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The AGB stars are a particularly prominent feature in he €CMDs of the LAIC and the SAIC. | The AGB stars are a particularly prominent feature in the CMDs of the LMC and the SMC. |
Following Nikolaev Weinberg (2000) the tail of AGB stars extending redware of J-Ix 521.4 is identified with carbon rich thermally pulsing. or TP-AGD stars. having presumed ages of 1.4 Cyr. | Following Nikolaev Weinberg (2000) the tail of AGB stars extending redward of $_S$$>$ 1.4 is identified with carbon rich thermally pulsing, or TP-AGB stars, having presumed ages of 1—4 Gyr. |
On CALDs in 2" fields centered on the LMC and the SAIC. we find that the brightest such AGB stars occur at Ixz10 in re LMC. and at Ixz10.5 in the SAIC. | On CMDs in $^o$ fields centered on the LMC and the SMC, we find that the brightest such AGB stars occur at $\approx$ 10 in the LMC, and at $\approx$ 10.5 in the SMC. |
Adopting the LMC SAIC distances given in Mateo (1998). 18.5 18.8 (or 49 58 Ixpc). respectively. the Ix-band absolute magnitudes of these C-rich TP-ACGD stars are about 8.5. | Adopting the LMC SMC distances given in Mateo (1998), 18.5 18.8 (or 49 58 Kpc), respectively, the K-band absolute magnitudes of these C-rich TP-AGB stars are about –8.5. |
We would have cleary seen such an AGB feature to within about half a magnitude of the detection limit. | We would have cleary seen such an AGB feature to within about half a magnitude of the detection limit. |
Since a prominent AGB is not seen to Ixz14. then the CLIVCs are either more cistant than m-M—22.5 or d=0.3 Mpe. or. if they are within about 0.3 Mpe. they do not contain a population of C-rich AGB stars similar to that seen in the LMC and the SAIC. | Since a prominent AGB is not seen to $\approx$ 14, then the CHVCs are either more distant than m-M=22.5 or d=0.3 Mpc, or, if they are within about 0.3 Mpc, they do not contain a population of C-rich TP-AGB stars similar to that seen in the LMC and the SMC. |
Another Milkv Way companion with a well-known population of bright. and intermediate-age AGB stars is the dSph Fornax. | Another Milky Way companion with a well-known population of bright, and intermediate-age AGB stars is the dSph Fornax. |
Saviane. Held AXortelli/ (2000) analvzed. it most recently. and emphasised. the presence. of bright AGB stars. which they attribute to an intermeciate-age (210 Gyr) stellar population. | Saviane, Held Bertelli (2000) analyzed it most recently, and emphasised the presence of bright AGB stars, which they attribute to an intermediate-age (2–10 Gyr) stellar population. |
Fornax is at quite a high Galactic latitudo. of almost -66". and sullers less from Galactic foreground contamination than do our CIIVCs. the LAIC ancl the SAIC. | Fornax is at quite a high Galactic latitude, of almost $^o$, and suffers less from Galactic foreground contamination than do our CHVCs, the LMC and the SMC. |
Fornax has a distance modulus of 20.70. or a distance of 138 Ίνρο. a V-baned central surface. brightness of 23.4 magnitude 27. and an exponential scale length of 10/22 (Mateo 1998). | Fornax has a distance modulus of 20.70, or a distance of 138 Kpc, a V-band central surface brightness of 23.4 magnitude $^{-2}$, and an exponential scale length of 2 (Mateo 1998). |
The 2ALASS data of Fornax. shown in Figure 9. exhibit a population of AGB stars with J-Ixz«1. ancl possibly of up to 2. near Ixz13.5. | The 2MASS data of Fornax, shown in Figure 9, exhibit a population of AGB stars with $>$ 1, and possibly of up to 2, near $\approx$ 13.5. |
Phe C-rich. TP-AGB stars with colours of J-WKs> 14 identified in the MCSs are very sparse in Fornax. | The C-rich TP-AGB stars with colours of $_S$$>$ 1.4 identified in the MCs are very sparse in Fornax. |
Fhese stars extend to the red from a more strongly populated AGB feature seen between 1.0«J-Ix2« 1.4. | These stars extend to the red from a more strongly populated AGB feature seen between $<$ $_S$$<$ 1.4. |
The magnitudes of these AGD stars. most probably the voung AGB component of Fornax. are near Ixz213.5. | The magnitudes of these AGB stars, most probably the young AGB component of Fornax, are near $\approx$ 13.5. |
This implies absolute Ix. magnitudes of zz 7.2. | This implies absolute K magnitudes of $\approx$ –7.2. |
