A&A 403, 93-103 (2003)
DOI: 10.1051/0004-6361:20030282
W. Nowotny1 - F. Kerschbaum1 - H. Olofsson2 - H. E. Schwarz3
1 - Institut für Astronomie der Universität Wien,
Türkenschanzstraße 17, 1180 Wien, Austria
2 - SCFAB, Stockholm Observatory / Department of Astronomy,
106 91 Stockholm, Sweden
3 - Cerro Tololo Inter-American Observatory, Casilla 603, La Serena,
Chile
Received 13 December 2002 / Accepted 25 February 2003
Abstract
We present results of our ongoing photometric survey of Local Group
galaxies, using a four filter technique based on the method of Wing
(1971) to identify and characterise the late-type stellar content. Two
narrow band filters centred on spectral features of TiO and CN allow us to
distinguish between AGB stars of different chemistries [M-type (O-rich) and
C-type (C-rich)]. The major parts of two dwarf galaxies of the M 31 subgroup
- NGC 185 and NGC 147 - were observed. From photometry in V and i we
estimate the tip of the RGB, and derive distance moduli
respectively. With additional photometric data in the narrow band filters
TiO and CN we identify 154 new AGB carbon stars in NGC 185 and 146 in
NGC 147. C/M ratios are derived, as well as mean absolute magnitudes
Mi
,
bolometric magnitudes
,
luminosity functions,
and the spatial/radial distributions of the C stars in both galaxies.
Key words: stars: AGB and post-AGB - stars: carbon - galaxies: Local Group - galaxies: individual: NGC 185, NGC 147 - surveys
Our first paper on this topic, Nowotny et al. (2001, in the following called Paper I), presented a search for Asymptotic Giant Branch (AGB) stars in M 31, and gave a motivation as to why it is important to study such stars in galaxies of the Local Group. Here we briefly summarise the most important points.
AGB stars are evolved stars with low to intermediate masses
( 0.8-8
), which have already experienced core
helium-burning and appear as late-type giants in colour-magnitude diagrams
(CMDs) of stellar populations. They are characterised by long-period
pulsations, strong stellar winds (leading to high mass loss rates), and a
change of their chemical atmospheric compositions from O-rich to C-rich
(spectral type M
C) under certain circumstances (thermal pulses,
convection) during their evolution. Because of their intrinsically high
luminosities of up to a few
,
and their enormous sizes, AGB
stars are among the brightest and coolest (reddest) stars in old stellar
populations. They contribute significantly to the integrated light of distant
galaxies, especially in the NIR (Mouhcine & Lançon 2002).
Within the galaxies of the Local Group that can be resolved into single stars
due to their proximity, they form an important subsample and can easily be
detected individually. On the one hand these AGB stars can give us clues about
the populations (extensions, distance estimates, and kinematics of
extragalactic systems, etc.), and on the other hand, by observing whole
samples of stars in extragalactic systems rather than single galactic
stars, one has the advantage of large numbers and the known properties of the
system (type, distance, metallicity) to get information about global AGB
properties and evolution. Large samples of extragalactic AGB stars are needed
for this in general.
Of special interest are AGB stars of spectral type C (that is stars with an
abundance of C higher than that of O, C/O 1), which are located close
to the tip of the AGB and have very characteristic spectra (dominated by carbon
species, such as CN), that differ strongly from M-type spectra (stars with an
abundance of C lower than that of O, C/O
1, and hence their spectra
are dominated by oxides, such as TiO). These very bright and red stars are
typical representatives of the intermediate-age population, and can serve as
standard candles because of their narrow luminosity function or can be used as
probes of galaxy dynamics (Dejonghe & van Caelenberg 1999), as has
been done e.g. for the Magellanic Clouds (van der Marel et al. 2002;
Hatzidimitriou et al. 1997 or Kunkel et al. 1997). Also,
they can be used as a tracer for the extent of a galaxy, as no contamination by
foreground stars is possible as for extragalactic M giants by galactic M
dwarfs; an important advantage. The evolution of carbon star properties and
their role in whole populations has recently been investigated by Mouhcine &
Lancon (2003). As these C stars have approximately the same magnitude
and colour as the much more numerous M stars, one needs some additional (at
least low-resolution) spectral information to separate the two groups.
