adsabs/FOCAL huggingface dataset:
Bright arc-like structures extending from the Sun, when seen in extreme-ultraviolet (EUV) wavelengths, are called coronal loops. Being a characteristic feature of the corona, it is important to assimilate high-resolution information on these structures in order to understand how plasma in the outer atmosphere is heated to multi-million Kelvin temperatures. The plasma temperatures and pressures are known to vary along the lengths of the loops, whose gradients provide some insight on the underlying heating mechanisms (Rosner et al. 1978), but the main attribute of critical importance is the cross-field morphology (Patsourakos & Klimchuk 2007). It has been suggested that coronal loops consist of multiple unresolved thin strands (Reale 2014; Klimchuk 2015). Observational evidence demonstrating low filling factors (Di Matteo et al. 1999), different degrees of “fuzziness” observed in loops at different temperatures (Tripathi et al. 2009; Reale et al. 2011), and plasma dynamics at small spatial scales (Antolin & Rouppe van der Voort 2012) indicate that this scenario is likely to be true. The typical cross-sections of the strands were estimated to be a few tens to a few hundreds of kilometers (DeForest 2007; Antolin & Rouppe van der Voort 2012; Brooks et al. 2013; Klimchuk 2015). Possible braiding of these strands can instigate magnetic reconnection, thus releasing significant energy to directly heat the plasma inside the loop (Cirtain et al. 2013; Klimchuk 2015). Each strand is believed to be impulsively and independently heated, leading to a multithermal configuration across the loop (Cargill 1994; Reale et al. 2005; Klimchuk et al. 2008; Tripathi et al. 2011). However, not all studies are consistent with this scenario since some loops demonstrate isothermal structuring (Del Zanna & Mason 2003; Noglik et al. 2008; Schmelz et al. 2009). One of the main issues plaguing observations of coronal loops is the intrinsic optically thin emission (Del Zanna & Mason 2003; Terzo & Reale 2010). Since the observed intensities are integrated along the line of sight, it becomes difficult to exclude emission from overlapping structures, which can be formed at different temperatures and inadvertently imply a multithermal structure. In this letter, we use a unique set of observations to reveal the multi-stranded and multithermal nature of a coronal loop, and isolate two of its components through the application of MHD seismology. More importantly, our results are free from the adverse effects arising from line of sight integrations.
Citation: Di Matteo et al. 1999
Citation function: Similarities; Similarities
Citation: Tripathi et al. 2011
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Citation: Terzo & Reale 2010
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Citation: Reale 2014
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In the context of a standard flare and coronal mass ejection (CME) model (e.g., Shibata et al. 1995; Lin et al. 2004), there are two parallel flare ribbons where nonthermal electrons collide and heat the chromosphere, which are observed in the Ca ii H, Hα, UV, and EUV wavelengths. Apart from the two ribbons, a particular type of flare ribbons, i.e., circular ribbons, exist (Masson et al. 2009; Reid et al. 2012; Wang & Liu 2012; Jiang et al. 2013; Liu et al. 2013; Sun et al. 2013; Yang et al. 2015; Zhang et al. 2015). They are always associated with the spine-fan configuration in the presence of the magnetic null point, which is a singular point where the magnetic field vanishes (





 = 0; Lau & Finn 1990). The magnetic field 





 near the null point can be expressed as the linear term 





, where 





 is a Jacobian matrix with elements 





 and 





 is the position vector 





 centered at the null point (Parnell et al. 1996). The divergence-free condition (





