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In agreement with the hardness ratio map (Fig. 5)). | In agreement with the hardness ratio map (Fig. \ref{cold.fig}) ), |
our spectral analysis therefore supports the presence of multiphase gas along the filaments. | our spectral analysis therefore supports the presence of multiphase gas along the filaments. |
Interestingly. such cool filaments follow the morphology of the powerful central radio source nicely. although the western part of the southern radio lobe appears devoid of cool gas. | Interestingly, such cool filaments follow the morphology of the powerful central radio source nicely, although the western part of the southern radio lobe appears devoid of cool gas. |
The SSW sector is indeed the hottest and is also the region where the presence of a second thermal component is least significant. | The SSW sector is indeed the hottest and is also the region where the presence of a second thermal component is least significant. |
This sector les at the location where the southern radio lobe appears to fold back on itself (Lane et al. | This sector lies at the location where the southern radio lobe appears to fold back on itself (Lane et al. |
2004). | 2004). |
The properties of the cavity are also consistent with a sharp bend in the southern jet there (Wise et al. | The properties of the cavity are also consistent with a sharp bend in the southern jet there (Wise et al. |
2007). | 2007). |
Finally. we also attempted to map the emission measure distribution by fitting more complicated multi-phase spectral models (such as with fixedtemperatures and free normalizations). | Finally, we also attempted to map the emission measure distribution by fitting more complicated multi-phase spectral models (such as , with fixedtemperatures and free normalizations). |
However. wefound our data are inadequate due to limited statistics. and the spectral resolution of Chandra. | However, wefound our data are inadequate due to limited statistics and the spectral resolution of . |
These detailed spectral studies will. hopefully. be possible in the future with the spectral capabilities of the (IXO). | These detailed spectral studies will, hopefully, be possible in the future with the spectral capabilities of the . |
Before concluding that the simulations are correct and the model erroneous in this case, however, we must add a final caution. | Before concluding that the simulations are correct and the model erroneous in this case, however, we must add a final caution. |
The timescale we compute is inversely proportional to the subgrid clumping factor assumed in the simulations. | The timescale we compute is inversely proportional to the subgrid clumping factor assumed in the simulations. |
While this number has been estimated for metallicities down to Z'z0.2 in the Small Magellanic Cloud, we have no direct knowledge of its value in lower metallicity galaxies. | While this number has been estimated for metallicities down to $Z'\approx 0.2$ in the Small Magellanic Cloud, we have no direct knowledge of its value in lower metallicity galaxies. |
Nonetheless, we tentatively conclude that the time-dependence of the Hg fraction is not a major effect at metallicities Z’> 1073. so a independent model is adequate in thisregime. | Nonetheless, we tentatively conclude that the time-dependence of the $_2$ fraction is not a major effect at metallicities $Z' \ga 10^{-2}$ , so a time-independent model is adequate in thisregime. |
around the Na I D line. which is dominated by high- ancl low-pressure sodium lamps. | around the Na I D line, which is dominated by high- and low-pressure sodium lamps. |
The (lux in this region is about 1 to 3x10P eres ? | 1! for the Tucson ancl Phoenix directions. and smaller for the other directions. | The flux in this region is about 1 to $\times 10^{-15}$ ergs $^{-2}$ $^{-1}$ $^{-1}$ for the Tucson and Phoenix directions, and smaller for the other directions. |
The current data show no more flux than the 19838 or 1999 data. and. if anything. is somewhat less. | The current data show no more flux than the 1988 or 1999 data, and, if anything, is somewhat less. |
However. the variations within a single night of observations towards Tucson are more significant than the changes over the past two decades. | However, the variations within a single night of observations towards Tucson are more significant than the changes over the past two decades. |
Overall. the sky brightness at Kitt Peak has remained remarkably constant over the past 20 vears. | Overall, the sky brightness at Kitt Peak has remained remarkably constant over the past 20 years. |