Since similar ACD. stars are not seen to Welt in the CLIVCS. then the CLIVCSs are either more distant than m-M-21.2 or d=0.17 Alpe. or. if hey are within 0.17 Alpe. they do not contain a population of AGB stars similar to that seen in Fornax. | Since similar AGB stars are not seen to $\approx$ 14 in the CHVCs, then the CHVCs are either more distant than m-M=21.2 or d=0.17 Mpc, or, if they are within 0.17 Mpc, they do not contain a population of AGB stars similar to that seen in Fornax. |
We also culled data for the dSph galaxies. Sextans (bzz42"). Sculptor (he 83°) and. Ursa Minor (bzz45") from. he 2ALASS database. | We also culled data for the dSph galaxies Sextans $\approx$ $^o$ ), Sculptor $\approx$ $^o$ ) and Ursa Minor $\approx$ $^o$ ) from the 2MASS database. |
Data for Carina and Draco are not available in the second incremental clata release. | Data for Carina and Draco are not available in the second incremental data release. |
All our dSph observed. have exponential scale lengths of the order of 10. aremin. | All four dSph observed have exponential scale lengths of the order of 10 arcmin. |
Fornax and Seulptor have rather ugh V-hand central surface. brightnesses. of 23.4 and 23.7 mag acrsec7 (see Mateo 1998). and they are easily seen as concentrations in the position plots shown in Fig. | Fornax and Sculptor have rather high V-band central surface brightnesses, of 23.4 and 23.7 mag $^{-2}$ (see Mateo 1998), and they are easily seen as concentrations in the position plots shown in Fig. |
10. | 10. |
The surface brightnesses Ursa. Minor ancl Sextans are only 25.5 and 26.2 magnitude: respectively. | The surface brightnesses Ursa Minor and Sextans are only 25.5 and 26.2 magnitude $^{-2}$ respectively. |
They. are not distinguishable on position plots against their respective Galactic foreground contaminations. although their L-band luminosity functions. exhibit a feature where we would expect to find their PROBs (compare Schulte-Ladbeck. et al. | They are not distinguishable on position plots against their respective Galactic foreground contaminations, although their H-band luminosity functions exhibit a feature where we would expect to find their TRGBs (compare Schulte-Ladbeck et al. |
1999). | 1999). |
We performed KS tests using the Wk-band ROB LE of Sculptor (whose distance is about SO Ixpe. see Mateo 1998). | We performed KS tests using the K-band RGB LF of Sculptor (whose distance is about 80 Kpc, see Mateo 1998). |
We first selected stars by position to occur near the centre of the field. of view. | We first selected stars by position to occur near the centre of the field of view. |
This emphasises true. member over Galactic foreground stars. | This emphasises true member over Galactic foreground stars. |
We then constrained the colour and luminosity range for the RGB to be 0.7«J-Ix «« 1.4. and ]x«215. | We then constrained the colour and luminosity range for the RGB to be $<$ $_S$$<$ 1.4, and $_S$$>$ 15. |
The Galactic foreground. of the Πλος is much stronger than that of Sculptor. since the former are situated at about 50° lower Galactic latitude. | The Galactic foreground of the CHVCs is much stronger than that of Sculptor, since the former are situated at about $^o$ lower Galactic latitude. |
We needed: το add multiples of the Sculptor RGB stars to the J-Ixs. Ks} CMDs of the οΗΧΟΣ before we could discriminate the additional ROB stars against the strong Galactic foreground of the CIIVC fields. | We needed to add multiples of the Sculptor RGB stars to the $_S$, $_S$ ] CMDs of the CHVCs before we could discriminate the additional RGB stars against the strong Galactic foreground of the CHVC fields. |
The data start to suggest the presence of a Sculptor-like RGB when four times the Sculptor stars are added to the data. | The data start to suggest the presence of a Sculptor-like RGB when four times the Sculptor stars are added to the data. |
Significant dillerences with probabilities of less than for a chance event occur when about eight times the amount of Sculptor ROB stars is added to the data. | Significant differences with probabilities of less than for a chance event occur when about eight times the amount of Sculptor RGB stars is added to the data. |
