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Choosing the models whose γ΄ values fall below 3.9 (assuming that the errors are Gaussian and that we want a the radius and mass are determined by calculating the average value of each. and the uncertainty is given by the dispersion in the values divided by 6 (see?.fordetails)...
Choosing the models whose $\chi^2$ values fall below $^2$ (assuming that the errors are Gaussian and that we want a the radius and mass are determined by calculating the average value of each, and the uncertainty is given by the dispersion in the values divided by 6 \citep[see][for details]{2010A&A...511A..46M}.
We estimate the age of the star by simply matching the “fitted mass” with the closest age from the model parameter results. we therefore have no quantitative measure of its uncertainty.
We estimate the age of the star by simply matching the “fitted mass” with the closest age from the model parameter results, we therefore have no quantitative measure of its uncertainty.
leading to a very preliminary value.
leading to a very preliminary value.
The results are given in Table 2..
The results are given in Table \ref{tab:rm}.
The estimate of the age implies a star coming close to the end of the main sequence phase.
The estimate of the age implies a star coming close to the end of the main sequence phase.
Thanks to spectroscopic observations from NARVAL. we have been able to estimate the effective temperature as well as the chemical abundances of HD 170987 showing that this star has alow metallicity + ddex) and a high abundance of lithium.
Thanks to spectroscopic observations from NARVAL, we have been able to estimate the effective temperature as well as the chemical abundances of HD 170987 showing that this star has a low metallicity $\pm$ dex) and a high abundance of lithium.
The surface light elements abundances could giveus an interesting clue about the internal rotation of HD 170987.
The surface light elements abundances could giveus an interesting clue about the internal rotation of HD 170987.
The effective temperature of 6540K sets this star in the so- lithium dip where both “Li and "Be are depleted (?)..
The effective temperature of $\rm 6540\,K$ sets this star in the so-called lithium dip where both $\rm ^7Li$ and $\rm ^9Be$ are depleted \citep{2002ApJ...565..587B}. .
Here we emphasize that we have also successfully tested our algorithm at the higher frequencies (~350 MHz) by simulating Galactic emission towards the cluster Abel 2255.
Here we emphasize that we have also successfully tested our algorithm at the higher frequencies $\sim350~{\rm MHz}$ ) by simulating Galactic emission towards the cluster Abel 2255.
The simulated maps have been compared with the observations obtained by R.F. Pizzo et al. communication)).
The simulated maps have been compared with the observations obtained by R.F. Pizzo et al. ).
Details of these simulations will be presented in a separate paper (Jeli¢ et al., prep-)).
Details of these simulations will be presented in a separate paper (Jelić et al., ).
To summarize, simulated maps show all observed characteristics of the Galactic emission, e.g. presence of the structures at different scales, spatial and frequency variations of the brightness temperature and its spectral index, complex Faraday structures, and depolarization.
To summarize, simulated maps show all observed characteristics of the Galactic emission, e.g. presence of the structures at different scales, spatial and frequency variations of the brightness temperature and its spectral index, complex Faraday structures, and depolarization.
Based on these results, we conclude that our model is able to simulate realistic maps of the Galactic emission both in total and polarized intensity, and can be used as an realistic foreground model in simulations of the EoR experiments.
Based on these results, we conclude that our model is able to simulate realistic maps of the Galactic emission both in total and polarized intensity, and can be used as an realistic foreground model in simulations of the EoR experiments.
The LOFAR-EoÓR project relies on a detailed understanding of astrophysical and non-astrophysical contaminations that can contaminate the EoR signal: the Galactic and extragalactic foregrounds, ionosphere, instrumental effects and systematics.
The LOFAR-EoR project relies on a detailed understanding of astrophysical and non-astrophysical contaminations that can contaminate the EoR signal: the Galactic and extragalactic foregrounds, ionosphere, instrumental effects and systematics.
In order to study these components and their influence on the detection of the EoR signal, a LOFAR-EoR simulation pipeline is being developed by the LOFAR-EoR team.
In order to study these components and their influence on the detection of the EoR signal, a LOFAR-EoR simulation pipeline is being developed by the LOFAR-EoR team.
The pipeline consists of the three main modules: the EoR signal (basedonsimulationsdescribedin ?),, the foregrounds (basedonthispaperand?) and the instrumentalresponse (describedon ?)..