While the magnitudes of many artificial sources have increased. (his should have little effect on astronomical spectroscopy unless the observer is particularly interested in the Na I D lines. | While the magnitudes of many artificial sources have increased, this should have little effect on astronomical spectroscopy unless the observer is particularly interested in the Na I D lines. |
Other aspects besides artificial sources. such as (he solar cvele. also contribute to the sky brightness (Pilachowski οἱ 11989). | Other aspects besides artificial sources, such as the solar cycle, also contribute to the sky brightness (Pilachowski et 1989). |
As discussed in both Massey. et ((1990) and Alassev Foltz (2000). previous solar maximnuns occurred in July 1989 and March. 2000. corresponding almost perfectly with the first two studies. | As discussed in both Massey et (1990) and Massey Foltz (2000), previous solar maximums occurred in July 1989 and March 2000, corresponding almost perfectly with the first two studies. |
Even though these investigations have been completed almost one solar evcle (11 vears) later. Fieure G shows that the solar flix changed quite drastically between our observations. | Even though these investigations have been completed almost one solar cycle (11 years) later, Figure \ref{fig:sunPlot} shows that the solar flux changed quite drastically between our observations. |
While (he 1999 data was taken al a period close to the solar maximum. the 2009/10 data was taken al a clear solar minimum while the 1988 dala was taken at an intermediate stage. | While the 1999 data was taken at a period close to the solar maximum, the 2009/10 data was taken at a clear solar minimum while the 1988 data was taken at an intermediate stage. |
This is a likely explanation for why the 1999 data is significantly brighter than the 2009/10 data. | This is a likely explanation for why the 1999 data is significantly brighter than the 2009/10 data. |
The solar evcle should affect the atmospheric lines (such as NI A5199. OI A5577 and OI A6300-64) as well as portions οἱ the blue and ved "pseudo-continuum. | The solar cycle should affect the atmospheric lines (such as NI $\lambda$ 5199, OI $\lambda$ 5577 and OI $\lambda$ 6300-64) as well as portions of the blue and red “pseudo-continuum." |
However. the Nal or Hel lisht pollution lines should not be affected. | However, the NaI or HgI light pollution lines should not be affected. |
In an attempt (o correct for the effects of the solar cvele on our data. we sealed our V and 2 magnitudes to the solar minimum using the knowledge that the skv gets brighter by c-0.4 magnitudes between the solar minimum and the solar maximum (Benn Ellison 1993. Peclani 2009). and (he “solar phases” (based upon the solar flix) when our observations were taken. | In an attempt to correct for the effects of the solar cycle on our data, we scaled our $V$ and $B$ magnitudes to the solar minimum using the knowledge that the sky gets brighter by $\sim$ 0.4 magnitudes between the solar minimum and the solar maximum (Benn Ellison 1998, Pedani 2009), and the “solar phases" (based upon the solar flux) when our observations were taken. |
These results are shown in Table 3. and suggest that while the observed sky brightness is comparable to what it was in 1999 (especially considering our 0.04 magnitude error). the skv has brightened at Zenith by approximately 0.1 magnitude since 1988. | These results are shown in Table \ref{tab:sFlux} and suggest that while the observed sky brightness is comparable to what it was in 1999 (especially considering our 0.04 magnitude error), the sky has brightened at Zenith by approximately 0.1 magnitude since 1988. |
Similarly. (hese corrected data suggest that while the skv towards both Tucson and “nowhere” has increased in brightness by approximately 0.3 magnitudes in V. and 2 since 1988. thev are actually ζω)... magnitude darker than thev were in 1999. | Similarly, these corrected data suggest that while the sky towards both Tucson and “nowhere" has increased in brightness by approximately 0.3 magnitudes in $V$ and $B$ since 1988, they are actually $\sim$ 0.1 magnitude darker than they were in 1999. |
Once needs to recall that these corrections are approximate. but thev help place (he measurements on a uniform (if hypothetical) basis. | Once needs to recall that these corrections are approximate, but they help place the measurements on a uniform (if hypothetical) basis. |
CE(B—V)«0.01:Chandaretal.2004) so we do not correct for it. | \citep[$E(B-V)<0.01$;][]{cwl04} so we do not correct for it. |
