adsabs/FOCAL huggingface dataset:
The third scenario to consider is emission from the quasar host galaxy. The approach of using the residual flux in DLA troughs to observe the Lyα emission associated with star formation in the host galaxies has been studied extensively (e.g., Kulkarni et al. 2006; Fynbo et al. 2010; Finley et al. 2013; Noterdaeme et al. 2014). Using the Kennicutt et al. (2008) calibration 





 and assuming an intensity ratio of 





 for the case B recombination, we estimate a star formation ratio of 





 based on the luminosity of the residual Lyα emission. This is much higher than estimates of the SFR in quasar host galaxies (





 e.g., Ho 2005; Cai et al. 2014) and in the high-redshift Lyman break galaxies and Lyα emitters (





 e.g., Shapley et al. 2003; Erb et al. 2006; Gronwall et al. 2007). To further check whether the residual Lyα emission originated from the star formation, we compare the colors (the flux density ratios: FUV/NUV and NUV/Lyα) of the residues with those of the galaxies with active star formation. The comparison galaxies are selected from the starburst original 1999 data set (Leitherer et al. 1999). The spectral energy distributions (SEDs) with “Topic—stellar continua” and “Quantity—stellar emission only” (Figures 7–12 on the Starburst99 website http://www.stsci.edu/science/starburst99/docs/table-index.html) are used to compute the galaxy’s colors. The starlight models present in a homogeneous way for five metallicities between Z = 0.040 and 0.001 and three choices of the initial mass function. The age coverage is from 1 Myr to 1 Gyr, and both star formation laws (instantaneous and continuous) are contained. The SEDs placed at the quasar’s redshift are extracted from the flux densities at the effective wavelengths of the FUV/NUV filters and 1215.6 Å. Figure 3(a) shows that SDSS J1259+6212 seriously deviates from the star-forming galaxy group. At wavelengths of a few hundred Angstroms, SDSS J1259+6212 has much bluer slopes than the bluest galaxy continuum. In Fathivavsari et al. (2018), strong and narrow Lyα emission is reported in 155 eclipsing damped Lyα systems, which is revealed with certainty as the narrow emission line from the host galaxy. However, the residual Lyα emission of SDSS J1259+6212 fills in the sub-DLA trough rather, forming a narrow peaked profile, and its luminosity is also higher than that from Fathivavsari et al., eclipsing the DLA sample. Thus, the host galaxy scenario is questionable.
Citation: Shapley et al. 2003
Citation function: Differences; Differences
In both models, as we will see, the thick disc scale length is about a factor of two shorter than that of the thin disc, in agreement with the results by Bovy et al. (2012a). The choice of presenting two mass models for the mass distribution of our Galaxy is mainly dictated by two reasons. First, the need to add a central bulge to the global gravitational potential to reproduce the rotation curve in the inner kpcs of the Milky Way strongly depends on the observational data with which one compares the theoretical curve: to reproduce the rise observed in the inner kpcs (see the observational data adopted by Caldwell & Ostriker 1981), Allen & Santillan (1991) introduced a central mass concentration, whose mass is about 15% of the disc mass. However, the central rise observed in the rotation of the molecular gas in the inner Galaxy (for more recent estimates see, for example,  Sofue 2012) may be an effect of non circular motions generated by large scale asymmetries like the bar, as has been shown recently by Chemin et al. (2015). Moreover, this feature is not reported in all the observational studies (see, for example, Reid et al. 2014). Secondly, there is growing evidence that the mass of any classical bulge, if present in the Milky Way, must be small (Shen et al. 2010; Kunder et al. 2012, 2016; Di Matteo et al. 2014, 2015). For these reasons, we prefer to present a second model, our Model II, which does not include any spherical central component, and which is still compatible with the rotation curve of the Galaxy, as given by Reid et al. (2014). Because it has been widely used in the last decades, and due to the facility of its implementation, we explicitly aim at generating Galactic models similar to the Allen & Santillan (1991) model, so to make any implementation of these new models, and any comparison with Allen & Santillan (1991), straightforward. As for the model proposed by Allen & Santillan (1991), Models I and II are axisymmetric and time-independent, and do not include stellar asymmetries such as a bar or spiral arms. No truncation is assumed for the discs, while the halo is truncated at 100 kpc, in agreement with the choice of Allen & Santillan (1991). As we describe in the following section, the analytic forms for the discs, halo, and bulge potentials are the same as those adopted by Allen & Santillan (1991). To allow an easy comparison with the Allen & Santillan (1991) model, in the following we will make use of the same system of units adopted by these authors: the potential is given in units of 100 km2/ s2, lengths are in kpc, masses in units of 2.32 × 107M⊙, time in units of 0.1 Gyr, velocities in units of 10 km s-1 and the vertical force in units of 10-9 cm s2. In these units, the gravitational constant G is equal to 1 and the mass volume density is in units of 2.32 × 107M⊙/ kpc3.
Citation: Chemin et al. (2015)
Citation function: Motivation
21 cm Intensity Mapping (IM) is a promising technique to map the cosmological large scale distribution of matter through the observation of 21 cm radio emission/absorption of neutral hydrogen gas (H i), while not requiring the detection of individual sources (Bharadwaj, Nath & Sethi 2001; Battye, Davies & Weller 2004) and has been largely explored in the context of the search for the EoR (Epoch of Reionization) signal (Pritchard & Loeb 2008; Morales & Wyithe 2010). Subsequently, it was suggested that post-EoR 21 cm Intensity Mapping surveys could be used to constrain dark energy through the measurement of the Baryon Acoustic Oscillations (BAO) scale (Chang et al. 2008; Seo et al. 2010; Ansari et al. 2012) in the large scale structure (LSS) distribution, over a broad redshift range (z ≲ 6). Such surveys require instruments with large instantaneous bandwidth and field of view and several groups have built dense interferometric arrays to explore IM, such as CHIME (Bandura et al. 2014) or Tianlai (Chen 2012) in the last decade. Smaller instruments such as PAON4 (Ansari et al. 2020) and BMX (O’Connor et al. 2021) have also been built to explore specific technical aspects of these arrays, as well as transit mode operation and calibration. CHIME has proved to be a powerful fast radio burst and pulsar observation machine (The CHIME/FRB Collaboration 2021) and has motivated the design and construction of larger, dish-based, dense interferometric arrays such as HIRAX (Newburgh et al. 2016) and CHORD (Vanderlinde et al. 2019). Non-interferometric surveys have also been considered, such as BINGO (Battye et al. 2016; Wuensche et al. 2022), FAST (Hu et al. 2020) or using the SKA precursor MeerKAT (Wang et al. 2021). Intensity Mapping is also being used to search for signals from the EoR and the cosmic dawn, at redshifts above z ≳ 10, by several large radio-interferometers, such as LOFAR (van Haarlem et al. 2013), MWA (Tingay et al. 2013), HERA (DeBoer et al. 2017), and LWA (Eastwood et al. 2018). It is also planned to be used with SKA-low (Mondal et al. 2020).
Citation: Seo et al. 2010
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Citation: Bharadwaj, Nath & Sethi 2001
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Citation: Bandura et al. 2014
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Citation: O’Connor et al. 2021
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Citation: DeBoer et al. 2017
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The shape of the bow shock nose can be directly inferred from the Hα images and used to constrain the bow shock standoff radius. In the thin-shell limit, the radial shape of the bow shock can be expressed as (Wilkin 1996)
7




