
v
r
∂
v
z
∂
r
+
v
z
∂
v
z
∂
z
=
0
,
where p is the pressure, ρ is the density of the gas, and ψ is the gravitational potential (Abramowicz et al. 1997). As we focus on how the general properties of the disk structure is affected by the radiation-driven outflows, we simplify the expression of the hydrostatic balance in the vertical direction as
20
1
ρ
∂
p
∂
z
+
∂
ψ
∂
z
=
0
,
where we assume the terms of ∂vz/∂r and ∂vz/∂z to be negligible. It has been pointed out that the widely used approximation for the vertical component of gravity, GMz/R3, is only valid for the thin disk with H/R ≪ 1, and a more accurate expression of vertical gravity should be adopted for a slim disk (McClintock et al. 2006; Gu & Lu 2007; Jiao et al. 2009; Gu 2012; Cao & Gu 2015). As discussed in Cao & Gu (2015), there are upper limits on the radiation flux, frad, and half-thickness, H, for a slim disk, above which the radiation force will overwhelm the vertical gravity. The maximal flux can be calculated with
21
f
rad
=
q
(
H
)
=
GM
ρ
H
r
2
+
H
2
r
2
+
H
2
−
r
s
2
.
How the gas is blown away from the disk surface by the radiation force is still quite unclear, and, therefore, it is assumed that the gas at the disk surface will be accelerated into outflows by the radiation force when the radiation force overwhelms the vertical gravity. Then, the mass accretion rate in the disk decreases due to outflows. This self-adjustment mechanism leads to a maximal radiation flux/thickness of the disk (see the detailed discussion in Cao & Gu 2015). Thus, we assume that the outflows are triggered in the disk where the condition
22
Q
rad
−
≥
2
f
rad
max
,
is satisfied, where
f
rad
max
is the maximal radiation flux and then the mass accretion rate is self-adjusted by the outflows to maintain
23
Q
rad
−
≡
2
f
rad
max
.
This relation is adopted in the energy of Equation (15) for the region in the disk where outflows are driven. Otherwise, the outflows are suppressed, and the disk is the same as the conventional slim disk without outflows, which is described by a set of equations for a slim disk (see Chapter 7 in Kato et al. 2008), i.e., the continuity equation:
24
M
˙
=
−
2
π
r
Σ
v
r
,
where
M
˙
remains constant radially, and the angular momentum equation is
25
v
r
(
Ω
r
2
−
j
in
)
=
−
α
rc
s
2
.
Except for Equations (23)–(25), the radial-component of the momentum equation, the equation of state, and energy equation are the same as the disk with outflows (see Equations (4) and (14)–(18)).
|
[
"McClintock et al. 2006"
] |
[
"It has been pointed out that the widely used approximation for the vertical component of gravity, GMz/R3, is only valid for the thin disk with H/R ≪ 1, and a more accurate expression of vertical gravity should be adopted for a slim disk"
] |
[
"Compare/Contrast"
] |
[
[
1152,
1174
]
] |
[
[
914,
1150
]
] |
2017AandA...605A..88L__Ceccarelli_et_al._2000_Instance_1
|
Altogether, the approximately thirty molecules recently detected have confirmed the chemical complexity in the nebula, and generated our interest for the present study. Of these species, we will focus our attention on the seventeen species listed by molecular families in Table 1. As can be seen in this table, the WHISPER survey allowed the detection of some organic molecules in the Horsehead nebula, such as formaldehyde (H2CO) and methanol (CH3OH), which constitute key species in the likely synthesis of more complex organic molecules such as some prebiotic molecules (Bernstein et al. 2002; Muñoz Caro et al. 2002; Garrod et al. 2008). Because they are detected in a wide variety of interstellar sources – in hot cores (Sutton et al. 1995; Ceccarelli et al. 2000), dark clouds (Bergman et al. 2011), shocked regions (e.g. Sakai et al. 2012; Codella et al. 2012; Tafalla et al. 2010) and even in comets (Mumma & Charnley 2011; Cordiner et al. 2015) – it is of prime importance to understand well how these precursor molecules form. H2CO is commonly thought to form both in the gas-phase and on grain surfaces, while CH3OH is believed to be only formed on grain surfaces (Garrod et al. 2006; Geppert et al. 2006). Guzman et al. (2013) reported the observations of these two molecules toward the Horsehead nebula in both the PDR and Core positions. Unable to reproduce the observed abundances of either H2CO or CH3OH at the PDR position with only pure gas-phase models, they concluded that, for this region, both species are formed on grain surfaces and then photodesorbed into the gas phase. On the other hand, at the Core position, a pure gas-phase model can reproduce the observed H2CO abundance, while photodesorption of ices is still needed to explain the observed abundance of CH3OH. Other organic molecules were reported in the Horsehead nebula as first detections in a PDR environment, including HCOOH (formic acid), CH2CO (ketene), CH3CHO (acetaldehyde), and CH3CCH (propyne) (Guzman et al. 2014). Their abundances were found to be higher at the PDR position than at the Core, revealing that complex organic chemistry is also occurring in UV-illuminated neutral gas (Guzman et al. 2014). Of these molecules, some – HCOOH, CH2CO, and CH3CHO – have now also been detected in the Orion bar PDR (Cuadrado et al. 2016, 2017).
|
[
"Ceccarelli et al. 2000"
] |
[
"Because they are detected in a wide variety of interstellar sources – in hot cores",
"it is of prime importance to understand well how these precursor molecules form."
] |
[
"Motivation",
"Motivation"
] |
[
[
746,
768
]
] |
[
[
642,
724
],
[
956,
1036
]
] |
2016MNRAS.458.3760S__Pettini_et_al._1994_Instance_1
|
Plots of various characteristics of sub-DLAs and DLAs in the literature. Blue star points mark the detections presented in this work. The Xs mark detections from Péroux et al. (2011a, 2012). The points mark detections from previous studies in the literature. References: Pettini et al. 2000, Lacey et al. 2003, Junkkarinen et al. 2004, Chen et al. 2005, Burbidge et al. 1996, Chun et al. 2006, Rao et al. 2005, Boisse et al. 1998, Le Brun et al. 1997, Gharanfoli et al. 2007, Rao et al. 2006, Steidel et al. 2002, Lanzetta et al. 1997, Lanzetta et al. 1995, Bergeron & Boisse 1991, Deharveng et al. 1995, Meiring et al. 2007, Cristiani et al. 1987, Bowen et al. 2005, Schulte-Ladbeck et al. 2005, Bowen et al. 2001, Rosenberg et al. 2006, Steidel et al. 1997, Bergeron 1986, Prochaska & Wolfe 1997, Weatherly et al. 2005, Lu et al. 1993, Lu et al. 1997, Djorgovski et al. 1996, Christensen et al. 2004, Rao & Turnshek 2000, Ellison et al. 2005, Pettini et al. 1994, Francis et al. 1996, D'Odorico et al. 2002, Francis & Williger 2004, Francis et al. 2001, Kulkarni et al. 2005, Cherinka et al. 2009, Noterdaeme et al. 2009 and Fynbo et al. 2010. See table 4 of Péroux et al. (2011a) for the full list of measurements. For the SFR plots, colour indicates the emission from which the value was determined: blue for [O ii], red for Hα, green for Hβ, purple for [O iii], and black for Lyα. SFRs have not been dust corrected. The solid red line in the SFR versus zabs plot is an arbitrarily normalized Madau relationship from Madau & Dickinson (2014). Qualitatively, the trend in host galaxy SFR with absorber redshift is consistent with this trend. The blue outliers are the interacting pair in the field of Q1436-0051. The dashed vertical lines indicate the column density boundary between sub-DLAs and DLAs. The outlier in b at ∼180 kpc is a host galaxy that is part of a cluster environment reported in Francis et al. (2004) and references therein. Grey arrows indicate limits for non-detections. The error bars on Q1436-0051 zabs = 0.9281 correspond to the range of possible H i column densities and Zn metallicities. See the text for details.
|
[
"Pettini et al. 1994"
] |
[
"The points mark detections from previous studies in the literature. References:"
] |
[
"Uses"
] |
[
[
945,
964
]
] |
[
[
191,
270
]
] |
2020AandA...639A..48S__Shulyak_et_al._2004_Instance_1
|
In order to predict emission and transmission spectra of HJs we utilized the τ-REx (Tau Retrieval for Exoplanets) software package (Waldmann et al. 2015b,a). This package uses up-to-date molecular cross sections based on line lists provided by EXOMOL6 project (Tennyson & Yurchenko 2012) and HITEMP (Rothman et al. 2010). We additionally used the HCN line list after Harris et al. (2006) provided by the EXOCLIME project7 and generated with the HELIOS-K code (Grimm & Heng 2015). These cross sections are precomputed on a grid of temperatures and pressures and are stored in binary opacity tables that are available for a number of spectral resolutions. The continuum opacity includes Rayleigh scattering on molecules and collisionally-induced absorption due to H2 –H2 and H2–He either after Abel et al. (2011, 2012) or Borysow et al. (2001); Borysow (2002); Borysow & Frommhold (1989), respectively. We extended the public version of τ-REx8 by incorporating additional opacity sources essential for the atmospheres of UHJs. In particular, bound-free and free–free transitions of H− become one of the major continuum opacity contributors for the temperatures hotter than about 2000 K. The cross sections of H− are from John (1988). We also included opacity due to free–free transitions of He− as well as Rayleigh scattering on H I atoms and Thomson scattering on free electrons. The He− cross sections are those originally from John (1968) using the polynomial fit by Carbon et al. (1969). Rayleigh scattering on H I is calculated after Dalgarno (1962). All relevant numerical routines were extracted from the LLMODELS stellar model atmosphere code (Shulyak et al. 2004). At the temperatures of HJs the H− is the dominant continuum opacity source that impacts the observed spectra of these planets (Arcangeli et al. 2018). However, at millibar pressures the He− opacity can become comparable or even stronger than that of H− for wavelengths longer than 1.6 μm, as shown in Fig. 3 (third panel from the bottom), where we display examples of continuum opacity coefficient at different altitudes in the atmosphere of a HJ with Teq = 3000 K. At even smaller pressures, electron scattering and Rayleigh scattering on H I atoms also become important contributors to the continuum opacity at particular wavelengths (top panel in Fig. 3). However, the contribution of H I and e− on the transmission spectra is marginal because their opacity is only strong at low pressures that are hardly probed by transmission spectroscopy. Thus, among all continuum opacity sources, only H− and He− significantly contribute to the predicted amplitude of the transmission spectra, as shown in the bottom plot of Fig. 3. When both continuum and line opacity are included, the effect of He− on predicted spectra is diluted by a much stronger opacity in molecular lines while H− contribution is still significant. Nevertheless, as can be seen from the second plot (from bottom) in Fig. 3, in optically thick layers the He− opacity could still be stronger than, for example, collision-induced absorption due to H2 -H2 and H2-He. We thus conclude that accurate calculation of atmospheric opacity requires He−, H I, and e− opacity included especially at low pressures, following modern stellar atmosphere codes. However, observed transmission and emission spectra of HJs are hardly affected by these three opacity sources. Finally, we updated HELIOS with He− and e− opacity and found out that this has little impact on the atmospheric temperature structure (with a modification in local temperature of at most ΔT ≈ 10 K) and thus can be ignored; the original version of HELIOS already includes H I Rayleigh scattering.
|
[
"Shulyak et al. 2004"
] |
[
"All relevant numerical routines were extracted from the LLMODELS stellar model atmosphere code"
] |
[
"Uses"
] |
[
[
1650,
1669
]
] |
[
[
1554,
1648
]
] |
2021MNRAS.507.4389G__Masters_et_al._2011_Instance_1
|
Erwin (2018) showed that, in a sample drawn from the Spitzer Survey of Stellar Structure in Galaxies (S4G), the bar fraction is constant over a range of (g −r) colours and gas fractions. Their bar fraction does not increase, but rather decreases for stellar masses higher than ∼ 109.7M⊙. These results are in contrast to many SDSS-based studies cited above. Erwin (2018) argues that this apparent contradiction can be explained if SDSS-based studies miss bars in low-mass blue galaxies. In Figs 5 and 6, we showed that the newly detected bars in GZD (compared to GZ2) are weak bars in low-mass blue galaxies. Nevertheless, the ‘combined’ bar fraction in Fig. 6 is not constant over (g −r) colour and agrees well with Masters et al. (2011) for redder colours [(g −r) colour > 0.5]. Additionally, our ‘combined’ bar fraction remains roughly constant over stellar mass. As mentioned before, we conclude that strong bars drive the trends of bar fraction with (g −r) colour, stellar mass, and SFR observed in other studies (Nair & Abraham 2010b; Masters et al. 2011, 2012; Vera et al. 2016; Cervantes Sodi 2017). However, the addition of weak bars in low-mass blue galaxies is insufficient to resolve the apparent disagreement between Erwin (2018) and many SDSS-based studies (Masters et al. 2011, 2012; Vera et al. 2016; Cervantes Sodi 2017; Kruk et al. 2018), which instead seems likely to be due to the very different sample selection of the S4G and SDSS galaxy samples. For example, the median stellar mass of the sample used in Erwin (2018) is ∼109.6M⊙ (based on their Fig. 4 and the bins in the top left-hand panel of their Fig. 5). However, the median stellar mass of our sample is 1010.6M⊙. As stellar mass correlates with many parameters (including bar length), this can have major consequences. Additionally, as Erwin (2018) notes, there is also the issue of resolution to consider. With an r-band FWHM of 1.18 arcsec from DECaLS (Dey et al. 2019) and a mean redshift of 0.036, the mean linear resolution of our sample is approximately 834 pc, which is higher than the 165 pc of Erwin (2018). This explains why they observe many sub-kpc bars, while we do not. These differences in stellar mass and resolution will manifest themselves in the conclusions, so a more detailed analysis is needed for a proper comparison with Erwin (2018).
|
[
"Masters et al. (2011)"
] |
[
"Nevertheless, the ‘combined’ bar fraction in Fig. 6 is not constant over (g −r) colour and agrees well with",
"for redder colours [(g −r) colour > 0.5]"
] |
[
"Similarities",
"Similarities"
] |
[
[
717,
738
]
] |
[
[
609,
716
],
[
739,
779
]
] |
2015ApJ...814...73S__Soria_et_al._2004_Instance_1
|
At a distance of 4.8 Mpc (Karachentsev et al. 2002), NGC 5408 X-1 is one of the best studied ULXs. It has been observed by several of the current generation of X-ray satellites, on a multitude of occasions. Observations include: an XMM-Newton large programme (e.g., Pasham & Strohmayer 2012); Swift XRT monitoring (Kaaret & Feng 2009; Grisé et al. 2013); and eight Chandra exposures, which we re-analyze here. The flux variability of the source rules out an X-ray supernova remnant and confirms that it is powered by accretion onto a compact object (Kaaret et al. 2003; Soria et al. 2004). It persistently displays a distinct soft ultraluminous two component X-ray spectrum in XMM-Newton data (Sutton et al. 2013) at an average 0.3–10 keV unabsorbed luminosity of
(Strohmayer 2009; although we note that they fit the high energy spectrum with a soft power-law, so may over-estimate the intrinsic luminosity, cf. Middleton et al. 2014). Additional soft residuals have been detected in the XMM-Newton spectra which can be well modeled as thermal plasma emission (Strohmayer & Mushotzky 2009; Miller et al. 2013; Middleton et al. 2014). It has previously been assumed that these were the result of diffuse star formation related emission in the host galaxy (Strohmayer et al. 2007; Strohmayer & Mushotzky 2009; Miller et al. 2013). However, we know from observational studies of galaxies that the X-ray luminosity of such emission is correlated with star formation rate (
/
; Mineo et al. 2012), and Middleton et al. (2014) contend that the luminosity of the putative thermal plasma emission (
calculated from Miller et al. 2013) greatly exceeds that inferred from star formation, even over the entirety of NGC 5408 (
calculated from a 24 μm flux density of
, Dale et al. 2005). Instead, Middleton et al. (2014) show that the putative plasma emission features could actually be commensurate with broadened, blueshifted absorption in a partially ionized, optically thin medium. Such a medium would be expected to occur in a super-Eddington wind, as it becomes optically thin at large distances (
) from the central black hole.
