Long-slit spectra centered on the 5170 Å Mg triplet of the galaxies
were taken along their major axes. Table 2 describes
the setups used during the 6 runs of observations at the
2.4 m telescope of the Michigan-Dartmouth-M.I.T. (MDM) observatory at
Kitt Peak, the 2.7 m telescope at Mc Donald observatory (McD) and
the 3.5 m
telescope of the German-Spanish Astronomical Center on Calar Alto (CA).
Table 3 gives the log of the observations. Depending on the
telescope
the seeing was either measured by a seeing monitor or by measuring the point
spread function (PSF) of a star image in the beginning of the night and
controlling it via the guiding probe. The seeing listed in
Table 3 is the worst value measured for the individual
exposures.
As we also
needed "velocity standard'' stars to determine the kinematical parameters of
the galaxies, we obtained spectra for several G9 to K1 - giant template
stars for all three setups.
To make sure that the illumination
across the slit was uniform and like in the galaxies
we trailed and wiggled the stars
across the slit. This ensured that the instrumental broadening was always
the same in the stars as in the galaxies. For the kinematic analysis itself we
finally only used spectra of HR 6817 (K1III) and HD 172401 (KOIII) both from
Gonzales (1993, G93).
Additionally we observed the flux standard star G191B2B or HD 19228
to calibrate the flux of the spectra before line strength indices were
derived.
The standard reduction (bias & dark subtraction,
flat fielding, wavelength calibration, cleaning for cosmics, sky subtraction,
correction for CCD misalignment) was carried out with our own programs running
under the image processing package MIDAS provided by ESO and which are
described in details in Bender et al. (1994, BSG94). Finally multiple
exposures for one galaxy were centered and summed up to increase the final
signal-to-noise.
The pixel-to-pixel noise in the normalized flat field is < 0.2%.
A sky subtraction better than 1% was achieved, while the errors in the
wavelength calibration were Å. The variation of the
spectral resolution with
was also
Å and stable during
each observing run (
Å).
The quality of the final spectrum
- and hence the resulting S/N - depends on the spectral resolution, the
efficiency of the spectrograph as well as on the seeing conditions.
Figure 4 shows examples of central spectra covering the 3 quality
classes listed in Table 3 for the different instrumental
resolutions
.
The spectra are averaged within the
radius
,
which
corresponds to a circular
standard aperture having a diameter of
(see Sect. 5).
For
= 67.9 km s-1 (run 2, 3, 5, 6),
the quality parameter is either 1 or 2, while for
=
129.4 km s-1 (run 1 & 4) it is only 2 or 3.
Because of the low instrumental resolution we excluded the latter ones from
having quality 1 by default.
In the next step the spectra were rebinned along the slit (spatial
direction) to achieve a minimum signal-to-noise (S/N)
40/Å at all
radii (for details of the procedure see BSG94).
Of course the central pixels have higher S/N (up to a factor of 5).
Monte Carlo simulations described by BSG94 showed that a minimum on
S/N = 30 - 40 is necessary to derive meaningful kinematical parameters with
negligible systematic errors.
As a final step of the pre-processing, the galaxy continuum was removed by
fitting a forth to sixth order polynomial.
We then determined the
line-of-sight-velocity-distributions (LOSVDs) by using the Fourier
Correlation Quotient (FCQ) method (Bender 1990), which provides
the stellar rotational velocities
,
velocity
dispersions
and first orders of asymmetric (H3) and
symmetric (H4) deviations of the LOSVDs from real Gaussian profiles
(van der Marel & Franx 1993; Gerhard 1993; BSG94).
As expected, we
find that the FCQ method is little influenced by template mismatching
(Bender 1990). Following BSG94, Monte Carlo
simulations were made to find the best fit-order for the continuum and to
check for systematic effects. It turns out that for runs 2, 3, 5, and 6, which
have high spectral resolution there are no systematic errors for all kinematic
parameters. Even for runs 1 & 4, that have rather low
spectral resolution (129.4 km s-1),
,
and H3 do not show
any systematic errors. Only the derived H4 amplitude for this setup
turns out to be
systematically negative for S/N
80. At S/N = 40 the systematic error
is -0.04 and hence is of the order of the error bars (see below).
We corrected for this systematic error using the dependence of H4 on S/N
delivered by the Monte Carlo simulations.
The error bars were derived from photon statistics and CCD read-out noise
and were calibrated via Monte Carlo simulations
described by Gerhard et al. (1998): Noise is added to template stars, broadened
following the observed values of H4 and
,
matching the power
spectrum noise to peak ratio of the galaxy. The accuracy of the estimated
error bars is about 20%.
Additionally we tested whether our estimated kinematic errors indeed
reproduce the RMS scatter between different exposures of the same galaxy.