The pipeline consists of the three main modules: the EoR signal \citep[based on simulations described in][]{thomas09}, , the foregrounds \citep[based on this paper and ][]{jelic08} and the instrumentalresponse \citep[described on][]{panos09}. .
The mass detection threshold was calculated assuming that clouds would be unresolved in the GBT beam. using the relation: where D is the average distance to the filament. 7, is the RMS noise in one channel. AV is the chamuel width. aud Nis the umber of channels required Or a sect detection.
The mass detection threshold was calculated assuming that clouds would be unresolved in the GBT beam, using the relation: where $D$ is the average distance to the filament, $\sigma_s$ is the RMS noise in one channel, $\Delta V$ is the channel width, and $N$ is the number of channels required for a secure detection.
For this calculation. we used an cloud linewidth of 20 kins |. corresponding oN=|.
For this calculation, we used an cloud linewidth of 20 km $^{-1}$, corresponding to $N=4$.
This asstuned Lnewidth strikes a valance between narrower IIVC lines (Wakker& aud the wider Luewidths of nassive clouds (Chynowethetal. 2008)..
This assumed linewidth strikes a balance between narrower HVC lines \citep{wvw97} and the wider linewidths of massive clouds \citep{chy08}. .
Adopting th AIS1 distance of 3.6 Mpc. the la nass detection threshold ranges from 2.5 to 12.9 ν 10° Mo. as shown in Figure 3..
Adopting the M81 distance of 3.6 Mpc, the $\sigma$ mass detection threshold ranges from 2.5 to 12.9 $\times$ $^5$ $_{\odot}$, as shown in Figure \ref{RMSfig}.
The average 55 detection threshold for mass was 32 «& 109 ND...
The average $\sigma$ detection threshold for mass was 3.2 $\times$ $^{6}$ $_{\odot}$.
With this threshold. we would be able to detect analogues to the most massive clouds around M31 aud M23. the M81/M82 clouds. the AILIOL cloud. aud large Milky. Way objects such as Complexes C aud II and the \lagellanic Stream.
With this threshold, we would be able to detect analogues to the most massive clouds around M31 and M33, the M81/M82 clouds, the M101 cloud, and large Milky Way objects such as Complexes C and H and the Magellanic Stream.
We would also be able to detect cold accretion clouds found in the simulations of I&eres&IIeru-quist (2009).
We would also be able to detect cold accretion clouds found in the simulations of \citet{Dusan:2009}.
We discovered 5 new clouds.
We discovered 5 new clouds.
We divide these clouds iuto two groups. those associated with the M81 Fibuueut aud those associated with the Milkv Way.
We divide these clouds into two groups, those associated with the M81 Filament and those associated with the Milky Way.
We determine the most likely association of cach cloud based on its position aud velocity.
We determine the most likely association of each cloud based on its position and velocity.
There is some overlap in velocity space between the ALSL Filament aud the Milkv War. and over the velocity range -85 to 25 lan | we can not discrinuuate between local aud distant cluission.
There is some overlap in velocity space between the M81 Filament and the Milky Way, and over the velocity range -85 to 25 km $^{-1}$, we can not discriminate between local and distant emission.
In addition. the Galactic IVC Complex A. which ranges im velocity from -120 lan | to -200 kin coincides with part of the mapped region (Wakker&vanWoorden1997).
In addition, the Galactic HVC Complex A, which ranges in velocity from -120 km $^{-1}$ to -200 km $^{-1}$, coincides with part of the mapped region \citep{wvw97}.
. Therefore we cannot distinguish clouds in this position and velocity range from Couples A. All of the clouds we previously detected (Chyuowethetal.2008.20090). are visible iu this extended tage.
Therefore we cannot distinguish clouds in this position and velocity range from Complex A. All of the clouds we previously detected \citep{chy08, chy09} are visible in this extended image.
Their locatious are marked iu Figure 1..
Their locations are marked in Figure \ref{mom0}.
In order to unity the clouds identified in this paper with those in Chynowethetal.(2008) and Chynowethetal.(2009).. and auv clouds identified with the GBT iun our current and future observations. we have re-designated the clouds with the ideutifier GBC for "Green Bank Cloud”. and the coordinates of the cloud. i. ον GBC Jhbuuuss.s|dadaiuuss.