|
[
"Soria et al. 2004"
] |
[
"The flux variability of the source rules out an X-ray supernova remnant and confirms that it is powered by accretion onto a compact object"
] |
[
"Compare/Contrast"
] |
[
[
570,
587
]
] |
[
[
410,
548
]
] |
2015AandA...579A..46M__Ilyushin_et_al._2010_Instance_1
|
The Hamiltonian used in the present work is the so-called RAM (rho axis method) internal-rotation Hamiltonian based on the work of Kirtman (Kirtman 1962), Lees and Baker (Lees & Baker 1968), and Herbst et al. (Herbst et al. 1984). Since rather complete descriptions of this method, which takes its name from the choice of axis system, have been presented several times (Hougen et al. 1994; Kleiner 2010) we do not repeat this general description here. The main advantage of the RAM Hamiltonian is its general approach that simultaneously takes into account the A- and E-symmetry species and all the torsional levels, intrinsically taking the intertorsional interactions into account within the rotation-torsion manifold of energy levels. This method was successfully applied to a number of molecules containing a C3v rotor and Cs frame, including the main isotopolog of acetaldehyde (Smirnov et al. 2014). As for the main isotopolog (Smirnov et al. 2014) we employed the RAM36 (rho-axis-method for 3- and 6-fold barriers) code that uses the RAM approach for the molecules with the C3v top attached to a molecular frame of Cs or C2v symmetry and having 3- or 6-fold barriers to internal rotation, respectively (Ilyushin et al. 2010, 2013). The Hamiltonian in the RAM36 program is presented by the following expression: (1)\begin{eqnarray} \nonumber H &=& (1/2) \sum_{knpqrs} B_{knpqrs0} \left[P^{2k}P_z^nP_x^pP_y^qp_{\alpha}^r \cos(3s{\alpha}) \right.\\[2mm] \nonumber &&\left. + \cos(3s\alpha) p_{\alpha}^rP_y^qP_x^pP_z^nP^{2k}\right] \\[2mm] \nonumber && + (1/2) \sum_{knpqrt} B_{knpqr0t} \left[P^{2k}P_z^nP_x^pP_y^qp_{\alpha}^r \sin(3t\alpha) \right.\\ [2mm] &&\left. + \sin(3t\alpha) p_{\alpha}^rP_y^qP_x^pP_z^nP^{2k}\right] \end{eqnarray}H=(1/2)∑knpqrsBknpqrs0[P2kPznPxpPyqpαrcos(3sα)+cos(3sα)pαrPyqPxpPznP2k]+(1/2)∑knpqrtBknpqr0t[P2kPznPxpPyqpαrsin(3tα)+sin(3tα)pαrPyqPxpPznP2k]where the Bknpqrst are fitting parameters; pα is the angular momentum conjugate to the internal rotation angle α; and Px,Py,Pz are projections on the x,y,z axes of the total angular momentum P. In the case of a C3v top and Cs frame (as is appropriate for acetaldehyde), the allowed terms in the torsion-rotation Hamiltonian must be totally symmetric in the group G6 (and also must be Hermitian and invariant to the time reversal operation). Since all individual operators pα,Px,Py,Pz,P2,cos(3sα) and sin(3tα) used in Eq. (1) are Hermitian, all possible terms provided by Eq. (1) will automatically be Hermitian. The particular term to be fit is represented in the input file with a set of k,n,p,q,r,s,t integer indices that are checked by the program for conformity with time reversal and symmetry requirements, to prevent accidental introduction of symmetry-forbidden terms into the Hamiltonian.
|
[
"Ilyushin et al. 2010"
] |
[
"As for the main isotopolog",
"we employed the RAM36 (rho-axis-method for 3- and 6-fold barriers) code that uses the RAM approach for the molecules with the C3v top attached to a molecular frame of Cs or C2v symmetry and having 3- or 6-fold barriers to internal rotation, respectively"
] |
[
"Uses",
"Uses"
] |
[
[
1210,
1230
]
] |
[
[
906,
932
],
[
955,
1208
]
] |
2021MNRAS.500.2336Y__Matonick_&_Fesen_1997_Instance_2
|
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.
|
[
"Matonick & Fesen 1997"
] |
[
"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"
] |
[
"Background"
] |
[
[
1264,
1285
]
] |
[
[
1041,
1171
]
] |
2022AandA...658A.194P__Khata_et_al._2020_Instance_3
|
The stellar photospheric parameters we collected from literature for the benchmark stars are summarized in Table A.1. Although most benchmark stars have v sini 2 km s−1 (Reiners et al. 2018), there are two stars with larger values: J07558+833 (12.1 km s−1) and J13005+056 (16.4 km s−1). These stars are useful to investigate the performance of the algorithms when dealing with higher rotational velocities. The literature values were derived with different methods. These methods include: interferometry to estimate the stellar radius and Teff (Boyajian et al. 2012; Ségransan et al. 2003; von Braun et al. 2014; Berger et al. 2006; Newton et al. 2015), synthetic model fitting using BT-Settl models to determine Teff (Gaidos et al. 2014; Lépine et al. 2013; Gaidos & Mann 2014; Mann et al. 2015) and log g (Lépine et al. 2013), empirical relations to derive stellar mass in the form of mass-luminosity relations (Mann et al. 2015; Khata et al. 2020; Boyajian et al. 2012; Berger et al. 2006; Ségransan et al. 2003), along with the mass-magnitude relations (Maldonado et al. 2015), mass-radius relations (von Braun et al. 2014), mass–Teff relations (Gaidos & Mann 2014; Gaidos et al. 2014), empirical relations to derive the stellar radius in the form of mass-radius relations (Maldonado et al. 2015) and Teff–radius relations (Gaidos & Mann 2014; Gaidos et al. 2014; Houdebine et al. 2019), pEW measurements to determine Teff (Maldonado et al. 2015; Neves et al. 2014; Newton et al. 2015) and [Fe/H] (Maldonado et al. 2015; Neves et al. 2014; Gaidos et al. 2014; Mann et al. 2015), the definition of spectral indices such as the H2O-K2 index to estimate Teff (Rojas-Ayala et al. 2012), as well as the combination of the H2O-K2 index with pEWs to derive [Fe/H] (Rojas-Ayala et al. 2012; Khata et al. 2020), the stellar radius and Teff (Khata et al. 2020), and spectral curvature indices for the determination of Teff (Gaidos & Mann 2014). Additionally, [Fe/H] was derived by using color-magnitude metallicity relations (Dittmann et al. 2016), atomic line strength relations (Gaidos & Mann 2014), and spectral feature relations (Terrien et al. 2015). Terrien et al. (2015) used K-band magnitudes and the Dartmouth Stellar Evolution Program (Dotter et al. 2008) to derive the stellar radius, whereas Mann et al. (2015) employed the Boltzmann equation with Teff determined from synthetic model fits. Last, but not least, Houdebine et al. (2019) derived Teff from photometric colors. For more details on the individual methods, we refer to the descriptions in the corresponding works.
|
[
"Khata et al. 2020"
] |
[
"the stellar radius and Teff"
] |
[
"Background"
] |
[
[
1837,
1854
]
] |
[
[
1808,
1835
]
] |
2020AandA...644A..97C__Leroy_et_al._2013_Instance_3
|
Major nearby galaxy cold gas mapping surveys (Regan et al. 2001; Wilson et al. 2009; Rahman et al. 2011; Leroy et al. 2009; Donovan Meyer et al. 2013; Bolatto et al. 2017; Sorai et al. 2019; Sun et al. 2018) have focused on observations of the molecular gas (through CO lines). Despite a few notable exceptions (e.g. Alatalo et al. 2013; Saintonge et al. 2017), these surveys observed mainly spiral or infrared-bright galaxies (i.e. galaxies with significant star formation) and have furthered our understanding of how star formation happens, rather than how it stops. This boils down to quantifying the relation between molecular gas and star formation rate (SFR), which appears nearly linear in nearby discs (Kennicutt 1998; Bigiel et al. 2008; Leroy et al. 2013; Lin et al. 2019). This relationship is often parametrised via the ratio between the SFR and the molecular gas mass (Mmol), which is called the molecular star formation efficiency (SFE = SFR/Mmol = 1∕τdep), where the inverse of the SFE is the depletion time, τdep. The depletion time indicates how much time is necessary to convert all the available molecular gas into stars at the current star formation rate. On kpc scales and in the discs of nearby star-forming galaxies, τdep is approximately constant around 1–2 Gyr (Bigiel et al. 2011; Rahman et al. 2012; Leroy et al. 2013; Utomo et al. 2017), and it appears to weakly correlate with many galactic properties such as stellar mass surface density or environmental hydrostatic pressure (Leroy et al. 2008; Rahman et al. 2012). Nevertheless, small but important deviations for a constant SFE have been noticed, which can be the first hints of star formation quenching. In some galaxies, the depletion time in the centres appear shorter (Leroy et al. 2013; Utomo et al. 2017) or longer (Utomo et al. 2017) with respectto their discs. These differences may correlate with the presence of a bar or with galaxy mergers (Utomo et al. 2017; see also Muraoka et al. 2019) and do not seem to be related to unaccounted variation in the CO-to-H2 conversion factor (Leroy et al. 2013; Utomo et al. 2017). Spiral arm streaming motions have also been observed to lengthen depletion times (Meidt et al. 2013; Leroy et al. 2015).
|
[
"Leroy et al. 2013"
] |
[
"In some galaxies, the depletion time in the centres appear shorter"
] |
[
"Background"
] |
[
[
1756,
1773
]
] |
[
[
1688,
1754
]
] |
2022MNRAS.515.2698A__In_2018_Instance_1
|
In the last few years, similar (though not identical) experimental developments to our method have been reported. In 2017, Wehres et al. (2017) described the design, construction, and operation of two laboratory broad-band emission spectrometers for gas-phase characterization of large molecules. The first is based on a Schottky-barrier diode heterodyne receiver spectrometer operating between 80 and 110 GHz, whilst the second uses cryo-cooled SIS technology and operates at higher frequency between 270 and 290 GHz (Wehres et al. 2018). Gas phase pyridine and methyl cyanide features were successfully detected and matching analytic simulations were presented for the observed spectral transitions. In 2018, a new technique combining a Terahertz radiometer (41–49) GHz and a vacuum chamber was developed to observe the generation of cold plasma and UV photochemistry in the gas phase under low pressure as described by Tanarro et al. (2018). However, in the context of the experiment described here, it is important to note that the spectrometer developments were not coupled to a system where the gas-phase molecules originated from the solid state. More recently, Yocum et al. (2019) reported on an experiment using a THz source and hot-electron bolometer to detect the gas-phase absorption spectra of simple molecular species such as H2O, D2O, and CH3OH desorbing into the gas-phase from ices grown in an ultra-high vacuum (UHV) chamber. This experimental set-up (SubLIME), combined with a Fourier-transform Infrared spectrometer, gave the spectroscopic insight on the ultraviolet photolysis and warm up of methanol ice sample, for the first time observed at those frequencies in a laboratory environment (Yocum et al. 2021). This showed the potential of the use of submillimetre/far-infrared technique to identify molecules in complex gas mixtures. One challenge of this technique was the sensitivity of absorption spectroscopy to the desorbing molecules. The combination of temperature programmed desorption and microwave spectroscopy techniques was discussed by Theulé et al. (2020). In these experiments, desorption of water, deuterium, methanol, and ammonia was measured using a hot electron bolometer and Si-diode technology combined with Fourier transform spectrometers. While Theule et al. observed the desorption through a waveguide cavity, Yocum et al. (2019) detected the desorption above a metal substrate as seen in standard TPD studies.
|
[
"Tanarro et al. (2018)"
] |
[
"In 2018, a new technique combining a Terahertz radiometer (41–49) GHz and a vacuum chamber was developed to observe the generation of cold plasma and UV photochemistry in the gas phase under low pressure as described by",
"However, in the context of the experiment described here, it is important to note that the spectrometer developments were not coupled to a system where the gas-phase molecules originated from the solid state."
] |
[
"Background",
"Differences"
] |
[
[
922,
943
]
] |
[
[
702,
921
],
[
945,
1153
]
] |
2022AandA...665A..25C__Lambrechts_&_Johansen_2012_Instance_1
|
Various theoretical studies employing numerical simulations have been performed to investigate planet formation through core accretion or gravitational instability in the low-stellar-mass regime. Payne & Lodato (2007) assessed planet formation around brown dwarfs adapting models for higher stellar masses based on core accretion. Through Monte Carlo simulations, they found that Earth-like planets can form in this condition and the planet mass depends strongly on the disk mass. However, none of their simulations showed a planetary rocky core accreting a gaseous envelope to form a giant planet. A way to overcome the radial drift barrier is a rapid rocky core growth. Pebble accretion is a mechanism able to speed up significantly the giant planet formation process (i.e., Lambrechts & Johansen 2012; Bitsch et al. 2015). Liu et al. (2020) carried out a theoretical study on planet formation driven by pebble accretion in the (sub)stellar mass range between 0.01 and 0.1 M⊙. First, they calculated the initial masses of protoplanets by extrapolating previous numerical simulations conducted in previous literature. Next, they performed a population synthesis study to track the growth and migration of a large sample of protoplanets under the influence of pebble accretion. Their results show that, around a 0.01 M⊙ brown dwarf, planets can grow up to 0.1−0.2 M⊕, while, around 0.1 M⊙ stars, planets can reach a maximum mass of 2−3 M⊕. Findings from this study show that even pebble accretion does not seem to be sufficient to form gas giants around VLM stars and brown dwarfs. Miguel et al. (2020) used a population synthesis approach based on planetesimal accretion to explore planet formation in the stellar mass range between 0.05 and 0.25 M⊙. They let the synthetic population of planetary systems evolve for 108 yr. The authors find that to form planets with masses higher than 0.1 M⊕ they need stars of at least 0.07 M⊙, implying that planet formation around brown dwarf may not be a usual outcome. Then, stars with masses higher than 0.15 M⊙ are necessary to form planets more massive than the Earth. Therefore, from all of these studies, we conclude that core accretion model currently cannot explain the presence of gas giants around VLM stars or brown dwarfs. Either the core accretion theory is incomplete, or another mechanism for planet formation is needed. This is confirmed by Lodato et al. (2005), who discussed the origin of the 5 MJup planet detected around the 25 MJup brown dwarf 2MASSW J1207334-393254 (Chauvin et al. 2005). They found that the core accretion mechanism is far too slow to generate such a planet in less than 107 yr, the estimated age of the system. Therefore, the authors proposed gravitational instabilities arising during the early phases of the disk lifetime as a viable possibility for the formation of the planet.