For all setups the expected errors - derived via our Monte Carlo simulations
and shown in Fig. 6 -
are of the order of the RMS. In fact for most galaxies the errors were
slightly larger than the RMS by a factor of
1.2. In a few cases the
estimated errors were overestimated up to a factor of 2.
Mg, Fe and H
line strength indices were derived
following Faber et al. (1985) and Worthey (1992) from flux calibrated
spectra, rebinned radially as before. For the 11 galaxies
observed with the MDM setup (run 1 and 4), that covers a larger
range, we could in addition derive the NaD
line index profiles.
We corrected all measured indices for
velocity dispersion broadening and calibrated our measurements to the
Lick system using stars from Faber et al. (1985). The errors are
derived from photon statistics and CCD read-out noise.
As for the kinematic parameters we tested whether our estimated errors indeed
reproduce the RMS scatter between different exposures of the same galaxy.
Again, for all setups the estimated
errors - derived via our Monte Carlo simulations
and shown in Fig. 6 -
are of the order of the RMS. For most galaxies the errors were
slightly larger than the RMS by a factor of
1.2 and
in a few cases the estimated errors were overestimated up to a factor of 2.
In the following we indicate the
average Iron index with
Fe
(Fe5270+Fe
5335)/2 (Gorgas et al. 1990) and
the usual combined Magnesium-Iron index with
[MgFe
(G93).
Some of the line index profiles needed further treatment. At
small radii some of them show unreal features caused by the varying
focus of the spectrograph in the dispersion direction.
Following G93, this effect is
only detectable if the variation of the focus' point spread function
(PSF) is dominant compared to the atmospherical seeing (FWHM). For
some of the galaxies the focus variation was the dominant effect and
we followed the procedure described in Mehlert et al. (1998, MSBW98)
to correct it. The Mg1and Mg2 indices defined by Faber et al. (1985) are affected most,
because their pseudo continua are 200 Å apart from the
line windows. For all the other line indices the continuum and line
windows are close to each other. H
and the Fe indices are in
the red and blue part of the spectra and thus slightly affected, while
the Mgb index at the central wavelength is almost unaffected.
Recent investigations showed evidence for small amount of ionized gas and
dust in the interstellar medium (ISM) of many
elliptical galaxies (e.g. Bregman et al. 1992;
Goudfrooij et al. 1994).
Consequently, the detection of some emission (e.g. H,
N, and O) in
their spectra is expected. Since the velocities of the stars and gas in the galaxies may
differ by up to 100 km s-1 (Bertola et al. 1995) asymmetric
contamination of line indices is possible. In addition the dust and ionized
gas show a wide range of distribution: Smoothly along their major/minor
axis or in rather patchy or filamentary features (Goudfrooij et al. 1994).
If one wants to investigate the stellar population of elliptical galaxies
(as we are going to do in Paper II) the contamination of H
by emission
is especially crucial.
Emission will weaken the age indicating H
absorption index and hence
lead to an overestimate of the age of the dominant stellar population
(see Worthey 1994, W94). Simulations with the kinematical template stars
we used showed that we can
detect emission in H
for equivalent width larger
than about 0.3 Å. Figure 5
shows examples of the strong (a) and weak (b) H
emission
seen in the our spectra. We will refer to these
two classes in the comments on the galaxies given in Sect. 5.
In particular
Fig. 5a demonstrates how asymmetric H
emission can be. Since no OIII emission is detected in this galaxy
(GMP 4315) at all,
these data illustrate that it is impossible to correct for H
emission via OIII as suggested by G93.
Additionally Goudfrooij & Emsellem (1996) showed that NI emission may
influence the measured Mg indices (Mgb up to 0.5 Å and
Mg2 up to 0.03 mag). The influence of asymmetric NI emission can
be seen in the Mgb profile of GMP 2390
(galaxy 6 in Fig. 6).
