A&A 367, 86-105 (2001)
DOI: 10.1051/0004-6361:20000479
S. J. Smartt1,4 - K. A. Venn2 - P. L.
Dufton3 -
D. J. Lennon4 - W. R. J. Rolleston3 - F. P. Keenan3
1 -
Institute of Astronomy, University of Cambridge, Madingley Road,
Cambridge CB3 0HA, UK
2 -
Department of Physics and Astronomy, Macalester College, St. Paul, MN
55101, USA
3 -
The Department of Pure and Applied Physics, The Queen's University
of Belfast, Belfast BT7 1NN, North Ireland
4 -
ING Telescopes, Apartado de Correos 368, Santa Cruz De
La Palma, 38780, Canary Islands, Spain
Received 2 August 2000 / Accepted 17 November 2000
Abstract
High-resolution, high signal-to-noise spectral data are presented
for four young B-type stars lying towards the Galactic Centre.
Determination of their atmospheric parameters from their
absorption line profiles, and uvby photometric
measurement of the continua
indicate that they are massive objects lying slightly out of the
plane, and were probably born in the disk between 2.5-5kpc from
the Centre.
We have carried out a detailed absolute and differential line-by-line
abundance analyses
of the four stars compared to two stars with very similar atmospheric
parameters in the solar neighbourhood. The stars appear to be
rich in all the well sampled chemical elements (C, N, Si, Mg, S, Al),
except for oxygen. Oxygen abundances derived in the atmospheres of
these four stars are very similar to that in the solar neighbourhood.
If the photospheric composition of these young stars is reflective of
the gaseous ISM in the inner Galaxy, then the values derived for the
enhanced metals are in excellent agreement with the extrapolation of
the Galactic abundance gradients previously derived by Rolleston
et al. (2000) and others. However, the data for oxygen suggests that the
inner Galaxy may not be richer than normal in this element, and the
physical reasons for such a scenario are unclear.
Key words: stars: early-type - stars: atmospheres - stars: abundances - Galaxy: evolution - Galaxy: center - Galaxy: abundances
Much work has been done on tracking the chemical evolution of the
solar neighbourhood as a function of time using age-metallicity
relations and element abundance ratios
in cool F and G-type stars. The seminal work by
Edvardsson et al. (1993) showed the potential of using samples of
nearby
stars with an age spread to trace the temporal evolution of the disk
and the halo, and such work has been extended by, for example, Fuhrmann
(1998) and Israelian et al. (1998).
These types of metallicity ratios
have been modelled by many authors (e.g. Carraro et al. 1998;
Pilyugin & Edmunds 1996 and references therein), in order to
constrain our Galaxy's history. The photospheric abundances of
massive stars (viz. stars with ZAMS
)
do not
have the potential to provide us with temporal information as
they are by nature short-lived objects (with typical lifetimes
of the order of
10-50Myrs). However, their high intrinsic
luminosities mean that we can spectroscopically observe them
at large distances from the solar neighbourhood; OB-type main-sequence
stars are typically
8
visually brighter than early G-type dwarfs. Their
photospheres are representative of the interstellar material from which they
were born due to their relative youth.
Hence, we can spatially and radially sample the metallicity of the
Galactic disk by determining abundances in distant, luminous blue
stars. The radial variation of different elements provides a further key
constraint on the formation and evolutionary mechanisms of the
Galactic disk (Matteucci & Francois 1989;
Matteucci & Chiappini 1999; Prantzos & Aubert 1995;
Portinari & Chiosi 1999; Pagel & Tautvaisine 1995).
These models require further observational constraints to restrict the
possible influence of their many free parameters such as initial mass function,
gas fraction variations, star formation rates, infall and outflow
rates and metallicity dependent stellar yields. Previous work has shown
that the metallicity of the Galactic disk decreases as Galactocentric
distance increases. This metallicity gradient has been
reproduced quantitatively using a number of different methods.
For example, extensive studies of optical
emission lines in H II regions by Shaver et al.
(1983) and Fich & Silkey (1991) have
derived an oxygen abundance gradient of
dex kpc-1;
within the region
kpc (where
refers
to Galactocentric distance throughout this paper). Good agreement
was also found by Maciel & Köppen (1994) and
Maciel & Quireza (1999) from
PN studies, and by Afflerbach et al. (1997) in a study of
ultra-compact H II regions. Several studies of B-type
main-sequence stars throughout the disk had given discrepant results, and
a much shallower gradient than that derived from the nebular studies
(e.g Fitzsimmons et al. 1992, and references therein).
However Smartt & Rolleston
(1997), in a consistent re-analysis of an extensive data-set of B-type
stars covering
kpc, showed that previous
B-type star studies produced spurious gradient determinations,
and derived an oxygen
abundance gradient of
dexkpc-1 in excellent
agreement with the nebular studies. As a follow up to this, Rolleston
et al. (2000; hereafter RSDR)
carried out a more detailed and exhaustive
survey of the large data-set gathered by the group at Queen's University
and have confirmed this result, while also producing abundance
ratios of different elements along the disk. Further work by
Gummersbach et al. (1998) also indicates that the nebular and
stellar results are in agreement within the errors of measurement.
However, most studies of the chemical composition of
the Galactic disk have concentrated on the solar neighbourhood,
and the anti-centre direction. Optical extinction rises steeply
in the direction toward the Galactic Centre, and restricts the
optical observation of disk stars more than a few kiloparsecs away.
In the anti-centre direction we have previously probed out to
kpc,
experiencing
(Smartt et al. 1996), but have
only managed to observe open clusters towards the Centre approximately
2-3kpc away (RSDR) with a reddening of
.
The metallicity within the solar Galactic radius is poorly
sampled, with only a few IR studies available.
Rudolph et al. (1997) and Afflerbach et al. (1997) have used similar observational methods for
far-IR emission lines in H II regions (of differing sizes)
as ISM abundance probes. This far-IR technique has allowed observation
of some inner Galactic regions. These studies suggest that the
abundance gradients for both nitrogen and sulphur do tend to increase
towards the Galactic Centre, but both indicate that their
oxygen abundance derivations may suffer some uncertainties (see
Sect.5.1). Carr et al. (2000) have
achieved a first determination of stellar abundances (in an M-type
supergiant) directly at the Galactic Centre in the spectacular
star forming region within the central few hundred parsecs.
They find an Fe abundance close to that of the solar neighbourhood,
which at first may seems surprising given the wealth of evidence
of a strong abundance gradient in the Milky Way (and also in other
spirals). However as we discuss in Sect. 4,
this region may be unrelated to the starforming Galactic disk
(
kpc) as a whole, and its evolution
may have proceeded independently.