) requires that the sum of the three eigenvalues equals zero (Zhang et al. 2012). The two eigenvectors corresponding to the two eigenvalues of the same sign determine the fan surface, which divides the space into two regions having a distinct connectivity. The third eigenvector corresponding to the third eigenvalue of the opposite sign determines the direction of spines passing through the null point. Magnetic reconnection and particle acceleration in null point reconnection regions have been explored in analytical studies (Priest & Titov 1996; Litvinenko 2004) and three-dimensional (3D) numerical simulations (Rosdahl & Galsgaard 2010; Baumann et al. 2013a, 2013b). The circular ribbons are believed to be intersections of the fan surfaces and the chromosphere. The central or inner ribbons within the circular ribbons are thought to be intersections of the inner spines and the chromosphere (Reid et al. 2012; Wang & Liu 2012). Sometimes, there are multiple flare ribbons owing to the extraordinarily complex magnetic topology of the active regions (ARs; Joshi et al. 2015; Liu et al. 2015).
Citation: Parnell et al. 1996
Citation function: Uses
To study the effect that assuming light traces mass has, we use a model of a cluster based on real data that contains 174 small scale galaxy-scale haloes (spectroscopically identified by Grillo et al. 2015) embedded in two large, cluster-scale DM haloes (Jauzac et al. 2014; Grillo et al. 2015). In this study, we act only to study the effect of changes to the galaxy-scale haloes on the positions of multiple images. We use the position of multiple images in the Hubble Frontier Field (HFF) galaxy cluster, MACS J0416 and the best-fitting cluster model that fits these multiple images as derived by the parametric strong lensing algorithm lenstool (Jullo et al. 2007; Jauzac et al. 2014). This equates to 140 multiple images and 174 cluster members. Each potential was originally fitted with a pseudo-isothermal elliptical mass distribution (PIEMD; Kneib et al. 1996; Natarajan & Kneib 1997; Elíasdóttir et al. 2007; Jullo et al. 2007), which follows the analytical density profile,

(1)

\begin{equation}
\frac{\rho (r)}{\rho _0} = \frac{1}{\sqrt{r^2+r_{\rm core}^2}} - \frac{1}{\sqrt{r^2+r_{\rm cut}^2}},
\end{equation}

where r is the radial distance from the centre of the halo, and the profile is parametrized by the core radius, rcore, and the cut radius, rcut. The cluster itself has a measured mass within a 200 kpc aperture of (1.6 ± 0.01) × 1014 M⊙ (Jauzac et al. 2014), and the mass of each galaxy has a lognormal distribution centred around log(M/M⊙) = 10.1 ± 0.6. The distribution of ellipticities peaks at 0.1 and decreases towards larger ellipticities. The galaxies span the entire range between 0 and 0.9. We split each galaxy halo potential into two components: a baryonic core and a DM halo. We separate the halo into two PIEMDs with the baryonic $r^{\rm B}_{\rm core} = r_{\rm core}$. For the baryonic cut radius, we conservatively cut it to a quarter of the original cut radius $r^{\rm B}_{\rm cut} = \frac {r_{\rm cut}}{4}$ (Velander et al. 2014), although this could be smaller. We also force the dark matter core radius to the baryonic cut radius, $r^{\rm DM}_{\rm core} = r^{\rm B}_{\rm cut}$ and the dark matter cut radius as the original cut radius $r^{\rm DM}_{\rm cut} = r_{\rm cut}$. Separating the halo like this means that the mass of each halo is conserved with respect to the pre-split halo. Given this two-component galaxy model, we project the images back to the source plane to create a list of background sources for the cluster. To do this, we use the same redshift information for the sources as that in Jauzac et al. (2014).
Citation: Elíasdóttir et al. 2007
Citation function: Uses; Uses
Observations of the lowest-lying rotational levels of water—those able to probe cold gas—have required space-based observatories. The Infrared Space Observatory-Short Wavelength Spectrometer (ISO-SWS; de Graauw et al. 1996; Kessler et al. 1996) covered the 





 ro-vibrational band (symmetric bending mode) of H2O centered near 6 μm, and absorption out of the lowest-lying levels of the ortho and para nuclear spin modifications (1





 and 0





, respectively) was detected toward several massive protostars (Boonman  van Dishoeck 2003). Due to the low spectral resolution of the observing configuration though (





 using SWS in AOT6 grating mode), these lines were significantly blended with absorption from other nearby H2O lines, making the determination of level-specific column densities impossible. Instead, the entire 





 band was fit simultaneously assuming a single temperature to determine the total water column density, 