In order to unify the clouds identified in this paper with those in \citet{chy08} and \citet{chy09}, and any clouds identified with the GBT in our current and future observations, we have re-designated the clouds with the identifier GBC for “Green Bank Cloud", and the coordinates of the cloud, i. e. GBC Jhhmmss.s+ddmmss.
These designations are found in Table 3..
These designations are found in Table \ref{fil_cloudprop}. .
To distinguish the newly discovered clouds frou. the clouds previously discovered. the new clouds are designated with CIO0-1 to C10-5.
To distinguish the newly discovered clouds from the clouds previously discovered, the new clouds are designated with C10-1 to C10-5.
Table 3/— lists their properties.
Table \ref{fil_cloudprop} lists their properties.
Spectra are shown in Figure L.
Spectra are shown in Figure \ref{cloud_ispec}.
Figure 2. shows the locatious of all clouds that we have determined to be located within AISI Filament.
Figure \ref{fil_analysis} shows the locations of all clouds that we have determined to be located within M81 Filament.
Cloud masses were calculated using where Fis the total flux of the cloud.
Cloud masses were calculated using where $F$ is the total flux of the cloud.
Exror estimates on the masses difficult to obtain. given the wide range in RAIS noise values over the map.
Error estimates on the masses difficult to obtain, given the wide range in RMS noise values over the map.
The dominant error contribution is the approximately uncertainty idu the absolute calibration for these observations.
The dominant error contribution is the approximately uncertainty in the absolute calibration for these observations.
Cloud CLlO-L (GBC J092635.5|702850) is to the northwest of the M81 eroup.
Cloud C10-1 (GBC J092635.8+702850) is to the northwest of the M81 group.
With a velocity of -178 Ian ft. Cloud C10-1 iav be part of Couples A. However. Complex A does not overlap in position with Cloud CIO0-1. so we include this cloud in further analysis of clouds in the M8SI Filuneut.
With a velocity of -178 km $^{-1}$, Cloud C10-1 may be part of Complex A. However, Complex A does not overlap in position with Cloud C10-1, so we include this cloud in further analysis of clouds in the M81 Filament.
Cloud C10-2 (GBC J101926.3|675222) is south of IC 2571 aud au adjacent IILJASS source. HILJASS J1021]|6510 (Dovcoetal.2001). ane has a velocity of 38 lan s|.
Cloud C10-2 (GBC J101926.3+675222) is south of IC 2574 and an adjacent HIJASS source, HIJASS J1021+6840 \citep{boy01} and has a velocity of 38 km $^{-1}$.
Cloud C10-3 (GBC J091952.7|680937) is to the soutlavest of the M81 eroup and has a velocity of -108 kin +.
Cloud C10-3 (GBC J091952.7+680937) is to the southwest of the M81 group and has a velocity of -108 km $^{-1}$.
Clo ClOL (OBC J1022:39.1|681057) is located near the previously reported IILTASS source.
Cloud C10-4 (GBC J1022:39.1+684057) is located near the previously reported HIJASS source.
This clouc is relatively brightaud is a few arcaumutes cast of the IILFASS source.
This cloud is relatively brightand is a few arc-minutes east of the HIJASS source.
Boveeetal.(2001). fux the ΠΙΑ source is extended in the directionof IC 257I.
\citet{boy01} find the HIJASS source is extended in the directionof IC 2574.
Our obscrvations have higher aneular resolutiou. aud in these observations Cloud
Our observations have higher angular resolution, and in these observations Cloud
we found that the ratio between the fluxes in the rest-frame wavelength ranges 4000-4100 aand 3850-3950 iis D,(4000)~1.06 (Baloghetal. 1999).
we found that the ratio between the fluxes in the rest-frame wavelength ranges 4000–4100 and 3850–3950 is $\rm D_{n}(4000) \sim 1.06$ \citep{bal99}. .
Assuming D,(4000)~2 for the host galaxy (Kauffmannetal.2003),, as typical of giant early-type galaxies, this implies a host galaxy contribution not greater than a few percent.
Assuming $\rm D_{n}(4000) \sim 2$ for the host galaxy \citep{kau03}, as typical of giant early-type galaxies, this implies a host galaxy contribution not greater than a few percent.
Thus we did not correct the spectra for the host galaxy contribution, because it is negligible for our purposes.
Thus we did not correct the spectra for the host galaxy contribution, because it is negligible for our purposes.
