Datasets:
ReadingTimeMachine
/
rtm-sgt-ocr-v1

Dataset card Data Studio Files Community
2
source
stringlengths
1
2.05k
⌀
target
stringlengths
1
11.7k
This equation essentially says that ve atomic or ionic temperature Z(m)=TyGn/mo)*.
This equation essentially says that the atomic or ionic temperature $T(m) = T_0 (m/m_0)^d$.
The form chosen for Equation (10) is μαelevant because. in the solar coronal case. the heating has been shown to be more pronounced ian "mass proporüonal healing” (Cranmer2002).. which corresponds to d>1.
The form chosen for Equation (10) is relevant because, in the solar coronal case, the heating has been shown to be more pronounced than “mass proportional heating” \citep{Cranmer02}, which corresponds to $d \geq 1$.
A fit of Equation (10) to the data introduces three model parameters €. (74.d). rather ian (he two parameters of Equation (1) utilized bv RedfieldandLinsky (2004)..
A fit of Equation (10) to the data introduces three model parameters $T_0, \xi, d$ ), rather than the two parameters of Equation (1) utilized by \cite{Redfield04}. .
This means jab there is a broader basin of acceptability in a X? sense.
This means that there is a broader basin of acceptability in a $\chi^2$ sense.
The following analvsis was undertaken.
The following analysis was undertaken.
Arp 2 is a globular cluster locaed at/= s.54° 5b=20.78" fa=19"28"4p. 8=30°21"14". J2000.0).
Arp 2 is a globular cluster located at $l=8.54^{o}$ , $b=-20.78^{o}$ $\alpha=19^h~28^m~44^s$, $\delta=-30^{o}21^{\prime}~14^{\prime\prime}$, J2000.0).
It is commonly believed to have formed inside the Ser dwarf spheroidal galaxy (Monaco et al.
It is commonly believed to have formed inside the Sgr dwarf spheroidal galaxy (Monaco et al.
2005). and then released into the Milky Way through tidal interaction.
2005), and then released into the Milky Way through tidal interaction.
With a metallicity of [Fe/H]=-1.77 (Mottini et al.
With a metallicity of [Fe/H]=-1.77 (Mottini et al.
2008). this cluster appears to be 3-4 Gyr younger than the old globulars. but ~1—2 Gyr older than the youngest globulars associated to Ser (Carraro et al.
2008), this cluster appears to be 3-4 Gyr younger than the old globulars, but $\sim1-2$ Gyr older than the youngest globulars associated to Sgr (Carraro et al.
2007. Layden Sarajedini 2000).
2007, Layden Sarajedini 2000).
The first photometric study of this cluster was performed by Buonanno et al. (
The first photometric study of this cluster was performed by Buonanno et al. (
1994).
1994).
The derived color-magnitude diagram (CMD) reveals an intriguing feature. namely that the horizontal branch (HB) is located entirely blue-ward of the RR-Lyrae instability strip. a fact that allowed the authors to assess its age through a differential comparison with 47 Tuc and Ruprecht 106.
The derived color-magnitude diagram (CMD) reveals an intriguing feature, namely that the horizontal branch (HB) is located entirely blue-ward of the RR-Lyrae instability strip, a fact that allowed the authors to assess its age through a differential comparison with 47 Tuc and Ruprecht 106.
A secondary. prominent feature. which the authors do not comment On. is a group of stars right above the turn-off (TO). which are probably the blue straggler stars (BSS) population in Arp 2.
A secondary, prominent feature, which the authors do not comment on, is a group of stars right above the turn-off (TO), which are probably the blue straggler stars (BSS) population in Arp 2.
BSS are anormal stellar population in clusters. since they are present in all of the properly observed Globular Clusters (GC. Ferraro 2006: Ferraro et al.
BSS are a normal stellar population in clusters, since they are present in all of the properly observed Globular Clusters (GC, Ferraro 2006; Ferraro et al.
2009. and reference therein).
2009, and reference therein).
Current scenarios for these stars in globulars are that either they are binary system with significant mass exchange. or stellar mergers resulting from direct collisions between two or more stars (Davies et al.
Current scenarios for these stars in globulars are that either they are binary system with significant mass exchange, or stellar mergers resulting from direct collisions between two or more stars (Davies et al.
2004: Knigge et al.