. The Submillimeter Wave Astronomy Satellite (SWAS; Melnick et al. 2000b) provided much higher spectral resolution (





 km s−1) and covered the 1





–1





 pure rotational transition of H2O at 557 GHz. This line was observed in both emission and absorption in multiple sources (e.g., Melnick et al. 2000a; Snell et al. 2000), demonstrating the ability to probe cold water. More recently, the study of low-lying rotational levels at high spectral resolution (∼0.5 km s−1) has been facilitated by the Heterodyne Instrument for the Far-Infrared (HIFI; de Graauw et al. 2010) on board the Herschel Space Observatory (Pilbratt et al. 2010). Water has been detected in both emission and absorption out of levels with 





 in several protostars (e.g., van Dishoeck et al. 2011; van der Tak et al. 2013), and in absorption out of the 0





 and 1





 levels in the molecular ISM (e.g., Sonnentrucker et al. 2010; Flagey et al. 2013). Observations that resolve the velocity structure of absorption lines are vital to both determining level-specific column densities and understanding the dynamics of the absorbing/emitting regions. This is especially important for protostars as such objects contain multiple dynamical components (e.g., disk, envelope, jets, outflows, shocks).
Citation: Melnick et al. 2000b
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Within this framework, Lovric et al. (2017) introduced a spectral color index in the solar UV that is linked to the ratio between the flux integrate over the Far-UV (FUV) and Middle-UV (MUV) spectral broadbands. Such a descriptor can be used to characterize UV stellar emission, which modulates the photochemistry of molecular species, e.g., oxygen, in the atmospheres of planets (e.g., Tian et al. 2014). By using solar irradiance measurements obtained with radiometers on board the Solar Radiation and Climate Experiment (SORCE) satellite (McClintock et al. 2005) for almost a solar cycle, the authors showed that the color index so defined is linearly correlated with the Bremen Magnesium II index, which is an excellent proxy of the magnetic activity (Viereck et al. 2004). Lovirc et a. showed that the correlation coefficient is slightly different for the descending phase of Cycle 23 and the ascending phase of the subsequent cycle. Such a difference was ascribed to residual instrumental effects, which, when compensated for, lead to a correlation coefficient constant with time. On the other hand, several solar photospheric and chromospheric indices present a clear asymmetry during different phases of a cycle (namely, an hysteresis pattern, e.g., Bachmann & White 1994; Criscuoli 2016; Salabert et al. 2017), whereas correlation coefficients between indices may vary from cycle to cycle (e.g., Bruevich & Yakunina 2014; Tapping & Morgan 2017), so that the question arises whether, and to what extent, the data corrections applied by Lovric et al. (2017) might include a physical variation of solar emission. The purposes of this paper are to answer this question and to investigate the physical mechanisms determining the high linear relation between the UV spectral color index and solar activity measured with the Mg ii index. To this aim, we compare the results of Lovric et al. (2017) with synthetic indices obtained with an irradiance reconstruction technique based on the use of semi-empirical atmosphere models and full-disk observations.
Citation: Viereck et al. 2004
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Citation: Lovric et al. (2017); Lovric et al. (2017)
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Citation: Lovric et al. (2017)
Citation function: Motivation; Motivation
The MF is an important statistical tool to understand the formation of stars (Sharma et al. 2017 and references therein). The MF is often expressed by a power law, 





 and the slope of the MF is given as 





, where N (log m) is the number of stars per unit logarithmic mass interval. We have used our deep optical data to generate the MF of the Sh 2-305 region. For this, we have utilized the optical V0 versus 





 CMDs of the sources in the target region and that of the nearby field region of equal area and decontaminated the previous sources of foreground/background stars and corrected for data incompleteness using a statistical subtraction method already described in detail in our previous papers (see Sharma et al. 2007, 2012, 2017; Pandey et al. 2008, 2013; Chauhan et al. 2011; Jose et al. 2013). As an example, in Figure 10 (left panel), we have shown the V0 versus 





 CMDs for the stars lying within the central cluster “Mayer 3” in panel (a) and for those in the reference field region selected as an annular region outside the boundary of the Sh 2-305 region (see Section 3.1) in panel (b). The magnitudes were corrected for the AV values derived from the reddening map (see Section 3.7). In panel (c), we have plotted the statistically cleaned V0 versus 





 CMD for the central cluster “Mayer 3,” which is showing the presence of PMS stars in the region. The ages and masses of the stars in this statistically cleaned CMD have been derived by applying the procedure described earlier in our previous papers (Chauhan et al. 2009; Sharma et al. 2017). For reference, the post-MS isochrone for 2 Myr calculated by Marigo et al. (2008; thick blue curve) along with the PMS isochrones of 0.1 and 3.5 Myr (purple curves) and evolutionary tracks of different masses (red curves) by Siess et al. (2000) are also shown in panel (c). These isochrones are corrected for the distance of Sh 2-305 (3.7 kpc; see Section 3.3). The corresponding MF has subsequently been plotted in Figure 10 (right panel) for the central cluster “Mayer 3.” For this, we have used only those sources that have ages equivalent to the average age of the optically identified YSOs combined with error (i.e., 





