Datasets:

ReadingTimeMachine
/

rtm-sgt-ocr-v1

Dataset card Data Studio Files Files and versions Community

Split (1)

train · 1.53M rows

Subsets and Splits

No community queries yet

The top public SQL queries from the community will appear here once available.

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
In total there were 10 regions which were selected for follow-up spectra which consisted of ss mmin) of ON source time.	In total there were 10 regions which were selected for follow-up spectra which consisted of s min) of ON source time.
Data were reduced. using the ATNE packageslivedata.gridzilla. and lor mapping. was used to perform a bandpass calibration for each row. using the preceding olf scan as a reference and applieda 1 order polvnomial fit (Le. linear) to the baseline.	Data were reduced using the ATNF packages, and For mapping, was used to perform a bandpass calibration for each row, using the preceding off scan as a reference and applieda $^{\rm st}$ order polynomial fit (i.e. linear) to the baseline.
re-gridded and combined all data fron all mapping scans onto a single cata cube with pixels 15"15	re-gridded and combined all data from all mapping scans onto a single data cube with pixels $^{\prime\prime}\times15^{\prime\prime}\times$
bypothesiis that TINCO could be a good traccor of duterstellar shocks.	is that HNCO could be a good tracer of interstellar shocks.
star formation, the dark matter has a roughly isothermal inner density profile with y~2.	star formation, the dark matter has a roughly isothermal inner density profile with $\gamma \approx 2$.
The central dark matter density increases by over an order of magnitude compared with the non-radiative case.	The central dark matter density increases by over an order of magnitude compared with the non-radiative case.
Figure 5 shows the local enclosed mass fraction M/Mx« of the different matter species in simulation A. 'The fractions sum up to unity at each radius.	Figure \ref{fig:fraction} shows the local enclosed mass fraction $M/M_\mathrm{tot}$ of the different matter species in simulation A. The fractions sum up to unity at each radius.
Generally, the stars dominate in inner of the virial radius 3 while the dark matter dominates in the outer (aboutparts of kpc),the halo.	Generally, the stars dominate in inner of the virial radius (about 3 kpc), while the dark matter dominates in the outer parts of the halo.
The gas mass near the center is typically higher than the dark matter mass.	The gas mass near the center is typically higher than the dark matter mass.
An interesting feature is the much smaller scatter of the combined baryon profile than of the gas or stellar component separately.	An interesting feature is the much smaller scatter of the combined baryon profile than of the gas or stellar component separately.
The star formation history varies from object to object, but whenever the stellar density is higher than the median value the gas density is lower by a similar amount, and vice versa.	The star formation history varies from object to object, but whenever the stellar density is higher than the median value the gas density is lower by a similar amount, and vice versa.
This is encouraging for our study.	This is encouraging for our study.
While our modeling of the star formation rate and feedback may not be exactly right, the effect of the baryon dissipation on the dark matter properties is calculated more robustly.	While our modeling of the star formation rate and feedback may not be exactly right, the effect of the baryon dissipation on the dark matter properties is calculated more robustly.
The universal baryon fraction in our simulations is =0.164.	The universal baryon fraction in our simulations is $\Omega_\mathrm{B,0}/\Omega_\mathrm{M,0} = 0.164$.
In general, we find that the median baryonQp,0/Qu,o fraction within the virial radius, fg(r2oop,a), is higher than universal at all times (by ~5—15%) for the runs with cooling and star formation but lower than universal (by &5%) for the non-radiative case.	In general, we find that the median baryon fraction within the virial radius, $f_\mathrm{B}(r_\mathrm{200b,A})$, is higher than universal at all times (by $\approx 5-15\%$ ) for the runs with cooling and star formation but lower than universal (by $\approx 5\%$ ) for the non-radiative case.
Also, fe(reoon,a) generally increases with time, up to above the universal fraction at z=2.	Also, $f_\mathrm{B}(r_\mathrm{200b,A})$ generally increases with time, up to above the universal fraction at $z=2$.
The halos in the simulations without supernova feedback can retain slightly more baryons than the halos in the simulations with supernova feedback.	The halos in the simulations without supernova feedback can retain slightly more baryons than the halos in the simulations with supernova feedback.
But the effect is relatively weak (between and 9%)) and within the scatter among individual halos.	But the effect is relatively weak (between and ) and within the scatter among individual halos.
These mass fractions should only be used for a qualitative comparison between the different simulations.	These mass fractions should only be used for a qualitative comparison between the different simulations.
There is nothing special about our choice of ragop,a as the length scale.	There is nothing special about our choice of $r_\mathrm{200b,A}$ as the length scale.