We can compare the results of our selection with the study of BresolinDINS).. | We can compare the results of our selection with the study of \citet[hereafter BKS]{bks96}. |
These authors performed a visual search for clusters in one IIST/WEDPC? field which is contained within our central ACS Their 43 cluster candidates have a median V. magnitude of 21.5. | These authors performed a visual search for clusters in one HST/WFPC2 field which is contained within our central ACS Their 43 cluster candidates have a median $V$ magnitude of 21.5. |
Two of their cancliclates fall in the gap between the ACS detectors. and two are not identifiable on the WEDPC? image. | Two of their candidates fall in the gap between the ACS detectors, and two are not identifiable on the WFPC2 image. |
In (he same area searched bv BINS. our cluster selection procedure picked 232 candidates with a median magnitude of V.=22.8. | In the same area searched by BKS, our cluster selection procedure picked 232 candidates with a median magnitude of $V=22.8$. |
Twenty-two of the remaining 39 BIS candidates (56%)) are in our candidate list. | Twenty-two of the remaining 39 BKS candidates ) are in our candidate list. |
Visual examination of the 17 BINS clusters not in our list shows them to have a range of morphologies. but thev are generally less round. and less isolated than the elusters on our list (see also Figure 2)). | Visual examination of the 17 BKS clusters not in our list shows them to have a range of morphologies, but they are generally less round and less isolated than the clusters on our list (see also Figure \ref{m101-detail}) ). |
As indicated by the median magnitudes. our candidates are eenerally Tainter than those of BINS. | As indicated by the median magnitudes, our candidates are generally fainter than those of BKS. |
Comparison of photometry for objects in common shows a small offset: median dillerences between their photometry and ours are AV=04620.04. ACD—V)=-0.1440.03. and AG—7)=40.08& 0.03. | Comparison of photometry for objects in common shows a small offset: median differences between their photometry and ours are $\Delta V = 0.16\pm 0.04$, $\Delta(B-V) = -0.14\pm 0.03$, and $\Delta(V-I) = +0.08\pm 0.03$ . |
BINS state that their magnitudes are uncertain by >0.1 mag. so we do not believe the offset indicates a serious problem in our photometry. | BKS state that their magnitudes are uncertain by $\geq 0.1$ mag, so we do not believe the offset indicates a serious problem in our photometry. |
SG cluster svstem studies have generally focused on (wo separate populations: the old. elobular clusters or the voung. massive clusters. | Star cluster system studies have generally focused on two separate populations: the old, globular clusters or the young, massive clusters. |
Elliptical galaxies. of course. have only the first (wpe of cluster. while the clusters studied in galaxy mergers exemplilv the second. | Elliptical galaxies, of course, have only the first type of cluster, while the clusters studied in galaxy mergers exemplify the second. |
The color distribution of the MIOLI cluster candidates is expected to rellect their age distribution: while colors of globular clusters are often used as indicators of metallicity 1998)... many of the MIOL candidates are [ar Coo blue to be old. metal-poor clusters. | The color distribution of the M101 cluster candidates is expected to reflect their age distribution: while colors of globular clusters are often used as indicators of metallicity \citep[e.g.][]{kw98}, many of the M101 candidates are far too blue to be old, metal-poor clusters. |
In the following analvsis. we will use color as a simple observational distinction between older and younger clusters. (o ease comparison wilh other star cluster studies. | In the following analysis, we will use color as a simple observational distinction between older and younger clusters, to ease comparison with other star cluster studies. |
5—V=0.45 is the red limit used by Larsen&Richtler(1999). for their studies of “Young Massive Clusters’. and is also close to the blue limit of the Milkv. Way globular cluster (2—V), distribution. so we adopt (2—V)y=0.45 as the dividing line between old and young clusters. | $B-V=0.45$ is the red limit used by \citet{lr99} for their studies of `Young Massive Clusters', and is also close to the blue limit of the Milky Way globular cluster $(B-V)_0$ distribution, so we adopt $(B-V)_0=0.45$ as the dividing line between old and young clusters. |