|
[
"Lambrechts & Johansen 2012"
] |
[
"Pebble accretion is a mechanism able to speed up significantly the giant planet formation process (i.e.,"
] |
[
"Background"
] |
[
[
777,
803
]
] |
[
[
672,
776
]
] |
2021MNRAS.505..435S__Deming_et_al._2013_Instance_1
|
Detections of ionic, atomic, and molecular species in exoplanetary atmospheres serve as a unique and strong diagnostic of those chemical and dynamical processes driving their formation and evolution. Their detection and abundance measurements could act as indicators of planetary formation scenarios and reveal connections to the primordial protoplanetary disc and the host star (Williams & Cieza 2011; Mordasini et al. 2016; Madhusudhan et al. 2017). Furthermore, discoveries of atmospheric chemical species allow us to better understand various thermodynamical processes and chemistry, winds in the upper atmosphere (Goodman 2009; Snellen et al. 2010; Brogi et al. 2016; Madhusudhan et al. 2016; Wyttenbach et al. 2020), and to probe planetary interiors and various bulk properties through their abundances (Kite et al. 2016; Thorngren & Fortney 2019; Madhusudhan et al. 2020). A whole host of ions, atoms, and molecules have been detected through a variety of, often complementary, techniques, such as differential spectrophotometry using low-to-mid resolution spectroscopy (e.g. Gibson et al. 2012, 2017; Deming et al. 2013; Kreidberg et al. 2014; Kirk et al. 2016; Nortmann et al. 2016), and high resolution spectroscopic techniques (e.g. Redfield et al. 2008; Snellen et al. 2008; Rodler, Lopez-Morales & Ribas 2012; Birkby et al. 2013; Hoeijmakers et al. 2015, 2018, 2020; Brogi et al. 2016; Birkby et al. 2017; Žák et al. 2019; Ehrenreich et al. 2020). To date, ionic species such as Fe ii and Ti ii (Hoeijmakers et al. 2019), atomic absorption from Na, K, H α, and He (e.g. Redfield et al. 2008; Sedaghati et al. 2016; Casasayas-Barris et al. 2017; Spake et al. 2018; Chen et al. 2020; Seidel et al. 2020), and molecules such as H2O, CH4, and CO (e.g. Konopacky et al. 2013; Brogi et al. 2014; Fraine et al. 2014; Barman et al. 2015; Sing et al. 2016) have been detected through the aforementioned techniques. Needless to say that this list of detected constituents is by no means exhaustive, nor that of methods employed to detect exoplanetary atmospheres. For instance, high-resolution imaging instruments such as SPHERE (Beuzit et al. 2019) and GRAVITY (Gravity Collaboration et al. 2017), both at the VLT (ESO’s Very Large Telescope), through combination with low-dispersion spectroscopy, have facilitated direct measurements of exoplanetary atmospheres (Samland et al. 2017; Gravity Collaboration et al. 2020).
|
[
"Deming et al. 2013"
] |
[
"A whole host of ions, atoms, and molecules have been detected through a variety of, often complementary, techniques, such as differential spectrophotometry using low-to-mid resolution spectroscopy (e.g."
] |
[
"Background"
] |
[
[
1109,
1127
]
] |
[
[
880,
1082
]
] |
2021AandA...655A.104P__Draine_&_Li_2007_Instance_1
|
This last, more technical aspect gave birth to a branch of research devoted to the problem of fitting a galaxy SED (Arnouts et al. 1999; Bolzonella et al. 2000; Cid Fernandes et al. 2005; Ilbert et al. 2006; Ocvirk et al. 2006; Tojeiro et al. 2007; Fritz et al. 2007, 2017; Franzetti et al. 2008; Pappalardo et al. 2010; Han & Han 2014; Leja et al. 2017; Weaver et al. 2021). Historically, models reproducing the evolution of stellar populations and dust emission have been developed separately, with reliable prescriptions to build synthetic spectra of different stellar populations on one side (Bruzual & Charlot 1993; Bressan et al. 1994; Worthey 1994; Fioc & Rocca-Volmerange 1997; Leitherer et al. 1999; Vazdekis 1999; Charlot & Fall 2000; Maraston 2005; Fioc & Rocca-Volmerange 2019), and a thorough comprehension of dust grains physics on the other (Draine & Lee 1984; Dale 2001; Takeuchi et al. 2003, 2005; Zubko et al. 2004; Draine & Li 2007; da Cunha et al. 2010; Silva et al. 2011; Asano et al. 2013; Calura et al. 2014; Zhukovska 2014; Mancini et al. 2015; Schneider et al. 2016; Popping et al. 2017; Aoyama et al. 2017; De Vis et al. 2017, 2019; Ginolfi et al. 2018; Graziani et al. 2019; Burgarella et al. 2019; Nanni et al. 2019; De Looze et al. 2020; Galliano et al. 2021). The two branches finally converged, performing simultaneous fitting of both components (Devriendt et al. 1999; Groves et al. 2008; da Cunha et al. 2008; Noll et al. 2009; Silva 2009; Graziani et al. 2019). Two of these methods, CIGALE (Noll et al. 2009; Boquien et al. 2019) and MAGPHYS (da Cunha et al. 2008), are based on the assumption that the radiation produced during the star formation process by the stellar and nebular components is partially absorbed by the dust and then re-emitted in the infrared part of the spectrum, satisfying the energy conservation. New families of codes have also been produced in an attempt to combine spectral and photometric analyses, recovering, through the implementation of photoionisation codes, emission lines feature; see for example Prospect (Robotham et al. 2020) and Prospector (Leja et al. 2017). Recently, even machine learning methods have been involved in this challenge, with neural network algorithms showing promising results (Simet et al. 2021).
|
[
"Draine & Li 2007"
] |
[
"Historically, models reproducing the evolution of stellar populations and dust emission have been developed separately,",
"and a thorough comprehension of dust grains physics on the other"
] |
[
"Background",
"Background"
] |
[
[
934,
950
]
] |
[
[
376,
495
],
[
791,
855
]
] |
2017ApJ...850...97B__Tamburro_et_al._2009_Instance_1
|
The H i mass fraction of every gas particle in the baryonic runs is calculated based on the particle’s temperature and density and the cosmic UV background radiation flux while including a prescription for self-shielding of H2 and dust shielding in both H i and H2 (Christensen et al. 2012). This allows for the straightforward calculation of the total H i mass of each simulated galaxy. We create mock H i data cubes only for the 42 halos that contain
. Specifically, we create mock data cubes that mimic ALFALFA observations (Haynes et al. 2011). After specifying a viewing angle (see below), our code considers the line-of-sight velocity of each gas particle. The velocity of each particle is tracked in the simulation by solving Newton’s equations of motion, but any turbulent velocity of the gas is not taken into account. Velocity dispersions in dwarf galaxies can be on the order of the rotational velocity, ∼10–15 km s−1 (e.g., Stanimirović et al. 2004; Tamburro et al. 2009; Oh et al. 2015). Dispersions are thought to be driven at least partially by thermal velocities or supernovae (Tamburro et al. 2009; Stilp et al. 2013a, 2013b). In our simulations, supernovae inject thermal energy, and the thermal state of the H i gas needs to be considered in the mock H i linewidth for a realistic comparison to observations. To account for the thermal velocity, the H i mass of each gas particle is assumed to be distributed along the line-of-sight in a Gaussian distribution with a standard deviation given by the thermal velocity dispersion,
, where T is the temperature of the gas particle. After this thermal broadening is calculated, a mock H i data cube can be generated by specifying the spatial and velocity resolution. For all of our galaxies, we adopt a spatial resolution of 54 pixels across 2Rvir. In practice, this corresponds to a range of ∼1 kpc resolution in our lowest-mass galaxies up to ∼9 kpc resolution in our most massive galaxies. However, the spatial resolution plays no role in our study, since measurements of the VF are based on spatially unresolved H i data. For the velocity resolution, we match the ALFALFA specification of 11.2 km s−1 (two-channel boxcar-smoothed).
|
[
"Tamburro et al. 2009"
] |
[
"Velocity dispersions in dwarf galaxies can be on the order of the rotational velocity, ∼10–15 km s−1 (e.g.,"
] |
[
"Compare/Contrast"
] |
[
[
968,
988
]
] |
[
[
834,
941
]
] |
2021ApJ...923L..22A__Dvorkin_&_Barausse_2017_Instance_1
|
Pulsar timing experiments (Sazhin 1978; Detweiler 1979) allow us to explore the low-frequency (∼1–100 nHz) part of the gravitational-wave (GW) spectrum. By measuring deviations from the expected arrival times of radio pulses from an array of millisecond pulsars, we can search for a variety of GW signals and their sources. The most promising sources in the nanohertz part of the GW spectrum are supermassive binary black holes (SMBHBs) that form via the mergers of massive galaxies. Orbiting SMBHBs produce a stochastic GW background (GWB; Lommen & Backer 2001; Jaffe & Backer 2003; Volonteri et al. 2003; Wyithe & Loeb 2003; Enoki et al. 2004; Sesana et al. 2008; McWilliams et al. 2012; Sesana 2013; Ravi et al. 2015; Rosado et al. 2015; Kelley et al. 2016; Sesana et al. 2016; Dvorkin & Barausse 2017; Kelley et al. 2017; Bonetti et al. 2018; Ryu et al. 2018), individual periodic signals or continuous waves (CWs; Sesana et al. 2009; Sesana & Vecchio 2010; Mingarelli et al. 2012; Roedig & Sesana 2012; Ravi et al. 2012, 2015; Rosado et al. 2015; Schutz & Ma 2016; Mingarelli et al. 2017; Kelley et al. 2018), and transient GW bursts (van Haasteren & Levin 2010; Cordes & Jenet 2012; Ravi et al. 2015; Madison et al. 2017; Islo et al. 2019; Bécsy & Cornish 2021). We expect to detect the GWB first, followed by detection of individual SMBHBs (Siemens et al. 2013; Rosado et al. 2015; Taylor et al. 2016; Mingarelli et al. 2017) that stand out above the GWB. Detection of GWs from SMBHBs will yield insights into galaxy mergers and evolution not possible through any other means. Other potential sources in the nanohertz band include cosmic strings (Damour & Vilenkin 2000, 2001; Berezinsky et al. 2004; Damour & Vilenkin 2005; Siemens et al. 2006, 2007; Ölmez et al. 2010; Sanidas et al. 2013; Blanco-Pillado et al. 2018; Chang & Cui 2021; Ghayour et al. 2021; Gorghetto et al. 2021; Wu et al. 2021a; Blanco-Pillado et al. 2021; Lin 2021; Chiang & Lu 2021; Lazarides et al. 2021; Chakrabortty et al. 2021; Ellis & Lewicki 2021), phase transitions in the early universe (Witten 1984; Caprini et al. 2010; Addazi et al. 2021; Arzoumanian et al. 2021; Di Bari et al.2021; Borah et al. 2021; Nakai et al. 2021; Brandenburg et al.2021; Neronov et al. 2021), and relic GWs from inflation (Starobinskiǐ 1979; Allen 1988; Lazarides et al. 2021; Ashoorioon et al. 2021; Yi & Zhu 2021; Li et al. 2021; Poletti 2021; Vagnozzi 2021; Sharma 2021), all of which would provide unique insights into high-energy and early-universe physics.
|
[
"Dvorkin & Barausse 2017"
] |
[
"Orbiting SMBHBs produce a stochastic GW background (GWB"
] |
[
"Background"
] |
[
[
781,
804
]
] |
[
[
484,
539
]
] |
2021ApJ...921..107B__Bellovary_et_al._2021_Instance_1
|
A black hole mass this large in Leo I is not expected from extrapolation of any of the standard black hole−host galaxy correlations. Of course, these small systems do not necessarily need to follow the trends seen in normal galaxies, but the black hole mass reported here does stand out. Lützgendorf et al. (2015) explore extrapolations of black hole correlations down to globular cluster scales, and using a velocity dispersion of 12 km s−1, Leo I has a black hole mass a factor of 100 more than the extrapolated trends. On the numerical side, van Wassenhove et al. (2010) consider different scenarios for formation of a black hole in Milky Way satellites and place the likelihood of one of them having a black hole around the size found here to be below 1%, but this result also depends on the initial seed mass (see also Bellovary et al. 2021). Runaway mergers of stellar mass black holes are unlikely to produce such a black hole in such a small galaxy, since the required initial mass function to reach the ratios seen in the models might be more top-heavy than what chemical abundances and star formation history studies suggest. An alternative explanation for the abnormally large central black hole may come from the recent study on Leo I’s star formation history from Ruiz-Lara et al. (2020). The authors identify a period of quenching from z = 1–2 followed by reignition until almost the present day, when ram pressure stripping may have shut it down as it fell into the Milky Way. While the authors speculate that this reignition at intermediate redshifts could be due to a past merger with a smaller dwarf, this could also be consistent with gas accretion and potential active galactic nuclei feedback, lending support to the high MBH values presented here. Amaro-Seoane et al. (2014) also suggest that dwarf systems may in fact have significantly larger black holes compared to the host galaxy–black hole relationships. Having a larger sample of black hole limits measured in dwarf galaxies will be important to explore.
|
[
"Bellovary et al. 2021"
] |
[
"On the numerical side, van Wassenhove et al. (2010) consider different scenarios for formation of a black hole in Milky Way satellites and place the likelihood of one of them having a black hole around the size found here to be below 1%, but this result also depends on the initial seed mass (see also"
] |
[
"Compare/Contrast"
] |
[
[
824,
845
]
] |
[
[
522,
823
]
] |
2016AandA...592A.157M__Mernier_et_al._2015_Instance_1
|
As seen in Fig. 6 (right), the large MOS-pn discrepancy in the Ni/Fe abundance ratio prevents us from deriving a precise measurement. This discrepancy is worrying, but can be explained by imperfections in the cross-calibration of the two instruments. Alternatively, and perhaps more likely, the high energy band around the Ni-K transitions is significanly affected by the instrumental background (as the flux of the cluster emission sharply decreases at high energies). This hard particle background (already mentioned in Sect. 3.1) has a different spectral shape in MOS and pn, which might even vary with time, thus between observations. In particular, an instrumental line (Cu Kα) is known to affect pn at a rest-frame energy of ~8 keV (Mernier et al. 2015). Despite our efforts to carefully estimate the background, that line might interfere with the Ni-K line in several observations, making a proper modelling of the Ni-K line impossible, and hence, boosting the Ni absolute abundance in pn. In this context, it can be instructive to compare our Ni/Fe measurements with those of Suzaku, which has a lower relative hard particle background. Sato et al. (2007b) (A 1060) and Tamura et al. (2009) (Perseus) reported ratios of ~1.3 ± 0.4 and ~1.11 ± 0.19, respectively (after rescaling to the proto-solar values). Although these measurements might be also be affected by further uncertainties (e.g. the choice of the spectral modelling, Sect. 4.3), they appear to be consistent with the Ni/Fe average ratio measured with MOS is this work, favouring our above supposition that MOS is more trustworthy than pn for measuring Ni/Fe. However, in order to be conservative, we prefer to retain the pn value as a possible result and, therefore, we keep large systematic uncertainties for Ni/Fe. We finally note that, unsurprisingly, Ni/Fe cannot be constrained in the cool objects (Fig. 6, left) because the gas temperature is too low to excite Ni-K transitions.
|
[
"Mernier et al. 2015"
] |
[
"In particular, an instrumental line (Cu Kα) is known to affect pn at a rest-frame energy of ~8 keV",
"Despite our efforts to carefully estimate the background, that line might interfere with the Ni-K line in several observations, making a proper modelling of the Ni-K line impossible, and hence, boosting the Ni absolute abundance in pn."
