ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

Several simulations of merging clusters have recently been presented in the literature that have similar characteristics with A1758N. Figure 3 in Bialek, Evrard, Mohr (2002) illustrates the trajectory of two subclusters that a deflection of about ~90? during their encounter.

Several simulations of merging clusters have recently been presented in the literature that have similar characteristics with A1758N. Figure 3 in Bialek, Evrard, Mohr (2002) illustrates the trajectory of two subclusters that undergo a deflection of about $\sim 90^{\circ}$ during their encounter.

Theundergo smaller cluster in their simulation has a steep X-ray surface brightness profile at its leading edge (i.e., a cold front), and is trailed by denser cooler gas that was displaced by ram pressure.

The smaller cluster in their simulation has a steep X-ray surface brightness profile at its leading edge (i.e., a cold front), and is trailed by denser cooler gas that was displaced by ram pressure.

A series of off-axis cluster mergers with different mass ratios was recently published by Ricker Sarazin (2001).

In their simulations with impact parameters of 2r, and 5r, (where r, is the scale radius in the NFW dark matter density profile), the clusters experience a significant deflection during their encounter.

In their simulations with impact parameters of $r_s$ and $r_s$ (where $r_s$ is the scale radius in the NFW dark matter density profile), the clusters experience a significant deflection during their encounter.

Their Figure 4 shows that both subclusters are with shocks as they separating.

Their Figure 4 shows that both subclusters are preceded with shocks as they begin separating.

The present state of precededthe gas in A1795N is beginprobably best reproduced by the simulations presented in Onura, Kay, Thomas (2002).

The present state of the gas in A1795N is probably best reproduced by the simulations presented in Onura, Kay, Thomas (2002).

Clusters 2 and 8 in their Figure 1 are bimodal with the two subclusters, identified by clumps of low entropy presently separating and preceded by cold fronts and shock gas,heated gas.

Clusters 2 and 8 in their Figure 1 are bimodal with the two subclusters, identified by clumps of low entropy gas, presently separating and preceded by cold fronts and shock heated gas.

There is no evidence for a shock front within 400 kpc of either subcluster in A1758N. The average gas temperature within this region has a significantly higher temperature than the gas in the two merging cores, indicating that any merger shocks induced during the encounter have already passed through most of the cluster.

As demonstrated by Vikhlinin, Markevitch, Murray (20014). the strength of a shock can

As demonstrated by Vikhlinin, Markevitch, Murray (2001a), the strength of a shock can

(D0515]705). D0531]6121. (D0539]|6200). D0538|7131 r)|6354 (D1526|670). BI53s|5920 (B1550|5815). B1600|7131. (B1531]|7122). 151|6707 and D1S41|6715 (DISA42|6851).

(B0518+705), B0531+6121 (B0539+6200), B0538+7131 (B0535+6743), B0755+6354 B1525+6801 (B1526+670), B1538+5920 (B1550+5815), B1600+7131 (B1531+722), B1819+6707 and B1841+6715 (B1842+681).

Data reduction of the phase referenced: observations is similar to that for the 5 Cllz data.

Data reduction of the phase referenced observations is similar to that for the 5 GHz data.

The twpical ve coverage obtained for a source at 15 CGllz is shown in figure 1..

The typical $u,v$ coverage obtained for a source at 15 GHz is shown in figure \ref{UV}.

The L6 CGllz data were obtained during two observing sessions. both involving the ten telescopes of the VLBA and l antennas of the EWN (see table 13).

The 1.6 GHz data were obtained during two observing sessions, both involving the ten telescopes of the VLBA and 4 antennas of the EVN (see table \ref{tel}) ).

The Westerbork data in the second. session was lost due to technical failure.

The Westerbork data in the second session was lost due to technical failure.

The data were recorded in 12842 mode (128 Mbits/sec. 4 IF channels. 2 bit/ssumple). with an cllective bandwidth of 32 MlIz centred at 1663 MIIz and 1655 during the first and second session respectively.

The data were recorded in $128-4-2$ mode (128 Mbits/sec, 4 IF channels, 2 bit/sample), with an effective bandwidth of 32 MHz centred at 1663 MHz and 1655 during the first and second session respectively.

In the first session. a subsample of 23 objects was observed. for 212 hours on 14 and 16 September 1997.

In the first session, a subsample of 23 objects was observed for $2\times12$ hours on 14 and 16 September 1997.

his subsample contained all sources with peak frequencies «5 Gllz. whieh were found to be extended in the 5 CGllz observations.

This subsample contained all sources with peak frequencies $<5$ GHz, which were found to be extended in the 5 GHz observations.