The VHR-R spectra were corrected for the telluric absorption bands, a particularly important step, because the 6870 ooxygen B-band affects the blue wing of the Ha line.
The VHR-R spectra were corrected for the telluric absorption bands, a particularly important step, because the 6870 oxygen $B$ -band affects the blue wing of the $\alpha$ line.
This was done by dividing the spectra by a template obtained from the spectrophotometric standard star.
This was done by dividing the spectra by a template obtained from the spectrophotometric standard star.
Figure 2 displays the low-resolution (left panels) as well as the high-resolution (right panels) spectra, the latter both before and after correction for atmospheric absorption around the Ha line.
Figure \ref{spectra} displays the low-resolution (left panels) as well as the high-resolution (right panels) spectra, the latter both before and after correction for atmospheric absorption around the $\alpha$ line.
We measured the emission line intensities by means of the package in IRAF (Kriss1994).
We measured the emission line intensities by means of the package in IRAF \citep{kri94}.
. To reduce the number of free parameters we fixed the relative wavelength distance between lines and required the FWHM to be the same for all the narrow lines.
To reduce the number of free parameters we fixed the relative wavelength distance between lines and required the FWHM to be the same for all the narrow lines.
From the low-resolution spectra we derived the [O 1I1]444959,5007 flux, and an upper limit to the H8 flux.
From the low-resolution spectra we derived the [O $\lambda\lambda$ 4959,5007 flux, and an upper limit to the $\beta$ flux.
In the high-resolution spectra, we fitted the 6700-7350 sspectral region with seven components: a Gaussian profile for each of the lines: Ha (broad and narrow), [N II]446548,84, and [S I1]446716,31, plus a linear component for the continuum.
In the high-resolution spectra, we fitted the 6700–7350 spectral region with seven components: a Gaussian profile for each of the lines: $\alpha$ (broad and narrow), [N $\lambda\lambda$ 6548,84, and [S $\lambda\lambda$ 6716,31, plus a linear component for the continuum.
The results of this spectral fitting are shown in refspecfit..
The results of this spectral fitting are shown in \\ref{specfit}.
The fit to the [S II]J446716,6731 doublet must be considered with caution, because this line is strongly affected by atmospheric absorption and thus its measurement strongly depends on the telluric correction (see refspectra)).
The fit to the [S $\lambda\lambda$ 6716,6731 doublet must be considered with caution, because this line is strongly affected by atmospheric absorption and thus its measurement strongly depends on the telluric correction (see \\ref{spectra}) ).
The FWHM of the narrow lines ranges from ~ to ~600kms", while that of the broad Ha component is 4200-5000 kms.
The FWHM of the narrow lines ranges from $\sim 500$ to $\sim 600 \rm \, km \, s^{-1}$, while that of the broad $\alpha$ component is 4200–5000 $\rm km \, s^{-1}$.
Including a further component to fit the [O 1]446300,64 line resulted in a marginal detection only.
Including a further component to fit the [O $\lambda\lambda$ 6300,64 line resulted in a marginal detection only.
Table 2 reports the line fluxes after correction for the Galactic extinction, adopting Ag=1.42, E(B—V)=0.35, and the extinction law of Cardellietal.(1989).
Table \ref{lines} reports the line fluxes after correction for the Galactic extinction, adopting $A_B = 1.42$, $E(B-V)$ =0.35, and the extinction law of \citet{car89}.
. The first detection of broad emission lines in the optical spectrum of BL Lacertae was reported by Vermeulenet (1995)..
The first detection of broad emission lines in the optical spectrum of BL Lacertae was reported by \citet{ver95}. .
They took two spectra in May and June 1995, and measured Ha with FWHM of 3400-3700kms! and flux of 4.4x107ergcm? s~!, while for Hf they obtained FWHM=4400kms! and F=1.3x107ergcm?s!.
They took two spectra in May and June 1995, and measured $\alpha$ with FWHM of $3700 \rm \, km \, s^{-1}$ and flux of $4.4 \times 10^{-14} \rm \, erg \, cm^{-2} \, s^{-1}$ , while for $\beta$ they obtained $ \rm FWHM = 4400 \, km \, s^{-1}$ and $F = 1.3 \times 10^{-14} \rm \, erg \, cm^{-2} \, s^{-1}$.
They stressed that such a luminous Ha line should have been observed before, while it is not recognizable in earlier spectra taken in 1976 by Miller&Hawley(1977),, in 1985 by Lawrenceetal. (1996),, and in 1989 by Stickeletal.(1993).