Various methods have been used to search for AGB stars in extragalactic systems. Out of these, narrow band photometry proved to be most powerful for an efficient identification and characterisation of AGB stars even in distant crowded fields (Paper I). Our method uses conventional Vi filters as an indicator of temperature or spectral type, whereas two narrow band Wing-type filters TiO and CN, centred on characteristic molecular features (around 800 nm) of O-rich and C-rich objects, respectively, provide low-resolution spectroscopic information to distinguish M from C stars. For further details about this method and a synthetic colour-colour diagram, the reader is referred to Paper I (Sect. 2) or Nowotny & Kerschbaum (2002).
Earlier work by e.g. Cook et al. (1986) showed that the method works
well, but the number of identified AGB stars was limited by the then available
small detectors. For global information on galaxies or AGB properties one needs
statistically meaningful star numbers. Apart from a few well studied galaxies,
information is incomplete or lacking for many of the systems of the Local
Group, see
Groenewegen (1999, 2002) or Mateo (1998). With
the advent of new technologies (large CCD detectors, mosaics, wide field focal
reducer instruments) and therefore larger fields of view (FOV), the Wing-filter
method was applied very successfully and comprehensively e.g. by Brewer et al.
(1995) and Letarte et al. (2002).
We have continued our investigations with two M 31 companions, NGC 185 and
NGC 147, using the same instrument. These two dwarf spheroidals were
discovered by W. Herschel on November 30, 1787, and by J. Herschel in the
1820s. Baade (1944) recognised them for the first time as members of
the Local Group, when he could resolve their stars. Being separated by only
1
(Fig. 12.2., van den Bergh 2000), they probably
form a gravitationally bound pair on the front side of the M 31 subgroup of
the Local Group (van den Bergh 1998).
The following values were measured for general parameters of NGC 185 and
NGC 147: tidal radii of 16
and 20
,
core radii of
60
and 67
,
ellipticities of 0.22 and 0.44,
respectively. References can be found in van den Bergh (2000) and
Mateo (1998). Total magnitudes for the galaxies can be derived from
their observed surface brightness profiles, which was done by Price
(1985), Kent (1987), Caldwell et al. (1992), and
Kim (1998). Using the values for V, (m-M), and EB-Vfrom Mateo (1998), one can derive absolute visual magnitudes
MV of -15.46 and -15.51 for NGC 185 and NGC 147, respectively.
Detailed population studies of NGC 185 have been carried out by Lee et al.
(1993b), Martínez-Delgado & Aparicio (1998), and
Martínez-Delgado et al. (1999). A number of stars above the RGB
indicate a significant intermediate-age stellar population, the AGB-stars
extend up to
-5.0
and an AGB luminosity
function (LF) based on about 1300 stars is given. The colours are distinctly
bluer in the central part. The stellar content of NGC 147 was investigated by
Davidge (1994) and Han et al. (1997). The existence of
RR Lyrae stars shows the presence of an old stellar population of about
10 Gyr, while the found AGB-stars (the tip occurs at
-5.0
)
suggest a stellar population of
5 Gyr age. The younger population seems to be more concentrated
toward the centre than the old population.
A few methods (tip of the RGB, HB stars brightness, and RR Lyrae stars) have
been used to derive distances to the two extragalactic systems. An overview and
references can be found in Tables 12.3 and 12.5 of van den Bergh
(2000). Following the review by Mateo (1998), we used
distance moduli of
and
for NGC 185 and NGC 147, respectively, which were derived from the data of
Lee et al. (1993b) and Han et al. (1997) by using the
magnitude of HB stars, and the I0 magnitude of the tip of the RGB (Lee et al. 1993a). We could confirm the latter with our own photometric
data (Sect. 4.1).