R

(
θ
)

=


R


0



csc

θ


3
(
1
−
θ

cot

θ
)


,


where R0 is the standoff radius and θ represents the angle between the pulsar’s velocity and a point R(θ) along the bow shock. The standoff radius is dictated by the pressure balance between the ambient ISM and the neutron star wind and directly related to the interstellar density ρA, the pulsar wind velocity vw and mass-loss rate 








m







w




, and the pulsar velocity vp (Wilkin 1996):
8






R


0


=











m







w




v


w




4
π


ρ


A




v


p


2









1

/

2


.


Equation (8) assumes an isotropic stellar wind, which is not generally true of a relativistic neutron star wind. When the pulsar rotation vector is skewed relative to the pulsar proper motion and magnetic field, the pulsar wind will be quasi-isotropic, and its exact behavior will depend on both the orientations of these vectors and the opening angle of the wind. For B2224+65, the light-cylinder radius is significantly smaller than the bow shock standoff radius, suggesting that an isotropic wind is an adequate assumption without additional evidence for anisotropy (for a detailed discussion, see Chatterjee & Cordes 2002). Assuming that the spin-down energy loss is entirely carried by the relativistic wind, 








m







w




v


w


=


E






/

c


. The standoff radius can be conveniently reformulated as an angle given by
9






θ


0


=
56.3

mas







sin


2


i




n


A


1

/

2
















E







33


1

/

2






μ


100




D


kpc


2







,


where i is the inclination angle, nA = ρA/mH is the total number density of the interstellar hydrogen and helium mixture in atomic mass units per cubic centimeter, 