A [actor of eight in [Lux corresponds to about 2.3 mag. | A factor of eight in flux corresponds to about 2.3 mag. |
Therefore. in order to detect a Seulptor-like RGB population in our CVLIC fields. it would have to have a V-band surface brightness of about 21-4 magnitude 7. | Therefore, in order to detect a Sculptor-like RGB population in our CVHC fields, it would have to have a V-band surface brightness of about 21.4 magnitude $^{-2}$. |
This is much higher than exhibited by any. of the Sph or dlrr/dSsph companions of the Milky. Way. | This is much higher than exhibited by any of the dSph or dIrr/dSph companions of the Milky Way. |
We note that M 31. however. has such high-surface-brightness companions. e.g.. NGC 147. | We note that M 31, however, has such high-surface-brightness companions, e.g., NGC 147. |
The 2NLASS data allow us to set additional limits on the absence of a dwarl-galaxy like stellar population for four CLIVCs. | The 2MASS data allow us to set additional limits on the absence of a dwarf-galaxy like stellar population for four CHVCs. |
of interest. | of interest. |
As all these tables are similar. we give a single description for all here. | As all these tables are similar, we give a single description for all here. |
Col. ( | Col. ( |
1) labels the region according to region selection shown in the accompanied figure. column (2) lists the temperature in keV. Derived quantities. that use an estimate of the projected length. as described below are reported in cols. ( | 1) labels the region according to region selection shown in the accompanied figure, column (2) lists the temperature in keV. Derived quantities, that use an estimate of the projected length, as described below are reported in cols. ( |
3-6). | 3–6). |
These are electron density. entropy. pressure. and the (local) gas mass. | These are electron density, entropy, pressure, and the (local) gas mass. |
Cols. | Cols. |
7-8 report the minimal (min) and maximal 675,40 radit of the extraction area. col. ( | 7–8 report the minimal $r_{min}$ ) and maximal $r_{max}$ ) radii of the extraction area, col. ( |
9) provides remarks on the region. | 9) provides remarks on the region. |
For this detailed analysis we also perform an estimate of the projection length of each analyzed region to obtain actual gas properties at these locations. as described at length in Henry et al. ( | For this detailed analysis we also perform an estimate of the projection length of each analyzed region to obtain actual gas properties at these locations, as described at length in Henry et al. ( |
2004) and Mahdavi et al. ( | 2004) and Mahdavi et al. ( |
2005). | 2005). |
To avoid the importance of the projection effects. we discard the regions having a ratio of the minimal to the maximal radii of values exceeding 0.8. | To avoid the importance of the projection effects, we discard the regions having a ratio of the minimal to the maximal radii of values exceeding 0.8. |
The statistics achieved in the observation of the REFLEX-DXL clusters allows us only to recognize the strongest fluctuations in either temperature. entropy. or pressure. | The statistics achieved in the observation of the REFLEX-DXL clusters allows us only to recognize the strongest fluctuations in either temperature, entropy, or pressure. |
In the presentation of the results we indicate the features seen in the hardness ratio based maps and discuss how much it 1s possible to confirm them through a direct spectroscopic analysis. | In the presentation of the results we indicate the features seen in the hardness ratio based maps and discuss how much it is possible to confirm them through a direct spectroscopic analysis. |
In selecting the regions according to their properties or according to their statistics. we implicitly perform an adaptive filtering of the signal. | In selecting the regions according to their properties or according to their statistics, we implicitly perform an adaptive filtering of the signal. |
It is therefore important. to characterize the spatial frequencies sampled in the analysis. which is also a way to characterize the analysis carried out and the cluster spatial scales sampled. | It is therefore important to characterize the spatial frequencies sampled in the analysis, which is also a way to characterize the analysis carried out and the cluster spatial scales sampled. |