The metallicity of the inner disk therefore has not been investigated
in any great detail, and lacks information on a variety of elemental
abundances. Spectra of B-type stars provide the potential
to sample many more chemical elements than is possible through
emission line studies. As discussed above, massive B-type main-sequence
stars are excellent probes of the current interstellar medium
and sample particularly well the -processed elements (O, Mg, and
Si) as well as C and N (Smartt et al. 1996).
This paper describes observations of metal-rich
massive stars lying towards the Galactic Centre together with
detailed model atmosphere and abundance analyses. This allows the
comparison of abundance ratios in this interesting regime with those
in the rest of the disk, as a key probe of the evolution of our Galaxy.
The combination of these criteria should produce targets which are
potentially normal early B-type dwarfs (or near main-sequence objects)
lying within 4kpc of the Galactic Centre and twelve such stars
were found.
The database of Reed & Beatty (1995) contains published spectral types for a subset of the objects in Reed (1993). However, the fainter objects of the earlier study (which were the most promising distant B-type stellar candidates) did not have an entry in Reed & Beatty (1995). As a result we had to rely on colours as a primary selector.
Star | SpT | V | B-V | U-B | ![]() |
![]() |
LS5130 | B2IV | 12.09 | 0.24 | -0.53 | 21.12 | -5.63 |
LS4419 | B2IV | 11.08 | 0.08 | -0.65 | 351.78 | -5.88 |
LS4784 | B2IV | 11.55 | 0.09 | -0.60 | 1.70 | -6.12 |
LS5127 | B3II-III | 11.96 | 0.15 | -0.55 | 16.13 | -7.34 |
During both observing programmes the observational routine was similar. Each star
was observed long enough to accumulate sufficient counts in the continuum
to yield a minimum signal-to-noise ratio of 70, and in most cases
more than 100. The observations were split into exposures of
1200-1500s to minimize the impact of cosmic-ray contamination and
were bracketed by observations of Cu-Ar and/or Cu-Ne wavelength calibration
lamps. Flat-field exposures and bias frames were generally taken at both the
beginning and the end of the night.
The single order spectra were reduced using the STARLINK package FIGARO (Shortridge et al. 1997) as discussed in Smartt et al. (1996), where further details can be found. The échelle data were reduced using the Image Reduction and Analysis Facility ( IRAF), using standard techniques. When in one dimensional format, all the spectra were transferred for further analysis to the STARLINK program DIPSO (Howarth et al. 1998). Normalization was achieved by carefully selecting continuum regions free from absorption lines, and fitting low order (of degree 3 or 4) polynomials through the noise. Equivalent widths for the metal lines and non-diffuse lines of neutral helium were measured by the non-linear least square fitting of single or multiple Gaussian profiles to the normalised spectra. Each line is assigned an error estimate reflecting the reliability of the equivalent width measurement, viz. a: error less than 10%, b: error less than 20%, c: error greater than 20%. These were assigned by considering the numerical error returned by the DIPSO line fitting computation and the qualitative accuracy of the profile fit. The hydrogen and diffuse helium lines were not measured in this manner, but the normalised profiles were extracted directly for comparison with Galactic standards and with theoretical profiles. The equivalent width measurements for each star can be found in Appendix A (available electronically).
In Table 2 we list the individual photometric indices
as derived from the SAAO observations. Typical uncertainies in the
SAAO photoelectric photometry are 0.01 and 0.005 magnitudes in colours
and the individual passbands respectively. The Strömgren reddening free
[u-b] index should be accurate to 0.03 mag. The ESO CCD
dataset comprised four independent observations per filter per field
containing each of the LS stars. Photometric errors were estimated using
two methods. First, each set of uvby instrumental magnitudes were
independently transformed to the standard system of y and (b-y),
m1, c1 indicies. These were used to deduce mean values and associated
standard errors of the mean. Secondly, values of the arithmetic mean and
standard error were deduced for the independent sets of uvby instrumental
measurements. The mean values were transformed to the standard magnitude
and colour indicies, while the associated photometric uncertainties were
deduced by propagating the standard error of the mean values through the
transformation equations.
For LS4419, both these methods inferred a photometric accuracy of 0.006
magnitudes in y and 0.011 magnitudes in the colour (b-y). The
reddening free [u-b] index should be accurate to
mag.
Thus, it is not surprising that we find excellent agreement between the
SAAO and ESO photometry for LS4419. Photometric uncertainties are
somewhat larger for the ESO data obtained for LS5127, viz. 0.015, 0.035
and 0.15 mag in y, colour and [u-b] index respectively. For this
star, the SAAO data should be the more reliable. However, it should be noted
that the SAAO and ESO measurements agree within the photometric errors -
which in turn lead to similar estimates of the stellar effective temperature
(see Sect. 3).
Star | y | (b-y) | m1 | c1 | [c1] | [u-b] (SAAO) | [u-b] (ESO) |
LS 5130 | 12.038 | 0.264 | -0.025 | 0.182 | 0.129 | 0.248 | - |
LS 4419 | 11.089 | 0.115 | 0.032 | 0.105 | 0.082 | 0.220 | 0.204 |
LS 4784 | 11.519 | 0.157 | 0.035 | 0.122 | 0.091 | 0.261 | - |
LS 5127 | 11.883 | 0.193 | 0.023 | 0.233 | 0.194 | 0.364 | 0.456 |
![]() |
2.84 | -0.106 | 0.093 | 0.116 | 0.137 | 0.255 | |
![]() |
3.80 | -0.064 | 0.078 | 0.294 | 0.307 | 0.423 |
The methods employed to derive stellar atmospheric parameters are similar to those described in Smartt et al. (1996) and Rolleston et al. (1997). All results are based on the ATLAS9 grid of line-blanketed model atmospheres of Kurucz (1991). LTE line formation codes were used to derive line profiles leading to determinations of atmospheric parameters and chemical compositions.
As a first attempt to constrain the effective temperatures (
)
of the four LS stars, which were all of spectral type B2/B3,
we used the ionization balance of the two stages of silicon;
the Si II lines at 4128Å and 4130Å, and the Si III
triplet at 4552-4574Å.
This method has been previously employed to analyse the
bright standard
Pegasi (HR39), and is known to produce
anomalous and unsatisfactory
results (Ryans et al. 1996; Gies & Lambert 1992; Peters 1976). Absolute
abundances (derived using this temperature diagnostic) of N II,
O II, Mg II, and S III show systematic
underabundances of around 0.4dex with respect to normal B-type
stellar values.