E







 is the spin-down luminosity in erg per second, D is the distance to the pulsar in kiloparsecs, and μ100 is the pulsar proper motion in 100 mas yr−1 (Chatterjee & Cordes 2002). The spin-down luminosity is 






E





=
4


π


2


I


P






/



P


3




, where I ≈ 1045 g cm2. The period derivative 






P







 can be corrected for the Schklovskii effect using
10






P





=




P







obs


−
2.43
×


10


−
21


P


μ


masy


2




D


kpc


,


where μmasy is the proper motion in milliarcseconds per year (Shklovskii 1970; Brownsberger & Romani 2014). This correction is negligible for B2224+65. The atomic number density nA is converted to an electron density ne assuming a cosmic abundance γH = 1.37, where nA = nHγH and nH ≈ ne. The standoff angle that is inferred from the outer edge of the Hα emission corresponds to a forward shock that lies slightly upstream of the contact discontinuity at an angular distance θa ≈ 1.3θ0 (Aldcroft et al. 1992; Bucciantini 2002). The period, period derivative, and spin-down luminosity for B2224+65 are shown in Table 1.
Citation: Chatterjee & Cordes 2002
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Citation: Chatterjee & Cordes 2002
Citation function: Uses
To constrain the number density of the clump, getting the size of the clump (rc) should be required. The size of the self-gravitating clump in the context of AGNs has been estimated in the literature (e.g., Krolik & Begelman 1988; Hönig & Beckert 2007; Kawaguchi & Mori 2011). According to the authors of these works, we have rc ≤ 0.05 pc at 1 pc from the central SMBH with MBH = 8 × 108M⊙ (Scharwächter et al. 2013). On the other hand, the size of clumps can be constrained by the observations of transient X-ray absorption events in nearby AGNs, i.e., the typical size is 0.002 pc (e.g., Markowitz et al. 2014; Tanimoto et al. 2019). By adopting Nc and rc, taking into account their uncertainties of 3 ≤ Nc ≤ 15 and 





, the lower limit of number density of each ionized cloud 





 can be given by
8





In the case of ionized gas clumps, the typical Thomson scattering opacity of each cloud can be estimated as 





 since the optical depth is given by 





, where σT is the Thomson cross section. If we adopt the opacity ratio of the dusty gas and the ionized gas is ≃103 at the UV band for the AGN radiation (e.g., Umemura et al. 1998; Ohsuga & Umemura 2001; Wada 2012), the optical depth of each cloud is 





, which is consistent with the lower value for nearby Seyfert galaxies (e.g., Table 10 in Ramos Almeida et al. 2011). This might indicate that the ionized gas is also clumpy within the 1 pc region of 3C 84. Wada et al. (2018) examine properties of the ionized gas irradiated by less luminous AGN such as Seyfert galaxies based on their “radiation-driven fountain” model (Wada 2012). They found that the ionized region show nonuniform internal structures, corresponding to the clumpy fountain flows caused by the radiation pressure on dusty gas, although the typical density (





) is smaller than our estimate. In addition, by the optical/NIR observations, the existence of dense clumps with 