In the second. session. all 9 remaining sources with peak frequencies <3 CGllz. which had not been imaged before at this frequency. were observed.

In the second session, all 9 remaining sources with peak frequencies $<3$ GHz, which had not been imaged before at this frequency, were observed.

The sources were tvpically observed for 4.11 minutes each. and an example of à s.c coverage is shown in figure L..

The sources were typically observed for $4\times11$ minutes each, and an example of a $u,v$ coverage is shown in figure \ref{UV}.

The data were correlated in Socorro.

No fringes were found for BO513|7129. B0537|6444. and. D0544|5847.

No fringes were found for B0513+7129, B0537+6444, and B0544+5847.

Several sources in the second session were observed. using phase referencing.

Several sources in the second session were observed using phase referencing.

These sources. with their calibrators in brackets. are D0537|6444 (D0535|677). D0830|5813 1205060|513). DB1557|6220 (B15581595). B1639|6711 (D1700|655). and Biso0s|6813 (D1749|TOL).

These sources, with their calibrators in brackets, are B0537+6444 (B0535+677), B0830+5813 (B0806+573), B1557+6220 (B1558+595), B1639+6711 (B1700+685), and B1808+6813 (B1749+701).

The data were reduced. in a similar way as the data at 5 Gllz.

The data were reduced in a similar way as the data at 5 GHz.

The parameters of the resulting 102 images (29 at 1.6 Cillz. 47 at 5 Cillz. and 26 at 15 (114) are given in table 2..

The parameters of the resulting 102 images (29 at 1.6 GHz, 47 at 5 GHz, and 26 at 15 GHz) are given in table \ref{mappar}.

ligure 2. shows the rms noise as function of the peak brightness in the images at the three observing frequencies.

Figure \ref{rmspeak} shows the rms noise as function of the peak brightness in the images at the three observing frequencies.

The dynamic ranges (defined as the ratio of the maximum brightness in the image to the rms noise in an area of blank sky) are between 125 and 2500 at 1.6 1. between 25 and 1700 at 5 Cllz. and between 30 and 500 at 15 Cllz.

The dynamic ranges (defined as the ratio of the maximum brightness in the image to the rms noise in an area of blank sky) are between 125 and 2500 at 1.6 GHz, between 25 and 1700 at 5 GHz, and between 30 and 500 at 15 GHz.

At 1.6 Gllz. two of the bright sources have higher rmis-noise levels than expected. which may indicate that the dynamic range is not limited. by the thermal noise.

At 1.6 GHz, two of the bright sources have higher rms-noise levels than expected, which may indicate that the dynamic range is not limited by the thermal noise.

To be able to compare the VLBI observations of this faint sample with those on bright GPS samples. it is important to determine whether components have been missed. due to the Iimited cdvnamic range for this faint sample.

To be able to compare the VLBI observations of this faint sample with those on bright GPS samples, it is important to determine whether components have been missed due to the limited dynamic range for this faint sample.

We therefore plotted the distribution of dvnamic range for the observations closest in frequency to the spectral peak (Fig. 3)).

We therefore plotted the distribution of dynamic range for the observations closest in frequency to the spectral peak (Fig. \ref{dynrange}) ).

Only 2 objects (D0755]6354 and D0544|5847) turn out not to have an image with a dynamic range οLOO.

Only 2 objects (B0755+6354 and B0544+5847) turn out not to have an image with a dynamic range $>100$.

In Figure 4. the ratio of total VLBI flux density in he images to the llux density in the NVSS at 1.6 CGllz. to he MERLLIN observations at 5 C€llz. and to the VLA 15 Cilz tus densities (from Snellen et al.

In Figure \ref{totvlbi} the ratio of total VLBI flux density in the images to the flux density in the NVSS at 1.6 GHz, to the MERLIN observations at 5 GHz, and to the VLA 15 GHz flux densities (from Snellen et al.

19982). are plotted.

1998a), are plotted.

This enables us to judge whether substantial structure has en resolved oul in the VLBI observations.

This enables us to judge whether substantial structure has been resolved out in the VLBI observations.

At 1.6 Cillz. vpicallv of the NVSS lux density. is recovered. in he VLBI observations. while at 5 CGllz the distribution »aks at.

At 1.6 GHz, typically of the NVSS flux density is recovered in the VLBI observations, while at 5 GHz the distribution peaks at.

10054... Only at 15 Cllz is the clistribution much oader and. peaks at about of the flux density in the VLA observations. and. hence provides some evidence that at this frequency some extended: structure may be missed.