Further, and more critically, large discrepancies then exist between the
abundances derived from other ionization equilibria
(C II/C III, S II/S III, and
Si III/Si IV) when using this Si II/Si III
temperature balance.
Using Strömgren photometry as an alternative
effective temperature indicator produces more satisfactory results;
typically, a value systematically
3000K less than that estimated from the Si II/Si III
balance has been obtained which agrees with the other ionization balances,
and the metal abundances give normal results.
The problem with the Si II lines historically stretches back
more than two decades, when Peters (1976) first reported the
problem. All stars in the Gies & Lambert (1992) study (which have
;
from Strömgren indices) appear to produce spurious
Si II results and a similar temperature problem has been found
by the authors in
at least two other stars - 22Ori and
Her.
In each case the temperature
found from the Si II/Si III lines is systematically
higher than that produced from the Strömgren indices.
Hence, it appears that the
Si II/Si III ionization balance is flawed as a temperature
diagnostic due to reasons which are, as yet, not well understood.
We find that including non-LTE corrections
(Brown 1987) in the line formation calculations fails to remove the discrepancy.
Cuhna & Lambert (1994) have also reported the same effect
when adopting Strömgren temperatures for Orion B-types and
their subsequent determination of Si abundances from both ionization
stages.
In the rest of this paper we shall adopt the effective temperatures
derived from the
Strömgren indices listed in Table 2.
The colour index [u-b] measures the Balmer jump as a function of
and we have primarily used the calibration of
Napiwotzki et al. (1993). We also compared the
effective temperatures derived using
the calibration of Lester et al. (1986),
which provides the [u-b] vs.
relation as
a function of surface gravity; consistent results were found
in each case. Table 3 lists the
effective temperatures that we have adopted. For the two
stars for which we have both SAAO and ESO data, we list the
mean of the two temperatures. Excellent agreement was found from the
two data sets for LS4419, with temperatures of 21900 and 22300K derived
from ESO and SAAO respectively; for LS5127 the temperatures were
in poorer agreement, 18700 and 17100K, but consistent
with a
K range.
Logarithmic surface gravities ()
were estimated by fitting
model profiles of the Balmer lines from our atmospheric analysis codes
to observed H
,
H
,
and H
depending on which lines were available in the
RGO spectra (see Sect.2). In all cases the results from each
line were in good agreement. A conservative estimate of
the probable random errors in our analyses, given the errors in the
measured data are
1000K and
dex.
A value for microturbulent velocity (
)
was determined for all the B-type
stars by requiring that the derived abundances from, primarily,
O II were independent of line strength. The O II ion
was chosen as it has the richest absorption line spectra in stars of
this type. However where a significant number of N II lines
were available,
was independently estimated and the mean taken.
Atmospheric parameters for all the stars are listed in Table 3.
As a final step in defining the model atmosphere appropriate to
each star, we determined the helium abundance within this iterative
procedure.
Helium abundances were calculated from the line strengths of the available
non-diffuse lines (3964, 4437, 4713, 5015 & 5047Å) and the profile
fitting of the diffuse lines (4009, 4026, 4387, & 4471Å). In each
star we found a normal helium composition of approximately 10.9dex
(corresponding to
by fraction). Hence, we did not need to
repeat the
iterations for models with
significantly abnormal helium fractions.
Star |
![]() |
![]() |
![]() |
[He] | ![]() |
(K) | (cgs) | (kms-1) | dex | (kms-1) | |
LS5130 | 21200 | 3.5 | 5 |
![]() |
35 |
LS4419 | 22100 | 3.7 | 7 |
![]() |
80 |
LS4784 | 20900 | 3.9 | 5 |
![]() |
93 |
LS5127 | 17900 | 3.2 | 7 |
![]() |
120 |
![]() |
21000 | 3.8 | 5 |
![]() |
5 |
![]() |
17700 | 3.9 | 5 |
![]() |
10 |
Species | LS5130 | ![]() |
Differential | Nearby | Other solar | Solar | |
B-stars1 | neighbourhood | ||||||
C II |
![]() |
![]() |
![]() |
8.20 | 8.13![]() |
8.55 | |
N II |
![]() |
![]() |
![]() |
7.81 | 7.94![]() |
7.97 | |
O II |
![]() |
![]() |
![]() |
8.68 | 8.70![]() |
8.87 | |
Mg II | 7.64 (1) | 7.22 (1) | +0.42 (1) | 7.38![]() |
- | 7.58 | |
Al III |
![]() |
![]() |
![]() |
6.45 | - | 6.47 | |
Si II |
![]() |
![]() |
![]() |
7.28![]() |
- | 7.55 | |
Si III |
![]() |
![]() |
![]() |
7.28![]() |
- | 7.55 | |
S II |
![]() |
![]() |
![]() |
7.21 | 7.12![]() |
7.21 | |
S III | 7.29 (1) |
![]() |
+0.23 (1) | 7.21 | 7.12![]() |
7.21 | |
Ar II |
![]() |
![]() |
![]() |
- | - | 6.58 | |
Fe III |
![]() |
![]() |
![]() |
7.36![]() |
- | 7.51 |
1. All typical B-type stellar values are taken from Gies & Lambert (1992), apart from |
a: Kilian (1994). |
2. Solar neighbourhood values taken from: |
c: Local ISM abundances from Meyer et al. (1998), Cardelli et al. (1996). |
d: Orion nebular abundances from Baldwin et al. (1991). |
3. Solar values taken from Grevesse & Noels (1993) and Anders & Grevesse (1989). |
Species | LS4419 (absolute) | LS4419 (differential) | LS4784 (absolute) | LS4784 (differential) | |
C II |
![]() |
![]() |
![]() |
![]() |
|
N II |
![]() |
![]() |
![]() |
![]() |
|
O II |
![]() |
![]() |
![]() |
+0.17 ![]() |
|
Mg II | 7.50 (1) | +0.28 (1) | 7.70 (1) | +0.48 (1) | |
Al III |
![]() |
+0.18 (1) |
![]() |
+0.47 (1) | |
Si II | 6.90 (1) | +0.23 (1) |
![]() |
![]() |
|
Si III |
![]() |
![]() |
![]() |
+0.58 ![]() |
|
S II | 7.17 (1) | +0.12 (1) | 7.51 (1) | +0.33 (1) | |
Ar II | 6.59 (1) | +0.09 (1) | - | - | |
Fe III | 6.45 (1) | -0.18 (1) | - | - |
Ion | LS5127 | ![]() |
Differential |
C II |
![]() |
![]() |
+0.55 (1) |
N II |
![]() |
![]() |
![]() |
O II |
![]() |
![]() |
![]() |
Mg II | 7.76 (1) |
![]() |
+0.69 (1) |