They stressed that such a luminous $\alpha$ line should have been observed before, while it is not recognizable in earlier spectra taken in 1975--1976 by \citet{mil77}, in 1985 by \citet{law96}, and in 1989 by \citet{sti93}.
. They estimated that the broad Ha flux may have increased by at least a factor 5 since 1989.
They estimated that the broad $\alpha$ flux may have increased by at least a factor 5 since 1989.
Soon after, Corbettetal.(1996) analysed two other optical spectra acquired in June 1995, and confirmed the results of Vermeulenetal.(1995).
Soon after, \citet{cor96} analysed two other optical spectra acquired in June 1995, and confirmed the results of \citet{ver95}.
. They discussed the appearance of the Ha line as due to an increase of either the amount of gas in the broad line region (BLR) or the strength of the photoionising source.
They discussed the appearance of the $\alpha$ line as due to an increase of either the amount of gas in the broad line region (BLR) or the strength of the photoionising source.
In the lattercase, the accretion disc seemed the most likely source of photoionising radiation.
In the lattercase, the accretion disc seemed the most likely source of photoionising radiation.
The signature of the accretion disc was later recognizedby Raiterietal.(2009) as a UV excess in the broad-band SEDs of BL Lacertae built with contemporaneous low-energy data from the WEBT and data from the XMM-Newtonsatellite.
The signature of the accretion disc was later recognizedby \citet{rai09} as a UV excess in the broad-band SEDs of BL Lacertae built with contemporaneous low-energy data from the WEBT and high-energy data from the XMM-Newtonsatellite.
both the @=(0 axis (upper panel) aud the Y=0 axis (lower panel).
both the ${\Phi}=0$ axis (upper panel) and the $W= 0$ axis (lower panel).
Given the orbit parameters. it is worthwhile to relate them. with other paramucters ofLyrs.
Given the orbit parameters, it is worthwhile to relate them with other parameters of.
. For that. we have divided our sample in subgroups iu various wavs BR rofRvol.tab)).
For that, we have divided our sample in subgroups in various ways \\ref{RRvel.tab}) ).
We used uon kinematic selectio1 criteria as inetallicity and period but the kinematic criterion eccentricity as well.
We used non kinematic selection criteria as metallicity and period but the kinematic criterion eccentricity as well.
We choose as cuts [Fe/TT]0.5. > d. o loauxb zco d.5.forceeem0.15 and 0.15. aud for period P«0.35 dd. 0.35«P< 0.55dd and P>0.55 d. is not a kincmatic parameter.
We choose as cuts $[{\rm Fe/H}]>-0.5$, $>-1$ , $\leq-1$, and $\leq-1.5$ , for $ecc \leq0.45$ and $>0.45$, and for period $P<0.35 $ d, $0.35<P<0.55$ d and $P>0.55$ d. is not a kinematic parameter.
Yet. ietallicitv eives no clear division between disk aud halo stars.
Yet, metallicity gives no clear division between disk and halo stars.
Stars with ο=Lo may be menibers of the disk eroup in view of their O 4ji particular those with [Fe/H|= 0.5).
Stars with $[{\rm Fe/H}] \geq -1$ may be members of the disk group in view of their $\bar{\Theta}$ (in particular those with $[{\rm Fe/H}] \geq -0.5$ ).
But the velocity dispersion σο=90 Hs lager than expected for a pure disk population.
But the velocity dispersion $\sigma_\Theta \simeq 90$ is larger than expected for a pure disk population.
So sole of the wwith |Fe/II] 2—1 cau be regarded to be halo stars.
So some of the with [Fe/H] $\geq -1$ can be regarded to be halo stars.
The more metal-poor wwith [Fe/TMJ<—1 are candidates for halo stars.
The more metal-poor with $\leq-1$ are candidates for halo stars.
This is corroborated by the result of O=11laus. go=110hans... οσο=0.69 and το=0.71 for this group.
This is corroborated by the result of $\bar{\Theta}=14$, $\sigma_\Theta=110$, $\overline{ecc} =0.69$ and $\overline{nze} =0.71$ for this group.
The velocity values of ©. V. aud Ü for the metal poor part of our sample are simular to several published halo samples (sce Martin Morrison (1998). them Table 5) even if in their works the cuts for laο stars are set at more metal poor levels.