The median colour of RGB stars was used to derive metallicities for NGC 185.
Lee et al. (1993b) found a mean value of
dex. Taking into account their smaller FOV, this
means that their metallicity estimate is consistent with the values found by
Martínez-Delgado & Aparicio (1998) in the corresponding area.
In addition, the latter derived decreasing metallicities for regions further
away from the galaxy centre - down to -1.76 for the outermost parts -
resulting in an average of
dex. Based on the results of
Martínez-Delgado & Aparicio (1998), we expect values around
-1.2 for the area covered by our frames. The same method was used for
NGC 147, by Han et al. (1997), resulting in a mean
[Fe/H] =
dex for the centre field and
[Fe/H] =
dex for the off-centre field (values of -0.85 for the
outermost parts of their field). We adopted a mean metallicity of -1, which
is also reported by Davidge (1994). Despite the not very well
determined metallicities and their dispersions, NGC 185 can be regarded as the
metal-poorer system of the two.
An interesting aspect of the two galaxies is that, while we can observe young stars, a significant amount of interstellar gas, and prominent dust patches in NGC 185, NGC 147 appears to be free of gas and dust. Also, NGC 185 is more concentrated in the central region, as can be seen in Figs. 1 and 6.
Table 1: Basic properties of the two galaxies surveyed, taken from: NASA Extragalactic Database (NED) - 1, van den Bergh (2000) - 2, Lee et al. (1993b) - 3, Han et al. (1997) - 4, and Mateo (1998) - 5.
Table 1 summarises the properties of the two extragalactic systems. For further details one is referred to the review by Mateo (1998), and to van den Bergh (2000). The two galaxies NGC 185 and NGC 147 are quite similar, almost like twins in terms of their stellar content (see the CMDs in Fig. 2). They have almost the same Star Formation History (SFH), Fig. 8 in Mateo (1998). A more detailed comparison between them is worthwhile because they have approximately the same absolute magnitude of -15.5
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Figure 1:
i-band CCD images of the two M 31 companions NGC 185 (upper)
and NGC 147 (lower), obtained at the 2.56 m Nordic Optical
Telescope. The field covers
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The observed CMDs of NGC 185 and NGC 147 (Martínez-Delgado & Aparicio
1998; Han et al. 1997) look promising for a further
investigation of their AGB populations using the Wing-filter method. In both
of them a number of AGB stars are found to be as red as
(V-i)0 4
.
Until now, no C stars have been
identified (Groenewegen 1999; Table 8 of Mateo 1998).
The observations were obtained with the 2.56 m Nordic Optical Telescope (NOT,
see http://www.not.iac.es) at La Palma (Spain) during the nights from August 30
to September 1, 2000. Using the ALFOSC focal reducer instrument with a
2k2k Loral-Lesser thinned CCD, we observed fields centred on NGC 185
and NGC 147 in the four filters described in Table 2. Also
listed are the chosen exposure times. All three nights were photometric, the
seeing varied between 0
7 and 1
1 on the combined frames. The two
fields were observed at air masses of between 1.1 and 1.8. For the photometric
calibration in V and i we additionally observed the Selected
Area Fields #92, #110, and #114 from Landolt (1992). Sky
flats for each filter were obtained, as were bias images. Figure 1
shows the i-filter images. With a pixel scale of 0
189/pixel
the CCD field is
,
i.e., the major
fractions of the galaxies could be observed on one frame. Although this will
include the majority of the stars, our survey is not totally complete in terms
of area. Approximate sizes of the galaxies are given in Sect. 2
(or compare with the surface brightness profiles of Kent 1987). In
addition, there are some restrictions due to vignetting and pairing, as
described in more detail in Sect. 3.3. The C stars found can be
used to give a rough estimate of the extent of the galaxies
(Sect. 4.4).