Al III | 6.68 (1) |
![]() |
+0.36 (1) |
Si II |
![]() |
![]() |
![]() |
Si III |
![]() |
![]() |
![]() |
S II |
![]() |
![]() |
![]() |
S III |
![]() |
![]() |
![]() |
Fe III | - (0) |
![]() |
- (0) |
Species |
![]() |
![]() |
![]() |
![]() |
![]() |
![]() |
C II |
0.01 | 0.00 | 0.02 | 0.03 | -0.08 | 0.11 |
N II | 0.07 | -0.05 | -0.06 | 0.12 | -0.08 | 0.11 |
O II | 0.12 | -0.11 | -0.18 | 0.23 | -0.11 | 0.15 |
Mg II | -0.05 | 0.05 | 0.13 | -0.11 | -0.22 | 0.32 |
Al III | 0.05 | -0.03 | -0.02 | 0.08 | -0.08 | 0.15 |
Si II | -0.06 | 0.08 | 0.20 | -0.15 | -0.05 | 0.11 |
Si III | 0.09 | -0.09 | -0.14 | 0.18 | -0.21 | 0.24 |
S II | -0.04 | 0.05 | 0.16 | -0.14 | -0.09 | 0.21 |
S III | 0.13 | -0.12 | -0.14 | 0.18 | -0.07 | 0.13 |
Ar II | 0.01 | 0.02 | 0.12 | -0.06 | -0.06 | 0.14 |
Fe III | 0.09 | -0.08 | -0.04 | 0.08 | -0.08 | 0.24 |
![]() |
Figure 1:
The Mg II doublet at 4481Å in LS5127 with
synthesised fits. The observed line profile is the solid histogram,
together with theoretical fits for magnesium abundances of
7.2, 7.5, 7.8 and 8.1dex (each with a microturbulence of ![]() ![]() |
Open with DEXTER |
Star | M |
![]() |
![]() |
z | PRV | ![]() |
![]() |
(![]() |
(Myrs) | (kpc) | (kpc) | (kms-1) | (kms-1) | (kpc) | |
LS5130 | ![]() |
![]() |
![]() |
0.7 | -7 | 64 | ![]() |
LS4419 | ![]() |
![]() |
![]() |
0.5 | -28 | 45 | ![]() |
LS4784 | ![]() |
![]() |
![]() |
0.4 | -30 | 36 | ![]() |
LS5127 | ![]() |
![]() |
![]() |
1.2 | -3 | 98 | ![]() |
The four B-type targets lie at significant distances below the plane
of the Galaxy (see Table 8), and it is unlikely that
they have formed at their current positions. The probable explanation
is that they were born in the Galactic disk, and were then
subsequently ejected. This class of early-type stars lying at, in some
cases, significant distances from their formation sites (i.e. runaway
stars) has been studied extensively, and their existence is well
accepted (see Rolleston et al. 1999).
In a previous paper (Smartt et al. 1997) we have discussed
the analysis of four blue supergiants lying at similar Galactic
latitudes, and a full discussion of the possible origins and ejection
mechanisms from the disk of these types of objects is
presented. Leonard (1993) has presented plausible mechanisms
for early-type stars suffering large ejection velocities within
massive binary evolution scenarios, or dynamical ejection from
clusters. These can result in initial velocity injections of up to
350 kms-1, easily allowing our objects to reach their
current estimated distance from the Galactic plane (z).
In order to constrain the regions of the disk where these stars were
formed, we follow the arguments of Smartt et al. (1997). For each
of the stellar spectra we have determined a radial velocity, and we
compare this velocity to that expected if the stars were co-rotating
with the disk at their current projected positions (assuming a flat
rotation curve and
= 8.5kpc,
kms-1 from
Kerr & Lynden-Bell 1986). We calculate a peculiar radial
velocity (PRV) which is the difference between these two, as a
representative measure of the dynamics of the stars with respect to
the Galaxy rotation. This PRV is only one component of the total
peculiar space velocity, with the other two components unknown
(directed along axes mutually perpendicular to the PRV). We assume
that the second component of velocity is perpendicular to the Galactic
plane (which, given the small Galactic latitudes of these stars is a
reasonable assumption). House & Kilkenny (1980) have given an
expression for the force field normal to the Galactic plane (Kz),
but this is probably only valid for
kpc. Given
that these stars have origins closer to the Galactic Centre, we have assumed
a simple relation which is an approximation to the shape of
the House & Kilkenny form of Kz, but which is always greater
in magnitude (with a peak of twice the House & Kilkenny maximum).
We assume that within
pc the force field
Kz is approximated by -kz (with
k = 2.29 10-11cms-2 pc-1) and for z > 700 pc then
Kz = -k. Analytical expressions for
the initial velocity (
)
and time required for the
stars to reach their present z distance, could then be deduced.
The initial ejection velocity perpendicular to the Galactic plane
needs to be no more than 100kms-1 (see Table 8)
for these stars to reach their
present positions, and significantly less for three of the objects.
Velocities of this order are well within those plausible for
common ejection methods (Leonard 1993). The PRV of each star
is also relatively low, indicating that the initial space
velocities of these stars were probably directed mostly in the z-direction.
However without proper motion information we cannot definitively constrain
their space velocity. If we assume that the
mean ejection velocity in the sample (
kms-1)
is indicative of the maximum value of any one velocity component, then
we can calculate a circular locus within which the star probably
formed. In 20Myrs, a star could have moved 1.7kpc within
(or parallel to) the Galactic disk, and this is a reasonable
error on the birth sites of the stars.
![]() |
Figure 2:
a) A schematic view of the Galaxy with the
Galactic Centre at (0, 0).
The inner dotted circle represents the possible bulge-disk
transition with which a rapid decline in the disk stellar
density is associated; the outer dashed circle represents the
stellar ring proposed by Bertelli et al. (1995) of radius 3.5kpc
(they estimate a Gaussian half-width of
![]() ![]() |
Open with DEXTER |
![]() |
Figure 3:
Galactic abundance gradients and the stellar sample (open circles)
of RSDR. Also shown are the four current targets (open stars) at
![]() ![]() ![]() ![]() |
Open with DEXTER |
Figure 2 illustrates the estimated stellar birth sites
and also schematically draws attention to features of
Galactic structure in this region.
Bertelli et al. (1995) have discussed the possibility that
a "stellar ring'' exists, which surrounds the Galactic Centre at a
distance of 3.5kpc, with a 0.5kpc half-width.