In Fig.l we present such an analysis. where it becomes clear that the choice of the regions corresponds to a grid in cylindrical coordinates. sampling the azimuthal angle with typically 3 sectors on radial scales from 0.1 to 1 Mpe. | In \ref{f:tf} we present such an analysis, where it becomes clear that the choice of the regions corresponds to a grid in cylindrical coordinates, sampling the azimuthal angle with typically 3 sectors on radial scales from 0.1 to 1 Mpc. |
Since there is a discussion on the reliability of the temperature determination. we have provided a plot for each cluster. where à comparison to the average cluster temperature profile of Vikhlinin et al. ( | Since there is a discussion on the reliability of the temperature determination, we have provided a plot for each cluster, where a comparison to the average cluster temperature profile of Vikhlinin et al. ( |
2005) is presented. | 2005) is presented. |
One can see in each case that there is à good agreement in the results. | One can see in each case that there is a good agreement in the results. |
We also point out that a few clusters where we probed the region outside rsoy. reveal strong asymmetries indicating aceretion from a filament. | We also point out that a few clusters where we probed the region outside $r_{500}$, reveal strong asymmetries indicating accretion from a filament. |
It may be that the presence or removal of such zones could be the underlying reason for some of the reported disagreement in temperature profiles. | It may be that the presence or removal of such zones could be the underlying reason for some of the reported disagreement in temperature profiles. |
We define a cool core of a cluster or its debris as the gas with entropy significantly below 200—300 keV enr. which according to Voit Bryan (2001) could cool in a Hubble time. | We define a cool core of a cluster or its debris as the gas with entropy significantly below 200–300 keV $^2$, which according to Voit Bryan (2001) could cool in a Hubble time. |
As local examples of cool cores have entropies lower than 100 keV enr. we have used the later criterion for the cooling core identification. | As local examples of cool cores have entropies lower than 100 keV $^2$, we have used the later criterion for the cooling core identification. |
Before proceeding with the description of individual systems. we summarize the results by presenting the fits to the entropy and pressure profiles. | Before proceeding with the description of individual systems, we summarize the results by presenting the fits to the entropy and pressure profiles. |
To define the shape of the entropy and pressure profiles. we applied the non-parametric locally weighted regression. following Sanderson et al. ( | To define the shape of the entropy and pressure profiles, we applied the non-parametric locally weighted regression, following Sanderson et al. ( |
2005 and references therein). | 2005 and references therein). |
This analysis results in. the non-parametric curve. which we approximate below with power laws. | This analysis results in the non-parametric curve, which we approximate below with power laws. |
Our analysis illustrated in. Fig.2 suggests a broken power law approximation to the entropy profile with an inner and an outer slopes of 0.78 and 0.52. respectively. and a break at 0.500. | Our analysis illustrated in \ref{f:comp} suggests a broken power law approximation to the entropy profile with an inner and an outer slopes of 0.78 and 0.52, respectively, and a break at $0.5r_{500}$. |
The amplitude of fluctuations around the best exceeds the effect of the statistics and is a measure of the important of substructure. as discussed below. | The amplitude of fluctuations around the best exceeds the effect of the statistics and is a measure of the important of substructure, as discussed below. |
The average level of fluctuations. which is (20%)) in. the case of entropy (pressure). could be taken as the accuracy to which the approximation to the entropy distribution could be determined. | The average level of fluctuations, which is ) in the case of entropy (pressure), could be taken as the accuracy to which the approximation to the entropy distribution could be determined. |