; see Table 6). Our photometry is more than 90% complete up to V = 21.5 mag, which corresponds to the detection limit of a 0.8 M⊙ (see Figure 10(c)) PMS star of ≃1.8 Myr age embedded in the nebulosity of AV ≃ 3.0 mag (i.e., the average values for the optically detected YSOs; see Table 6). We have applied a similar approach to derive the MF of the southern clustering and the whole region of Sh 2-305 (see Section 3.1), and the corresponding values of their MF slopes in the mass range 1.5  M⊙  6.6 are given in Table 7. The MF of the northern clustering cannot be determined due to an insignificant number of optically detected stars.
Citation: Jose et al. 2013
Citation function: Uses
The characteristic frequency of curvature radiation is
8






ν


curv


=



3
c


γ


3




4
π
ρ



=
2.4
×


10


3




γ


3




ρ


30


−
1



Hz
,


where 




ρ
=
30


ρ


30




 km is the curvature radius and 




γ


 is the Lorentz factor of an emitting electron. For a typical FRB, we have
9




γ
=
75


ρ


30


1

/

3




ν


curv,9


1

/

3


,


where 






ν


curv,9


=


ν


curv



/



10


9



Hz


. We then discuss whether or not a photon can propagate through the plasma in the magnetosphere by considering three effects. First, the plasma frequency in the emission region is
10






ν


p


=



1


2
π










4
π


n


e




e


2






m


e









1

/

2


.


In a highly magnetized magnetosphere, the plasma effect is negligible if 






ν


p





γ


1

/

2




ν


curv




 (Lyubarskii & Petrova 1998), which is further written as
11






a


30


>
1.0


B


*
,
12


2

/

9




γ


2


−
2

/

9




ν


curv,9


−
4

/

9


,


where 






γ


2


=
γ

/



10


2




 and we have taken the electron density of the emission region to be roughly that of the acceleration region. Second, the cyclotron absorption in the magnetosphere is generally considered in pulsar physics, whose optical depth is given by Lyubarskii & Petrova (1998):
12






τ


cyc


≃
2
×


10


−
3




B


*
,
12


3

/

5




P


*


−
9

/

5




ν


curv,9


−
3

/

5




γ


2


−
3

/

5


,


where 






P


*


=
2
π

/



Ω


*




 and the electron number density outside of the flux tube (see Figure 1) has been assumed to be described by the Goldreich–Julian density. Thus, the cyclotron absorption can be ignored if 






P


*


>
0.03


B


*
,
12


1

/

3




ν


curv,9


−
1

/

3




γ


2


−
1

/

3




 s. Third, the characteristic distance (l) of a photon due to the Thomson scattering is (Lyubarskii & Petrova 1998)
13




l
=



1


(
1
−


β


e


cos
θ
)




n


¯




e




σ


T





=



1.5
×


10


14



cm






n


¯




e
,
10


(
1
−


β


e


cos
θ
)



,


where 






β


e


c


 is the velocity of an electron, θ is the angle between the motion directions of the electron and photon, and 








n


¯




e
,
10


=




n


¯




e



/



10


10





cm


−
3




 has been assumed to be the mean electron number density of the magnetosphere. The light cylinder RL for a pulsar with period P* is 






R


L


=


cP


*



/

2
π
=
4.88
×


10


9


(


P


*



/

1

s
)

cm
≪
l


. Therefore, we conclude that a photon with frequency of order 1 GHz can propagate freely through the magnetosphere.
Citation: Lyubarskii & Petrova 1998
Citation function: Uses
Citation: Lyubarskii & Petrova (1998)
Citation function: Uses; Uses
Citation: Lyubarskii & Petrova 1998
Citation function: Uses; Uses
Our next goal is to estimate the density, spatial extent, and total baryonic mass contained in the CGM of NGC 3221, given that the detected signal is truly from NGC 3221. We calculate the average EM profile from the cumulative EI profile. As the EI of the larger annulus encompasses the EI of the smaller annulus, the uncertainties are not independent. We assume that the correlation coefficient of the consecutive measurements is 0.5, i.e., the covariance is half the geometric mean of their own variance: σi,i−1 = 0.5σiσi−1. We fit the EM profile with a β-model. The β-model has a density profile
3





where no is the central density and rc is the core radius, resulting in an EM profile similar to the surface brightness profile (Sarazin 1986)
4





where EMo is the central EM. Because of the small number of data points, large error bars, and the degeneracy between the parameters, we could not fit the three parameters of the β-model simultaneously. Therefore, as we did for the surface brightness profile in Section 3.1, we fit the EM profile in two different ways:

(1)
A β-model: We first fixed β = 0.5 and fit the radial profile for EMo and rc. Then, we fixed rc to the best-fit value (152.4 kpc) and fit for EMo and β. The resulting fit (χ2/dof = 3.3/4) yielded parameter values EMo =(2.8 ± 1.1) × 10−5 cm−6 kpc and β = 0.56 ± 0.36. We could not determine the range of rc from the fit, but we obtain the range of rc empirically for the best-fitted values of EMo and β that is consistent with the 1σ error span of the EM profile, with minimum rc = 110 kpc and maximum rc = 225 kpc. Interestingly, the best-fitted values and the confidence intervals are consistent with those obtained by fitting the radial profile of surface brightness in Section 3.1. Additionally, we fit the EM profile with Model B, a constant-density model.


(2)
A constant-density model: The truncated constant-density homogeneous medium (n = no(constant)) has an elliptical profile of projected EM,
5





where Rout is the spatial extent of the gaseous medium. We obtain best-fitted (χ2/dof = 1.4/4) values of n =





 = (2.7 ± 0.2) × 10−4 cm−3 and Rout = 175 ± 2 kpc (Figure 8).
Citation: Sarazin 1986
Citation function: Uses; Uses
We study here the Sco OB1 association (Figs. 1 and 2), using this and other techniques. The general properties of this large OB association, which spans almost 5° on the sky, and is surrounded by a ring-shaped HII region called Gum 55, are reviewed by Reipurth (2008). Its central cluster NGC 6231 contains several tens of OB stars, which have been extensively studied. On the other hand, many fewer studies, all recent, were devoted to the full mass spectrum, using optical photometry (Sung et al. 1998, 2013) and X-rays (Sana et al. 2006, 2007; Damiani et al. 2016; Kuhn et al. 2017a,b). The currently accepted distance of NGC 6231 is approximately 1580 pc, and its age is between 2and 8 Myr, with a significant intrinsic spread (Sung et al. 2013; Damiani et al. 2016). No ongoing star formation is known to occur therein, however. Approximately one degree North of the cluster, the loose cluster Trumpler 24 (Tr 24) also belongs to the association. There is little literature on this cluster (Seggewiss 1968; Heske & Wendker 1984, 1985; Fu et al. 2003, 2005) which unlike NGC 6231 lacks a well-defined center and covers about one square degree on the sky. Its age is  10 Myr according toHeske & Wendker (1984, 1985), who find several PMS stars, and its distance is 1570–1630 pc according to Seggewiss (1968). Other studies of the entire Sco OB1 association include MacConnell & Perry (1969 – Hα-emission stars), Schild et al. (1969 – spectroscopy), Crawford et al. (1971 – photometry), Laval (Laval 1972a,b – gas and star kinematics, respectively), van Genderen et al. (1984 – Walraven photometry), and Perry et al.(1991 – photometry). At the northern extreme of Sco OB1, the partially obscured HII region G345.45+1.50 and its less obscured neighbor IC4628 were studied by Laval (1972a), Caswell & Haynes (1987), López et al. (2011), and López-Calderón et al. (2016). They contain massive young stellar objects (YSOs; Mottram et al. 2007), maser sources (Avison et al. 2016), and the IRAS source 16562-3959 with its radio jet (Guzmán et al. 2010), outflow (Guzmán et al. 2011), and ionized wind (Guzmán et al. 2014), and are therefore extremely young (1 Myr or less). The distance of G345.45+1.50 was estimated as 1.9 kpc by Caswell & Haynes (1987), and 1.7 kpc by López et al. (2011), in fair agreement with distances of Sco OB1 stars. In Fig. 1 of Reipurth (2008) a strip of blue stars is visible, connecting NGC 6231 to the region of IC4628.
Citation: Heske & Wendker; 1985
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Citation: Heske & Wendker; 1985
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