Table 2: Observing log, the exposure times were the same for both galaxies.
For the basic reduction of the frames we used MIDAS (bias subtraction, cosmic ray removal, flat-fielding with sky flats, matching of frames, adding). The narrow band filters were used in the converging beam of the telescope to avoid differential wavelength shifts within the field due to the large angles of incidence in parallel beam instruments such as ALFOSC, which resulted in some vignetting in the outermost parts of the TiO and CN frames (Fig. 6). Photometry was done for all stars, the ones lying in the affected areas, were sorted out afterwards. For the photometry of all four added frames we used the PSF-fitting software written by Ch. Alard for the data reduction of the DENIS-project. This program calculates a number of quality flags (correlation coefficient of the fit "C1'', "error'', etc.), which can be used to sort out data of low quality. The photometry of the stars had to meet the following criteria to be considered good enough: C1 > 0.7, error < 0.5. Stars meeting weaker criteria only blur the CMDs, and none of the bright AGB stars will be lost by excluding these stars.
To measure the atmospheric extinction at La Palma, we observed Landolt's field
SA 114 in a sequence of different airmasses in all four filters. The
resulting extinction coefficients (
kV=0.174,
ki=0.067,
and
)
were taken into account.
The correction for interstellar reddening in V and i was done
according to the values from the NASA Extragalactic Database (NGC 185:
,
;
NGC 147:
,
).
The effect of reddening on (TiO-CN) is very small, due to the negligible
difference in central wavelength of these two filters (Battinelli & Demers
2000 found
).
Therefore, no correction was applied to the data.
Using several unsaturated but bright single stars with their neighbours subtracted from the frame, we established aperture corrections, between the PSF fitting magnitudes and the aperture photometry magnitudes for V and i. These corrections were applied to all stars.
Photometric zero-points for the filters V and i were obtained
from stars of the standard fields SA 92, 110, and 114 of Landolt
(1992). We used these, as we did in Paper I, because of the
similarity of Landolt's Cousins i-filter (
Å and
Å) and our i-filter (see Table 2).
For the two narrow band filters no absolute photometric calibration was done.
To be compatible with Paper I, we used the fact, that "early'' spectral types
lack TiO/CN-features, and therefore have (TiO-CN)
.
Stars having (V-i)0 < 0.7, C1 > 0.8,
error < 0.1, and only one counterpart in all frames during the
pairing process, were considered to be good enough for determining the special
offset for (TiO-CN)0. 31 stars in NGC 185 and 58 in NGC 147 led to
offsets of -0.079
0.05
and -0.076
0.05
,
respectively. These offsets were applied to all stars.
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Figure 2:
Colour-magnitude diagrams for all stars having good photometry in
V and i (sample 1; mean 1![]() |
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All single frames were shifted and added, the frames of different filters were
shifted to match each other (i as reference), which resulted in small losses
of covered area. The pairing of stars, detected in different filters, was done
with the DENIS-software "Cross_Colour''. All stars within a radius of 3
pixels (0
6) were considered to correspond to each other.
Only stars with one and only one definitive counterpart in all four filters
were used in the further analysis. No multiple matches of other nearby pairing
candidates were accepted. Using stars from the Guide Star Catalogue (GSC I),
that appear on our frames, and the MIDAS/ASTROMET-package we produced an
absolute astrometric calibration. This was then used to derive absolute
coordinates for identified stars or to calculate positions for the galaxy
centres in pixel coordinates. By comparing different methods to determine the
centres on the i frames, such as values from the literature, surface
brightness plots, average coordinates of all stars or C stars etc., we found
them to differ by only about 10
,
which we considered as
consistent. In the following we used the centre coordinates from Cotton et al.