The nature of this ring is still somewhat uncertain, but Bertelli et al.
claim that assuming it to be a region of increased stellar disk density
would produce synthetic colour-magnitude diagrams closely matching those
they have observed in fields towards the Galactic Centre. They suggest
that it is associated with active star formation, and that
the cross-section through this ring is at a maximum when viewed at
.
They further suggest that this is consistent with an
observed peak in OH maser sources (Blommaert
et al. 1994) around
,
a signature of
relatively recent star formation. Whether this is a complete ring
surrounding the inner Galaxy, or part of an inner spiral arm,
is still debatable (Ortiz & Lépine 1993).
Clemens et al. (1988) have mapped the distribution
of molecular gas in the first quadrant of the Galaxy, and found an H2ring at approximately 4kpc from the Centre. This would appear to
correlate with the stellar ring found by Bertelli et al., but the
latter suggest that the stellar and molecular ring are not necessarily
co-incident. Rather that the molecular ring maybe
the result of a "bow-shock of propagating star formation''.
Although some Galaxy models (Bahcall & Soneira 1980; Bahcall 1986) adopt
a relation for the number density of disk stars
which falls off exponentially with
increasing distance from the Galactic Centre,
it appears that at
kpc
from the Galactic Centre there is a sharp decrease in the stellar
disk density (Bertelli et al. 1995, Paczynski et al. 1994).
This has been interpreted as the transition between the disk and the bulge.
Figure 2 suggests that the four B-type stars
have their origins in this stellar ring, or where the molecular ring
dynamics are causing current star formation. Indeed if there is a
disk/bulge transition, early-type stars may be relatively rare within
this inner region, although current and
spectacular star formation is certainly on-going in the inner few
hundred parsecs (Serabyn et al. 1998; Najarro et al. 1997). The above
indicates that our B-type stars probably formed around
kpc
in the Galactic disk and were subsequently ejected.
![]() |
Figure 4:
a) Filled stellar symbols are oxygen
abundances of the Galactic Centre objects, together with stars
from RSDR with
![]() ![]() ![]() ![]() ![]() |
Open with DEXTER |
Carr et al. (2000) have presented
IR spectra of a red M-type supergiant (IRS7) lying
in the central cluster, and compared the metal line spectra with
that of a similar star in the solar neighbourhood. A differential
abundance analysis of the star indicates that it has an Fe
abundance very close to solar, and certainly does not appear
metal enhanced in any significant way (although it does appear to
show CNO-abundances typical of severe dredge-up after core H-burning).
At first this may seem a surprising
result given that one might expect a linear Galactic abundance gradient
in the thin disk to keep rising towards the Centre, and that the central
regions would thus show super-metal rich abundances. However
there is lack of molecular gas, young stars and H II regions
within the inner 3.5kpc, until one gets to the
central cluster. Hence, whatever the evolutionary history of the
massive stellar clusters at the Centre, it seems to have occurred
independently of the Milky Way disk (
kpc).
There is no reason therefore too assume that the disk and central cluster
are chemically linked, and one should consider them as separate
entities in evolutionary scenarios.
Afflerbach et al. (1997) have carried out an extensive
study of the N, S and O abundances in H II regions with
kpc. In Fig. 4, their O
abundances are plotted together with those for our stars
and the stellar sample of RSDR with
kpc.
The results from the H II
regions and the RSDR data-set in the outer Galaxy are
in excellent agreement (see RSDR and Smartt 2000 for
further details). However differences appear at smaller
Galactocentric distances. In particular while the
H II region data is consistent with a linear gradient
-0.06dexkpc-1, no significant variations are apparant
for the B-stars. One possible explanation is that the zero-points
for the two studies is somewhat different. For example,
if the Afflerbach et al.
data were shifted upwards by
0.2dex (the difference in the
mean results within
kpc), then
good agreement would be found within the solar neighbourhood,
but the B-stars towards the Centre would not reflect the
H II region results. The cause of this discrepancy is unclear
and appears to be limited to oxygen (see Sects. 5.2
and 5.3). We note that Afflerbach et al.
is the only H II region study of the inner Galaxy to quote
oxygen abundances. Rudolph et al. (1997) were unable to
independently determine S and O, while Shaver et al. (1983) point out
that their two inner Galaxy objects (
4kpc)
are somewhat peculiar, and that their radio and optical sources
may not be the same. Afflerbach et al. have
recalculated N abundances from the data of Simpson et al. (1995)
and find differences of a factor of
2
in some of the highest
metallicity regions. Additionally,
on average there seems to be a 50% discrepancy in the sulphur
abundances. Therefore it appears
that the nebular abundances in the inner Galaxy require further
study to constrain their absolute values, while a
more extensive sample of B-stars is required. Smartt et al.
(2000) have discussed the problem of determining
reliable absolute nebular abundances in regions of solar
abundance and above in M31, and find significant discrepancies
across different analysis methods.
![]() |
Figure 5:
Oxygen abundances in the Milky Way from our stellar
sample (open circles) and the H II region
oxygen abundances in NGC2403 from Garnett et al. (1997, solid
circles). Both are shown as a function of effective radius.
Within
![]() |
Open with DEXTER |
In terms of effective radius (where
kpc, from de
Vaucouleurs & Pence 1978), the break in the observed O abundance
gradient occurs at
.
The spiral galaxy
NGC2403 has been studied by Garnett et al. (1997) who estimated
H II region abundances at 12 points across the disk.
In Fig. 5 the oxygen abundance is plotted
as a function of effective radius for both these galaxies.
A similar gradient change may be present in NGC2403, although
the sampling is relatively sparse.
This would imply that a two component model for the
radial variation of oxygen is not exclusive
to our own Galaxy and suggests indirect evidence that it is a
real effect. Certainly the data suggests that for both
galaxies, at
,
there is little
evidence for increasing oxygen abundances.
However, the well studied galaxy M33 (Garnett et al.
1997; Henry & Howard 1995; Monteverde et al.
2000) which is similar in size, mass and metallicity
to NGC2403 does not appear to show such a break. Additionally,
M101 and NGC4303 are other examples of extensively studied galaxies
(Kennicutt & Garnett 1996; Martin & Roy 1992),
with 30-80 H II
region data points sampled in each. Neither show evidence
for anything other than a linear relation across the whole
disk. As these spirals are much better sampled than the Galaxy
or NG2403, we require more data points in the inner Milky Way
to confirm that this flattening is a real effect.
The magnesium and silicon abundances for the Galactic
Centre stars and the
corresponding RSDR dataset are plotted in Fig. 3.