2004; Knigge et al.
2009: Perets Frabrycky In all these studies a proper assessment of the membership of BSSs through comparison with Red Giant Branch (RGB) stars is routinely performed (Ferraro et al.
2009; Perets Frabrycky In all these studies a proper assessment of the membership of BSSs through comparison with Red Giant Branch (RGB) stars is routinely performed (Ferraro et al.
1993).
1993).
The comparison of their cumulative radial distribution may hint to a possible common origin. specifically to confirm or deny whether they belong to the same parent Additionaly. the radial distribution of BSS in a star cluster is the most effective tool to understand their origin and which is the dominant production channel (Ferraro 2006).
The comparison of their cumulative radial distribution may hint to a possible common origin, specifically to confirm or deny whether they belong to the same parent Additionaly, the radial distribution of BSS in a star cluster is the most effective tool to understand their origin and which is the dominant production channel (Ferraro 2006).
BSS. are routinely found - with the exception of Omega Cen and NGC 2419 (Dalessandro et al.
BSS are routinely found - with the exception of Omega Cen and NGC 2419 (Dalessandro et al.
2008)- to be centrally concentrated.
2008)- to be centrally concentrated.
Their radial protile then smooths down. while in the cluster periphery it shows again an increase of the BSS contribution (Lanzoni et al.
Their radial profile then smooths down, while in the cluster periphery it shows again an increase of the BSS contribution (Lanzoni et al.
2007. Dalessandro et al.
2007, Dalessandro et al.
In this paper we report on a new photometric data-set of Arp 2 obtained with the goal of analyzing the BSS populationin a
In this paper we report on a new photometric data-set of Arp 2 obtained with the goal of analyzing the BSS populationin a
Oguriοἱal.(2002).. fit the LCDAI model well.
\citet{ogu02}, fit the LCDM model well.
Comparison of Figure 9 with (2002)s Figure 6 indicates (hat our compound model fits the observations better.
Comparison of Figure \ref{fig8} with \citet{ogu02}' 's Figure 6 indicates that our compound model fits the observations better.
The single population model predicts a single (almost) straight line in the leA9—Mua space.
The single population model predicts a single (almost) straight line in the $\lg\Delta\theta - \lg\Delta t_{\rm med}$ space.
For the compound model. a 7step is produced at the point where the mass density. prolile changes.
For the compound model, a “step” is produced at the point where the mass density profile changes.
The “step” that we see in Figure 9. corresponds to the transition from population A (galaxies) to population B (galaxy eroups/clusters).
The “step” that we see in Figure \ref{fig8} corresponds to the transition from population A (galaxies) to population B (galaxy groups/clusters).
As an extension of our previous work (LOO2). we computed the lensing probability produced by a compound population of dark halos.
As an extension of our previous work (LO02), we computed the lensing probability produced by a compound population of dark halos.
We have caleulated the lensing probability or both image separation and time delay.
We have calculated the lensing probability for both image separation and time delay.
The calculations confirma our previous results (LOO?) that the lensing probability produced by GNEW halos with a<1.3 is lower than that produced. by SIS halos with same masses by orders of magnitudes. where —a is the inner slope of the halo mass density.
The calculations confirm our previous results (LO02) that the lensing probability produced by GNFW halos with $\alpha \la 1.3$ is lower than that produced by SIS halos with same masses by orders of magnitudes, where $-\alpha$ is the inner slope of the halo mass density.
So. for the compound population of halos. both the iunber of lenses with large image separation (AG25") and the number of lenses with small image separation (A@10 7”) are greatly suppressed.
So, for the compound population of halos, both the number of lenses with large image separation $\Delta\theta \ga 5^{\prime\prime}$ ) and the number of lenses with small image separation $\Delta\theta \la 10^{-2\,\prime\prime}$ ) are greatly suppressed.
The same conclusion holds also for the number of lenses with large time delay (Al10 vears) and the number of lenses with «mall time delay (AZ<10! vear). (
The same conclusion holds also for the number of lenses with large time delay $\Delta t \ga 10\,{\rm years}$ ) and the number of lenses with small time delay $\Delta t \la 10^{-4}\,{\rm year}$ ). (
See Figs.
See Figs.
1. and 6..
\ref{fig1} and \ref{fig5}.
This conclusion holds even when the ellect of magnification bias is considered. see Figs.