(1999). Stars lying in areas on the TiO and CN frames that were
affected by vignetting had to be identified and removed. For this, we measured
the centres and radii of the vignetted areas on original frames of both filters
and transformed it to the i reference coordinates.
This resulted in a total of 26 496 stars for NGC 185 and 18 300 stars for
NGC 147, having good photometry in V and i (sample
1). The maximum distance of a star in this sample to the centre of the galaxy
is 4
6 for both galaxies.
Having detections in the narrow band filters results in an even stronger
criterium. Despite the longer exposure times, also here mainly the bright red
giants were detected. But as we are interested in the AGB population, which is
made up of the brightest stars on the RGB/AGB sequence, this is no major
drawback. 8546 stars for NGC 185 and 6332 stars for NGC 147 have photometry
of good quality in all four filters (sample 2). The maximum central
distance for this sample is 3
7 for NGC 185 and 3
8 for NGC 147.
Coordinates and photometric properties for all stars (sample 1) are available in electronic form at the CDS, a short extract of the results for NGC 185 is shown in Table A.1.
A direct estimate of the foreground contamination from our frames seems difficult as the galaxy sizes are a bit larger than our FOV (Sect. 4.4) and no extra field off the galaxies could be observed. As far as the C stars are concerned, the contamination doesn't matter at all, but it may be quite important for the M giants. Galactic M dwarfs in the directions of the two galaxies can affect the sample significantly, see e.g. Albert et al. (2000). But from the low numbers in the outermost bins of Fig. 7, we assume no huge contamination in our case, for further discussions see Sect. 4.4.
As a check of the reliability of our photometry, we measured the tip of the RGB
from the photometric data. Figure 3 shows the i0 luminosity
function for all stars of sample 1 for the two galaxies. As
edge-detection algorithm we used (like Lee et al. 1993a) a zero-sum
Sobel kernel [-2, -1, 0, 1, 2], the resulting histogram is also plotted in
Fig. 3. To avoid "binning noise'', we calculated i0 for the
tip-RGB as weighted mean of a few bins (with significantly large star counts)
around the maximum bin. We obtained 19.96
for NGC 185 and
20.36
for NGC 147, which is in good agreement with the values of
19.92
0.1
by Lee et al. (1993b) and
20.3
0.04
by Han et al. (1997). This gives us
confidence in the correctness of our photometry. In Fig. 2, triangles
mark the corresponding RGB-tips.
Following the (tip-RGB) method as described in Lee et al. (1993a), we
calculated mean colour indices (V-i)0 for the RGB and derived
bolometric corrections BCI and metallicities [Fe/H] from it. Despite the
good agreement in i0 for the tip-RGB, we find considerable differences in
the determination of [Fe/H]. Our (V-i)-3.5 result in metallicities of -0.89
for NGC 185 and -1.11 for NGC 147, which differ from the values in the
literature as given in Sect. 2. Finally, we calculated distance
moduli of 24.04
for NGC 185 and 24.44
for NGC 147.
As the discrepant metallicities have only a weak influence on the determination
of the distance moduli (m-M)0 we derive from our photometry, the
latter still are in good agreement with the ones from the literature.
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Figure 3: i0-LF for all stars of sample 1 for both galaxies. Overplotted is the resulting histogram of the edge-detection algorithm (zero-sum Sobel kernel), from which we calculated the magnitude of the RGB-tip. |
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Figure 4 (left) shows colour-colour diagrams for all stars of
sample 2. Two well separated branches can be seen toward redder
colours. M and C stars differ in (TiO-CN)0 by up to 1.5
.
The identification of the different chemistries is based on the location of the
stars in this colour-colour diagram, the selection areas (also plotted in
Fig. 4) were defined as follows:
M: (V-i)0 > 1.6
,
(TiO-CN)0 > 0.15
,
C: (V-i)0 > 1.16
,
(TiO-CN)0 < -0.3
For M stars, the selection criteria in both colours were kept unchanged from those in Paper I (Sect. 4.2. and Fig. 2 in Paper I), where they were defined by the synthetic photometry of an M0 star. We also kept the same border in (TiO-CN)0 for the C stars.