For oxygen and silicon we note that the absolute abundances derived in our
two solar neighbourhood standard stars are similar to the mean of
the RSDR data set within 1kpc of the Sun. Hence the
absolute results of the Galactic Centre stars are on a consistent scale to
the rest of the data. However with Mg,
the abundance derived in Peg and
Her
is 7.2dex, whereas in RSDR the mean local value is 7.85 dex
(again estimated from the mean of stellar values
within 1kpc of the Sun).
The stars analysed in this paper have effective temperatures
in the range
K, while the RSDR sample have a mean
temperature of 25275K and it is possible that the modelling of this line is
temperature dependent. It is important to use homogeneous datasets when
looking for abundance trends and that is indeed what was done in RSDR.
They showed that although the zero-point of any particular
abundance may be in the error, any gradient biasing is removed by using
suitable sub-samples of stars. Hence to put the four Galactic Centre stars
onto the same scale as RSDR, we have added 7.85dex to
their differential Mg abundances.
This assumes that
Peg,
Her and the stars within 1kpc
of the Sun in RSDR have similar abundances and that any offset is
due to temperature dependent modelling inaccuracies. It appears
that the Galactic Centre stars are richer in Mg than
their local counterparts (see also Sect. 3).
From Table 7 an underestimate of the microturbulence
(
)
in each star could reduce this Mg overabundance as it is based on
a relatively strong line. However an increase in
by 5kms-1 or more
in each star would be required. While observational uncertainties
might cause a particular determination of
to have such an error,
there is no evidence that we have systematically underestimated
the microturbulence in each star. The Mg abundance
gradient in Fig. 3 hence appears to steadily increase
towards the Centre. A value of
dexkpc-1 is derived by fitting
all points, which is slightly larger the RSDR result, although they
agree within the uncertainties.
The silicon abundances of the Galactic centre stars are also plotted
in Fig. 3, and again appear compatible with
a steadily increasing abundance gradient towards the centre.
A fit through all the Si points gives a gradient of
dexkpc-1,
again in good agreement with the RSDR value for the outer Galaxy only.
The differential abundances derived for each star are between
0.3-0.6dex higher than their respective standards, strongly
supporting the idea that these stars are richer in Si than
solar neighbourhood material.
In Fig. 6 the Al abundances of the galactic
Centre stars are plotted again with the data from RSDR.
Aluminium appears to increase steadily toward the Centre, and a
gradient of
dexkpc-1 is derived from the whole
data set, similar to that produced by RSDR for the anti-centre direction.
Sulphur abundances for our stars are also available, although we
have no similar data at greater radii from RSDR. In
Fig. 6 our stellar abundances
are compared with with those from Afflerbach et al.
(1997). Reasonable agreement is found with
these stars appearing to be S-rich and having abundances similar to those
found in nebular studies towards the Centre (although the caveat
regarding the nebular absolute values
discussed in the previous section may be important).
![]() |
Figure 6:
a) Asterisk symbols are aluminium
abundances of the Galactic Centre objects and circular points are
Al abundances from RSDR stars.
Error bars for the Galactic Centre stars are taken as typically
![]() ![]() |
Open with DEXTER |
The correlation of of O, Mg, and Si abundances found in the outer Galaxy (see RSDR) and in the low metallicity Magellanic Clouds (Rolleston et al. 1996) is theoretically expected as these elements are all produced and returned to the ISM through supernovae typeII. One would expect that this trend should continue towards the inner Galaxy, and it is surprising that our O abundances show no sign of increase, while Mg and Si appear significantly higher than their comparison stars. The S abundances in these stars are also higher than normal, although the dataset is somewhat less robust given the weakness of the absorption features of this element.
There has been some investigation into the effects of metallicity
dependent yields from massive stars.
Prantzos et al. (1994) have taken the
yields of Maeder 1992) and investigated the the
galactic chemical evolution of C and O. These studies suggest
that oxygen yields from massive
stars tend to decrease towards higher metallicities. This is due to
mass loss being higher at high metallicities and
significantly more mass is ejected in the stellar wind (which is rich in
helium and carbon). This leaves
a smaller mass fraction which can be converted
into oxygen. However a detailed study of element yields other than C and O,
and their subsequent evolution has not yet been done.
Hence given the statistically small number of stars we have in the present
study, and the lack of metallicity dependent yield calculations
for the other elements we cannot interpret the data in any more
detail at present. Analysis of more stars in this region of
the Galaxy is desirable.
The observational findings presented here should provide a stimulus to
investigations of the evolution of the Galaxy at high metallicities,
given the
possibility that stellar yields are highly variable with metallicity.
Furthermore, we should attempt
to determine the true quantitative variation in the ratios of the
-elements in the inner Galaxy to investigate if this trend is
real. The recent chemical evolution models of Portinari & Chiosi
(1999, 2000) do indicate that the O abundance gradient
may indeed flatten off in the region
kpc,
which is encouraging. However their SN typeII yields would suggest that
the other
-processed elements (Mg and Si) should follow a similar
trend, which we do not observe. Boissier & Prantzos (1999)
also suggest that their evolutionary models for the Milky Way predict
a flattening of the metallicity gradient with time, and that
saturation is reached in the inner Galaxy. This is
in reasonable quantitative agreement with what we find for oxygen
abundances, but does not explain the behaviour of the other
elements.
Each of the four stars do appear to have nitrogen abundances significantly
higher than their solar neighbourhood counterparts. The result for LS5130
is more marginal than the others, but with 19 features measured an
error in the mean (
)
of 0.03 dex would suggest that
the enrichment is real although small. Figure 4b shows the nitrogen
abundances of the Galactic Centre stars plotted along with a sample of
stars from RSDR, and the H II regions from Afflerbach et al. (1997).
The gradient found in RSDR appears to continue towards the Centre;
certainly the abundance estimates do not appear to flatten as for oxygen.
Reasonably good agreement
with the absolute values of the H II region analyses is found, and
we calculate quite a significant gradient for
N/O. The
full sample of RSDR yielded a gradient in N/O of
dexkpc-1,
and this is increased slightly by the inclusion of the four Galactic Centre
stars. The latter effect is clearly due to the stars showing enhanced
nitrogen in their atmospheres, but normal oxygen.
We have previously discussed the possibility of N-rich core gas
contaminating the surface of young massive stars, and hence polluting
the natal photospheric material (RSDR). However due to the lack of
a strong anti-correlation of C and N, RSDR found no evidence for such
contamination and the same argument holds for the four stars analysed here.