] |
[
"Uses",
"Uses"
] |
[
[
739,
758
]
] |
[
[
639,
737
],
[
761,
996
]
] |
2020AandA...643A..93H__Lin_et_al._2016a_Instance_1
|
In recent years, type Ia supernovae (SNe Ia; Amanullah et al. 2010; Suzuki et al. 2012; Betoule et al. 2014; Scolnic et al. 2018) have been widely employed to test cosmic isotropy. Antoniou & Perivolaropoulos (2010) searched for the preferred direction of anisotropy for the Union2 sample by adopting the hemisphere comparison (HC) method (Schwarz & Weinhorst 2007). They found a maximum accelerating expansion rate, which corresponds to a preferred direction of anisotropy. After that, Mariano and Perivolaropoulos (Mariano & Perivolaropoulos 2012) found a possible preferred anisotropic direction at the 2σ level using the Union2 sample, but by employing the dipole fitting (DF) method. Since then, these two methods have been widely used to explore the cosmic anisotropy (Cai & Tuo 2012; Cai et al. 2013; Zhao et al. 2013; Li et al. 2013; Heneka et al. 2014; Bengaly et al. 2015; Andrade et al. 2018; Sun & Wang 2019) by investigating observational data of, for instance, the Union2.1 sample (Yang et al. 2014; Javanmardi et al. 2015; Lin et al. 2016a), the Joint Light-Curve Analysis (JLA) sample (Lin et al. 2016b; Chang et al. 2018a; Wang & Wang 2018), the Pantheon sample (Sun & Wang 2018), gamma-ray bursts (GRBs; Wang & Wang 2014), galaxies (Zhou et al. 2017), as well as gravitational wave and fast radio bursts (Qiang et al. 2019; Cai et al. 2019). Using HC and DF methods, Zhao et al. (2019) studied the cosmic anisotropy via the Pantheon sample. They found that the SDSS sample plays a decisive role in the Pantheon sample. It may imply that the inhomogeneous distribution has a significant effect on the cosmic anisotropy (Chang et al. 2018b). This opinion was also presented by Sun & Wang (2019). Their conclusions show that the effect of redshift on the result is weak and there is a negligible anisotropy when making a redshift tomography. Deng & Wei (2018a) tested the cosmic anisotropy with the Pantheon sample, but by using the following three methods: the HC method, the DF method, and Healpix1 (Górski et al. 2005). They also performed a cross check. There are two preferred directions from the HC method. In adopting the DF method and Healpix, they found no noticeable anisotropy. They also compared the HC method with the DF method by using the JLA sample (Deng & Wei 2018b) and found that the results of these two methods have not always been approximately coincident with each other. In order to better test the cosmic isotropy, the best way would be to add new samples with a relatively homogeneous distribution.
|
[
"Lin et al. 2016a"
] |
[
"Since then, these two methods have been widely used to explore the cosmic anisotropy",
"by investigating observational data of, for instance, the Union2.1 sample"
] |
[
"Background",
"Background"
] |
[
[
1038,
1054
]
] |
[
[
689,
773
],
[
921,
994
]
] |
2018MNRAS.473.3810Y__Mitrushchenkov_et_al._2017_Instance_1
|
The lack of data on inelastic processes due to collisions with neutral hydrogen atoms has been a major limitation on modelling of F-, G- and K-star spectra in statistical equilibrium, and thus to reliably proceeding beyond the assumption of local thermodynamic equilibrium (LTE) in analysis of stellar spectra and the determination of elemental abundances. This problem has been well documented, e.g. see Lambert (1993); Barklem (2016a) and references therein. Significant progress has been made in recent times through detailed full-quantum scattering calculations, based on quantum chemical data, for the cases of simple atoms such as Li, Na, Mg and Ca (Belyaev & Barklem 2003; Barklem, Belyaev & Asplund 2003; Belyaev et al. 2010; Barklem et al. 2010; Belyaev et al. 2012; Barklem et al. 2012; Mitrushchenkov et al. 2017). These calculations have demonstrated the importance of the ionic-covalent curve crossing mechanism leading naturally to charge transfer processes (mutual neutralization and ion-pair production), in addition to excitation and de-excitation processes. The importance of this mechanism has allowed various simplified model approaches to be developed, which may be used in cases where suitable quantum chemistry data are not been available. In particular a semi-empirical model has been employed for Al, Si, Be and Ca (Belyaev 2013a,b; Belyaev, Yakovleva & Barklem 2014b; Yakovleva, Voronov & Belyaev 2016; Belyaev et al. 2016), and a theoretical model based on a two-electron asymptotic linear combinations of atomic orbitals (LCAO) approach, has also been employed for Ca (Barklem 2016b, 2017). Comparisons of the two methods show quite good agreement and reasonable agreement with the full quantum calculations is found, particularly for the most important processes with the largest rates (Barklem 2016b, 2017; Mashonkina, Sitnova & Belyaev 2017; Mitrushchenkov et al. 2017). Thus, the model approaches provide a useful route for obtaining estimates of the rates for these processes for many elements of astrophysical interest.
|
[
"Mitrushchenkov et al. 2017"
] |
[
"Significant progress has been made in recent times through detailed full-quantum scattering calculations, based on quantum chemical data, for the cases of simple atoms such as Li, Na, Mg and Ca"
] |
[
"Background"
] |
[
[
797,
823
]
] |
[
[
461,
654
]
] |
2019MNRAS.482.3803P__Schaerer_&_Vacca_1998_Instance_1
|
The broad blue WR bump feature around 4686 Å was searched in all the star-forming regions in the galaxies. A clear detection of the broad blue WR bump was made only in two star-forming regions hosted by SBS 1222+614 as shown in Fig. 8. This detection is in good agreement with Shirazi & Brinchmann (2012), who have also reported the WR features in SBS 1222+614. The broad blue bump consists of a blend of C iii/C iv λ4650, 4658, N iii λ4634, 4640, [Ar iv] λ4711, 4740, and He ii λ4686 emission lines. This detection generally indicates a good number (102–105) of young WR stars in the galaxy (e.g. Kunth & Sargent 1981; Kunth & Schild 1986). The blue bump appears mainly due to the presence of late-type WN (WNL) and early-type WC (WCE) stars (Schaerer & Vacca 1998). The red WR bump around 5808 Å is also expected in the WR galaxies. We also possibly identified the red bump feature in the SBS 1222+614 (# a + b) region as shown in Fig. 9. The red WR bump appears mainly due to the presence of emission lines [N ii] λ5795 and He i λ5875 from the WCE stars. The red bump is rarely detected in WR galaxies as it contains very weak emission line and is expected in high-mettalicity region (Guseva, Izotov & Thuan 2000). Although we could not detect the WR features in the star-forming regions of other galaxies in our sample, the optical SDSS spectrum for the brightest regions of knot #b and #a in IC 3521 and CGCG 0410-023, respectively, shows a detection of WR features (Brinchmann et al. 2008). The SDSS spectrum also shows the detection of the broad WR blue bump having a relatively lower strength in the knot #b of CGCG 038-051. These detections are missed out in this study most likely due to the low SNR in the blue part of the spectra. The starburst ages estimated to be very young (≤6 Myr; see Table 3) for these star-forming regions including those present in SBS 1222+614 are consistent with the detection of the WR features, which appears only during the very early periods of star formation. Although a few other star-forming regions such as the knot #a in IC 3521 and CGCG 038-051 and knot #c in CGCG 041-023 showing a very young starburst of ≤6 Myr are expected to have the WR features, it is not clear from the present data if these regions also host WR stars or not. Some star-forming regions such as knot #c and b in CGCG 038-051 and CGCG 041-023, respectively, show the starburst of ages of ∼10 Myr, indicating that they have probably completed their WR phases. Overall, the detections of the WR features from our own observations and those from the SDSS data from at least one star-forming region in each galaxy in the sample suggest that the sample galaxies are undergoing young massive star formation phase having a significant population of WR stars.
|
[
"Schaerer & Vacca 1998"
] |
[
"The blue bump appears mainly due to the presence of late-type WN (WNL) and early-type WC (WCE) stars"
] |
[
"Background"
] |
[
[
744,
765
]
] |
[
[
642,
742
]
] |
2020AandA...640A.133R___et_al._2017_Instance_1
|
The full sample for which the circumstellar CO line emission will be modeled consists of the ∼180 C-, M-, and S-type AGB stars analyzed in Schöier & Olofsson (2001), González Delgado et al. (2003), and Ramstedt et al. (2006) together with additional sources presented in Danilovich et al. (2015). In this initial paper, the new interferometric data for the southern M- and C-type stars are presented. Some of the available sample statistics for the full ∼180 star DEATHSTAR sample are shown in Fig. 1. The distance distribution (Fig. 1, middle) is compared with the estimated distribution in the solar neighborhood (Jura 1990; Jura & Kleinmann 1992; Jura et al. 1993). The estimated distribution is derived from 2MASS and ground-based observations (Jura & Kleinmann 1990) and assumes a smooth distribution of 40 C-type stars kpc−2, a scale height of 200 pc, and that there are a third as many S-type as C-type stars. For the full ∼180 star DEATHSTAR sample, the C-type stars from Schöier & Olofsson (2001) are all brighter than K = 2 mag. The M-type sample consists of the non-Mira stars from the General Catalog of Variable Stars (GCVS; Samus’ et al. 2017) with quality flag 3 in the IRAS 12, 25, and 60 μm bands and 60 μm flux ≳3 Jy with the addition of the Mira stars in González Delgado et al. (2003). The S-type sample also consists of stars that have good quality flux measurements in the IRAS 12, 25, and 60 μm bands, that are found in the General Catalog of Galactic S stars, and that are detected in Tc and are hence intrinsic. The completeness of the S-type sample is discussed in Ramstedt et al. (2009) and it is thought to be complete out to 600 pc. Furthermore, stars of all three spectral types are only included in the sample if they are detected in CO radio line emission, which could be reproduced under the assumption of spherical symmetry. Sources that show strongly asymmetric line profiles when observed with single-dish telescopes, or with known CO-detached shells, are hence not included (e.g., R Scl, U Ant, EP Aqr, and π1 Gru). In this paper, we also exclude stars that have previously been observed with ALMA; however, they will be included in the future analysis. The lower panel of Fig. 1 shows that the stellar and wind parameters of the full ∼180 star sample cover the ranges expected for AGB stars. As expected, fewer stars are found at the high end. The sample is biased to mass-losing stars since only stars that are previously detected in CO radio line emission are included. It is also likely that the full range of AGB masses is not covered simply because higher mass stars are rare. It is our assessment that the full ∼180 star DEATHSTAR sample is representative of Galactic mass-losing AGB stars and covers the relevant ranges of wind and stellar parameters to provide the necessary constraints for theoretical models.
|
[
"Samus’ et al. 2017"
] |
[
"The M-type sample consists of the non-Mira stars from the General Catalog of Variable Stars (GCVS;",
"with quality flag 3 in the IRAS 12, 25, and 60 μm bands and 60 μm flux ≳3 Jy"
] |
[
"Uses",
"Uses"
] |
[
[
1138,
1156
]
] |
[
[
1039,
1137
],
[
1158,
1234
]
] |
2021MNRAS.504.2168G__Steiner_et_al._2011_Instance_1
|
Finally, we attempt to characterize the reflection component using the full 2–35 keV spectra with a sophisticated model [M4: ${\tt{\rm constant}}$*${\tt{\rm TBabs}}$*(${\tt{\rm simplr}}$*${\tt {\rm kerrbb2}}$+${\tt{\rm kerrconv}}$*(${\tt{\rm ireflect}}$*${\tt{\rm simplc}}$)), to evaluate the impact on the spin measurement and understand the changes of the accretion flow and the interaction between the disc and the corona. This model features a self-consistent treatment of the thermal, Compton scattering and the reflection component: ${\tt {\rm kerrbb2}}$ describes the thermal component and supplies the seed photons for ${\tt{\rm simplr}}$ (a modified version of ${\tt{\rm simpl}}$, Steiner et al. 2011) to generate the Compton component; while a portion of the Compton component will escape to reach an observer, the remains (refer as ${\tt{\rm simplc}}$, Steiner et al. 2011) will strike back to the disc to generate the reflected component. The reflection fraction Rref in ${\tt{\rm ireflect}}$ (Magdziarz & Zdziarski 1995), defined as the ratio of the Compton photons striking back to the disc to that escaping to infinity, is restricted to negative value thereby only the reflected component is returned by ${\tt{\rm ireflect}}$. It is linked to the reflection constant parameter x in ${\tt{\rm simplr}}$ via the relation x = 1 + |Rref| (Gou et al. 2011). We set the elemental abundance to unity and the iron abundance AFe to five times the solar abundance (Bharali et al. 2019; Buisson et al. 2019; Xu et al. 2020). The disc temperature Tin is fixed at the value returned by ${\tt{\rm diskbb}}$ (M1, refer to Gou et al. 2011). The ionization parameter ξ is fixed at 1000 (i.e. log(ξ) = 3, Xu et al. 2020; Buisson et al. 2019), as it is difficult to be constrained. Finally we use ${\tt{\rm kerrconv}}$ (Brenneman & Reynolds 2006) to apply relativistic effects assuming an unbroken emissivity profile with index q = 3. The key parameters in ${\tt {\rm kerrbb2}}$ and ${\tt{\rm kerrconv}}$ are linked together.
|
[
"Steiner et al. 2011"
] |
[
"This model features a self-consistent treatment of the thermal, Compton scattering and the reflection component: ${\\tt {\\rm kerrbb2}}$ describes the thermal component and supplies the seed photons for ${\\tt{\\rm simplr}}$ (a modified version of ${\\tt{\\rm simpl}}$,",
"to generate the Compton component;"
] |
[
"Uses",
"Uses"
] |
[
[
690,
709
]
] |
[
[
426,
689
],
[
711,
745
]
] |
2019MNRAS.484.1100M__Dekel_&_Birnboim_2006_Instance_1
|
The stream velocity is proportional to the halo virial velocity, which can be related to the sound speed of gas at the virial temperature, yielding Mb ∼ 1. The density contrast is obtained by assuming pressure equilibrium between hot gas at Tb ∼ Tv, the virial temperature of an NFW halo, and cold gas at Ts ∼ 104K, set by the steep drop in the cooling rate below that temperature (Sutherland & Dopita 1993). If both the halo and the stream are roughly isothermal, then this ratio is constant throughout the halo. In practice, the stream temperature can be as high as ∼3 × 104 K due to photoheating from the UV background (e.g. Goerdt et al. 2010), while the post-shock temperature in the hot halo near Rv may only be ∼0.5Tv (Dekel & Birnboim 2006, P18). Finally, the stream radius is related to its velocity and density through the mass accretion rate along the stream, ${\dot{M}}_{\rm s}\simeq \pi R_{\rm s}^2 \rho _{\rm s}V_{\rm s}$. Cosmological simulations suggest that this is typically a fixed fraction of the total accretion rate on to the halo virial radius (Danovich et al. 2012), which is constrained by cosmology (e.g. Dekel et al. 2013). The final expressions, including uncertainties in model parameters, are
(42)
\begin{eqnarray*}
M_{\rm b}\simeq 0.75-2.25,
\end{eqnarray*}
(43)
\begin{eqnarray*}
\delta \simeq (10-100)\times M_{12}^{2/3} (1+z)_3^{-1},
\end{eqnarray*}
(44)
\begin{eqnarray*}
\frac{R_{\rm s}}{R_{\rm v}} \simeq (0.01-0.1)\times \left(\delta _{75} M_{\rm b,1.5}\right)^{-1/2},
\end{eqnarray*}
where $M_{12}=M_{\rm v}/10^{12}\, \mathrm{M}_\odot$, (1 + $z$)3 = (1 + $z$)/3, δ75 = δ/75, and Mb,1.5 = Mb/1.5. We note that δ ∼ 10 implies that the background gas has a temperature of ${\sim } 3\times 10^5\, {\rm K}$, which would make it thermally unstable unless it is illuminated by a very hard ionizing background (e.g. Efstathiou 1992, though see also Binney, Nipoti & Fraternali 2009, who argue that buoyancy effects may stabilize the halo against thermal instability even without a photoionizing source). This implies that δ ≳ 30 is likely the most physically plausible regime.
|
[
"Dekel & Birnboim 2006"
] |
[
"while the post-shock temperature in the hot halo near Rv may only be ∼0.5Tv",
", P18)."