For two reasons, we changed the border in (V-i)0 for C stars. The
bluest C star in Fig. 2 of Paper I was approximately of spectral type C5.
Doing the same synthetic photometry for another spectral library, that of Silva
& Cornell (1992), we found their C0 star to have
(V-i)0 = 1.16
.
Figure 4 suggests
that there are good candidates for being a C star, i.e., with a clear
distinction in (TiO-CN)0, bluer than (V-i)0 = 1.6
. Taking into account these facts, we adopted the border in (V-i)0 for C stars to -1.16.
There may be some S type star candidates among the reddest stars, as was found
by Letarte et al. (2002) and spectroscopically confirmed by Brewer et al. (1996). While S stars still show TiO bands and will therefore lie
in the same area as the M stars, SC stars (C/O = 1) with no prominent feature
within the filters TiO and CN are expected to have
(TiO-CN)
and should lie between the two
selection areas (Nowotny & Kerschbaum 2002).
Using our criteria, we found 154 new C star candidates in NGC 185 and 146 in
NGC 147. They can be found above, as well as below, the tip of the RGB and are
- with a few exceptions - among the brightest red giants of the whole
population, as can be seen in Fig. 4 (right). 1732 M stars lie in
the corresponding area for NGC 185 and 950 for NGC 147. This leads to C/M
ratios for the observed fields of 0.089 and 0.154,
respectively. In general, one would expect
NGC 185 to have more C stars (larger C/M ratios), as it is the more metal-poor
system, has an almost identical SFH to NGC 147, and the limiting magnitudes
for both galaxies are roughly the same. This is not the case and the ratios are
contrary to the expectation. Note, that one would not face this discrepancy
with the metallicities we derived in Sect. 4.1. For a comparison with
Groenewegen (2002) we also calculated the ratio C/M5+, where
we counted all stars later than spectral type M5, using our synthetic
photometry to set the (V-i)0 limit (Paper I; we used the same
spectral library of Fluks et al. 1994). Again we find a
larger ratio for NGC 147, Table 3. Both ratios agree
within the uncertainties of the method with Fig. 5 of Groenewegen
(2002).
Absolute star numbers, C/M ratios, and global statements should be considered with caution, as they depend on the selection criteria (especially dramatic for the M star candidates), foreground contamination for M stars, photometric variability of AGB stars, and the limiting magnitude of the observations (Brewer et al. 1995). For a quantitative comparison of star counts for different galaxies, the selection should be done consistently.
Based on the observations of other Local Group galaxies, see Fig. 2 of Groenewegen (1999) and Fig. 6 of Battinelli & Demers (2000), one would expect more C stars to be found in NGC 185 and NGC 147. The limited field is not the main reason why we detect relatively few of them. As there are only very few stars in the outermost bins of Fig. 7, we do not expect to miss many C stars. The selection in quality of the photometry can not have a large influence, as C stars are bright and easily detectable.
Table 3: Star counts for sample 2 of both galaxies and the mean properties of the C star candidates identified in Fig. 4.
In Paper I, we found 515 red stars, that were not detected in V because of the short exposure time (Sect. 4.5 in Paper I). For these we plotted a special colour-magnitude diagram (Fig. 8 in Paper I), which allowed us to sort out some more C star candidates. For the observations of NGC 185 and NGC 147 we reached a fainter limiting magnitude in V and did not miss this kind of stars. As a check, we plotted the same special CMD for the stars of NGC 185 and NGC 147. The C star candidates found in this way are almost identical to the ones found using the colour-colour diagram of Fig. 4. This confirms our conclusions in Sect. 4.5 of Paper I.
