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The C/O ratio obviously varies significantly from star to star. indicating different amounts of enriched material dredged up to the surface.
The C/O ratio obviously varies significantly from star to star, indicating different amounts of enriched material dredged up to the surface.
It would be expected that an increased C/O ratio due to the dredge up of C would be accompanied by an increased isotopic abundance ratio.
It would be expected that an increased C/O ratio due to the dredge up of $^{12}$ C would be accompanied by an increased isotopic abundance ratio.
Indeed we find a nice correlation between C/O and C/C for the O-rich stars as can be seen from Table 2..
Indeed we find a nice correlation between C/O and $^{12}$ $^{13}$ C for the O-rich stars as can be seen from Table \ref{t:coetc}.
The values for C/O and "C/C ratio found are in good agreement with findings from Smith Lambert (1990)) for field stars.
The values for C/O and $^{12}$ $^{13}$ C ratio found are in good agreement with findings from Smith Lambert \cite{SL90}) ) for field stars.
The minimum C/O ratio in our sample of 0.2 agrees closely with expectations from first dredge-up.
The minimum C/O ratio in our sample of 0.2 agrees closely with expectations from first dredge-up.
As described above. one blend including a line of HF (1- R7) is covered by our K band observations.
As described above, one blend including a line of HF (1-0 R7) is covered by our $K$ band observations.
After limiting the various parameters affecting the shape and depth of this blend. namely mainly the temperature and the carbon isotopic ratio. from other parts of the stellar spectra we fitted the remaining difference between the observed and the calculated blend profile by a change in the F abundance.
After limiting the various parameters affecting the shape and depth of this blend, namely mainly the temperature and the carbon isotopic ratio, from other parts of the stellar spectra we fitted the remaining difference between the observed and the calculated blend profile by a change in the F abundance.
Due to the rather large uncertainties of the C-star fitting we decided to derive the F abundance only for the O-rich stars.
Due to the rather large uncertainties of the C-star fitting we decided to derive the F abundance only for the O-rich stars.
Figure 2 shows as an example part of the spectrum of LES including the HF line (marked by àn arrow).
Figure \ref{hfplot} shows as an example part of the spectrum of LE8 including the HF line (marked by an arrow).
Overplotted are three models with different F abundances.
Overplotted are three models with different F abundances.
Obviously. Fluorine is underabundant in LES relative to the other metals.
Obviously, Fluorine is underabundant in LE8 relative to the other metals.
It turns out that the F abundance has to be changed from star to star to allow for a reasonable fit of the observed blend.
It turns out that the F abundance has to be changed from star to star to allow for a reasonable fit of the observed blend.
Uncertainty in determining the F abundance comes from both the temperature and the C/" C ratio (via a "CO line in the blend).
Uncertainty in determining the F abundance comes from both the temperature and the $^{12}$ $^{13}$ C ratio (via a $^{13}$ CO line in the blend).
We made an estimate of the uncertainty by varying these two parameters around the values determined from other parts of the spectra and then summing in quadrature the respective uncertainties.
We made an estimate of the uncertainty by varying these two parameters around the values determined from other parts of the spectra and then summing in quadrature the respective uncertainties.
Taking mto account this uncertainty and the fact that we have only one HF line to derive the abundance the results of course have to be taken with some caution.
Taking into account this uncertainty and the fact that we have only one HF line to derive the abundance the results of course have to be taken with some caution.
Confirmatio from other HF lines is needed to limit the possible effect of unidentified lines on the profile of the blend and the derived abundances.
Confirmation from other HF lines is needed to limit the possible effect of unidentified lines on the profile of the blend and the derived abundances.
The C/O ratio was derived for two C-rich stars. LE]1] and LE2.
The C/O ratio was derived for two C-rich stars, LE11 and LE2.
For the other C-rich stars. either no H band spectra were available (as a result of limited observing time) or no acceptable fit of the spectrum could be achieved.
For the other C-rich stars, either no $H$ band spectra were available (as a result of limited observing time) or no acceptable fit of the spectrum could be achieved.
For the latter case. the starting influence of stellar variability on the stellar spectrum may play some role besides the above-mentioned problems.
For the latter case, the starting influence of stellar variability on the stellar spectrum may play some role besides the above-mentioned problems.