A smooth halo profile extends much further (we plot it out to two virial radii), and in general an edge of a halo is ill defined (seealso,e.g.,??)..	A smooth halo profile extends much further (we plot it out to two virial radii), and in general an edge of a halo is ill defined \citep[see also, e.g.,][]{2006ApJ...645.1001P,2008MNRAS.389..385C}.
In order to describe the more concentrated matter distribution in the simulations with cooling and star formation, we use an intrinsic and general measure for the concentration of a halo.	In order to describe the more concentrated matter distribution in the simulations with cooling and star formation, we use an intrinsic and general measure for the concentration of a halo.
It is the enclosed density within the radius rmax (location of the peak of the circular velocity curve Vmax) in units of the critical density at z—0: This concentration measure has the advantage that it is well defined both for isolated halos and subhalos (as long as the peak of the circular velocity curve can be found) and it does not make any assumptions about a specific shape of the density profile (??)..	It is the enclosed density within the radius $r_{\mathrm{max}}$ (location of the peak of the circular velocity curve $V_\mathrm{max}$ ) in units of the critical density at $z=0$: This concentration measure has the advantage that it is well defined both for isolated halos and subhalos (as long as the peak of the circular velocity curve can be found) and it does not make any assumptions about a specific shape of the density profile \citep{2002ApJ...572...34A, 2007ApJ...667..859D}.
In principle, ον can be derived from observable quantities.	In principle, $c_\mathrm{V}$ can be derived from observable quantities.
An alternative interpretation is that cy is related to the number of rotations, Nrot, at rax per Hubble time, 1/Ho, by Another common measure for a halo concentration is the virial concentration cyir= where rs is a characteristic scale radius.	An alternative interpretation is that $c_\mathrm{V}$ is related to the number of rotations, $N_\mathrm{rot}$ , at $r_{\mathrm{max}}$ per Hubble time, $1/H_0$, by Another common measure for a halo concentration is the virial concentration $c_\mathrm{vir} \equiv r_\mathrm{vir}/r_\mathrm{s}$, where $r_\mathrm{s}$ is a characteristic scale radius.
The virial Tvir/1s,concentration has two main drawbacks: i) cy, grows even when the inner mass distribution remains constant, due to the comoving definition of the virial radius, and ii) cyir is not well defined for subhalos.	The virial concentration has two main drawbacks: i) $c_\mathrm{vir}$ grows even when the inner mass distribution remains constant, due to the comoving definition of the virial radius, and ii) $c_\mathrm{vir}$ is not well defined for subhalos.
If an analytical halo density profile is known, it is straightforward to calculate the mapping between cy and cyiy.	If an analytical halo density profile is known, it is straightforward to calculate the mapping between $c_\mathrm{V}$ and $c_\mathrm{vir}$.
In the case of an NFW profile (?) see for example Figure 5 in ?..	In the case of an NFW profile \citep{1996ApJ...462..563N} see for example Figure 5 in \cite{2007ApJ...667..859D}.
However, in dissipative simulations dark matter no longer follows the NFW profile (the inner parts are modified more than the outer parts) and cy is a more useful measure of halo compactness.	However, in dissipative simulations dark matter no longer follows the NFW profile (the inner parts are modified more than the outer parts) and $c_\mathrm{V}$ is a more useful measure of halo compactness.
Figure 6 shows the concentrations of the objectsin run A versus the concentrations in run B. Here we determine	Figure \ref{fig:concentration} shows the concentrations of the objectsin run A versus the concentrations in run B. Here we determine
Galaxy properties (morphology. color. star formation. gas content. etc.)	Galaxy properties (morphology, color, star formation, gas content, etc.)
have a strong dependence on the local environment where galaxies reside.	have a strong dependence on the local environment where galaxies reside.
Several works found evidence that galaxies in high density local environment such as groups and clusters. show different properties than their isolated counterparts (e.g. Dressler 1980: Balogh et al.	Several works found evidence that galaxies in high density local environment such as groups and clusters, show different properties than their isolated counterparts (e.g. Dressler 1980; Balogh et al.
2004: Baldry et al.	2004; Baldry et al.
2006: Skibba Sheth 2009).	2006; Skibba Sheth 2009).
Groups and clusters are dominated by spheroidal and gas poor galaxies. because when spiral galaxies traverse the dense local environments of clusters. stripping removes their interstellar gas making them loose their ability to form new stars (e.g Gunn Gott 1972: Dressler 1980).	Groups and clusters are dominated by spheroidal and gas poor galaxies, because when spiral galaxies traverse the dense local environments of clusters, stripping removes their interstellar gas making them loose their ability to form new stars (e.g Gunn Gott 1972; Dressler 1980).
In addition. recent observational and theoretical works support the idea that galaxy properties also depend on the global environment.	In addition, recent observational and theoretical works support the idea that galaxy properties also depend on the global environment.
Regarding this topic. some authors (e.g. Balogh et al.	Regarding this topic, some authors (e.g. Balogh et al.
2004: Cecearelli et al.	2004; Ceccarelli et al.
2008: Park Chor 2009. Padilla et al.	2008; Park Choi 2009, Padilla et al.