The entropy profiles with exclusion of the substructure has been analyzed in Zhang et al. ( | The entropy profiles with exclusion of the substructure has been analyzed in Zhang et al. ( |
2005). yielding a steeper index of 0.95. | 2005), yielding a steeper index of 0.95. |
The characterization of the pressure profile is more complex. requiring three power laws. with slopes —0.64 at FOXθα. —2.47 at r>O.Srsoq and a slope of —1.50 in between. | The characterization of the pressure profile is more complex, requiring three power laws, with slopes $-0.64$ at $r<0.3r_{500}$, $-2.47$ at $r>0.5r_{500}$ and a slope of $-1.50$ in between. |
For completeness. we present in Tab.2 à standard approach of using the orthogonal regression and assuming a power law shape to approximate the shape of the entropy and pressure profiles. | For completeness, we present in \ref{t:cl-sp-all} a standard approach of using the orthogonal regression and assuming a power law shape to approximate the shape of the entropy and pressure profiles. |
We present the results obtained using different masks and also combine the clusters with and without the rescaling. described above. | We present the results obtained using different masks and also combine the clusters with and without the rescaling, described above. |
As could be easily seen from Tab.3.. this approach results in a much larger amplitude of the residuals. even taking into account that the power law shape was fitted separately for each cluster. while in the non-parametric approach. one shape is used to approximate all clusters. | As could be easily seen from \ref{t:sc0}, this approach results in a much larger amplitude of the residuals, even taking into account that the power law shape was fitted separately for each cluster, while in the non-parametric approach, one shape is used to approximate all clusters. |
One of the most important results is an observation of a flattening in the entropy profile at outer radii in DXL clusters. changing from 0.78 within the 0.500 to 0.54 outside. | One of the most important results is an observation of a flattening in the entropy profile at outer radii in DXL clusters, changing from 0.78 within the $0.5 r_{500}$ to 0.54 outside. |
As the sample consists of the most massive clusters in the Universe. we believe that the explanation of the observed trend should be searched in the details of the accretion. | As the sample consists of the most massive clusters in the Universe, we believe that the explanation of the observed trend should be searched in the details of the accretion. |
As summarized in Voit (2004). the index of the entropy profile is driven by the effects of mass growth as well as evolution of the virial density. | As summarized in Voit (2004), the index of the entropy profile is driven by the effects of mass growth as well as evolution of the virial density. |
Under the assumption of a smooth accretion. the entropy grows with radius às $~Mi* (Voit 2004). | Under the assumption of a smooth accretion, the entropy grows with radius as $S \sim
M_{gas}^{1-4/3}$ (Voit 2004). |
With a canonical cluster characteristic of the surface brightness profile. 8=2/3. Meyor. where M, is enclosed gas mass. | With a canonical cluster characteristic of the surface brightness profile, $\beta=2/3$, $M_{\rm gas} \sim r$, where $M_{\rm
gas}$ is enclosed gas mass. |
However. as gas mass fraction tends to level off at high radii. the cumulative gas mass starts to follow the mass of the dark matter and so Mas77. | However, as gas mass fraction tends to level off at high radii, the cumulative gas mass starts to follow the mass of the dark matter and so $M_{\rm gas} \sim r^{0.5}$. |
A similar flattening in the entropy distribution is then expected and is observed in our data at 7>O.4rsoo. | A similar flattening in the entropy distribution is then expected and is observed in our data at $r\ge
0.4 r_{500}$. |
If this indeed is the explanation to our data. one would not necessarily expect the same trend to be observed in low-mass clusters. where the baryor fraction is growing with the radius even at FS500- | If this indeed is the explanation to our data, one would not necessarily expect the same trend to be observed in low-mass clusters, where the baryon fraction is growing with the radius even at $r_{500}$ . |