The derived 7. and C/O for the C-rich star LEII come with a much larger uncertainty.
The derived $T_{\rm eff}$ and C/O for the C-rich star LE11 come with a much larger uncertainty.
From the strength of the CO 3-0 band head. we exclude a C/O of less than 1.4.
From the strength of the CO 3-0 band head, we exclude a C/O of less than 1.4.
A macroturbulent velocity of ss! is needed to both fit the width of the features and their strength.
A macroturbulent velocity of $^{-1}$ is needed to both fit the width of the features and their strength.
The temperature can be constrained even less.
The temperature can be constrained even less.
If we use the value from near infrared photometry. we get a Tay of KK. As noted by Lebzelter Wood (2007)). this temperature is likely to be too low as pulsation properties instead suggest a value close to KK. We produced model spectra around both temperatures.
If we use the value from near infrared photometry, we get a $T_{\rm eff}$ of K. As noted by Lebzelter Wood \cite{LW07}) ), this temperature is likely to be too low as pulsation properties instead suggest a value close to K. We produced model spectra around both temperatures.
A value of 7,;22950KK seems to give the better fit. but the sensitivity of the synthetic spectrum on this parameter is not very high.
A value of $T_{\rm eff}$ K seems to give the better fit, but the sensitivity of the synthetic spectrum on this parameter is not very high.
Concerning an upper limit of the C/O ratio an overall fit of the spectrum of similar quality can be reached for various combinations of C/O and Τομ with a higher temperature
Concerning an upper limit of the C/O ratio an overall fit of the spectrum of similar quality can be reached for various combinations of C/O and $T_{\rm eff}$ with a higher temperature
cutoff. 1024-2048. MIIz. anti-aliasing filter.
cutoff, 1024-2048 MHz, anti-aliasing filter.
This is the input to the 8 bit cligilizers in the DTS module. which operate on a 2045 MlIz clock.
This is the input to the 8 bit digitizers in the DTS module, which operate on a 2048 MHz clock.
Gain slope equalization is not performed on (his signal due to the greater cdvnamic range of the 8-bit digitizer.
Gain slope equalization is not performed on this signal due to the greater dynamic range of the 8-bit digitizer.
The local oscillator svstem (Figures 3. ancl 4)) is based on 123 MIIz and 512 MIIZ master reference signals. both. [rom a crystal oscillator locked to a livdrogen maser and a 1 llz. GPS-based. master timing signal.
The local oscillator system (Figures \ref{fig:if} and \ref{fig:lo}) ) is based on 128 MHz and 512 MHz master reference signals, both from a crystal oscillator locked to a hydrogen maser and a 1 Hz, GPS-based, master timing signal.
These signals are used to generate all other reference and timing information used by the EVLA svstem.
These signals are used to generate all other reference and timing information used by the EVLA system.
With the exception of the maser. the master relerence svstem is fully redundant (to ensure high reliability and continued operations during maintenance.
With the exception of the maser, the master reference system is fully redundant to ensure high reliability and continued operations during maintenance.
Relerence signals are distributed (ο an antenna over a single mode. fiber optic cable operating al the 1310 nm zero-dispersion wavelength.
Reference signals are distributed to an antenna over a single mode, fiber optic cable operating at the 1310 nm zero-dispersion wavelength.
The relerence signal consists of a 512 MIIz sine wave with a phase inversion of GO ns duration occurring once every 10 seconds as a timing reference.
The reference signal consists of a 512 MHz sine wave with a phase inversion of 60 ns duration occurring once every 10 seconds as a timing reference.
Absolute time is passed to the central and antenna electronics using the network time protocol (NTP) over the gigabit Ethernet emploved bx the EVLÀ's monitor and control svstem.
Absolute time is passed to the central and antenna electronics using the network time protocol (NTP) over the gigabit Ethernet employed by the EVLA's monitor and control system.
Precise time is obtained in the antenna by svnchronizing the NTP time with the timing signals in the local electronics.
Precise time is obtained in the antenna by synchronizing the NTP time with the timing signals in the local electronics.
The phase of the reference signal to each antenna can change due to temperature and mechanical effects in (he fiber.
The phase of the reference signal to each antenna can change due to temperature and mechanical effects in the fiber.