Run | Date | Telescope | Detector | ![]() |
Scale | Slit- | spectral |
Spectrograph | [Å] | [''/pix] | width | resolution ![]() |
|||
1 | 3/95 | MDM | TI: 1024 ![]() |
4300-6540 | 0.777 | 1.7'' | 2.23 Å |
4 | 3/96 | 2.4 m | Mark III | 129.4 km s-1 | |||
2 |
4/95 | McD | TI: 800 ![]() |
4850-5560 | 0.635 | 2.5'' | 1.17 Å |
5 | 4/96 | 2.7 m | LCS | 67.9 km s-1 | |||
3 | 5/95 | CA | TI: 1024 ![]() |
4730-5700 | 0.896 | 3.6'' | 1.17 Å |
6 | 5/96 | 3.5 m | TWIN/R | 67.9 km s-1 |
No. | GMP No. | NGC/IC No. | DN | Run | Single exp. times | Total exp. times | PA | Seeing | Q |
[s] | [s] | (N ![]() |
(FWHM) | ||||||
1 | 3329 | NGC 4874 | 129 | 2 |
![]() |
10800 | 45![]() |
1.8'' | 2 |
2 | 2921 | NGC 4889 | 148 | 2 |
![]() |
9000 | 81![]() |
2.2'' | 2 |
3 | 4928 | NGC 4839 | 31 | 1 |
![]() |
7200 | 63![]() |
1.8'' | 3 |
2 |
![]() |
8750 | 63![]() |
2.0'' | 3 | ||||
3 | 1800 | 63![]() |
1 | ||||||
6 | 1800 | 3600 | 63![]() |
2.8'' | 1 | ||||
4 | 4822 | NGC 4841A | 240 | 1 |
![]() |
9600 | 105![]() |
1.5'' | 2 |
5 | 1750 | NGC 4926 | 49 | 1 | 3600 + 4500 + 2700 | 10800 | 60![]() |
1.8'' | 2 |
6 | 2390 | IC 4051 | 143 | 1 |
![]() |
10800 | 106![]() |
1.8'' | 2 |
7 | 2795 | NGC 4895 | 206 | 1 |
![]() |
10800 | 156![]() |
1.5'' | 2 |
8 | 3792 | NGC 4860 | 194 | 1 | 3600 + 1520 + 3000 + 2400 | 10520 | 127![]() |
1.4'' | 2 |
9 | 2629 | NGC 4896 | 232 | 1 |
![]() |
7![]() |
2 | ||
4 | 3000 + 2850 | 15060 | 7![]() |
1.8'' | 2 | ||||
10 | 3561 | NGC 4865 | 179 | 1 |
![]() |
10800 | 115![]() |
1.8'' | 2 |
11 | 2000 | NGC 4923 | 78 | 3 |
![]() |
7200 | 77![]() |
2.2'' | 1 |
12 | 2413 | - | 230 | 1 | 2700 + 2050 + 3000 | 25![]() |
2 | ||
4 |
![]() |
13750 | 25![]() |
1.8'' | 2 | ||||
13 | 4829 | NGC 4840 | 46 | 1 |
![]() |
9900 | 106![]() |
1.5'' | 2 |
14 | 3510 | NGC 4869 | 105 | 3 |
![]() |
7200 | 169![]() |
2.5'' | 1 |
15 | 2417 | NGC 4908 | 167 | 3 |
![]() |
7200 | 55![]() |
2.5'' | 1 |
16 | 2440 | IC 4045 | 168 | 3 | 2300 + 2500 + 2000 | 6800 | 108![]() |
2.2'' | 1 |
17 | 3414 | NGC 4871 | 131 | 3 | 3120 + 3200 + 2400 | 8700 | 178![]() |
2.8'' | 1 |
18 | 4315 | NGC 4850 | 137 | 3 |
![]() |
7200 | 153![]() |
2.2'' | 1 |
19 | 3073 | NGC 4883 | 175 | 3 |
![]() |
7200 | 106![]() |
2.2'' | 1 |
20 | 1853 | - | 190 | 4 |
![]() |
10800 | 88![]() |
1.6'' | 2 |
21 | 3201 | NGC 4876 | 124 | 4 |
![]() |
12600 | 24![]() |
1.7'' | 2 |
22 | 3661 | - | 13 | 5 |
![]() |
16800 | 139![]() |
1.4'' | 1 |
23 | 4679 | - | 75 | 5 |
![]() |
15300 | 114![]() |
1.6'' | 1 |
24 | 3352 | NGC 4872 | 130 | 5 |
![]() |
14400 | 111![]() |
1.5'' | 1 |
25 | 2535 | IC 4041 | 145 | 5 |
![]() |
14400 | 48![]() |
1.3'' | 1 |
26 | 3958 | IC 3947 | 72 | 5 |
![]() |
14400 | 102![]() |
1.5'' | 1 |
27 | 2776 | - | 39 | 5 |
![]() |
14400 | 77![]() |
2.3'' | 1 |
28 | 0144 | NGC 4957 | - | 6 | 4300 + 4500 | 8800 | 91![]() |
1.5'' | 1 |
29 | 0282 | NGC 4952 | - | 6 | 4100 + 5000 | 9100 | 135![]() |
2.5'' | 1 |
30 | 0756 | NGC 4944 | - | 6 |
![]() |
9000 | 88![]() |
2.0'' | 1 |
31 | 1176 | NGC 4931 | - | 6 |
![]() |
9000 | 78![]() |
2.0'' | 1 |
32 | 1990 | - | - | 6 | 5100 + 4500 | 9600 | 135![]() |
2.2'' | 1 |
33 | 5279 | NGC 4827 | - | 6 | 4500 + 5330 | 9830 | 56![]() |
1.5'' | 1 |
34 | 5568 | NGC 4816 | - | 6 | 3600 + 3400 | 7000 | 78![]() |
1.2'' | 1 |
35 | 5975 | NGC 4807 | - | 6 | 4500 + 3000 | 7500 | 23![]() |
2.0'' | 1 |
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