] |
[
"Uses",
"Uses"
] |
[
[
726,
747
]
] |
[
[
649,
724
],
[
747,
754
]
] |
2021AandA...653A..36M__Goulding_&_Alexander_(2009)_Instance_3
|
The SFG sample was constructed using the Great Observatories All-Sky LIRG Survey (GOALS sample, Armus et al. 2009), from which we extracted 158 galaxies, with data from Inami et al. (2013), who report the fine-structure lines at high resolution in the 10 − 36 μm interval, and Stierwalt et al. (2014), who include the detections of the H2 molecular lines and the PAH features at low spectral resolution. For those galaxies in the GOALS sample that have a single IRAS counterpart, but more than one source detected in the emission lines, we have added together the line or feature fluxes of all components, to consistently associate the correct line or feature emission to the total IR luminosity computed from the IRAS fluxes. To also cover lower luminosity galaxies, as the GOALS sample only includes luminous IR galaxies (LIRGs) and ultra-luminous IR galaxies (ULIRGs), we included 38 galaxies from Bernard-Salas et al. (2009) and Goulding & Alexander (2009), to reach the total sample of 196 galaxies with IR line fluxes in the 5.5 − 35 μm interval in which an AGN component is not detected. For the Bernard-Salas et al. (2009), Goulding & Alexander (2009), and the GOALS samples, we excluded all the composite starburst-AGN objects identified as those with a detection of [NeV] either at 14.3 or 24.3 μm. It is worth noting that the original samples from Goulding & Alexander (2009) and Bernard-Salas et al. (2009) have spectra solely covering the central region of the galaxies. To estimate the global SFR, we corrected the published line fluxes of the Spitzer spectra by multiplying them by the ratio of the continuum reported in the IRAS point source catalogue to the continuum measured on the Spitzer spectra extracted from the CASSIS database (Lebouteiller et al. 2015). We assumed here that the line emission scales (at first order) with the IR brightness distribution. In particular, we considered the continuum at 12 μm for the [NeII]12.8 μm and [NeIII]15.6 μm lines, and the continuum at 25 μm for the [OIV]25.9 μm, [FeII]26 μm, [SIII]33.5 μm, and [SiII]34.8 μm lines. This correction was not needed for the AGN sample and the GOALS sample because of the greater average redshift of the galaxies in the 12MGS and GOALS samples. In particular, the 12MGS active galaxy sample has a mean redshift of 0.028 (Rush et al. 1993), while the GOALS sample has a mean redshift of 0.026. The galaxies presented by Bernard-Salas et al. (2009) have instead an average redshift of 0.0074, while the sample by Goulding & Alexander (2009) has an average redshift of 0.0044. For the other lines in the 10 − 36 μm interval, Goulding & Alexander (2009) did not report a detection, and we used the data presented in Bernard-Salas et al. (2009) for a total of 15 objects. Both Bernard-Salas et al. (2009) and Goulding & Alexander (2009) reported data from the high-resolution Spitzer-IRS spectra. Data in the 50 − 205 μm interval were taken from Díaz-Santos et al. (2017). For the GOALS sample, 20 starburst galaxies were taken from Fernández-Ontiveros et al. (2016), and 23 objects were taken from the ISO-LWS observations of Negishi et al. (2001). As a result, we obtained a total sample of 193 objects. Lastly, the PAH features’ fluxes were measured from the low-resolution Spitzer-IRS spectra by Brandl et al. (2006), including 12 objects from the sample of Bernard-Salas et al. (2009) and 179 objects from Stierwalt et al. (2014).
|
[
"Goulding & Alexander (2009)"
] |
[
"It is worth noting that the original samples from",
"and Bernard-Salas et al. (2009) have spectra solely covering the central region of the galaxies."
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
1359,
1386
]
] |
[
[
1309,
1358
],
[
1387,
1483
]
] |
2018AandA...613A...7Y__Kostogryz_et_al._(2016)_Instance_1
|
where a is the area of the i, j pixel (constant for a regular grid), μ is the anglebetween the surface normal and the line of sight to the observer, ϕ is the polar angle of a system with the origin at the disk center, F is the total stellar flux, and q and u are normalized Stokes parameters. We note that the relative flux is normalized to the total flux of an unspotted photosphere in our analysis. The center-to-limb variations ofintensity I(μij) and polarization P(μij) were found through trilinear interpolation using the look-up tables from Kostogryz & Berdyugina (2015) and Kostogryz et al. (2016), according to the selected wavelength, surface gravity, and temperature of a star. Kostogryz & Berdyugina (2015) calculated the tables for continuum spectra of FGK stars (Teff = 4500 K–6900 K, log g = 2.0–5.0, and wavelength range 4000–7000 Å) for the Phoenix grid of plane-parallel models, and in Kostogryz et al. (2016) similar calculations were made for a wider range of models (Teff = 4000 K–7000 K and log g = 1.0–5.5) assuming a spherical stellar atmosphere. In addition, we used unpublished data on intensity and polarization variations for both plane-parallel and spherical cooler models with temperatures down to 3000 K, obtained with the same code by Kostogryz (2016, priv. comm.). It should be noted that these data for cooler stars do not include atomic and molecular absorption lines, which can lead to overestimated polarization values, especially for blue wavelengths and for dwarfs with higher surface gravities (depending on the selected wavelength). Evidently, a proper spectrum synthesis code is needed to calculate the intrinsic polarization for these cases. However, the depolarizing effect from absorption lines may not be very significant because when the absorption line suppresses the polarization, the intensity is also reduced. In turn, this will affect the normalized polarization parameters that we calculate less strongly. Additionally, coherent scattering processes, especially in molecular bands, as seen on the Sun (e.g., Berdyugina et al. 2002) probably increase polarization further, which counteracts the effect of the absorption lines.
|
[
"Kostogryz et al. (2016)"
] |
[
"The center-to-limb variations ofintensity I(μij) and polarization P(μij) were found through trilinear interpolation using the look-up tables from Kostogryz & Berdyugina (2015) and",
", according to the selected wavelength, surface gravity, and temperature of a star."
] |
[
"Uses",
"Uses"
] |
[
[
581,
604
]
] |
[
[
401,
580
],
[
604,
687
]
] |
2019MNRAS.484..814G__Tápai,_Zoltán_&_László_2015_Instance_1
|
There are many methods, beyond those already employed here, that can be used to identify, and probe the masses of, black holes. This includes reverberation mappings of AGNs (e.g. Bahcall, Kozlovsky & Salpeter 1972; Blandford & McKee 1982; Netzer & Peterson 1997), the ‘fundamental plane of black hole activity’ (Merloni et al. 2003; Falcke, Körding & Markoff 2004), spectral modelling of the high-energy X-ray photon coming from the hot accretion discs around IMBHs (Pringle & Rees 1972; Narayan & Yi 1995), high-ionization optical emission lines (Baldwin et al. 1981; Kewley et al. 2001), and high-spatial-resolution observations of maser emission using radio and millimetre/submillimetre interferometry (e.g. Miyoshi et al. 1995; Greenhill et al. 2003; Humphreys et al. 2016; Asada et al. 2017). In addition, the merging of black holes is now quite famously known to produce gravitational radiation during their orbital decay (Abbott et al. 2016). The merging of galaxies containing their own central IMBH is similarly expected to result in the eventual merging of these black holes. The Kamioka Gravitational Wave Detector (KAGRA; Aso et al. 2013) will be a 3-km-long underground interferometer in Japan capable of detecting the gravitational radiation emanating from collisions involving black holes with masses up to 200 M⊙ (Tápai, Zoltán & László 2015). The planned Deci-Hertz Interferometer Gravitational wave Observatory (DECIGO; Kawamura et al. 2011) and the European, Laser Interferometer Space Antenna (LISA) Pathfinder mission15 (Anza et al. 2005; McNamara 2013), with their greater separation of mirrors, will be able to detect longer wavelength gravitational waves and thus better reach into the domain of intermediate-mass and supermassive black hole mergers, the latter of which are currently being searched for via ‘pulsar timing arrays’ (PTAs) (e.g. Hobbs et al. 2010; Kramer & Champion 2013; Shannon et al. 2015). A key constraint to the expected detection threshold of such signals from PTAs – in particular the background of cosmic ripples from the merger of massive black holes (themselves arising from the merger of galaxies) – is the (black hole)-to-(host galaxy/bulge) mass ratio (see equation 4 for spiral galaxies). An additional source of long-wavelength gravitational radiation will arise from the inspiral of compact stellar-mass objects, such as neutron stars and black holes, around these IMBHs (Mapelli et al. 2012). It is reasonable to expect that the densely packed nuclear star clusters, which coexist with low-mass SMBHs (e.g. González Delgado et al. 2008; Seth et al. 2008; Graham & Spitler 2009), will similarly surround many IMBHs. Gravitational radiation and the gravitational tidal disruption of ill-fated stars that venture too close to these black holes (Komossa et al. 2009, Komossa 2013, and references therein; Zhong, Berczik & Spurzem 2015; Stone & Metzger 2016; Lin et al. 2018) are therefore expected from these astrophysical entities. There is, therefore, an array of future observations that could yield further confidence and insight into the realm of IMBHs.
|
[
"Tápai, Zoltán & László 2015"
] |
[
"The Kamioka Gravitational Wave Detector",
"will be a 3-km-long underground interferometer in Japan capable of detecting the gravitational radiation emanating from collisions involving black holes with masses up to 200 M⊙"
] |
[
"Background",
"Background"
] |
[
[
1330,
1357
]
] |
[
[
1086,
1125
],
[
1151,
1328
]
] |
2016ApJ...831...11S__Honeycutt_1992_Instance_1
|
The secondary eclipses of the binary are 0.03–0.04 mag deep, and so relatively small data reduction issues can have large effects on the fidelity of the light curve. Our brightness measurements were derived from aperture photometry using DAOPHOT (Stetson 1987), although we took several additional steps to increase the precision of the results. We conducted a curve-of-growth analysis of 12 apertures photometered per star using DAOGROW (Stetson 1990) in order to correct all measured stars to a uniform large aperture. We then attempted to unify the data for each filter to a consistent zero point by using ensemble photometry (Sandquist et al. 2003; Honeycutt 1992). This essentially uses all measured stars on the frame to determine magnitude offsets resulting from differences in exposure time, airmass, atmospheric transparency, and the like. Our implementation iteratively fits for position-dependent corrections that result from variations from point-spread function across the frame. These steps each brought noticeable reductions in the amount of scatter in the light curves. However, we still found that the shape of the ground-based light curves of the secondary eclipse did not match what was expected from the Kepler observations. After some investigation, we found that features in the light curve were correlated with those for the brightest star near KIC 9777062 in our images. (The star was about 70 pixels or 28″ distant on the sky, so there was no significant overlap of the point-spread function when the seeing was generally 4–8 pixels FWHM.) We found that after subtracting the light curve of this star the out-of-eclipse light levels and the shapes of the primary and secondary eclipse light curves were much more consistent from night to night. When more distant stars on the images were tested, we found a much poorer degree of correlation. We conclude that short length-scale variations were not being corrected for (and could not be corrected for, due to lack of star sampling on the image) with whole-image zero-point corrections or image-scale position-dependent corrections. The variations were of around 0.01 mag size, but could significantly affect the secondary eclipse light curves. We could not identify any features on our flat-field images that could explain the light-curve variations, and they occurred whether dome flats, twilight flats, or hybrid flats (combining the smoothed large-scale variations from the twilight flats and small-scale variations from dome flats) were used. The ground-based eclipse observations are shown in the bottom rows of Figure 1.
|
[
"Honeycutt 1992"
] |
[
"We then attempted to unify the data for each filter to a consistent zero point by using ensemble photometry",
"This essentially uses all measured stars on the frame to determine magnitude offsets resulting from differences in exposure time, airmass, atmospheric transparency, and the like. Our implementation iteratively fits for position-dependent corrections that result from variations from point-spread function across the frame."
] |
[
"Uses",
"Background"
] |
[
[
653,
667
]
] |
[
[
521,
628
],
[
670,
992
]
] |
2018AandA...615A.161M__Hartmann_1904_Instance_1
|
δ Ori Aa+Ab = Mintaka Aa+Ab = 34 Ori Aa+Ab = HD 36486 A+B = BD −00 983 A+B = ALS 14 779 A+B. δ Ori is Orion’s Belt westernmost star and a multiple system at the center of the Mintaka cluster (Caballero & Solano 2008). It has two close bright visual components (Aa and Ab) and two distant dim ones (B and C) that will not be considered here. Aa is itself composed of two spectroscopic components (Aa1+Aa2) in an eclipsing orbit with a5.7325 d period, while Ab is two magnitudes fainter than the two Aa components and located 0.′′ 267 away in 1993 and moving away from the primary (Hartmann 1904; Harvin et al. 2002). We observed the system on two consecutive nights and in both cases we were able to spatially resolve Ab from Aa, this being one of the two systems in this paper (σ Ori is the other one) where the power of Lucky Spectroscopy becomes more apparent. The two spectra for Aa yield the same spectral type as for the combined Aa+Ab value from GOSSS I, O9.5 IINwk, but the lines are slightly narrower. When analyzing the data for GOSSS I, the combined spectrum was close to being (n), something that does not happen for the spatially resolved spectra (this has the additional advantage of giving us a more purespectrum for a classical O9.5 II classification standard). There are very small variations between the two epochs for Aa, likely due to the motion of Aa1 and Aa2 but the signature of the latter in the combined spectrum is very weak (Shenar et al. 2015). The two epochs for Ab yield the same spectral type, O9.7 III:(n), and in both cases there are indications of only a slight residual contamination from Aa, with the second epoch being noisier. The (n) qualifier indicates that Ab is a fast rotator, as previously noted by Harvin et al. (2002) and Shenar et al. (2015). The differences between the spectral types of Aa (Aa1+Aa2) and Ab are consistent with the Teff and log g differences measured by Shenar et al. (2015) using spectral disentangling (as opposed to the spatial deconvolution used here). Richardson et al. (2015) also spatially resolved Aa and Ab in the UV using HST/STIS and obtained very similar values of Teff around 31 kK for Aa1 and Ab, with error bars close to 2000 K, which is also consistent with our spectral classifications. That paper using UV data is the only one we have found where δ Ori Aa+Ab was spatially deconvolved, making our result the first time it has been done in the optical.