![]() |
Figure 7:
a) The Galactic abundance gradient of carbon, composed of the
four programme stars and the data from RSDR. A gradient of
![]() ![]() ![]() |
Open with DEXTER |
An explanation of the N/O gradient was discussed in RSDR, based on the
original ideas of Vila-Costas & Edmunds (1993). At metallicities
higher than
O/H
8.3,
N/O would be proportional
to
O/H, because N enrichment would be dominated by secondary
production, i.e. from C and O seed nuclei in the
natal interstellar material through the CNO-cycle.
Primary production would play a minimal role in contributing to the ISM
abundance of N, and simple closed-box models in which the production of
N is dominated by secondary production do predict a linear
trend.
At low metallicities the secondary component
becomes less important (due to its inherent dependence on metallicity),
and the dominant mechanism for the production of nitrogen
is primary. The metallicity at which this happens is around
8.3 dex, close to the value found in the outer most regions of
the Galactic disk. These ideas stem from the fact that
the N/O ratios in low metallicity dwarf galaxies show
a very large scatter, with little discernable trend. As
primary nitrogen orginates in low-intermediate mass stars
there is a time delay between is enrichment and that of
primary oxygen (from short-lived massive stars), hence the
scatter at low metallicitities is often interpreted as
time dependent delays between oxygen and nitrogen
enrichments. However at the high metallicities we are dealing
with, it is likely that secondary production (in massive
stars is dominant), and there would thus be no delay between
N and O production.
Our N/O ratios toward the Galactic Centre tend to support the
conclusions of RSDR that the N/O gradient in the Galaxy is real, and
due to secondary N production dominating across the disk.
In Fig. 7 carbon abundances for the four Galactic Centre stars and from the RSDR compilation are plotted. As three of our new points are C-rich, it is not surprising that the abundances appears to steadily increase toward the Centre. As discussed for nitrogen, because the stars are not rich in oxygen the C/O ratio is greater than normal, hence producing a larger (and statistically more significant) gradient than that given in RSDR. As discussed by Maeder (1992) this would be consistent with C being produced at a greater rate than O in metal-rich stars (Sect. 5.2), due to the increased mass-loss in the post He-burning phase.
Acknowledgements
Spectroscopic data were obtained at the Anglo-Australian Observatory in Siding Spring, New South Wales and the ESO 3.6m Telescope on La Silla. We also obtained photoelectric photometry on the 0.5m in South Africa, and CCD photometry on the 1.5m Danish Telescope in La Silla. We are grateful to the staff at all observatories for their assistance. We are particularly grateful to Dave Kilkenny at SAAO who provided urgent initial photometry, which allowed some of this work to be incorporated into S. Smartt's Ph.D. Thesis. Data reduction was performed on the PPARC funded Northern Ireland STARLINK node, and some of the model atmosphere programs were made available through the PPARC supported Collaborative Computational Project No. 7. We made use of the SIMBAD database maintained at CDS, Strasbourg. We acknowledge financial support for this work from the British Council and the Deutscher Akademischer Austauschdienst in the context of the British-German Academic Research Collaboration initiative, and also the Visiting Fellowship program at Queen's University which allowed travel for SJS to further discussions. SJS and WRJR acknowledge funding from the PPARC. KAV would like to acknowledge travel and observing support by ESO, and to thank the Henry Luce Foundation for research funds through a Clare Boothe Luce professorship award.
Species and Line | LS5130 | LS4419 | LS4784 | ![]() |
O II 3911.96 | 57a | 96c | 76b | 50a |
C II 3918.98 | 138a | 130a | 261a7 | - |
N II 3919.01 | ||||
O II 3919.28 | ||||
C II 3920.69 | 127a | 113a | - | - |
O II 3945.04 | 30b | 63b | 58c | - |
O II 3954.37 | 51a | 63b | 74b | 48a |
He I 3964.73 | 177a | 187a | 188a | 155a |
O II 3982.72 | 29c | 72b1 | 60a1 | 31a |
S III 3983.77 | 28c | - | - | 14a |
Species and Line | LS5130 | LS4419 | LS4784 | ![]() |
N II 3995.00 | 102a | 104a | 93a | 78a |
N II 4035.08 | 35b | 46a | 42c | 25a |
N II 4041.31 | 44b | 51a | 60b | 37a |
O II 4069.62 | 95a | 115a | 97b | 94 |
O II 4069.89 | ||||
O II 4072.16 | 62a | 75a | 43c | 65a |
C II 4074.48 | 48b | 147a2 | 141b2 | 42a |
C II 4074.85 | ||||
C II 4075.85 | ||||
O II 4075.86 | 124a | - | - | - |
C II 4076.53 | ||||
N II 4076.91 | ||||
O II 4078.84 | 24c | 19c | - | 26a |
O II 4085.11 | 32c | 39b | 34b | 26a |
O II 4087.15 | 22c | - | - | 12a |
Si IV 4088.85 | 37c | 44b | 61a | - |
O II 4089.28 | ||||
Si IV 4116.10 | 20c | - | - | - |
O II 4119.22 | 62a | - | - | - |
O II 4120.28 | 291a | - | - | - |
He I 4120.81 | ||||
He I 4120.99 | ||||
O II 4121.46 | ||||