Only at 15 GHz is the distribution much broader and peaks at about of the flux density in the VLA observations, and hence provides some evidence that at this frequency some extended structure may be missed.

The broadness of the peak is probably also influenced by variability.

Figures 7.. ὃν, 12. give the maps of the individual sources with observations at three. two and one frequency respectively.

Figures \ref{fig3}, \ref{fig2}, \ref{fig1} give the maps of the individual sources with observations at three, two and one frequency respectively.

For cach source. the images have the same size at cach frequency. ancl are centred in such a way

For each source, the images have the same size at each frequency, and are centred in such a way

galaxies to have brighter bulge central surface brightnesses with increasing disk central surface brightnesses.

The color profiles (Figs. 1..

The color profiles (Figs. \ref{profiles},

bottom panels) show that in most galaxies the outer parts are bluer than the inner parts.

It was shown by dB95 that this is also true for late-type LSB galaxies.

It is hard to draw conclusions about the colors of the bulges judging from the color profiles alone.

But one can reveal differences between disk and bulge colors by comparing bulge-to-disk (B/D) ratios in different wavelength bands.

We determined B/D ratios by comparing the total light output of bulge and disk per passband.

We noticed a clear tendency for the B/D ratios to increase towards the redder wavelengths (Table 4)).

We noticed a clear tendency for the B/D ratios to increase towards the redder wavelengths (Table \ref{bdratios}) ).

This means that the bulges of LSB galaxies are than their disks. confirming a trend also observed for HSB galaxies (4193).

This means that the bulges of LSB galaxies are than their disks, confirming a trend also observed for HSB galaxies (dJ95).

In Table 5 we present the integrated colors of the galaxies in our sample and those of disk dominated LSB (de Blok 1997)) and giant LSB galaxies (Sprayberry et al. 1995)).

In Table \ref{total_colors} we present the integrated colors of the galaxies in our sample and those of disk dominated LSB (de Blok \cite{de blok}) ) and giant LSB galaxies (Sprayberry et al. \cite{sprayberry}) ).

Prior to determining the mean nuclear color of the galaxies in our sample galaxies without a clear bulge were excluded.

For the determination of the mean area weighted colors the galaxies without a clear bulge were included.

The systems in our sample have redder area weighted colors than disk LSB galaxies and this would be more pronounced if the more or less bulgeless galaxies were left out of the sample.

Figure 9 shows the distribution of color with disk scale length (top) as well as the distribution of color (center) and color (bottom) with disk central surface brightness.

Figure \ref{mubv} shows the distribution of color with disk scale length (top) as well as the distribution of color (center) and color (bottom) with disk central surface brightness.

There is no trend of central surface brightness with color anc the colors do not depend on size.

There is no trend of central surface brightness with color and the colors do not depend on size.

HSB galaxies cannot be the progenitors of LSB galaxies since galaxies fade andredden.

HSB galaxies cannot be the progenitors of LSB galaxies since galaxies fade and.

. The bluest galaxies are the disk dominated LSB galaxies anc are concentrated in a rather small area whereas the HSB galaxies scatter over the entire color range towards the redder colors.

The bluest galaxies are the disk dominated LSB galaxies and are concentrated in a rather small area whereas the HSB galaxies scatter over the entire color range towards the redder colors.

The bulge dominated LSB galaxies fill up the region between these two samples.

The large scatter in color for the bulge dominated LSB sample is due to the wider range im morphological types.

The large scatter in color for the bulge dominated LSB sample is due to the wider range in morphological types.

The disk dominated LSB galaxies are quiescent and form a fairly uniform sample.

The galaxies i our sample have red bulges. but some of them also have bars and (blue) rings making the spread in color larger than for a more uniform sample.

The galaxies in our sample have red bulges, but some of them also have bars and (blue) rings making the spread in color larger than for a more uniform sample.

Some of the large bulge dominated LSB galaxies in our sample could easily be classified as giant LSB galaxies. judging from their sizes and luminosities.

Some of the large bulge dominated LSB galaxies in our sample could easily be classified as giant LSB galaxies, judging from their sizes and luminosities.

It is therefore instructive to compare the colors with those of giant LSB galaxies.

The disk colors of the biggest galaxies in our sample are significantly than giant disks (Table 5)) and the B/D ratios are smaller than for typical giants (Sprayberry et al. 1995)).

The disk colors of the biggest galaxies in our sample are significantly than giant disks (Table \ref{total_colors}) ) and the B/D ratios are smaller than for typical giants (Sprayberry et al. \cite{sprayberry}) ).

So although they have bulges and very large low surface brightness disks. they are not as evolved as the giants.