|
[
"Hartmann 1904"
] |
[
"Aa is itself composed of two spectroscopic components (Aa1+Aa2) in an eclipsing orbit with a5.7325 d period, while Ab is two magnitudes fainter than the two Aa components and located 0.′′ 267 away in 1993 and moving away from the primary"
] |
[
"Background"
] |
[
[
580,
593
]
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[
[
341,
578
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2019MNRAS.489.4669S__Bigiel_et_al._2010_Instance_1
|
In Fig. 7 we compare UGC 1378’s SFR density versus gas surface density (the Schmidt–Kennicutt relation) to data in the literature. The gas surface density corresponds to H i calculated from the 0th moment map from Mishra et al. (2017) in the same areas as SFR density. Points for the HSB and LSB discs are plotted as black and grey circles, respectively. We plot the mean SFR and H i surface density for the entire galaxy with a large open circle. The black line corresponds to the relation with an exponent of 1.4 found by Kennicutt (1998). Triangles give results for LSB galaxies published by Wyder et al. (2009), and bright and faint crosses show normal spiral galaxies from Kennicutt (1998) – total and H isurface densities. A blue line shows the best-fitting relation for the H isurface density of Bluedisk galaxies from Roychowdhury et al. (2015). We also plot the SFR in the outer regions of spiral galaxies (Bigiel et al. 2010, square symbols). In Fig. 7 the UGC 1378 measurements lie between normal spirals and LSB galaxies. The HSB disc data lie above the relation plotted for normal spirals, possibly indicating that the SFR is boosted by the bar driving gas to the star-forming rings. We cannot account for molecular gas since there are no available measurements for UGC 1378. The contribution of molecular gas would likely move the HSB disc of UGC 1378 towards the locus of normal galaxies. Because the HSB SFR of UGC 1378 is close to the predicted SFR from the Kennicutt (1998) relation obtained from H i densities (faint crosses in Fig. 7). The LSB disc of UGC 1378 lies below the correlation and accounting for molecular gas would only increase the deviation from the normal Schmidt–Kennicutt relation. Similar deviations are observed in ‘classical’ LSB galaxies, Bluedisk galaxies (Roychowdhury et al. 2015), outer parts of HSB spiral galaxies (Bigiel et al. 2010), and H idiscs in early-type galaxies (Yıldız et al. 2017). These deviations for LSB galaxies are at least partially explained by their lower gas densities leading to lower SFRs (Abramova & Zasov 2011). A recent episode of gas accretion on to the disc of UGC 1378 may also contribute to a lower SFR if the gas is not yet fully participating in the star formation. Lutz et al. (2017) studied a sample of very H i rich galaxies and proposed that very high specific angular momentum in H irich galaxies prevents the accreted gas from being transported to the mid-plane of the disc and being converted into stars. This mechanism may act to preserve giant gaseous discs.
|
[
"Bigiel et al. 2010"
] |
[
"We also plot the SFR in the outer regions of spiral galaxies",
"square symbols)."
] |
[
"Uses",
"Uses"
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[
[
916,
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854,
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936,
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2021MNRAS.508.2019B__Gürkan_et_al._2018_Instance_1
|
Assuming that non-jetted H ii galaxies are powered by SF, the relation between radio luminosities in such galaxies and MBH can be deciphered as a link between thermal/non-thermal SF luminosity and BH mass, in the form $L_{\rm SF} \propto M_{\rm BH}^{0.61}$, which has the largest scatter of the observed linear regressions, due to the large uncertainties of MBH estimates at low values and numerous radio non-detections. This unprecedented relation is likely the result of the observed link found between SF indicators such as radio, far-infrared, optical or line luminosities, and galaxy mass (e.g. Mauch & Sadler 2007; Gürkan et al. 2018). Moreover, the SF contribution in the radio band with respect to the AGN activity is expected to increase with galaxy mass (Aird, Coil & Georgakakis 2019). In fact, the observed correlation suggests that the radio production due to SF broadly increases with BH mass, as the latter quantity is approximately a constant proportion of the galaxy mass (from a third down to a fifth in the lowest MBH ∼ 105 M⊙, Häring & Rix 2004; Reines & Volonteri 2015; Martín-Navarro & Mezcua 2018), i.e. the stellar content: more stars likely produce more radio emission. This is consistent with the idea that SF, although mostly included in the disc (Bluck et al. 2020) and known to scale with galaxy mass (Speagle et al. 2014), would possibly also scale with MBH. More precisely, the nuclear starburst is plausibly set by the available gas mass present in the NSC (MNSC, Fernández-Ontiveros, Prieto & Acosta-Pulido 2009; Neumayer, Seth & Böker 2020), which scales with the galaxy dynamical mass ($M_{\rm NSC} \propto M _{\mathrm{ dyn}}^{0.55}$, Scott & Graham 2013), which in turn respond to the SMBH gravitational well (Cen 2015; Pitchford et al. 2016). Such a sequence of links eventually enacts the radio-MBH and radio-optical relations observed for star-forming galaxies, e.g setting the fraction of radio emission with respect to the photoionizing energy from young stellar populations and SN products.
|
[
"Gürkan et al. 2018"
] |
[
"This unprecedented relation is likely the result of the observed link found between SF indicators such as radio, far-infrared, optical or line luminosities, and galaxy mass"
] |
[
"Uses"
] |
[
[
621,
639
]
] |
[
[
421,
593
]
] |
2020ApJ...896..169P__Dullemond_&_Dominik_2004_Instance_1
|
Pre-main-sequence stars are typically surrounded by protoplanetary disks, and since such disks are generally highly optically thick at optical/near-infrared (NIR) wavelengths, they may cast shadows on their surroundings. Such disk shadows have been observed on a wide range of angular scales; recently, high-contrast imaging has revealed shadows cast on outer disks by misaligned inner disks on subarcsecond scales (e.g., Marino et al. 2015; Benisty et al. 2018; Casassus et al. 2018). Such shadows reveal that angular distortions are common in inner disks. This is consistent with the existence of a class of self-shadowed disks in which the inner disk scale height is higher than that of the outer disk, leading to a general shadowing and related cooling of the outer disk (Dullemond & Dominik 2004; Dong 2015). A related type of disk shadow, but on a vastly larger angular scale of arcminutes, are shadows cast on large-scale reflection nebulosity. These great disk shadows can occur if a young, usually spatially unresolved, star-disk system is illuminating a reflection nebula, and are especially apparent for systems viewed close to edge-on (Hodapp et al. 2004; Pontoppidan & Dullemond 2005). Neither of these two types of disk shadow should be confused with silhouette disks that obscure background nebulosity, such as the Orion proplyds (O’dell 1998), or dark dust lanes in isolated edge-on disks (Burrows et al. 1996; Stapelfeldt et al. 1998; Duchêne et al. 2010). The projection of the disk onto a large reflection nebula can greatly magnify a small structure in the obscuring disk. Indeed, the apparent angular size of great disk shadows is only limited by the size of the reflection nebula illuminated by the central star, and may be orders of magnitude larger than the protoplanetary disk itself. Great disk shadows therefore present a unique opportunity to explore the geometry of disks on scales otherwise not resolved by direct imaging, primarily the disk scale height, inclination, and position angle.
|
[
"Dullemond & Dominik 2004"
] |
[
"Such shadows reveal that angular distortions are common in inner disks. This is consistent with the existence of a class of self-shadowed disks in which the inner disk scale height is higher than that of the outer disk, leading to a general shadowing and related cooling of the outer disk"
] |
[
"Similarities"
] |
[
[
776,
800
]
] |
[
[
486,
774
]
] |
2021MNRAS.501.3781R___2017_Instance_2
|
While spatially extended optical jets and bipolar CO molecular outflows have been observed in numerous Class 0/I protostars (e.g. Reipurth & Bally 2001; Bally 2016, and references therein), near-infrared high-resolution spectroscopy and spectroimaging observations in the past two decades have made it possible to study the kinematics of the outflowing gas and physical properties at the base of the jet within a few hundred au of the driving source in Class 0/I protostars (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These microjets are bright in [Fe ii] forbidden and H2 rovibrational emission lines, hence showing the presence of forbidden emission-line (FEL) regions and molecular hydrogen emission-line (MHEL) regions in low-mass Class 0/I protostars. While multiple low- and high-velocity components are observed in both MHELs and FELs, the higher velocity gas is slightly further offset from the driving source than the slower gas, and the kinematics of the H2 emission differs from [Fe ii] emission, revealing complicated kinematic structures. Evidence of H2 emission from cavity walls is also seen in some protostars, suggesting the presence of a wide-angled wind. Strong emission in the well-known accretion diagnostics of Paschen and Brackett hydrogen recombination lines is observed in protostars, with the ratio of the accretion luminosity to bolometric luminosity spanning from ∼0.1 to ∼1. The mass accretion and loss rates for Class 0/I low-mass protostars span the range of 10−6–10−8 M⊙ yr−1, and the derived jet efficiencies (ratio between mass ejection and accretion rates) range between ∼1 per cent and 10 per cent (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These measurements are within the range predicted by the magnetohydrodynamic jet launching models (e.g. Frank et al. 2014).
|
[
"Antoniucci et al.",
"2017"
] |
[
"The mass accretion and loss rates for Class 0/I low-mass protostars span the range of 10−6–10−8 M⊙ yr−1, and the derived jet efficiencies (ratio between mass ejection and accretion rates) range between ∼1 per cent and 10 per cent (e.g.",
"These measurements are within the range predicted by the magnetohydrodynamic jet launching models"
] |
[
"Background",
"Similarities"
] |
[
[
1887,
1904
],
[
1917,
1921
]
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[
[
1543,
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[
1956,
2053
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2021ApJ...916...57G__Gedalin_2017_Instance_1
|
Collisionless shocks are one of the most ubiquitous phenomena in space plasmas. The directed flow energy is converted into other forms at the shock front: ion and electron heating, particle acceleration, and magnetic field enhancement. One of the most important problems of the collisionless shock physics is prediction of the post-shock (downstream) state of the plasma and fields given the corresponding plasma and magnetic field state before the shock (upstream). In the absence of dissipative processes the relations between the upstream and downstream parameters (Rankine–Hugoniot relations, boundary conditions, and jump conditions) are just conservation laws of particle number, momentum, and energy. There is a complete hierarchy of scales in collisionless shocks and, accordingly, the boundary conditions depend on the scale at which the boundaries are placed. The standard Rankine–Hugoniot relations (RH) are formulated on the largest, magnetohydrodynamic (MHD) scale, at which the distributions thermalize (Kennel 1994). In many space plasma environments the ambient conditions change on spatial or temporal scales, which are too short to achieve such thermalization. Observational comparison of the upstream and downstream plasma is often done in various regions that do not satisfy the conditions for establishing the MHD RH. In the close vicinity of the shock front, ion distributions are significantly gyrophase dependent, and the jump conditions at the very transition should take this into account (Gedalin & Balikhin 2008; Gedalin 2016a). On larger scales or for measurements invoking substantial temporal and/or spatial averaging, the distributions become gyrotropic but can still remain anisotropic, since isotropization may be slow. The corresponding RH should take into account this anisotropy (Lyu & Kan 1986; Gedalin 2017; Gedalin et al. 2020). Additional complications arise when there are different populations of ions that undergo gyrotropization and isotropization on different scales. Eventual equilibration of temperatures may never happen. Even for different populations of ions of the same kind eventual merging into one thermal population may never be observed. The latter situation is expected to be typical when pickup ions (PUIs) are important constituents, as at the termination shock (Zank et al. 1996; Li et al. 2008; Richardson et al. 2008; Burrows et al. 2010; Matsukiyo & Scholer 2011, 2014; Ariad & Gedalin 2013; Jokipii 2013; Mostafavi et al. 2017, 2018; Kumar et al. 2018). In this case, MHD has to be replaced with a multispecies model, where each population obeys the corresponding conservation laws separately, while the magnetic field jump is obtained by combining all species together. In this approach one has to know the equations of state for each species. Typically, some form of thermodynamics motivated state equations are assumed (Florinski et al. 2009; Borovikov et al. 2011; Pogorelov et al. 2013; Zank et al. 2014; Mostafavi et al. 2018). This approach implicitly assumes that the population is in a kind of local thermodynamic equilibrium. However, observations show that this is not the case. In fact, the downstream ion distribution is determined primarily by ion dynamics in the macroscopic fields of the shock front. Interaction with the electromagnetic fluctuations plays a secondary role causing slow relaxation toward equilibrium, while binary collisions are negligible. The resolution of particle observations in the vicinity of the heliospheric termination shock corresponds to the scale on which the solar wind protons are possibly (but not certainly) isotropic but PUI are still anisotropic. Yet most studies of RH at the termination shock aim at the determination of jump conditions on the MHD scale, assuming isotropy of the distributions. Some previous analyses within a two-fluid model (solar wind and PUI) introduced heat flux and collisionless viscosity of PUI (Zank et al. 2014; Mostafavi et al. 2016, 2017, 2018; Zank 2016). This approach explicitly assumes that ion scattering within the shock transition is substantial so that the focused transport equation is valid throughout the shock (Zank et al. 2014). However, the shock transition is scatter-free (Toptyghin 1980; Drury 1983), so that the approach loses the physics of relation between the kinetic and MHD scales. Accordingly, the obtained magnetic profiles are monotonic and cannot explain observed overshoots (Mostafavi et al. 2018). Previous attempts to take into account the ion kinetics at a shock front assumed magnetic moment conservation (Fahr & Chalov 2008; Fahr et al. 2012; Fahr & Siewert 2013, 2015). However, it has been shown that magnetic moment is, in general, not conserved (Terasawa 1979; Gedalin 2020), and such approximation may be satisfactory only in the perpendicular regime. Proper establishment of the boundary conditions on MHD scales requires establishing a relation of the ultimately isotropic distributions to the collisionless ion dynamics within a shock front. This objective requires proper determination of downstream gyrophase averaged distributions, which are gyrotropic but anisotropic, with further relating these to the isotropic pressure (Gedalin et al. 2020). In the present paper, we numerically determine the parameters of these distributions, which will be used for modifying the Rankine–Hugoniot relation.
|
[
"Gedalin 2017"
] |
[
"On larger scales or for measurements invoking substantial temporal and/or spatial averaging, the distributions become gyrotropic but can still remain anisotropic, since isotropization may be slow. The corresponding RH should take into account this anisotropy"
] |
[
"Background"
] |
[
[
1833,
1845
]
] |
[
[
1557,
1815
]
] |
2020ApJ...899....4X__Zhou_et_al._2019_Instance_1
|
It is also very encouraging to note that recent analyses of heavy-ion reaction experiments in terrestrial laboratories and properties of NSs from multiple messengers have led to some new progress in constraining the Esym(ρ) up to about twice the saturation density. For example, shown in Figure 1 are the values of symmetry energy at 2ρ0, that is, Esym(2ρ0), from (1) the FOPI-LAND (Russotto et al. 2011) and (2) the ASY-EOS (Russotto et al. 2016) Collaborations by analyzing the relative flows and yields of light mirror nuclei as well as neutrons and protons in heavy-ion collisions at beam energies of 400 MeV/nucleon, (3) (Chen) an extrapolation of the systematics of low-density symmetry energy (Chen 2015), (4) (Zhang & Li) direct inversions of observed NS radii, tidal deformability, and maximum mass in the high-density EOS space (Zhang et al. 2018; Zhang & Li 2019a, 2020a), (5) (Xie & Li) a Bayesian inference from the radii of canonical NSs observed by using X-rays and gravitational waves from GW170817 (Xie & Li 2019), (6) (Zhou et al.) analyses of NS radii, tidal deformability, and maximum mass within an extended Skyrme Hartree–Fock approach (eSHF) (Zhou & Chen 2019; Zhou et al. 2019), (7) (Nakazato & Suzuki) analyzing cooling timescales of protoneutron stars as well as the radius and tidal deformability of GW170817 (Nakazato & Suzuki 2019), and (8) a Bayesian inference directly from the X-ray data of seven quiescent low-mass X-ray binaries in globular clusters (Baillot d’Etivaux et al. 2019). Despite the rather different assumptions and methods used in analyzing the different types of laboratory and observational data, it is very interesting to see that they all together are consistent with a fiducial value of Esym(2ρ0) = 47 MeV within the still relatively large error bars of the individual analyses. Moreover, several recent theoretical studies also predicted values of Esym(2ρ0) consistent with its fiducial value of 47 MeV. For example, an upper bound of Esym(2ρ0) ≤ 53.2 MeV was derived recently by Tong et al. (2020) by studying the radii of neutron drops using the state-of-the-art nuclear energy density functional theories. Quantum Monte Carlo calculations using local interactions derived from chiral EFT up to next-to-next-to-leading order predicted a value of Esym(2ρ0) ≈ 46 ± 4 MeV (Lonardoni et al. 2020), while the latest many-body perturbation theory calculations with consistent nucleon-nucleon and three-nucleon interactions up to fourth order in the EFT expansion predicted a value of Esym(2ρ0) ≈ 45 ± 3 MeV (Drischler et al. 2020). They are both consistent with the fiducial value of
and have much smaller error bars. It is worth noting that the chiral EFT is currently applicable to a maximum density of about 2ρ0.
|
[
"Zhou et al. 2019"
] |
[
"For example, shown in Figure 1 are the values of symmetry energy at 2ρ0, that is, Esym(2ρ0), from",
"(6)",
"analyses of NS radii, tidal deformability, and maximum mass within an extended Skyrme Hartree–Fock approach (eSHF)",
"Despite the rather different assumptions and methods used in analyzing the different types of laboratory and observational data, it is very interesting to see that they all together are consistent with a fiducial value of Esym(2ρ0) = 47 MeV within the still relatively large error bars of the individual analyses."