This conclusion holds even when the effect of magnification bias is considered, see Figs.
4. and 5..)
\ref{fig4a} and \ref{fig4}. .)
We have also tested the dependence of the lensing probability on the redshilt of the source object (Figs. 2... 3.. 7..
We have also tested the dependence of the lensing probability on the redshift of the source object (Figs. \ref{fig2}, \ref{fig3}, \ref{fig6},
and 8)).
and \ref{fig7}) ).
The results show that. the lensing probability is «uite sensitive to the change in the redshilt of the source object.
The results show that, the lensing probability is quite sensitive to the change in the redshift of the source object.
The number of lenses significantly increases as (he source redshift increases.
The number of lenses significantly increases as the source redshift increases.
ILowever. (he rate of the increase decreases as the source redshift becomes large. which is caused bv (he [act that the proper cosmological distance approaches a finite limit when ος—o.
However, the rate of the increase decreases as the source redshift becomes large, which is caused by the fact that the proper cosmological distance approaches a finite limit when $z_S \rightarrow \infty$.
Another interesting result is that. the peak of the lensing probability for each population moves toward large image separation or time delay. as the source redshift increases.
Another interesting result is that, the peak of the lensing probability for each population moves toward large image separation or time delay, as the source redshift increases.
We see (hal population C (dwarl halos) in an LCDAM model has a unique signature in the time domain. ο,
We see that population C (dwarf halos) in an LCDM model has a unique signature in the time domain, c.f.
Figures ο and 7..
Figures \ref{fig5} and \ref{fig6}.
Time delays of less than 10 seconds and greater than 0.1 second are predicted and should be found in ganuna-ray burst sources which are al cosmological distances and have the requisite temporal substructure.
Time delays of less than $10$ seconds and greater than 0.1 second are predicted and should be found in gamma-ray burst sources which are at cosmological distances and have the requisite temporal substructure.
Variants of CDM. such as warm dark matter (Bode.Ostriker.&Turok2001).. repulsive dark matter (Goodman 2000).. or collisional dark matter (Spergel&Steinhard2000) would not produce Chis feature.
Variants of CDM, such as warm dark matter \citep{bod01}, repulsive dark matter \citep{goo00}, or collisional dark matter \citep{spe00} would not produce this feature.
llowever. current. survevs do not go deep enough to provide a sulliciently large sample to test (he prediction.
However, current surveys do not go deep enough to provide a sufficiently large sample to test the prediction.
When more observational data on gamnma-ray burst tme delay and small
When more observational data on gamma-ray burst time delay and small
then application of the new Alixmaster algorithm to the Cowdy Tamiltonian (203) with 77=II,|Πω fox aud from the variation of {1 (where £. αν and 7 are functions of 0)aud
then application of the new Mixmaster algorithm to the Gowdy Hamiltonian \ref{ber-gowdywaveh}) ) with $H = H_1 + H_2$ for and These yield the exact solutions from the variation of $H_1$ (where $\xi$, $\kappa$ , and $\tau_0$ are functions of $\theta$ )and
It should be noted that about half of the Seyferts in the Maiolino and Rieke's (1995)) sample have [OIII] luminosities lower than NGC 4941, and have not been studied in the X rays.
It should be noted that about half of the Seyferts in the Maiolino and Rieke's \cite{maiolino_a}) ) sample have [OIII] luminosities lower than NGC 4941, and have not been studied in the X rays.
Therefore, the hard X-ray spectral properties of the lowest luminosity AGN population have still to be probed.
Therefore, the hard X–ray spectral properties of the lowest luminosity AGN population have still to be probed.
The most remarkable result of our survey is that all the objects are heavily obscured with Ny> 10??9cm-? and, in particular, 6 out of 8 objects are Compton thick.
The most remarkable result of our survey is that all the objects are heavily obscured with $_H > 10^{23.6}$ $^{-2}$ and, in particular, 6 out of 8 objects are Compton thick.
More specifically NGC2273, NGC3393, NGC4939, NGC5643 and MCG-05-18-002 were identified as Compton thick with Ny> 102cm-?, NGC1386 is Compton thick with Ng> 107cm-?, while NGC3081 and NGC 4941 were identified as Compton thin with Ng~5x 10?8cm7?.