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Figure 4: Left: colour-colour diagrams for all stars of sample 2. Also plotted are the selection areas for M and C stars, as described in the text. Right: corresponding colour-magnitude diagrams for all stars of sample 2. The identified C stars from the selection areas, marked with open circles, lie on the uppermost parts of the RGB/AGB-sequences. |
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Figure 5:
Left: luminosity function (apparent and absolute i magnitude) for all C stars, identified in Fig. 4, of
NGC 185 and NGC 147. Using the distance moduli from
Table 1, one can calculate a ![]() ![]() ![]() ![]() ![]() ![]() ![]() ![]() |
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The influence of the chosen criteria for C stars on the shape of the LF was already mentioned in Sect. 4.2, Fig. 6 of Letarte (2002). This can introduce an additional error to the derived distances. By selecting only the "redder'' more luminous C stars, the LF becomes better defined.
The two diagrams on the left in Fig. 5 show the i0-LFs
for all C stars in the selection areas of Fig. 4. Their
mean photometric properties can be found in Table 3.
The upper axes show our apparent magnitudes, the dotted line marks the
RGB-tip as it was found in Sect. 4.1. The lower axes shows
the corresponding absolute magnitudes calculated with the distance
moduli from Table 1. In good agreement is the absolute
magnitude of the RGB-tip for both galaxies. Mean values
Mi
of all C stars were found to be
-4.24
0.46 for NGC 185 and -4.24
0.42 for NGC 147.
These values are low compared to the ones found for other nearby
galaxies, see Fig. 5 of Demers & Battinelli (2002).
Mi
was found to be -4.75
0.47
for the LMC by Richer (1981), -4.69
0.28
for IC 1613 by Albert et al. (2000) and -4.7
for NGC 6822 by Letarte et al. (2002). On the other hand,
the check with the RGB-tip in Sect. 4.1 assures us that we do not
have a large error in our i0-magnitudes.
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Figure 6: Distribution of all stars of sample 2 of NGC 185 (upper) and NGC 147 (lower) on the i-band frames, overplotted are the identified C stars and the galaxy centres (Sect. 3.3). One can see the effect of vignetting in the corners, the strips without stars at the lower ends are due to a shift of the telescope between observations of frames of different filters. |
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Figure 7: Radial distributions of all stars of sample 2 and for the found C stars. |
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Also plotted in Fig. 5 are the bolometric LFs for the C stars
of both galaxies. For the calculation of
,
we used a
different bolometric correction (BC) than the one from Bessell & Wood
(1984) that we used in Paper I. Note, that there can be
relatively large differences in the BCs of different authors,
especially for stars with (V-i)
,
which can be seen
in Fig. 11 of Montegriffo et al. (1998). The BC of
Montegriffo et al. (1998) is based on a larger set of
photometric data than that of Bessell & Wood (1984). The
BC which they designate "metal-rich'', fits the expected
metallicities for our observations better and reaches values of
(V-i)
.
As it is derived from photometry
of globular clusters, it can be used for our M type giants
. The
clear deviations of the very few C stars in Fig. 1 of Bessell & Wood
(1984) suggest a different BC for C giants.
Bolometric corrections for C type giants have been published by Frogel
et al. (1980), Costa & Frogel (1996), Groenewegen
(1997), and Bergeat et al. (2002), but not for the colour
index (V-i)0. From the collection of data on galactic C stars of
Bergeat et al. (2002) a BC based on (V-i)0 was derived.
Bergeat (priv. comm.) kindly provided us with data containing stars of their
groups HC5 and CV1-CV7, which are variable carbon giants, and typical members
of the thin disk. Although there remain uncertainties, we calculated the
BCV by a polynomial fit. With the distance moduli from Table 1,
the bolometric magnitude can be calculated in the following way:
As expected, this results in luminosities considerably different from those derived by adopting the M star BC.