So although they have bulges and very large low surface brightness disks, they are not as evolved as the giants.

The comparison of our sample with a large sample of typical "Freeman galaxies" (de Jong van der Kruit 1994)) not only shows that on average the bulge dominated LSB galaxies are bluer. but also that one can distinguish à LSB from a HSB galaxy of the same type by its bluer color.

The comparison of our sample with a large sample of typical “Freeman galaxies” (de Jong van der Kruit \cite{de jong_vdkruit}) ) not only shows that on average the bulge dominated LSB galaxies are bluer, but also that one can distinguish a LSB from a HSB galaxy of the same type by its bluer color.

where. we using from the fluid equation. AMI| wp.

where, we using from the fluid equation, $\dot{\rho}=-3H(1+\omega)\rho$ .

Since. uw is Eos parameter for IDE is eiveu bx 1l Substituting the Eqs. (11)). (12))

Since, $\omega$ is The Eos parameter for HDE is given by Substituting the Eqs. \ref{7g}) ), \ref{8g}) )

ito the Eq. (103)

into the Eq. \ref{6g}) )

elves where OL,=JO,thfarID. is the curvature paraleter and G=G'/G Tere. we should follow up au expression related to he state parameter of equation at the present time.

gives where $\Omega_{ke}=\beta \Omega_k=-\beta k/a^2H^2,$ is the curvature parameter and ${\cal G}=G'/G$ Here, we should follow up an expression related to the state parameter of equation at the present time.

Since we have extracted the expressions for OA. we can calculate w(t) for simall redshifts :. performing he standard expansions of the literature.

Since we have extracted the expressions for $\Omega_{\Lambda}'$, we can calculate $w(z)$ for small redshifts $z$ , performing the standard expansions of the literature.

Iu particular. since pa~aOE wee acquire Expanding pa we have: here. the derivatives are taken at the present time ag=1.

In particular, since $\rho_\Lambda \sim a^{-3(1+w)}$ we acquire Expanding $\rho_{\Lambda}$ we have: here, the derivatives are taken at the present time $a_0=1$.

Then. (62) is giveu im the simall red shifts |2)ξ2 up to second order. as: We cau rewrite (16)) as: Simce after some calculation. aud some simplification we achieve to wy.wy as follows: Now. substituting Eq.(11)) iuto Eqs.(19)). obtain: ↖↖↽∐↸∖↥⋅↸∖∙∖∙∟⋝∙≦⋜⋯∖≼∐∖∐↕∐∖⋜↧↴∖↴↕⋟∪∐∪↖↖↰∖↴∶ Iu the following. we want compare the diagram of the state parameter equation at the preseut time versus À on both flat space and non-flat space for differcut redshifts.

Then, $w(z)$ is given in the small red shifts $\ln a=-\ln (1+z)=-z$ up to second order, as: We can rewrite \ref{2s}) ) as: Since after some calculation, and some simplification we achieve to $\omega_0, \omega_1$ as follows: Now, substituting \ref{10g}) ) into \ref{6s}) ), \ref{7s}) ) we obtain: where, $\chi, \zeta, \xi$ are define as follows: In the following, we want compare the diagram of the state parameter equation at the present time versus $\lambda$ on both flat space and non-flat space for different redshifts.

In the Fie.l we have plotted theequation of state parameter iz. versus A on both dat aud non flat space time for τ=0.01 and G=0.2 and small OQ...

In the Fig.1 we have plotted theequation of state parameter $\omega$, versus $\lambda$ on both flat and non flat space time for $z=0.01$ and ${\cal G}=0.2$ and small $\Omega_{ke}$.

We have obtained interval AA,=(0.32.0.93) aud AA»=(0.21.1.132) for fat space time (ας nou flat space time (0) respectivly. in which x accept the allows values between (€1.1)

We have obtained interval $\Delta\lambda_1=(0.32, 0.93)$ and $\Delta\lambda_2=(0.24, 1.132)$ for flat space time $(a)$ and non flat space time $(b)$ respectivly, in which $\omega$ accept the allows values between $(-1, 1)$.

as well as this action in the Fig.2 aud Fig. also have plotted for :=0.5. aud >=0.9on both flat (a) and nonflat (b) space time respectivlv.

as well as this action in the Fig.2 and Fig.3, also have plotted for $z=0.5$, and $z=0.9$on both flat (a) and nonflat (b) space time respectivly.

It is clearly που that is AA»> Άλι.

It is clearly seen that is $\Delta\lambda_2>\Delta \lambda_1$ .