 cm−3 has been reported from the detection of coronal lines within NLR (e.g., Murayama & Taniguchi 1998). These high density clouds in NLRs might contribute the absorption feature of N1.
Citation: Tanimoto et al. 2019
Citation function: Uses; Uses
The origin of the Milky Way Galaxy presumably commenced within the initial one billion years (Gyr) subsequent to the Big Bang origin of the Universe around 13.7 Gyr ago. The formation of the Galaxy initiated with the merging of the initial diffuse neutral hydrogen gas clouds. This was followed by the gravitational merging of the colliding initial protogalaxies, thereby leading to the gradual accretional evolution of the galaxies in a hierarchical manner. The formation and evolution of the successive generations of stars inside the Galaxy led to the gradual abundance evolution of the elemental (and isotopic) inventories of the Galaxy. Galactic chemical evolution (GCE) models deal with understanding the origin of the distribution and gradual evolution of the elemental (and isotopic) abundances across the Galaxy over the Galactic time-scale of ∼13.7 Gyr (Matteucci and Franҫois 1989; Rana 1991; Chiappini, Matteucci & Gratton 1997; Pagel 1997; Chang et al. 1999; Goswami & Prantzos 2000; Alibés, Labay & Canal 2001; Chiappini, Matteucci & Romano 2001; Matteucci 2003, 2014; Franҫois et al. 2004; Kobayashi et al. 2006; Kobayashi and Nomoto 2009; Matteucci et al. 2009; Kobayashi and Nakasato 2011; Kobayashi, Karakas & Umeda 2011; Micali, Matteucci & Romano 2013; Minchev, Chiappini & Martig, 2013, 2014; Mott, Spitoni & Matteucci 2013; Sahijpal 2013, 2014; Sahijpal & Gupta 2013; Sahijpal 2014; Minchev et al. 2015; Spitoni et al. 2015). The astronomical observed trends in the elemental abundance distribution across the Galaxy serve as a major constraint on the GCE models in order to understand the accretional growth rate and the dynamical evolution of the Galaxy with its three main components, i.e. the Galactic halo, the thick disc, and the thin disc (Marochnik & Suchkov 1996; Matteucci 2003; Sparke & Gallagher 2007; Chiappini et al. 2015). The GCE model provides a detailed evolutionary account of the star formation rate (SFR), the average prevalent initial mass function (IMF) for star formation, the supernova (SN Ia, Ib/c, II) rates, and the stellar nucleosynthetic contributions of stars with distinct mass and metallicity.
Citation: Chiappini, Matteucci & Romano 2001
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Citation: Matteucci 2003
Citation function: Motivation
Most of these OH absorbers show an OH column density about the order of $10^{15} \rm\, cm^{-2}$. Gupta et al. (2018) show that galaxies with z  0.4 could have much lower OH column densities, i.e. below $10^{14} \rm\, cm^{-2}$, using a detected OH absorber and upper limits of eight non-detections and assuming the OH excitation temperature $T_x^{\rm OH} = 3.5\,$ K, however, a more commonly used $T_x^{\rm OH}$ for OH absorbers at cosmological distances is 10 K (e.g. Kanekar & Chengalur 2002; Kanekar et al. 2003; Curran et al. 2011). The Gupta et al. (2018) sample contains only intervening absorbers with relatively low redshifts. Grasha et al. (2020) searched 11 intervening and five associated H i absorbers for OH absorptions and report no new detections except for a re-detection of the 1667 MHz OH absorption towards PKS 1830−211. Grasha et al. (2019) searched a large sample of compact radio sources for associated H i and OH absorptions but only six of them, which have relatively larger redshifts (z ∼ 0.6), have H i absorption detections and none of them have OH absorption detection. Intervening absorbers are mostly from foreground galaxies or H i dark clouds. Associated absorbers can be found in AGN outflows, circumnuclear discs, and cold gas clouds in the host galaxy or merger relics (Maccagni et al. 2017; Allison et al. 2019; Oosterloo et al. 2019). The cold molecular gas in these systems is not well studied due to a low detection number (Allison et al. 2019) and they could potentially have very different OH column densities and different excitation temperatures (Curran et al. 2016). Therefore, we observe a sample of low-redshift (z  0.3) galaxies with associated H i 21-cm absorption detections, with an extra target with a deep intervening H i 21-cm absorption for comparison, using the Five-hundred-meter Aperture Spherical radio Telescope (FAST; Nan et al. 2011; Li & Pan 2016; Li et al. 2018a; Jiang et al. 2020) aiming to study their cold gas parameters through OH absorptions.
Citation: Kanekar et al. 2003
Citation function: Compare/Contrast
In general, the time evolution of the PSR is governed by,
(3)$$\begin{eqnarray*}
{T^{\mu \nu }}_{;\nu } = 0,