Although our i-magnitudes appear to be somewhat fainter than
expected, the bolometric LFs for our C stars are in good agreement
with the one for the SMC (comparable [Fe/H] and statistically
meaningful star numbers) as given by Groenewegen (1999) in
his Fig. 5. The same goes for the mean
,
which also fit well into his Fig. 6. A possible
influence of the metallicity on
cannot
be deduced, considering the large uncertainties in the BC. From the
differences of LFs for different galaxies (Groenewegen 2002)
it seems that the LF and the corresponding
depend on metallicity and the SFH. Due to a large
scatter and uncertainties for the low-luminosity end of the LF, it
appears to be more suited for a rough estimation of the distance than
as a standard candle.
Figure 6 shows the distributions of detected sample 2 stars on the i-frames, marked are the identified C star candidates and the galaxy centres (derived as described in Sect. 3.3). The influence of vignetting of the TiO- and CN-frames can be seen for both groups, as can the concentration of stars toward the centre (the slightly inhomogeneous distribution of C stars in NGC 147 can probably be explained by small number statistics). The galaxy centres from Cotton et al. (1999) fit well. NGC 185 is more highly concentrated toward the centre.
Figure 7 shows the radial distributions of star counts for all stars and for the selected C stars in our FOV, where the number density is scaled by the area and total number. For the sake of simplicity, the star counts where made in circular annuli around the centre.
As it was stated by Albert et al. (2000), see their Figs. 6 and 8, while the distribution of M stars can be severely contaminated by foreground M dwarfs, the "clean'' C star distribution is a good measure of the extent of extragalactic systems. As there are C stars all over the frames in Fig. 6, our FOV cannot cover the whole area of the galaxies on the sky, which can also be concluded from a simple visual inspection of the frames. A quanitative foreground star count from the outer regions of the frames is not possible.
From the very small numbers (on the order of 10 stars per square arcmin)
in the outermost bins we conclude that we detected the majority of the (AGB)
stars of these galaxies. Only few such stars lie outside our FOV, because the
stellar density decreases strongly (compare the luminosity profiles of Kent
1987). This number densities can also be regarded as a crude estimate
for the foreground contamination and align well with the values given by
Ratnatunga & Bahcall (1985), which find 11 stars per square
arcmin for NGC 147 and for V down to 25
(our detection limit).
The distributions for all stars and for only C stars are almost the same for
all galactocentric distances. The distribution for
M stars is probably also similar. Therefore, the C/M ratio will be
approximately constant as a function of distance from the centre for both
galaxies.
The results from our four-colour CCD photometric observations of NGC 185 and NGC 147 are:
Acknowledgements
The authors wish to thank J. Bergeat very much for extracting the bolometric correction for C stars from his database and also for fruitful discussions on the topic. This work was supported by the "Fonds zur Förderung der Wissenschaftlichen Forschung'' under project numbers S7308-AST and P14365-PHY and the Austrian Federal Ministry of Transport, Innovation and Technology in the course of the Project FIRST-PACS. We used the Simbad database operated at CDS, Strasbourg, France.
Table A.1:
Example for the NGC 185 results, extracted from the full table, which
is available via CDS. Column 1 gives an ID number for every star, while Cols. 2 and 3 list the coordinates (J2000.0) in the form hhmmss.sss and
ddmmss.ss, respectively. A special flag F is given in Col. 4, which
denotes in which region of Fig. 4 the star falls ( c for
carbon stars, o for oxygen-rich stars, r for the rest), if it is
member of sample 2. If the flag F is u (for
unclassified), the star is from sample 1 and not included in the
smaller sample 2, which means it has only Vi-photometry. Columns
5-7 list the photometric results, i, V-i, TiO-CN.
The astrometric data may be afflicted with a systematic error of
0.5-1.2
introduced by the accuracy limits of the GSC reference stars
(Taff et al. 1990), the mean uncertainty of the MIDAS astrometric
solution is 0.05
rms (with a maximum of 0.11
).