The change is measured by a round trip phase (RED) svstem (Figure 4)).
The change is measured by a round trip phase (RTP) system (Figure \ref{fig:lo}) ).
The measurement is mace by splitting the optical reference signal at the antenna and sending part of it back to the central LO svstem on a second single mode fiber in the same cable.
The measurement is made by splitting the optical reference signal at the antenna and sending part of it back to the central LO system on a second single mode fiber in the same cable.
This signal is (hen compared (o a master reference signal to get the total phase change in the fiber.
This signal is then compared to a master reference signal to get the total phase change in the fiber.
The result is divided by (wo to get the phase change at the antenna.
The result is divided by two to get the phase change at the antenna.
Although there is no guarantee that the two fibers will behave exactly the same. experience indicates that thev behave sufficiently similarly that the factor of two mentioned above is adequate for the needed accuracy.
Although there is no guarantee that the two fibers will behave exactly the same, experience indicates that they behave sufficiently similarly that the factor of two mentioned above is adequate for the needed accuracy.
The accuracy of the phase measuring equipment is of order 100 Is. which is smaller than the relative clelavs introduced into the astronomical signals by the varving atmosphere above the antennas for most VLA baselines.
The accuracy of the phase measuring equipment is of order 100 fs, which is smaller than the relative delays introduced into the astronomical signals by the varying atmosphere above the antennas for most VLA baselines.
At the antenna. the optical signal is converted back to an electrical signal and routed io the antenna reference generator module (Figure 3)).
At the antenna, the optical signal is converted back to an electrical signal and routed to the antenna reference generator module (Figure \ref{fig:if}) ).
In this module. a ervstal oscillator is phase locked to the incoming reference signal.
In this module, a crystal oscillator is phase locked to the incoming reference signal.
The 128 MIIz and 512 MlIZz signals from the oscillator are fed to (wo. step recovery. diode-based. comb generators used to derive all of the higher frequency reference ancl LO signals in the antenna.
The 128 MHz and 512 MHz signals from the oscillator are fed to two, step recovery diode-based, comb generators used to derive all of the higher frequency reference and LO signals in the antenna.
In addition. a GO ns timing
In addition, a 60 ns timing
Figure 2.
Figure 2.
A composit¢of three infrared maps of NGC E118 shown in false colors.
A composite of three infrared maps of NGC 4418 shown in false colors.
The 1.1 pau. 1.6 pou. 2.2 jn tages are shown as blue. erce1. aud red. respectively.
The 1.1 $\mu$ m, 1.6 $\mu$ m, 2.2 $\mu$ m images are shown as blue, green, and red, respectively.
The field of view is ~11.1"<11.1".
The field of view is $\sim 11.4\arcsec \times 11.4\arcsec$.
The logarithin of cach image was taken before combining he three wavelengths to compress the dynamic range.
The logarithm of each image was taken before combining the three wavelengths to compress the dynamic range.
values of {νους=(3.16.3.44.5.52.3.16.3.16) for w=0.50.0.75.100.1.25. 1.50). respecivelv.
values of $\tilde{K}_{\rm 200}= (3.76, 3.44, 3.32, 3.16, 3.16)$ for $\wQ=(-0.50, -0.75, -1.00, -1.25, -1.50)$ , respecively.
This range in Koon is bot consisten with one cosmology-independent value of Kou.
This range in $\tilde{K}_{200}$ is not consistent with one cosmology-independent value of $\tilde{K}_{200}$.
The restuts of this analysis suggest. that moclels similar to the Bol model. in which the halo concentration is cdelinec in terms of JA; and Avy. are more reaclily generalizable to alternative cosmologies as eiis related to afa; via a constant of proportionality. Avi.
The results of this analysis suggest that models similar to the B01 model, in which the halo concentration is defined in terms of $R_{\rm vir}$ and $\DeltaVir$, are more readily generalizable to alternative cosmologies as $\cvir$is related to $a/a_c$ via a constant of proportionality, $K_{\rm vir}$.