The velocity values of $\bar{U}$, $\bar{V}$ and $\bar{U}$ for the metal poor part of our sample are similar to several published halo samples (see Martin Morrison (1998), their Table 5) even if in their works the cuts for halo stars are set at more metal poor levels.
Therefore we give the results for wwith Γον—1.5. too.
Therefore we give the results for with $ \leq-1.5$, too.
For this group our results shows to be iu good agreement with results of previous investigations (soo Martin Morrison 1998. Chen 1999).
For this group our results shows to be in good agreement with results of previous investigations (see Martin Morrison 1998, Chen 1999).
This metal-poor group of αμα ο—13bo. a ümiean retrograde rotation with respect to the
This metal-poor group of has $\bar{\Theta}=-13$, i.e., a mean retrograde rotation with respect to the.
LSR?..Orbit divides between vounecr and older stars (nore gravitational interactions of disk stars lead to larger deviations frou the originally circular orbits).
supposedly divides between younger and older stars (more gravitational interactions of disk stars lead to larger deviations from the originally circular orbits).
The subdisisiou at ecce=0.15 shows the expected correlations with kincmatic parameters (see refRRvel.tab:: notablv £L).
The subdivision at $ecc=0.45$ shows the expected correlations with kinematic parameters (see \\ref{RRvel.tab}; notably $\bar{I_{\rm z}}$ ).
This subclivision does not correlate iu à prouounced way with metallicity.
This subdivision does not correlate in a pronounced way with metallicity.
Note that we did not include 7 stars with retrograde orbits iu the subgroup with eccx:0.15 refRBR vel.tab)).
Note that we did not include 7 stars with retrograde orbits in the subgroup with $ecc\leq 0.45$ \\ref{RRvel.tab}) ).
These 7 have aiean circular velocity wowith dispersion of σο=56kins..
These 7 have a mean circular velocity $\bar{\Theta}=-165$ with dispersion of $\sigma_\Theta =56$.
Also Chiba Yoshii (1997) report the presence of halo stars with low ecceutricityorbits.
Also Chiba Yoshii (1997) report the presence of halo stars with low eccentricityorbits.
Bothresults show that low eccentricity does not always mean that a star is a disk member.
Bothresults show that low eccentricity does not always mean that a star is a disk member.
The 58 aare good candidates for disk stars (thin aud thick disk).
The 58 are good candidates for disk stars (thin and thick disk).
density for each velocity component assuming an isotopic ratio of ΟΡΟ=150 and a CO/Ib=10 |.
density for each velocity component assuming an isotopic ratio of $\rm^{16}O/^{18}O=150$ \citep{Harrison99} and a $_2$ $10^{-4}$.
The relative abundances derived for all molecules are presented in Table 2..
The relative abundances derived for all molecules are presented in Table \ref{tab:abunRatios}.
Additionally. abundance ratio o£ CO — with respect all species is also presented in Table 2..
Additionally, abundance ratio of $^{13}$ $^+$ with respect all species is also presented in Table \ref{tab:abunRatios}.
The errors in (he derived column densities take into account the statistical error of integrated intensities and the uncertainty in (he assumed excitation temperature.
The errors in the derived column densities take into account the statistical error of integrated intensities and the uncertainty in the assumed excitation temperature.
These errors are subsequently propagated to (he abundance ratios.
These errors are subsequently propagated to the abundance ratios.
The abundance has been derived [rom that of IMCO —(o avoid the opacity effects affecting the main isotopologue.
The $^{+}$ abundance has been derived from that of $^{13}$ $^{+}$to avoid the opacity effects affecting the main isotopologue.
Indeed. if the C/C ratio of ~40 derived for 2253 and 44945(Ilenkeletal.1993.1994:Martín2005) applies to all these ealaxies. an average opacily of TeoyoyoZi lis derived from the ICO/IICO.J=1—0 line ratio (hisworkandUseroetal. 2004).
Indeed, if the $^{12}$ $^{13}$ C ratio of $\sim40$ derived for 253 and 4945\citep{Henkel93,Henkel94,Martin05} applies to all these galaxies, an average opacity of $\tau_{HCO^+\,J=1-0}\gtrsim1$ is derived from the $\rm HCO^+/H^{13}CO^+ \,J=1-0$ line ratio \citep[this work and][]{Usero04}. .