] |
[
"Uses",
"Uses",
"Uses",
"Similarities"
] |
[
[
1184,
1200
]
] |
[
[
266,
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[
1032,
1035
],
[
1050,
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[
1517,
1830
]
] |
2017ApJ...834..135T__Zolotov_et_al._2015_Instance_1
|
To explain the morphological transformation, two main evolutionary paths have been proposed in the literature. A slow cosmological path naturally follows from the strong redshift evolution of galaxy sizes,
(Mosleh et al. 2012; Newman et al. 2012; van der Wel et al. 2014b; Shibuya et al. 2015). Star-forming galaxies quench star formation and add to the passive population with approximately the same size in a later epoch (van Dokkum et al. 2015; Lilly & Carollo 2016). A second, fast path involves a downward transition in the mass–size plane, at approximately constant redshift (Barro et al. 2013, 2014; Dekel & Burkert 2014; Zolotov et al. 2015). This process requires a substantial “compaction” of the formally extended star-forming galaxies. One possible mechanism would be a major merger, which is known from observations and simulations to lead to substantial angular momentum redistribution, orbit reconfiguration, and mixing (Mihos & Hernquist 1996; Wuyts et al. 2010). Another possibility is an internal angular momentum redistribution within the star-forming disk. This process has been considered to be effective at high redshift (Noguchi 1999; Immeli et al. 2004a, 2004b; Elmegreen et al. 2008; Genzel et al. 2008; Bournaud et al. 2011), when galaxies are gas rich (Tacconi et al. 2013) and effective viscous dissipation leads to radial inward transport of gas and stars with a timescale of a few hundred megayears (Dekel et al. 2009) and buildup of a central dense core (bulge component) through circumnuclear concentration of gas. Nelson et al. (2016b) find in massive galaxies at z ∼ 1.4 that central 1 kpc regions are highly attenuated by dust and are responsible for half of the total star formation rate (SFR). In conjunction with morphological quenching (Martig et al. 2009; Genzel et al. 2014a), and powerful AGN outflows (Bower et al. 2006; Croton et al. 2006; Förster Schreiber et al. 2014; Genzel et al. 2014b), the compaction process may then lead to an inside-out quenching near the Schechter mass (Tacchella et al. 2015, 2016).
|
[
"Zolotov et al. 2015"
] |
[
"A second, fast path involves a downward transition in the mass–size plane, at approximately constant redshift",
"This process requires a substantial “compaction” of the formally extended star-forming galaxies."
] |
[
"Background",
"Background"
] |
[
[
636,
655
]
] |
[
[
478,
587
],
[
658,
754
]
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2018ApJ...865...22C__Criscuoli_et_al._2013_Instance_1
|
The strong linearity between the two indices shown in Figure 5 must be ascribed to their strong temperature dependence and to the proximity of the Mg ii core and FUV formation heights on one side, and Mg ii continuum and MUV on the other. Most of the radiation contributing to the FUV and MUV originates in the higher photosphere and chromosphere (Thuillier et al. 2012), where atmosphere models of quiet and magnetic structures present similar (and rather small) temperature and density gradients (c.f.r. Figure 1 in Fontenla et al. 2011). This means that, at least at first order, an increase of the activity corresponds to an increase of equal amount of the temperature and density from which both MUV and FUV, as well as Mg ii wings and core, originate. Under the assumption that the lambda-integrated fluxes can be described by a Plank function, it is easy to show that the color index must scale linearly with the temperature, at least for small temperature variations. Therefore, since the spectral brightness variations over a magnetic cycle do not exceed a few tens of K (Fontenla et al. 2011), the [FUV–MUV] color index must also scale linearly with the activity. In LTE, a linear dependence of the core/wing ratio is expected (e.g., Criscuoli et al. 2013), but such dependence is probably less obvious for the Mg ii h and k line, whose cores form in NLTE conditions. Nevertheless, collisions play an important role in the formation of the line (Linsky & Avrett 1970), making the core of this doublet strongly sensitive to temperature (e.g., Carlsson et al. 2015). Indeed, numerical magnetohydrodynamic simulations indicate that the brightness temperature derived from the Mg ii h and k cores is a good approximation for the plasma temperature at the corresponding formation heights (Leenaarts et al. 2013b). Moreover, the Bremen Mg ii index is derived from data of relatively modest spectral resolution, so that the cores are not well resolved. This means that the core intensities used to construct the Mg ii index must have a non-negligible contribution from several layers of the atmosphere, from the lower photosphere, where the response function of the spectral region between the two h and k peaks reaches its maximum (Uitenbroek 1997), to the base of the transition region (Thuillier et al. 2012; Leenaarts et al. 2013b; Carlsson et al. 2015), where the cores form. Employing again a Plank function to describe both the core and the wing intensities of the Mg ii doublet it is easy to show that the Mg ii index must scale linearly with temperature, at least for small temperature perturbations.
|
[
"Criscuoli et al. 2013"
] |
[
"In LTE, a linear dependence of the core/wing ratio is expected (e.g.,",
"but such dependence is probably less obvious for the Mg ii h and k line, whose cores form in NLTE conditions."
] |
[
"Compare/Contrast",
"Compare/Contrast"
] |
[
[
1244,
1265
]
] |
[
[
1174,
1243
],
[
1268,
1377
]
] |
2021ApJ...923..105L__Owens_et_al._2013_Instance_2
|
For the purpose of this study, we adopt the general concept that the global solar wind is generally composed of relatively steady high-speed solar winds from long-lived high latitude open coronal field regions. The high-speed solar winds from the northern and southern hemispheres are separated by a layer of—on average—slower, denser, more complicated, and variable outflows and structures (Gosling 1996; McComas et al. 2008). This boundary layer between the polar outflows has a latitudinal extent that is roughly defined by the heliospheric current sheet (HCS) maximum excursion. In the solar wind, there also exist various anomalous episodes, as consequences of transient and/or energetic activities that may be remote (in the corona) or heliospheric in origin. The following phenomena or transients relating to anomalous solar wind parameters have been studied and have inspired our study. (1) False magnetic field polarity reversals—evoking images of switchbacks—that have been argued as being interchange reconnection debris (Owens et al. 2013, 2017, 2020), and recently also observed much closer to the Sun by the Parker Solar Probe mission (Laker et al. 2021), and the origin is being debated. (2) Heliospheric magnetic field (HMF) departures from the Parker spiral to nearly radial field at 1 au (Murphy et al. 2002; Watari et al. 2005). The great majority of the initial mass function orientation statistics fall in the vicinity of Parker spiral fields for typical solar wind speeds. However, blocks of time occasionally stand out when the HMF is nearly radial. (3) Interplanetary coronal mass ejections (ICMEs), magnetic clouds (MCs), and other flux ropes, whose structures are foreign to the ambient solar wind and whose plasma differs from the background plasma, often bring anomalously low temperatures and plasma beta (Burlaga et al. 1981; Klein & Burlaga 1982; Burlaga 1988; Richardson & Cane 1995; Kilpua et al. 2009). (4) The intervals of counterstreaming suprathermal electrons (CSEs), which are often interpreted as closed loops of magnetic fields with both ends being connected to a hot electron source within the corona (Bame et al. 1981; Gosling et al. 1987, 1992). (5) Heat flux dropouts, suggesting complete disconnection from hot electron sources (McComas et al. 1989). Many studies have been devoted to these phenomena, including their relationship to the location of the HCS/coronal streamer belt and their solar cycle variations (McComas et al. 1989; Fitzenreiter & Ogilvie 1992; Gosling et al. 1992, 1993; Shodhan et al. 2000; Murphy et al. 2002; Watari et al. 2005; Lavraud et al. 2010; Yu et al. 2014, 2016; Owens et al. 2017, 2020). Weaker solar cycles and advanced imaging and modeling capabilities have raised the profile of pseudostreamers and their boundaries as contributors to the outflows within this slow wind layer (Wang et al. 2012; Owens et al. 2013, 2014). Similarly, the density blobs that are released from the ends of the coronal streamers are another source of slow solar wind (Sheeley & Wang 2007; Rouillard et al. 2010a, 2010b) that belong to this collection of features. These have been associated with in situ observations of small flux rope-like structures in the solar wind at 1 au (Kilpua et al. 2009; Rouillard et al. 2010a, 2010b; Viall & Vourlidas 2015; Sanchez-Diaz et al. 2017; Lavraud et al. 2020). Clearly, there are many reasons and sources of evidence for the occurrence of anomalous characteristics that differ from the typical ambient solar wind. The question is whether they occur in organized ways that provide insights into their origins.
|
[
"Owens et al. 2013"
] |
[
"Weaker solar cycles and advanced imaging and modeling capabilities have raised the profile of pseudostreamers and their boundaries as contributors to the outflows within this slow wind layer"
] |
[
"Motivation"
] |
[
[
2877,
2894
]
] |
[
[
2667,
2857
]
] |
2015ApJ...799...55G__Rosenvinge_et_al._2008_Instance_1
|
We use data from instruments on board STEREO-A, STEREO-B, ACE, SOHO, Wind, the Geostationary Operational Environmental Satellites (GOES), MESSENGER and the Solar Dynamics Observatory (SDO). Remote-sensing observations of the CME and the activity phenomena over the solar surface were provided by the Sun Earth Connection Coronal and Heliospheric Investigation (SECCHI) instrument suite on board STEREO (Howard et al. 2008), the Large Angle and Spectrometric Coronagraph experiment (LASCO) on board SOHO (Brueckner et al. 1995), the Atmospheric Imaging Assembly (AIA; Lemen et al. 2012) on board SDO, and the X-ray telescopes on board GOES. Synoptic maps including PFSS model results provided by the Global Oscillation Network Group (GONG, http://gong.nso.edu/) and by Solarsoft PFSS package were also examined. Radio observations were provided by the S/WAVES (Bougeret et al. 2008) investigation on board STEREO and by the WAVES (Bougeret et al. 1995) experiment on board Wind. In situ energetic particle observations were provided by the Solar Electron and Proton Telescope (SEPT; Müller-Mellin et al. 2008), the Low Energy Telescope (LET; Mewaldt et al. 2008), and the High Energy Telescope (HET; von Rosenvinge et al. 2008) on board STEREO (all of them part of the IMPACT instrument suite; Luhmann et al. 2008), the Comprehensive Suprathermal and Energetic Particle Analyzer (COSTEP; Müller-Mellin et al. 1995), and the Energetic and Relativistic Nuclei and Electron (ERNE; Torsti et al. 1995) instrument on board SOHO, the Electron, Proton, and Alpha Monitor (EPAM; Gold et al. 1998) on board ACE, the 3D Plasma and Energetic Particle Investigation (3DP) on board Wind, and the Energetic Particle and Plasma Spectrometer (EPPS; Andrews et al. 2007) on board MESSENGER. Additionally, Fe/O ratios at ACE were obtained using data from the Solar Isotope Spectrometer (SIS; Stone et al. 1998). Finally, solar wind plasma and magnetic field data were obtained from the Plasma and Suprathermal Ion Composition (PLASTIC; Galvin et al. 2008) investigation on board STEREO, the STEREO Magnetic Field Experiment (Acuña et al. 2008), the ACE Magnetic Field Experiment (Smith et al. 1998), and the ACE/Solar Wind Electron Proton Alpha Monitor (SWEPAM; McComas et al. 1998). Due to gaps in ACE/SWEPAM data, plasma data from the Charge, Element, and Isotope Analysis System (CELIAS; Hovestadt et al. 1995) instrument on board SOHO were also used. Magnetic field data from MESSENGER were obtained from the Magnetometer Instrument on board this s/c (Anderson et al. 2007). Interplanetary disturbance identifications were cross-checked using the STEREO and ACE level 3 event list maintained by L. Jian (http://www-ssc.igpp.ucla.edu/forms/stereo/stereo_level_3.html, http://www.srl.caltech.edu/ACE/ASC/DATA/level3/) and the near-Earth ICME list maintained by I. Richardson and H. Cane (http://www.srl.caltech.edu/ACE/ASC/DATA/level3/icmetable2.htm; see also Cane & Richardson 2003).