Put another way. defining the radius of a halo. and thus its concentration. using a fixedixeclL overdensity|t criterion[ necessitatestat usingising a cosmologv-logdependent. proportionality constant in σα, (10))
Put another way, defining the radius of a halo, and thus its concentration, using a fixed overdensity criterion necessitates using a cosmology-dependent proportionality constant in Eq. \ref{eq:B01_vir}) )
while the cosmologv-dependent virial overdensity definition seems to account for these dillerences. so that Ava is independent of cosmologv.
while the cosmology-dependent virial overdensity definition seems to account for these differences, so that $K_{\rm vir}$ is independent of cosmology.
As in previous studies (BOL: Jing 2000: Jing Suto 2002). we also find that haloes of a given mass have a broad distribution of concentrations.
As in previous studies (B01; Jing 2000; Jing Suto 2002), we also find that haloes of a given mass have a broad distribution of concentrations.
To determine the inherent scatter in the cec-AM relation it is important to account for the artificial scatter introduced by uncertainties in the fit to an NEW profile and by the Poisson noise in each bin.
To determine the inherent scatter in the $\cvir$ $\Mvir$ relation it is important to account for the artificial scatter introduced by uncertainties in the fit to an NFW profile and by the Poisson noise in each bin.
Following the BOL analysis. we corrected for the former by determining 500 one-sided Gaussian cleviates for cach halo with a standard deviation equal to the error in the eg fit returned. by the halo finder.
Following the B01 analysis, we corrected for the former by determining 500 one-sided Gaussian deviates for each halo with a standard deviation equal to the error in the $\cvir$ fit returned by the halo finder.
Phe deviates are. positive (negative) i c is less (greater) than the median in that bin.
The deviates are positive (negative) if $\cvir$ is less (greater) than the median in that bin.
We then determined. the 16 and SA"HL percentiles. in. Ιοσ(ὅνιν) ancl subtract olf the Poisson noise from. cach in quadrature.
We then determined the $16^{th}$ and $84^{th}$ percentiles in $\log(c_{\rm vir})$ and subtract off the Poisson noise from each in quadrature.
Phe resulting estimates of the intrinsic scatter are shown as the dashed. lines in Figure 7..
The resulting estimates of the intrinsic scatter are shown as the dashed lines in Figure \ref{fig:cofM}. .
The scatter is consistent with being independent of w ancl Aa. and we find hat taking the BOL propetionality constant to be Nias2.28 and Aui4.52 its the lower and. upper lines weI.
The scatter is consistent with being independent of $\wQ$ and $\Mvir$ , and we find that taking the B01 proportionality constant to be $K_{\rm low}=2.28$ and $K_{\rm high}=4.52$ fits the lower and upper lines well.
These values correspond to 9,4,=0.18 dex and Chosetial0.11 dex.
These values correspond to $\sigc$$_{,\rm low}=0.18$ dex and $\sigc$$_{,\rm high}=0.11$ dex.
Although these are similar in magnitude to the scatter reporte in previous studies. our clistribuions are skeweel away from log-normal toward lower concentrations.
Although these are similar in magnitude to the scatter reported in previous studies, our distributions are skewed away from log-normal toward lower concentrations.
We note that the skewness may likely be caused. by the lower resolution of our simulations. which tends to result in lower concentration haloes.
We note that the skewness may likely be caused by the lower resolution of our simulations, which tends to result in lower concentration haloes.
bor a Dixed mass the DOl model predicts. that concentration should decrease with redshift as L/(1|2).
For a fixed mass the B01 model predicts that concentration should decrease with redshift as $1/(1+z)$.
The haloes in our simulation also satisvo this relation. as shown bv Figure &.. in which we plot the redshift’ dependence of concentration for haloes of mass AM,=ilothtA.
The haloes in our simulation also satisfy this relation, as shown by Figure \ref{fig:cofz}, in which we plot the redshift dependence of concentration for haloes of mass $\Mvir=7 \times 10^{11} \hinv\msun$.
This figure shows that the concenrations follow the ecx(l|2) relation that is cmbocied in the BOL analytic model.
This figure shows that the concentrations follow the $c_{\rm vir} \propto (1+z)^{-1}$ relation that is embodied in the B01 analytic model.
At redshifts greater than 2.5. our catalogues of haloes in this mass bin with [itt«d NEW.profiles becomeincomplete.
At redshifts greater than $\sim 2.5$ , our catalogues of haloes in this mass bin with fitted NFWprofiles becomeincomplete.