\end{eqnarray*}$$where Tμν is the energy–momentum tensor. In the MPD formulation, a multipole expansion of the energy–momentum tensor is undertaken. This constructs the ‘gravitational skeleton’. In the extreme mass ratio limit, the pole and dipole terms are dominant and terms greater than the quadrupole can be neglected. The PSR dynamics are then entirely determined by the background gravitational field and the dynamical spin interaction with this field. The mass monopole is described by the 4-momentum pμ whilst the spin dipole is sμν. The corresponding equations of motion are (Mathisson 1937; Papapetrou 1951; Dixon 1974),
(4)$$\begin{eqnarray*}
\frac{D p^{\mu }}{\mathrm{ d} \tau } = - \frac{1}{2} {R^{\mu }}_{\nu \alpha \beta } u^{\nu } s^{\alpha \beta },
\end{eqnarray*}$$(5)$$\begin{eqnarray*}
\frac{D s^{\mu \nu }}{\mathrm{ d} \tau } = p^{\mu } u^{\nu } -p^{\nu } u^{\mu },
\end{eqnarray*}$$for proper time parametrization τ along the PSR worldline zα(τ), D/dτ denotes a covariant derivative whilst uν is the PSR 4-velocity and Rμναβ the Riemann curvature tensor. In order to first construct the gravitational skeleton requires first specifying a reference worldline with which to define the expansion. Typically, such a choice would be the worldline described by the centroid (centre of mass) of the body. However, in GR the centroid choice is observer-dependent. This uncertainty is evidenced in the system of equations (4) and (5), which is not a determinate set since there exist more unknowns than equations. This is related to the uncertainty in choosing a reference world line for the multipole expansion. Explicitly choosing an observer with respect to which the centre of mass is defined renders the system of equations determinate. Such a choice is known as the spin supplementary condition (SSC). For this work we adopt the Tulczyjew-Dixon (TD) condition which specifies the centroid to be that measured in zero 3-momentum frame:
(6)$$\begin{eqnarray*}
s^{\mu \nu } p_{\nu } = 0,
\end{eqnarray*}$$(Tulczyjew 1959; Dixon 1964). A key advantage of this SSC is that it specifies a unique worldline. In contrast, other choices of SSC are infinitely degenerate (for discussion, including application to EMRIs, see Babak, Gair & Cole 2014; Filipe Costa & Natário 2014). With this choice of SSC both the mass of the MSP,
(7)$$\begin{eqnarray*}
m = \sqrt{- p^{\mu } p_{\mu }}
\end{eqnarray*}$$and the scalar contraction of the spin vector
(8)$$\begin{eqnarray*}
s = s^{\mu } s_{\mu }
\end{eqnarray*}$$are constants of the motion. In turn, the spin vector is a contraction of the spin tensor,
(9)$$\begin{eqnarray*}
s_{\mu } = -\frac{1}{2m} \epsilon _{\alpha \beta \mu \nu } p^{\nu } s^{\alpha \beta }
.
\end{eqnarray*}$$
Citation: Dixon 1974
Citation function: Uses; Uses
To understand the physics of sunspots, one has to study their temporal evolution. Formation and decaying phases play an important role in sunspot evolution. We refer to Martínez Pillet (2002) for a review of decaying sunspot evolution. For a long time, it was only possible to study the decay of the morphological changes of the area of sunspots because of the lack of inversion codes to interpret full-Stokes measurements, although the magnetic flux is the more important parameter. Bumba (1963) found a linear decrease of the area of the sunspot with time. Martinez Pillet et al. (1993) confirmed this linear decay, but obtained a different coeffcient. In contrast, Petrovay & van Driel-Gesztelyi (1997) found a parabolic decay with a rate proportional to 





A
(
t
)




$ \sqrt{A(t)} $


, where A is the area of the spot. Linear decay rates of the areas of 32 sunspots are found by Chapman et al. (2003). Baumann & Solanki (2005) investigate the Greenwich sunspot group record, but they were not able to distinguish between a linear and a quadratic decay law. Hathaway & Choudhary (2008) publish an almost constant decay rate of 3.65μHemispheres day−1. Gafeira et al. (2014) study four sunspots and find an approximately linear decay of the areas similar for umbra and penumbra. In both studies, the decay rates are different for individual spots. Case studies resulting in a linear decrease of the magnetic flux during the decay phase are presented by Deng et al. (2007) and Verma et al. (2016). In the latter publication, the development of the area is non-monotonic. Sheeley et al. (2017) investigated the development of the magnetic flux in 36 sunspots, but not all of them were in the decaying phase. Some spots showed a nearly linear decay, but they found also indications of a bursty decay. In a 100 h numerical simulation of a mature sunspot, Rempel (2015) find a linear decay of the magnetic flux in the umbra for the last 80 h. The penumbral magnetic flux remained almost constant during this period.
Citation: Petrovay & van Driel-Gesztelyi (1997)
Citation function: Differences; Differences