[It should be remembered. however. that theage distribution of star clusters in MIOL is not necessarily | It should be remembered, however, that theage distribution of star clusters in M101 is not necessarily |
iree. degrees of freedom (for the flux aud position of the ποσο), | three degrees of freedom (for the flux and position of the source). |
Tn performing the fits. we use the analytical result vat the best fit has a uuuber of model plotous equal o the ummber of detected yhotons. | In performing the fits, we use the analytical result that the best fit has a number of model photons equal to the number of detected photons. |
Thus. in the model of a constant backeround. an constant value of 1, 1e optimum value of i; is found directly by dividiug the observed ΠΙΟ of photous by the uuuber of pixels under consideration. | Thus, in the model of a constant background, a constant value of $m_i$, the optimum value of $m_i$ is found directly by dividing the observed number of photons by the number of pixels under consideration. |
Iu fitting a conustaut plus oue source. we distribute he source counts around the source position according o the ROSAT TRI point PAspread function at the center of the detector (David et 11995). | In fitting a constant plus one source, we distribute the source counts around the source position according to the ROSAT HRI point spread function at the center of the detector (David et 1995). |
We then vary the vackerouud and source counts. aud the source position. to uiuimuize ln£. | We then vary the background and source counts, and the source position, to minimize $\ln L$. |
Iu doing so we keep the sumi of the source and backerouud counts fixed at the observed umber. | In doing so we keep the sum of the source and background counts fixed at the observed number. |
Next. we fit a coustaut backeround plus two or three »oimt sources. | Next, we fit a constant background plus two or three point sources. |
Iu this fit. the parameters of tle first source are allowed to vary: we use a Downhill Simplex method as miplemiented by Press et ((1992) to minimize lu£L witli respec to the 7 or |0 variables. | In this fit, the parameters of the first source are allowed to vary; we use a Downhill Simplex method as implemented by Press et (1992) to minimize $\ln L$ with respect to the 7 or 10 variables. |
Here again. the sua of the model counts for the backerouud aud the two/three sources is kept constant. at the observed τν, | Here again, the sum of the model counts for the background and the two/three sources is kept constant, at the observed number. |
The results of this fitting procedure are sunumnarized iu roftabiul: 1e resulting source positions are shown iu roffüenml aud listed in Table 3.. | The results of this fitting procedure are summarized in \\ref{tabml}; the resulting source positions are shown in \\ref{figml} and listed in Table \ref{tabpos}. |
A third source in the ceuter is nominally significant at the 3-0 level: however. our analvsis doesu't take iuto account auy remaiiue jitter iu the Point Spread Function. aud we consider the existence of this source not proven. | A third source in the center is nominally significant at the $\sigma$ level; however, our analysis doesn't take into account any remaining jitter in the Point Spread Function, and we consider the existence of this source not proven. |
Therefore. we do not list the source in roftabpos:: its position is shown in rofüeual.. | Therefore, we do not list the source in \\ref{tabpos}; its position is shown in \\ref{figml}. |
Optical images while the N-vav source was still on were taken for us in Service Time in the night of 1998 Aueust 26 to 27 at the 3.5au New Techologv Telescope (NTT) at La Silla. using the Superb Seeing Imager SUSI2. aud with the San Unit Telescope #11 (Autu) of the Verv Laree Telescope at Paranal. using the VLT Test Camera. | Optical images while the X-ray source was still on were taken for us in Service Time in the night of 1998 August 26 to 27 at the 3.5-m New Techology Telescope (NTT) at La Silla, using the Superb Seeing Imager SUSI2, and with the 8-m Unit Telescope 1 (Antu) of the Very Large Telescope at Paranal, using the VLT Test Camera. |
We will only discuss the NTT observations here. as these iad better secius. | We will only discuss the NTT observations here, as these had better seeing. |
With SUSI2. one 10-s aud two LO0-s exposures were taken through a Bessell Ro filter. as well as two 900-8 exposures through a Dessell D filter. | With SUSI2, one 10-s and two 100-s exposures were taken through a Bessell R filter, as well as two 900-s exposures through a Bessell B filter. |