This incompleteness oeferentiallv allects low concentrations haloes. causing the Coo2) relation to Hatten
This incompleteness preferentially affects low concentrations haloes, causing the $\cvir(z)$ relation to flatten
in the spaces between the IFU fibers.
in the spaces between the IFU fibers.
Three of the fields had four sets of dithers. but the July observing rum was less successful due to weather aud ouly a single dither set was taken for the final field.
Three of the fields had four sets of dithers, but the July observing run was less successful due to weather and only a single dither set was taken for the final field.
A custom software pipeline developed for the iustrunient (Adams 2010. in preparation) was used in all reductions.
A custom software pipeline developed for the instrument (Adams 2010, in preparation) was used in all reductions.
The notable features are that no interpolation was performed on the data to avoid correlated noise and that background subtraction was done between differently sampled fibers bv fitting B-spline 1nodols (Dierckx 1993)) in a runuing boxcar of 3l chip-adjaceut fibers iu a manner simular to optimal loneslit skv subtraction methods (Ikelsoun. 2003)).
The notable features are that no interpolation was performed on the data to avoid correlated noise and that background subtraction was done between differently sampled fibers by fitting B-spline models (Dierckx ) in a running boxcar of 31 chip-adjacent fibers in a manner similar to optimal longslit sky subtraction methods (Kelson ).
The NGC 5166 pointing was sparse enough to allow selfskv-ubtraction.
The NGC 5466 pointing was sparse enough to allow self-sky-subtraction.
The spectra were fux calibrated by observing a single flux standard on each night using the same dithering pattern used for the science targets.
The spectra were flux calibrated by observing a single flux standard on each night using the same dithering pattern used for the science targets.
The fiux solution was then applied to every fiber ii every exposure for that night.
The flux solution was then applied to every fiber in every exposure for that night.
Based ou many tens of flux staudards taken during other VIRUS-P projects the absolute flus calibration precision is and wavelength iudepeudeut for full dither sets
Based on many tens of flux standards taken during other VIRUS-P projects the absolute flux calibration precision is and wavelength independent for full dither sets.
The data reduction code does not correct for differeutial atmospheric refraction (DAR) because there is very Little iupact from DAR with these larec fibers when thev are stunned over full dither sets
The data reduction code does not correct for differential atmospheric refraction (DAR) because there is very little impact from DAR with these large fibers when they are summed over full dither sets.
Fiber spectra were then combined using the IRAF taskseomdbine to form a single spectrum for cach star.
Fiber spectra were then combined using the IRAF task to form a single spectrum for each star.
For some stars. this was done by combining a suele fiber frou all visits (usually four) to a specific dither position: for other stars. those falling near the edge of a fiber. it nav have required combining differcut fibers on differcut dithers (perhaps & or 12 spectra).
For some stars, this was done by combining a single fiber from all visits (usually four) to a specific dither position; for other stars, those falling near the edge of a fiber, it may have required combining different fibers on different dithers (perhaps 8 or 12 spectra).
Figure l is a color-magnitude diagram for NGC 5166. with different svinbols denoting various parts of our data set: stars uot observed im this particular VIRUS-P pointing are simall points. and the spectroscopic BHuple (inblended asviuptotie giant branch (ACB) aud RGB stars with high enough signal-to-noise for iudex nuüeasuremenuts) are shown as filled squares.
Figure \ref{fig1} is a color-magnitude diagram for NGC 5466, with different symbols denoting various parts of our data set: stars not observed in this particular VIRUS-P pointing are small points, and the spectroscopic sample (unblended asymptotic giant branch (AGB) and RGB stars with high enough signal-to-noise for index measurements) are shown as filled squares.
Our staudared for a spectroscopically unblended star is that there are uo other stars within 5 arcsec down to 3 mag below the sample stars D imaguitude.
Our standard for a spectroscopically unblended star is that there are no other stars within 5 arcsec down to 3 mag below the sample star's B magnitude.
Open squares represent stars with sufficient signal-to-noise (S/N) for spectroscopic analysis but possible spectroscopic blends. aud open triangles represent IID stars and RGD stars with low S/N. One strone-lined radial velocity nommember (star 1980) is shown as an inverted triangle.