|
[
"von Rosenvinge et al. 2008"
] |
[
"Radio observations were provided by",
"and the High Energy Telescope"
] |
[
"Uses",
"Uses"
] |
[
[
1200,
1226
]
] |
[
[
811,
846
],
[
1164,
1193
]
] |
2021MNRAS.507.1623C__Lesgourgues_2011_Instance_1
|
To describe the H i bispectrum in redshift space, we apply the standard redshift space kernels (see Heavens et al. 1998; Scoccimarro et al. 1999 for derivations) such that
(10)$$\begin{eqnarray*}
B_\rm {H\,\small {I}}(\boldsymbol{k}_{1},\boldsymbol{k}_{2},\boldsymbol{k}_{3}) &=& 2\, \overline{T}_\rm {H\,\small {I}}^2\left[Z_1(\boldsymbol{k}_{1}) Z_1(\boldsymbol{k}_{2}) Z_2(\boldsymbol{k}_{1}, \boldsymbol{k}_{2}) P_\text{lin}(k_{1}) P_\text{lin}(k_{2})\right.\nonumber\\
&&\left. +\, \text{cycl.}\right] D_\text{FoG}(\boldsymbol{k}_{1},\boldsymbol{k}_{2},\boldsymbol{k}_{3},\sigma _\text{B}),
\end{eqnarray*}$$where cycl. represents cyclic permutations which run over all possible pairs of $\boldsymbol{k}_{1},\boldsymbol{k}_{2}$ and $\boldsymbol{k}_{3}$. Plin represents the real-space, linear matter power spectrum for which we use the CLASS Boltzmann solver (Blas, Lesgourgues & Tram 2011; Lesgourgues 2011). Z1 is given in equation (7) and Z2 denotes the second-order kernel and is given by
(11)$$\begin{eqnarray*}
Z_2(\boldsymbol{k}_{i}, \boldsymbol{k}_{j}) &=& b_1 F_2(\boldsymbol{k}_{i}, \boldsymbol{k}_{j})+f \mu _{ij}^2 G_2(\boldsymbol{k}_{i}, \boldsymbol{k}_{j})+\frac{f \mu _{ij} k_{ij}}{2}\nonumber\\
&&\times \,\left[\frac{\mu _i}{k_i} Z_1(\boldsymbol{k}_{j})+\frac{\mu _j}{k_{j}} Z_1(\boldsymbol{k}_{i})\right] + \frac{b_2}{2},
\end{eqnarray*}$$where $\boldsymbol{k}_{ij} = \boldsymbol{k}_{i} + \boldsymbol{k}_{j}$ and $\mu _{ij} = \boldsymbol{k}_{ij}\cdot \hat{{\bf z}}/k_{ij}$. F2 and G2 denote the second-order kernels for the real-space density and velocity fields and are given by
(12)$$\begin{eqnarray*}
F_2(\boldsymbol{k}_{i}, \boldsymbol{k}_{j}) = \frac{5}{7} + \frac{m_{ij}}{2}\left(\frac{k_i}{k_{j}} + \frac{k_{j}}{k_i}\right) + \frac{2}{7} m_{ij}^2,
\end{eqnarray*}$$(13)$$\begin{eqnarray*}
G_2(\boldsymbol{k}_{i}, \boldsymbol{k}_{j}) = \frac{3}{7} + \frac{m_{ij}}{2}\left(\frac{k_i}{k_{j}} + \frac{k_{j}}{k_i}\right)+\frac{4}{7} m_{ij}^2 ,
\end{eqnarray*}$$where $m_{ij} = (\boldsymbol{k}_{i}\cdot \boldsymbol{k}_{j})/(k_i k_{j})$. The final term in equation (10), DFoG, is a phenomenological factor to address some non-linear RSD effects not sufficiently modelled by the redshift kernels alone. On smaller scales, internal motion inside virialized structures produces a radial smearing to the density field in redshift space, known as the fingers-of-god (FoG) effect (Jackson 1972). It is common to include a term which describes the FoG (Taruya, Nishimichi & Saito 2010), even when including higher order perturbation theory terms and should be seen as a phenomenological damping required to correct for non-linear effects (Verde et al. 1998; Gil-Marín et al. 2014). For our choice of model, this factor is given by (Gil-Marín et al. 2015)
(14)$$\begin{eqnarray*}
D_\text{FoG}(\boldsymbol{k}_{1},\boldsymbol{k}_{2},\boldsymbol{k}_{3},\sigma _\text{B}) = \left[1 + \frac{1}{2}\left(k_{1}^2 \mu _1^2 + k_{2}^2\mu _2^2 + k_{3}^2\mu _3^2\right)^2 \sigma _\text{B}^2\right]^{-2} . \nonumber\\
\end{eqnarray*}$$It is worth noting that it is necessary for galaxy surveys to also include modelling of shot noise caused by discreteness effects in their bispectra analyses. However, for H i IM, where unresolved signal is integrated over, this should not be a limiting factor (Spinelli et al. 2020). We therefore do not consider any treatment of shot noise in our analysis, making the assumption that this should be very low in H i IM observations.
|
[
"Lesgourgues 2011"
] |
[
"Plin represents the real-space, linear matter power spectrum for which we use the CLASS Boltzmann solver"
] |
[
"Uses"
] |
[
[
915,
931
]
] |
[
[
778,
882
]
] |
2017ApJ...840...91Y__Shatsky_&_Tokovinin_2002_Instance_1
|
Another binary parameter is the stellar mass ratio distribution. The choice of the mass ratio can be called the pairing function, which is combining stars into binary systems. When random pairing is used (i.e., binary companions are randomly chosen from a given IMF), we find similar mass segregation results to those for star clusters containing only single stars. However, random pairing is ruled out observationally, and there is also a lack of theoretical backing for a random pairing function (Shatsky & Tokovinin 2002; Kouwenhoven et al. 2005, 2007a, 2007b). Indeed, massive stars preferentially choose other massive stars as their binary companions (Kouwenhoven et al. 2010). This implies that both the geometric average and the median MST method may be biased by the short MST edges from binary systems, although the LnND method is less influenced. An alternative binary pairing function is the so-called ordered pairing (i.e., mass ratio of the binary star components ∼1, in particular for massive stars; see e.g., Kouwenhoven et al. 2009; Oh & Kroupa 2012), which matches stars in order of their mass distribution. We re-pair our test model using ordered pairing and find some extremely high values of the mass segregation degree when using the MST method. This is illustrated in Figure 5, where we see that the mass segregation degree for ordered pairing functions is biased toward very high values for relatively small massive-star samples. These extremely high degrees of mass segregation are not physical, because they are mainly produced by the short MST edges in the binary systems. This is illustrated by the sharp steps in the solid curves, which greatly deviates from the non-mass-segregation distribution equal to unity. Therefore, the gmMST and mMST methods should not be considered as a global compactness measure of massive stars in cases where binary fractions are significant, i.e., >10%. Massive stars may have much shorter MST edges than low-mass stars, due to the preferential presence of binaries among massive stars (e.g., Sana et al. 2014).
|
[
"Shatsky & Tokovinin 2002"
] |
[
"When random pairing is used (i.e., binary companions are randomly chosen from a given IMF), we find similar mass segregation results to those for star clusters containing only single stars. However, random pairing is ruled out observationally, and there is also a lack of theoretical backing for a random pairing function"
] |
[
"Compare/Contrast"
] |
[
[
499,
523
]
] |
[
[
176,
497
]
] |
2020MNRAS.492..708M__Davidson_&_Fesen_1985_Instance_1
|
Due to the high luminosity and seemingly long-term flux stability, the Crab pulsar and its pulsar wind nebula (PWN) is one of the most studied sources in the very high energy (VHE, $\rm {E} \gt 100$ GeV) regime. For many years, Crab has been used as a standard candle in X- and gamma-ray astronomy (Hester 2008). The Crab nebula was the first TeV gamma-ray source discovered (in 1989 by the Whipple 10-m telescope; Weekes et al. 1989), and soon after detected by numerous facilities above 100 GeV (Smith et al. 2000; Aharonian et al. 2006; Abdo et al. 2012; Aleksić et al. 2015; Meagher & VERITAS Collaboration 2015; Abeysekara et al. 2017). It has a (energy-dependent) angular size of ∼0.1° and its distance has been estimated to be ≈2.2 kpc, corresponding to a physical size of ≈3.8 pc (Trimble 1973; Davidson & Fesen 1985; Kaplan et al. 2008). The nebula non-thermal spectrum can be described by two components, a synchrotron component extending from radio to high energy gamma-rays and a second component emerging above 1 GeV (Atoyan & Aharonian 1996). The latter is interpreted as inverse Compton scattering (IC) of the same particles against soft background photons: cosmic microwave background (CMB), far-infrared (FIR), and near-infrared (NIR) background, and the synchrotron photons of the nebula itself or Synchrotron Self-Compton (SSC). The pulsar has a spin period of $\rm {P} = 33$ ms, a spin-down rate of $\rm {\dot{P}} = 4.21\times 10^{-13}$ and a spin-down luminosity of $\rm {L_{spin}} = 3.8\times 10^{38}\ \rm {erg}\ \rm {s^{-1}}$. The pulsed emission between 0.1 and 100 GeV is believed to be due to synchrotron-curvature radiation (Abdo et al. 2010; Ansoldi et al. 2016) and its spectrum is well parametrized by a power law with a sub-exponential cut-off function of spectral index γP = 1.59, the break located at an energy of about 500 MeV and curvature index of κ = 0.43. In addition, a power-law component emerges above the cut-off extending above 100 GeV (Aliu et al. 2011; Aleksić et al. 2012; Ansoldi et al. 2016).
|
[
"Davidson & Fesen 1985"
] |
[
"It has a (energy-dependent) angular size of ∼0.1° and its distance has been estimated to be ≈2.2 kpc, corresponding to a physical size of ≈3.8 pc"
] |
[
"Background"
] |
[
[
803,
824
]
] |
[
[
642,
787
]
] |
2022ApJ...935..135B__Thomas_et_al._2019_Instance_1
|
Disk galaxies typically reveal out-of-equilibrium features due to incomplete equilibration. These may appear in the form of bars and spiral arms, which are large-scale perturbations in the radial and azimuthal directions, responsible for a slow, secular evolution of the disk. In the vertical direction, disks often reveal warps (Binney 1992). In the case of the Milky Way (MW) disk, which can be studied in much greater detail than any other system, recent data from astrometric and radial velocity surveys such as SEGUE (Yanny et al. 2009), RAVE (Steinmetz et al. 2006), GALAH (Bland-Hawthorn et al. 2019), LAMOST (Cui et al. 2012), and above all Gaia (Gaia Collaboration et al. 2016, 2018a, 2018b) have revealed a variety of additional vertical distortions. At large galactocentric radii (>10 kpc) this includes, among others, oscillations and corrugations (Xu et al. 2015; Schönrich & Dehnen 2018), and streams of stars kicked up from the disk that undergo phase mixing, sometimes referred to as “feathers” (e.g., Price-Whelan et al. 2015; Thomas et al. 2019; Laporte et al. 2022). Similar oscillations and vertical asymmetries have also been reported in the solar vicinity (e.g., Widrow et al. 2012; Williams et al. 2013; Yanny & Gardner 2013; Gaia Collaboration et al. 2018b; Quillen et al. 2018; Bennett & Bovy 2019; Carrillo et al. 2019). One of the most intriguing structures is the phase-space spiral discovered by Antoja et al. (2018) and studied in more detail in subsequent studies (e.g., Bland-Hawthorn et al. 2019; Li 2021; Li & Widrow 2021; Gandhi et al. 2022). Using data from Gaia DR2 (Gaia Collaboration et al. 2018a), Antoja et al. (2018) selected ∼900,000 stars within a narrow range of galactocentric radius and azimuthal angle centered around the Sun. When plotting the density of stars in the (z, v
z
)-plane of vertical position, z, and vertical velocity, v
z
, they noticed a faint, unexpected spiral pattern, which became more enhanced when color-coding the (z, v
z
)-“pixels” by the median radial or azimuthal velocities. The one-armed spiral makes two to three complete wraps, resembling a snail shell, and is interpreted as a signature of phase mixing in the vertical direction following a perturbation, which Antoja et al. (2018) estimate to have occurred between 300 and 900 Myr ago. More careful analyses in later studies (e.g., Bland-Hawthorn et al. 2019; Li 2021) have nailed down the interaction time to ∼500 Myr ago.
|
[
"Thomas et al. 2019"
] |
[
"At large galactocentric radii (>10 kpc) this includes, among others,",
"and streams of stars kicked up from the disk that undergo phase mixing, sometimes referred to as “feathers” (e.g.,"
] |
[
"Background",
"Background"
] |
[
[
1044,
1062
]
] |
[
[
761,
829
],
[
903,
1017
]
] |
2015AandA...574A..50J__Takeda_et_al._(2008)_Instance_1
|
Several studies that analyzed the metallic content in the atmospheres of evolved stars hosting planets have been published in the last decade. The first results were based on small samples and/or the abundances were not obtained with a homogeneous technique. Schuler et al. (2005) derived the metallicity for one GWP and gathered abundances from the literature for another seven. They reported that, on average, GWP were metal-poor compared with planet hosting dwarfs. Similar results were found by Sadakane et al. (2005) analyzing 4 GWP. In 2007, Pasquini et al. studying 14 GWP (4 from the literature), concluded that in contrast to the distribution of main-sequence stars with planets, the GWP distribution does not favor high metallicity objects. Conversely, Hekker & Meléndez (2007), analyzing a sample of 380 GK giant stars including 20 with planets (15 from the literature), found an enhancement for GWP of 0.13 dex compared with stars without planets. Takeda et al. (2008), with a sample of 322 giants, including 10 planet-hosts, did not find any metallicity offset. Ghezzi et al. (2010a) found that the metallicity distribution of 16 GWP displays an average that is 0.17 dex more metal-poor than the sample of 117 planet-hosting dwarfs and, furthermore, that the subgiant sample is more metal-rich by 0.12 dex. Johnson et al. (2010) ruled out a flat metallicity relationship among their sample of 246 subgiants (36 with planets, including unpublished candidates). More recently, with larger samples, Mortier et al. (2013) did not find any metallicity enhancement in 71 evolved stars with planets in comparison with 733 evolved stars without planets, with metallicity values gathered from the literature. Finally, Maldonado et al. (2013) found a metallicity enhancement in 16 subgiants with planets relative to 55 without planets (50 from literature). These authors did not find evidence of a metallicity offset between giants with and without planets, analyzing 43 GWP and 67 GWOP. However, for stars with masses above 1.5 M⊙, they reported a slight metallicity enhancement for giants with planets relative to the control sample.
|
[
"Takeda et al. (2008)"
] |
[
"with a sample of 322 giants, including 10 planet-hosts, did not find any metallicity offset."
] |
[
"Compare/Contrast"
] |
[
[
960,
980
]
] |
[
[
982,
1074
]
] |
2022MNRAS.509..567S__Spinoso_et_al._2017_Instance_1
|
When a galaxy is isolated, star formation quenching is mostly driven by internal processes (Larson 1974; Dekel & Silk 1986; Vulcani et al. 2021). Active galatic nuclei (AGNs) feedback create an outflow of gas preventing further hot gas accretion from the circumgalactic medium (Dalla Vecchia & Schaye 2008; Bongiorno et al. 2016; Trussler et al. 2020). Star formation depends on gravitational instabilities, which may be prevented due to a transition from a stellar disc to a spheroid (‘morphological quenching’; Martig et al. 2009). Bars in spiral galaxies may drive gas inflows, which enhance central star formation (bar-driven evolution; Spinoso et al. 2017). However, in the local Universe most of the galaxies live in groups/clusters (Geller & Huchra 1983). Even before crossing the virial radius, infalling galaxies can stop accreting new gas (‘Starvation’; Larson, Tinsley & Caldwell 1980; Balogh, Navarro & Morris 2000; van de Voort et al. 2017). For instance, Trussler et al. (2020) suggest galaxy quenching has an extended phase (∼5 Gyr) of starvation. In addition, infalling galaxies can lose gas, stars and even dark matter via gravitational tides (‘tidal mass-loss’ - TML; Johnston, Sigurdsson & Hernquist 1999; Read et al. 2006). Within the virial radius, the hot gas in the intracluster medium (ICM) exerts pressure on galaxies moving within the cluster and may remove gas via Ram Pressure Stripping (RPS; Gunn & Gott 1972; Abadi, Moore & Bower 1999). Clusters provide a suitable environment for both direct and indirect galaxy interactions, especially in its core. Direct encounters may lead to galaxy mergers and cause a starburst event over a short time-scale and quickly exhaust a galaxy’s gas supply (Springel & Hernquist 2005; Cox et al. 2008; Teyssier, Chapon & Bournaud 2010), whereas repeated indirect encounters may leave interacting galaxies with disturbed morphologies. At last, it is important to note that clusters are built up by the accretion of galaxy groups. Cluster galaxies may therefore be affected by ‘pre-processing’, in which galaxy properties were altered even before entering the cluster (Fujita 2004; Mahajan 2013; Sarron et al. 2019).
|
[
"Spinoso et al. 2017"
] |
[
"Bars in spiral galaxies may drive gas inflows, which enhance central star formation (bar-driven evolution;"
] |
[
"Background"
] |
[
[
641,
660
]
] |
[
[
534,
640
]
] |
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