More specifically NGC2273, NGC3393, NGC4939, NGC5643 and MCG-05-18-002 were identified as Compton thick with $_H > 10^{25}$ $^{-2}$, NGC1386 is Compton thick with $_H > 10^{24}$ $^{-2}$, while NGC3081 and NGC 4941 were identified as Compton thin with $_H \sim 5\times 10^{23}$ $^{-2}$.
However, the nature of NGC4939, NGC4941 and 05-18-002 is still questionable.
However, the nature of NGC4939, NGC4941 and MGC-05-18-002 is still questionable.
'This result has to be compared with former spectral surveys.
This result has to be compared with former spectral surveys.
Smith Done (1996)) studied the spectra of a sample of type 2 and 1.9 Seyferts observed with Ginga, that were probably selected amongst bright X—ray sources.
Smith Done \cite{smith}) ) studied the spectra of a sample of type 2 and 1.9 Seyferts observed with Ginga, that were probably selected amongst bright X–ray sources.
The distribution of Ng in their sample is shown in Fig. 5..
The distribution of $_H$ in their sample is shown in Fig. \ref{fig_nh}.
The average absorbing column density is about 1022cm-?.
The average absorbing column density is about $10^{22.6}$ $^{-2}$.
Turner et al. (1997a,,1997b))
Turner et al. \cite{turner_a}, \cite{turner_b}) )
analyzed the spectra of a sample of Sy2s from the ASCA archive.
analyzed the spectra of a sample of Sy2s from the ASCA archive.
The latter sample includes a larger fraction of weak Sy2s with respect to the Ginga survey (see discussion in the former section).
The latter sample includes a larger fraction of weak Sy2s with respect to the Ginga survey (see discussion in the former section).
As a consequence, it contains a larger fraction of heavily absorbed Sy2s, as shown in Fig. 5..
As a consequence, it contains a larger fraction of heavily absorbed Sy2s, as shown in Fig. \ref{fig_nh}.
Our sample was selected by means of an isotropic indicator of the nuclear luminosity, the dereddened [ΟΠΠ
Our sample was selected by means of an isotropic indicator of the nuclear luminosity, the dereddened [OIII]
The main differences between the different WIMP models as regards the luminosity evolution in Fig.2 are in the time the "pure" dark star phase lasts.
The main differences between the different WIMP models as regards the luminosity evolution in \ref{evol} are in the time the “pure” dark star phase lasts.
The higher the boost factor, the shorter this phase.
The higher the boost factor, the shorter this phase.
Conversely a larger concentrationparameter ¢ prolongs the DS phase, since more DM is available.
Conversely a larger concentrationparameter $c$ prolongs the DS phase, since more DM is available.
In Fig.
In Fig.
3 we plot Hertzsprung-Russell (H-R) diagrams for the four cases.
\ref{HR} we plot Hertzsprung-Russell (H-R) diagrams for the four cases.
One can see two distinct phases.
One can see two distinct phases.
First, the DS goes up the Hayashi track with a very steep increase of the luminosity yet relatively cool surface temperature. 7;;;<10! K. At the end of the Hayashi track the star enters the Henyey track.
First, the DS goes up the Hayashi track with a very steep increase of the luminosity yet relatively cool surface temperature, $T_{eff}\leq 10^4$ K. At the end of the Hayashi track the star enters the Henyey track.
This path corresponds to an almost constant luminosity while the temperature increases fast, mostly due to the KH contraction phase.
This path corresponds to an almost constant luminosity while the temperature increases fast, mostly due to the KH contraction phase.
As a rule of thumb once a star is on the Heyney track its core should be fully radiative.
As a rule of thumb once a star is on the Heyney track its core should be fully radiative.
The graphs end at a temperature of ~10° K when the star reaches the main sequence.
The graphs end at a temperature of $\sim 10^5$ K when the star reaches the main sequence.
From the left panel of Fig.
From the left panel of Fig.
3 one can see that the boosted AH4 case has the highest luminosity, due to the extremely efficient DM heating (Qi;~(ov)/im, forming a luminosity peak.
\ref{HR} one can see that the boosted AH4 case has the highest luminosity, due to the extremely efficient DM heating $Q_{DM}\sim \sv/m_{\chi}$ forming a luminosity peak.
However, as the AH4 case burns up its DM, the its luminosity falls.
However, as the AH4 case burns up its DM, the its luminosity falls.