During he observations. the secing varied between 077 from the frst R-band tage to 172 in last B-baud nuage. | During the observations, the seeing varied between $0\farcs7$ from the first R-band image to $1\farcs2$ in last B-band image. |
The üeht was not photometric. | The night was not photometric. |
The detector was a mosaic of two EEV CCDs. cach composed of 20181006 square xxels of 15 un ou the side. | The detector was a mosaic of two EEV CCDs, each composed of $2048\times4096$ square pixels of $15\,\mu$ m on the side. |
They were road out biuned jr 2 in cach direction. as the plate scale of OVOSpix! would substantially oversaunple the secing. | They were read out binned by 2 in each direction, as the plate scale of $0\farcs08{\rm\,pix^{-1}}$ would substantially oversample the seeing. |
For all but he first. l0-« R-band image. the telescope was offset such that the core of the cluster was not too close to he eap between the two CCDs. | For all but the first, 10-s R-band image, the telescope was offset such that the core of the cluster was not too close to the gap between the two CCDs. |
The data reduction was done using standard procedures. determining the jas from the overscan regious (after verification on yas frames) and correcting for pixelto-pixel scusitivity variations using dome flats taken in the morning following he observations. | The data reduction was done using standard procedures, determining the bias from the overscan regions (after verification on bias frames) and correcting for pixel-to-pixel sensitivity variations using dome flats taken in the morning following the observations. |
On 1999 July 15. when the XN-rav source was off. we took images with the Focal Reducer/Low Dispersion Spectrograph FORSI on Autu through Besscll R. D. and U filters. | On 1999 July 15, when the X-ray source was off, we took images with the Focal Reducer/Low Dispersion Spectrograph FORS1 on Antu through Bessell R, B, and U filters. |
The might was no photometric. aud the seeing varied from 078 to 172. | The night was not photometric, and the seeing varied from $0\farcs8$ to $1\farcs2$. |
Two 10-8 and one 100-5 exposures were taken in R. two 30-5 and one 300-8 in D. aud oue and oue 600-3 in U. The detector was a Telstronix CCD with 2018« pixels of 21412. The standard resolution collimator was used. for whic ithe plate scale is (72pix1 | Two 10-s and one 100-s exposures were taken in R, two 30-s and one 300-s in B, and one 100-s and one 600-s in U. The detector was a Tektronix CCD with $2048\times2048$ pixels of $24\,\mu$ m. The standard resolution collimator was used, for which the plate scale is $0\farcs2{\rm\,pix^{-1}}$. |
The detector was read ou through all four amplifiers. using the low-eain setting. of about 3¢ADU!. | The detector was read out through all four amplifiers, using the low-gain setting, of about $3{\rm\,e^-\,ADU^{-1}}$. |
The data reduction was done usiue standard procedures. | The data reduction was done using standard procedures. |
From bias frames taken before and after the welt. it was found that the level was somewhat variable. both in time aud iu position on the detector. but that the | From bias frames taken before and after the night, it was found that the level was somewhat variable, both in time and in position on the detector, but that the |
(see NBWO3. Figure 3). | (see NBW03, Figure 8). |
It is clear that extva-mixing on the RGB alone does not account for the measurements in the most O-poor grains. | It is clear that extra-mixing on the RGB alone does not account for the measurements in the most $^{18}$ O-poor grains. |
Indeed. it would require a much higher content of O than observed. | Indeed, it would require a much higher content of $^{17}$ O than observed. |
The right panel of Figure 1. adds more evidence by including the Al isotopic ratio. | The right panel of Figure \ref{one} adds more evidence by including the Al isotopic ratio. |
The grid of dotted curves refers to extra-mixing calculations on the RGB fora 2 M. model: the results are however tvpical of the whole range of low mass stars. | The grid of dotted curves refers to extra-mixing calculations on the RGB for a 2 $M_{\odot}$ model: the results are however typical of the whole range of low mass stars. |
The curves roughly going from left to right are for M. values of 0.03. 0.1. 0.3. 1 in units of 5 M. /vr. | The curves roughly going from left to right are for $\dot M$ values of 0.03, 0.1, 0.3, 1 in units of $^{-6}$ $M_{\odot}/$ yr. |
The higher is the AM value. the larger is the O depletion. as indicated by the labels. | The higher is the $\dot M$ value, the larger is the $^{18}$ O depletion, as indicated by the labels. |