Open squares represent stars with sufficient signal-to-noise (S/N) for spectroscopic analysis but possible spectroscopic blends, and open triangles represent HB stars and RGB stars with low S/N. One strong-lined radial velocity nonmember (star 1980) is shown as an inverted triangle.
The two filled squares outlined with laree circles are stars 1398 aud 2[83. and are likely to he ACB stars based on their positions to the blue of the eiaut branch.
The two filled squares outlined with large circles are stars 1398 and 2483, and are likely to be AGB stars based on their positions to the blue of the giant branch.
The filled square outlined with a large star is star 1839. which las strong CN and CII features and is likely a CIT star.
The filled square outlined with a large star is star 1839, which has strong CN and CH features and is likely a CH star.
Although L839 lies on the elaut branch in the (WW 7) colormagnitude diagram. the strong absorption iu the UV. due to extra CN opacity aud the Doud-Neff effect (Boud Neff 1969)). males its (2 V) color distinctly red relative to theelaut branch.
Although 1839 lies on the giant branch in the $V, V-I$ ) color-magnitude diagram, the strong absorption in the UV, due to extra CN opacity and the Bond-Neff effect (Bond Neff ), makes its $B-V$ ) color distinctly red relative to thegiant branch.
The distance modulus aud reddening for NGC 5166 are (Qm Af),=16.15 aud E(B W)=0.023 (Dotter et al. 2010)).
The distance modulus and reddening for NGC 5466 are $m-M$ $_{V}=16.15$ and $B-V$ )=0.023 (Dotter et al. ).
Four typical spectra are shown in Figure 2.. to provide a sense for the quality and characteristics of the data.
Four typical spectra are shown in Figure \ref{fig2}, , to provide a sense for the quality and characteristics of the data.
There does seem {ο be a smaller velocity dispersion larther [rom the giants. but it’s nol even of one-siema signilicance outside the Local Group. and not much better within. (
There does seem to be a smaller velocity dispersion farther from the giants, but it's not even of one-sigma significance outside the Local Group, and not much better within. (
These should probably be taken as upper limits on the significance. since the satellite ealaxy dispersions are not really Gaussian.)
These should probably be taken as upper limits on the significance, since the satellite galaxy dispersions are not really Gaussian.)
How much should we see?
How much should we see?
Pradaetal.(2003). examined isolated. giant ancl satellite galaxy pairs in the SDSS data. caleulating the line of sieht velocity dispersion for 3000 satellites al various distances from the primary.
\citet{PVK03} examined isolated giant and satellite galaxy pairs in the SDSS data, calculating the line of sight velocity dispersion for 3000 satellites at various distances from the primary.
For an L* galaxy. (which is comparable to those examined above). the dispersion fell from 120 km tat 20 kpe to 60 km Fat 350 kpe.
For an L* galaxy (which is comparable to those examined above), the dispersion fell from 120 km $^{-1}$ at 20 kpc to 60 km $^{-1}$ at 350 kpc.
At this rate the dispersion would be a minor part of the dispersion as observed in the Local Volume al megaparsec scales.
At this rate the giant-satellite dispersion would be a minor part of the dispersion as observed in the Local Volume at megaparsec scales.
In all. it is unlikely that we have observed any inlluence of large galaxies on the velocity dispersion of satellites between 1. and 2 Mpe and much more likely that what has been observed is due to some other οδοί, (
In all, it is unlikely that we have observed any influence of large galaxies on the velocity dispersion of satellites between 1 and 2 Mpc and much more likely that what has been observed is due to some other effect. (
1 we separate the sample from bevond the Local Group into "behind and before" and "beside". (he inner aud outer "beside" samples have the same dispersion to within one km +: which. at 120 km !.ds larger (han the "before and behind” inner sample.
If we separate the sample from beyond the Local Group into “behind and before” and “beside”, the inner and outer “beside” samples have the same dispersion to within one km $^{-1}$; which, at 120 km $^{-1}$, is than the “before and behind” inner sample.
Bul selecting data in (his wav can have a pernicious elect on statistics. and at the very least increases the chance that small-number fIuctuations will give misleacling results.)
But selecting data in this way can have a pernicious effect on statistics, and at the very least increases the chance that small-number fluctuations will give misleading results.)