The boosted and unboosted cases eventually cross over at a temperature of ~104 Κ. and henceforth the unboosted case has a higher luminosity.
The boosted and unboosted cases eventually cross over at a temperature of $\sim 10^4$ K, and henceforth the unboosted case has a higher luminosity.
Consequently, the boosted AH4 case has the lowest luminosity as the star moves onto the main sequence, as discussed above.
Consequently, the boosted AH4 case has the lowest luminosity as the star moves onto the main sequence, as discussed above.
Inthe right panel ofFig.
In the right panel of Fig.
3 the trend is uniform: an increase in the concentration parameter leads to an inerease in the luminosity.
\ref{HR} the trend is uniform: an increase in the concentration parameter leads to an increase in the luminosity.
The difference is relatively small in the early stages of the evolution, at low temperatures.
The difference is relatively small in the early stages of the evolution, at low temperatures.
This is due to the fact that the adiabatically contracted DM density profile is not very sensitive to the concentration parameter, therefore about the same amount of DM heating will be generated in each case.
This is due to the fact that the adiabatically contracted DM density profile is not very sensitive to the concentration parameter, therefore about the same amount of DM heating will be generated in each case.
However, for à lower value of the concentration parameter, the adiabatically contracted DM runs out faster. as there is less DM available.
However, for a lower value of the concentration parameter, the adiabatically contracted DM runs out faster, as there is less DM available.
This leads to a shorter “pure” DS phase, as can also be seen from Fig. 2..
This leads to a shorter “pure” DS phase, as can also be seen from Fig. \ref{evol}, ,
and consequently to slightly lower final mass and luminosities.
and consequently to slightly lower final mass and luminosities.
choice of nominal masses of MW and MBL. 1.5x10775, and 3x10P5..
choice of nominal masses of MW and M31, $1.5\times 10^{12}m_\odot$ and $3\times 10^{12}m_\odot$.
This approximates conventional wisdom. but the choice may have led to a local mininnun of 4? with similar model masses. 1.6x1075, for MW and 2.4x10M. for M31.
This approximates conventional wisdom, but the choice may have led to a local minimum of $\chi^2$ with similar model masses, $1.6\times 10^{12}m_\odot$ for MW and $2.4\times 10^{12}m_\odot$ for M31.
An attempt to adjust the solution to bring these (vo masses closer by adding to 4? a penalty [or a significant mass difference had little effect.
An attempt to adjust the solution to bring these two masses closer by adding to $\chi^2$ a penalty for a significant mass difference had little effect.
This might be expected because changing the masses of these (wo ealaxies requires consistent. adjustments of the redshifts and distances of many LG galaxies. a slow operation by the present numerical method.
This might be expected because changing the masses of these two galaxies requires consistent adjustments of the redshifts and distances of many LG galaxies, a slow operation by the present numerical method.
The third point to consider is that our analvsis has allowed masses considerable freedom to float to aid the fit to distances and velocities.
The third point to consider is that our analysis has allowed masses considerable freedom to float to aid the fit to distances and velocities.
This means that some erroneous choices of orbits may have been be made to fit (he measured redshifts aud distances by the choice of erroneous masses.
This means that some erroneous choices of orbits may have been be made to fit the measured redshifts and distances by the choice of erroneous masses.
Further investigation of the last two issues will require reconstruction of the model. which we may hope will be aided by future still lighter constraints on distances aud proper motions that reduce the chance of including erroneous orbits.
Further investigation of the last two issues will require reconstruction of the model, which we may hope will be aided by future still tighter constraints on distances and proper motions that reduce the chance of including erroneous orbits.
The fit to constraints on LG. galaxy. positions and velocities relaxes the MW. circular velocity to ce,=256 km +.
The fit to constraints on LG galaxy positions and velocities relaxes the MW circular velocity to $v_c = 256$ km $^{-1}$.
This is 36 km s! larger than a conventional value and 26 km ! larger the nominal central value used in the computation of 47. but it is consistent with the Reid et al. (
This is 36 km $^{-1}$ larger than a conventional value and 26 km $^{-1}$ larger the nominal central value used in the computation of $\chi^2$, but it is consistent with the Reid et al. (
2009a) measurement. 254416 km !.
2009a) measurement, $254 \pm 16$ km $^{-1}$.