The almost vertical lines. instead. refer to different values of Aloe7 (0.2. 0.15. 0.1. | The almost vertical lines, instead, refer to different values of $\Delta \log~T$ (0.2, 0.15, 0.1, |
range of cases. | range of cases. |
An improved version of the formula with fi=V(d/4.0)(1— and fa=0.4 was presented in ?.. | An improved version of the formula with $f_1=\sqrt{(d/4.0)(1-2.3/k)}$ and $f_2=0.4$ was presented in \cite{2010ApJ...721..582B}. |
But these formulas 2.3/k)were derived for the case where the data is generated by a homogeneous Poisson process and are reasonably accurate for most applications. | But these formulas were derived for the case where the data is generated by a homogeneous Poisson process and are reasonably accurate for most applications. |
For non-homogeneous data the of on large scales makes the formula presenceslightly densityinaccurate. | For non-homogeneous data the presence of density gradients on large scales makes the formula slightly inaccurate. |
gradientsHence for greater precision one should derive the values of f, and f» for a model data whose large scale distribution is similar to the data being studied but otherwise does not have any substructures in it. | Hence for greater precision one should derive the values of $f_1$ and $f_2$ for a model data whose large scale distribution is similar to the data being studied but otherwise does not have any substructures in it. |
We consider two models here a) a Gaussian sphere b) a Hernquist sphere with a scale radius of 15 kpc, which was shown to be an appropriate description of the stellar halo on large scales by ?.. | We consider two models here a) a Gaussian sphere b) a Hernquist sphere with a scale radius of 15 kpc, which was shown to be an appropriate description of the stellar halo on large scales by \cite{2005ApJ...635..931B}. |
The sample size of models was set to 10". | The sample size of models was set to $10^7$. |
A value of fi=0.93 and fo=0.4 was found to fit the distribution of significance of spurious groups and this is shown in2. | A value of $f_1=0.93$ and $f_2=0.4$ was found to fit the distribution of significance of spurious groups and this is shown in. |
. The distribution of significance averaged over 11 ACDM Figurehalos for the rest of the data sets is also shown. | The distribution of significance averaged over $11$ $\Lambda$ CDM halos for the rest of the data sets is also shown. |
At small S the number of groups are dominated by the spurious groups and the curves lie on the predicted relationship. | At small $S$ the number of groups are dominated by the spurious groups and the curves lie on the predicted relationship. |
At large S due to presence of real structures the curves flatten out. | At large $S$ due to presence of real structures the curves flatten out. |
The symbols in the figure denote the mean adopted value of ST; for each data set. | The symbols in the figure denote the mean adopted value of $S_{Th}$ for each data set. |
The distribution of stars in data set S5 is very similar to data set S3 except for the fact that it is not over all sky. | The distribution of stars in data set S5 is very similar to data set S3 except for the fact that it is not over all sky. |
Hence we revise the value of 571, for the data set S5 slightly from 4.55 as shown in the plot to 4.75. | Hence we revise the value of $S_{Th}$ for the data set S5 slightly from 4.55 as shown in the plot to 4.75. |
Note, for data sets with sample size less than 10? the adopted relation is found to slightly underestimate the number of spurious groups but we ignore this. | Note, for data sets with sample size less than $10^5$ the adopted relation is found to slightly underestimate the number of spurious groups but we ignore this. |
Even ifthis fact is taken into account this would only lead to a very minor revision in values of S1 for data sets S2 and S4 and since the corresponding curves are quite flat in this region the number of detected groups would also not change much. | Even ifthis fact is taken into account this would only lead to a very minor revision in values of $S_{\rm Th}$ for data sets S2 and S4 and since the corresponding curves are quite flat in this region the number of detected groups would also not change much. |
An example application of the group finder to one of the halos from data set S1 is shown in1. | An example application of the group finder to one of the halos from data set S1 is shown in. |
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