a double peaked structure. | a double peaked structure. |
The observed time variation of ρω is shown in Figure 1. | The observed time variation of $L_{\rm bol}$ is shown in Figure 1. |
It is seen (hat the bolometric Iuminosity actually peaks ching the main radio observing period. | It is seen that the bolometric luminosity actually peaks during the main radio observing period. |
Although the observations span a rather limited range in time and. furthermore. awe likely to be affected by interstellar scattering and seintillation (BNC). it is noteworthy that no pronounced minim is apparent in (he light curves for the optically thin frequencies. | Although the observations span a rather limited range in time and, furthermore, are likely to be affected by interstellar scattering and scintillation (BKC), it is noteworthy that no pronounced minimum is apparent in the light curves for the optically thin frequencies. |
This can be understood if [ωμρ is roughly equal to /. | This can be understood if $t_{\rm comp}$ is roughly equal to $t$. |
Belore making a more detailed numerical fit to the data. we will show that if Aic/ we can obtain a consistent and plausible set of values [or vj. ἐν ancl εκ. | Before making a more detailed numerical fit to the data, we will show that if $t_{\rm comp}\sim t$ we can obtain a consistent and plausible set of values for $v_{\rm sh}$, $\epsilon_{\rm B}$ and $\epsilon_{\rm rel}$. |
The Gime scale for Compton cooling is given by With the use of (he condition fegmpc/ results in where Lii»=Lü4/107ergs!. | The time scale for Compton cooling is given by With the use of the condition $t_{\rm comp}\sim t$ results in where $L_{\rm bol,42}\equiv L_{\rm bol}/10^{42} \ergs$. |
Furthermore. rjj=1 and /=6 days have been used. | Furthermore, $\nu_{10} = 1$ and $t = 6$ days have been used. |
As an illustration we show in Figure 1. the ratio of the Compton cooling time scale to the adiabatic time scale for the specilic model in 39. with M=10°M,wrhn f. va(/4=I0)70.000kms!. and ey=10 7. | As an illustration we show in Figure \ref{fig1} the ratio of the Compton cooling time scale to the adiabatic time scale for the specific model in \ref{sec_3} with $\Mdot=10^{-5} \ml$, $v_{\rm w}=1000 \kms$ , $v_{\rm sh}(t_{\rm d}=10)=70,000 \kms$, and $\epsilon_{\rm B}=10^{-3}$ . |
For other values a?Liao1/2 | For other values $ t_{\rm comp}
\propto (\epsilon_{\rm B} \Mdot/v_{\rm w})^{1/4} v_{\rm sh}^2 t^{3/2}
L_{\rm bol}^{-1} \nu_{10}^{-1/2}$ . |
An X-ray [lux was detected [rom SN 2002ap by XMM-Newton on 2002. 33 (1226 davs) (Soria&Itong2002:Sutaria.Chandra.Bhatnagar. 2004). | An X-ray flux was detected from SN 2002ap by XMM-Newton on 2002, 3 $t\approx 6$ days) \citep{SK02,SCB03,SPM03}. |
. Due to the weak signal neither the total flux nor the spectral shape could be well determined. | Due to the weak signal neither the total flux nor the spectral shape could be well determined. |
The uncertainty in the high energy. [αν is affected by Che subtraction of the Εαν from a strong. nearby source with a hard spectrum. while at low energy. the inability to constrain absorption im excess of the galactic value makes the observed [ας a lower Himil to the intrinsic one. | The uncertainty in the high energy flux is affected by the subtraction of the flux from a strong, nearby source with a hard spectrum, while at low energy the inability to constrain absorption in excess of the galactic value makes the observed flux a lower limit to the intrinsic one. |
The observations can be fitted either with a (thermal or a power law spectrum. | The observations can be fitted either with a thermal or a power law spectrum. |
In the latter case. the deduced spectral index is consistent with that observed in (he racio. | In the latter case, the deduced spectral index is consistent with that observed in the radio. |
It is therefore likely that the X-ray flux is the Compton scattered optical radiation from the supernova. | It is therefore likely that the X-ray flux is the Compton scattered optical radiation from the supernova. |
If this is thecase. independent estimates of the energv densities in | If this is thecase, independent estimates of the energy densities in |
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