This result show that the suitable interval for A which w accept the values between (.1.1) iu non flat space time

This result show that the suitable interval for $\lambda$ which $\omega$ accept the values between $(-1, 1)$ in non flat space time

Because the vast majority of scatterings occur near the line center. where the color temperature and kinetic temperature are very nearly (he same. (his generalization does not significantly change the result for the spin temperature.

Because the vast majority of scatterings occur near the line center, where the color temperature and kinetic temperature are very nearly the same, this generalization does not significantly change the result for the spin temperature.

We have found by a numerical test that (his correction affects the spin temperature by only a small fraction of a percent.

We have found by a numerical test that this correction affects the spin temperature by only a small fraction of a percent.

We will now consider a simple model where the emissivitv of photons (turns on al a redshift z; and increases linearly with redshift thereafter.

We will now consider a simple model where the emissivity of photons turns on at a redshift $z_i$ and increases linearly with redshift thereafter.

We will also consider the effect of X-ray emission [rom the same star formation regions. which dominates (he heating rate of the eas.

We will also consider the effect of X-ray emission from the same star formation regions, which dominates the heating rate of the gas.

This will illustrate a plausible thermal history of the eas and the expected strength of the absorption or emission signal in (he redshilted 21cm line.

This will illustrate a plausible thermal history of the gas and the expected strength of the absorption or emission signal in the redshifted 21cm line.

The kinetic temperature of the gas in the expanding Universe evolves as where DL;,; aud A, are the total heating auc cooling rates. respectively. and nv is the total gas munber density.

The kinetic temperature of the gas in the expanding Universe evolves as where $\Gamma_{tot}$ and $\Lambda_{tot}$ are the total heating and cooling rates, respectively, and $n$ is the total gas number density.

In the absence of heating aud cooling. the eas temperature decreases acliabatically with Tx(1zy.

In the absence of heating and cooling, the gas temperature decreases adiabatically with $T \propto (1+z)^2$.

We use the code RECFAST (Seager.Sasselov.&Scott1999:Seager.Sasselov2000) to caleulate the temperature evolution of the eas before the first stars ancl quasars turn on.

We use the code RECFAST \citep{sss99,sss00} to calculate the temperature evolution of the gas before the first stars and quasars turn on.

We note that al 2Z40 the spin temperature drops below the CMD temperature due to the collisional coupling (in Fie.

We note that at $z\gtrsim 40$ the spin temperature drops below the CMB temperature due to the collisional coupling (in Fig.

6. (his drop is seen at (he highest redshilt end).

6, this drop is seen at the highest redshift end).

The corresponding absorption is. however. difficult to observe at (he very long wavelengths corresponding to Chis redshilt.

The corresponding absorption is, however, difficult to observe at the very long wavelengths corresponding to this redshift.

At lower redshifts. the spin temperature is practically equal to the CMD temperature until the first photons are produced.

At lower redshifts, the spin temperature is practically equal to the CMB temperature until the first photons are produced.

If the comoving photon emissivity (defined as the number of photons emitted per unit

with results trom the literature.

with results from the literature.

We also derived firmer speeds for a smaller number of objects based on three or more epochs. bv combining our results wil other model fits in the literature.

We also derived firmer speeds for a smaller number of objects based on three or more epochs, by combining our results with other model fits in the literature.

The collected apparent speeds for these aud other WBLs from the literature are tabulated in Table 7.. where Col. (

The collected apparent speeds for these and other HBLs from the literature are tabulated in Table \ref{tab:speed}, where Col. (

3) gives the jet component speeds. aud (4) the number of epochs used for estimating these speeds.

3) gives the jet component speeds, and (4) the number of epochs used for estimating these speeds.

LDLs tvpically show superluminal motions with apparent speeds in the range 1—5c (7? )..

LBLs typically show superluminal motions with apparent speeds in the range $1-5c$ \citep{Gabuzda00}. .

In Fig.

21 we have plotted the apparent speeds of these. along with the LBL 0829-046 [rom the saanple. shaded in black. which exhibits component speeds of zz5e.

\ref{fig:speeds} we have plotted the apparent speeds of these, along with the LBL 0829+046 from the sample, shaded in black, which exhibits component speeds of $\approx5c$.

Figure also shows the distribution of apparent speeds for our HBL sample.

Figure \ref{fig:speeds} also shows the distribution of apparent speeds for our HBL sample.

We find that v4, 15 ἱνρισα]]ν less than 2 for the IIDLs.

We find that $\beta_{app}$ is typically less than 2 for the HBLs.

A two-sided IX.S test indicates that the IIDBL and LBL speeds are different at the signilicance level.

A two-sided K–S test indicates that the HBL and LBL speeds are different at the significance level.