Si II 4128.07 | 53a | 38c | 65b | 33a |
Si II 4130.89 | 54a | 86c3 | 70c3 | 34a |
Ar II 4131.71 | 38b | - | - | 17a |
O II 4132.80 | 33b | - | 54c3 | 30a |
Al III 4149.50 | - | 56b | 64a | - |
Al III 4150.14 | - | |||
S II 4153.10 | 60a | 125b | 105a | - |
O II 4153.30 | ||||
S II 4162.70 | 47b | - | - | - |
Fe III 4164.79 | 45c | 32b | - | |
Fe III 4166.86 | 26c | - | - | 13a |
He I 4168.97 | 99a | 99a | 104a | - |
O II 4169.22 |
Species and Line | LS5130 | LS4419 | LS4784 | ![]() |
S II 4174.24 | 40b | 28c | - | - |
N II 4179.67 | 14c | 21c | - | - |
O II 4185.45 | 36b | 43b | - | 30a |
O II 4189.79 | 46b | 54b | 65b | - |
P III 4222.15 | 21c | 20c | - | 21a |
N II 4227.74 | 30b | 55b | 61b | 18a |
N II 4236.86 | 51a | 55a | 66c | 28a |
N II 4236.98 | ||||
N II 4241.78 | 49a | 86a | 105b | - |
S III 4253.59 | - | 118a | - | - |
O II 4253.90 | ||||
O II 4254.13 | ||||
C II 4267.02 | 262a | 245a | 305a | - |
C II 4267.27 | ||||
O II 4267.71 | ||||
O II 4275.56 | - | 97b | 123a8 | - |
O II 4277.40 | - | 72b | - | 16a |
S III 4284.99 | - | - | 47c | 28a |
S II 4294.43 | 35b | 50b | 37c | - |
O II 4294.79 | ||||
O II 4303.84 | 30c | 62a | 54c | 26a |
C II 4313.30 | 17c | - | - | - |
O II 4317.14 | 58a | 76a | 64b | 56a |
C II 4317.26 | ||||
O II 4319.63 | 48a | 78a | 63b | 52a |
O II 4325.76 | 40b | 23b | 29c | 20a |
O II 4345.56 | 30c | - | - | 24a |
O II 4347.42 | 26c | - | - | 21a |
Ar II 4348.11 | 32c | - | - | - |
O II 4349.43 | 66b | - | - | 63a |
O II 4351.26 | 31c | - | - | 41a |
O II 4366.89 | 60a | 61a | 32c | 48a |
O II 4369.27 | - | - | - | 14a |
Species and Line | LS5130 | LS4419 | LS4784 | ![]() |
C II 4372.49 | 52a | 91a | 67c9 | - |
C II 4374.27 | 32b | - | - | - |
Fe III 4395.78 | 40b | 60a | 58b | - |
O II 4395.94 | ||||
C II 4411.20 | 38b | 63b | - | - |
C II 4411.52 | ||||
O II 4414.90 | 84a | 95a | 52c | 75a |
O II 4416.97 | 64a | 76a | 28c | 61a |
Fe III 4419.59 | 44b | 32c | 35c | 34a |
Fe III 4430.95 | 33c | - | 39c | 24a |
N II 4432.74 | 38c | - | 51c | 20a |
He I 4437.55 | 126a | 103a | 135a | 118a |
N II 4442.02 | - | - | 38c | 9a |
N II 4447.03 | 60a | - | 72b | 15a |
O II 4448.19 | ||||
O II 4452.37 | 31c | - | 29c | 25a |
Al III 4479.89 | 55a | - | - | 38a |
Al III 4479.89 | ||||
Mg II 4481.13 | 202a | 196a | 265a | 160a |
Mg II 4481.33 | ||||
Al III 4512.54 | 57a | 65a | 63a | 39a |
Al III 4528.91 | 73a | 144a10 | 135a10 | 66a |
Al III 4529.20 | ||||
N II 4530.40 | 29c | - | - | 25a |
N II 4552.53 | 174a | 199a | 187a | 114a |
Si III 4552.62 | ||||
Si III 4567.82 | 150a | 152a | 151a | 102a |
Si III 4574.76 | 81a | 97a | 90a | 69a |
O II 4590.97 | 52a | 87a | 67a | - |
O II 4596.18 | 59a | 74a | 52a | 47a |
N II 4601.48 | 62b | 107a | 90a | 37a |
O II 4602.13 | 15a | |||
N II 4607.16 | 56a | 64a | 68a | 35a |
O II 4609.44 | 41b | 63a | 50a | 30a |
Species and Line | LS5130 | LS4419 | LS4784 | ![]() |
N II 4613.87 | 55a | 79a | 56a | 27a |
C II 4618.40 | 137c | 82a | 213a11 | - |
C II 4619.23 | ||||
N II 4621.29 | 41b | 73a | - | 26a |
N II 4630.54 | 92a | 122a | 123a | |
O II 4638.86 | 61b | 81a | 71a | 54a |
O II 4641.82 | 74b | 171a4 | 172a4 | 76a |
N II 4643.09 | 58b | - | - | 42a |
C III 4647.42 | - | 230a 5 | 192a5 | - |
O II 4649.14 | 90a | - | - | 92a |
O II 4650.84 | 64b | - | - | 52a |
O II 4661.63 | 67a | 94a | 75a | 54a |
O II 4673.74 | 26c | - | - | 19a |
O II 4676.24 | 44c | 45c | - | 46a |
N II 4678.10 | 24c | - | - | - |
N II 4694.70 | 15c | - | - | 9a |
O II 4696.35 | 15c | - | - | 12a |
O II 4699.00 | 42a | 69c | 37a | - |
O II 4699.22 | ||||
O II 4703.16 | - | 66a5 | 53a5 | - |
O II 4705.35 | 47b | - | - | - |
O II 4710.01 | 30b | 40c | - | 20a |
He I 4713.14 | 289a | 283a | 250a | 271a |
He I 4713.37 | ||||
Si III 4716.65 | 41b | 52b | - | 12a |
N II 4779.67 | 18c | 36c | 44b | 10a |
C II 4802.70 | 28c | - | - | 15a |
N II 4803.29 | 37c | 61a | 74a | 18a |
Ar II 4806.07 | 29c | 21c | - | 18a |
Si III 4813.30 | 31b | 41a | 75b12 | 21a |
S II 4815.52 | 54b | 27b | - | 27a |
Si III 4819.72 | 55b | 55a | 50a | 40a |
Species and Line | LS5130 | LS4419 | LS4784 | ![]() |
S II 4824.07 | 12c | - | 18c | - |
Si III 4828.96 | 44b | 64a | 59a | - |
Ar II 4879.14 | 28b | - | 19c | - |
O II 4890.93 | 21c | 22b | - | 11a |
O II 4906.83 | 38b | - | 29b | - |
O II 4941.07 | 25b | 38b | 14c | 15a |
O II 4943.00 | 39b | 52b | 45c | 28a |
S II 4991.94 | 16c | - | 24c | 22a |
N II 4994.36 | 31b | 55a | 59b | - |
S II 5009.62 | 32b | - | - | - |
N II 5010.62 | 40b | 40c | - | - |
S II 5014.03 | 52b | - | - | - |
He I 5015.68 | 322a | 319a | 354a | - |
S II 5032.41 | 66b | 56a | 91a | - |
Si II 5041.03 | 24c | - | 30c | 15a |
N II 5045.09 | 54b | 59a | 47c | - |
He I 5047.74 | 207a | 163a | 179b | - |
Si II 5055.98 | 35c | - | 56b | 22a |
Si II 5056.31 | ||||
S II 5103.30 | - | - | 20c | - |
C II 5132.96 | 97a | 74a | 99a | - |
C II 5133.29 |
Notes: |
1. Blend of O II 3982.72 and S III 3983.77. |
2. Includes the O II, N II and C II lines at 4075-4076AA. |
3. Blended with Ar II 4131.71. |
4. Blended with N II 4643.09. |
5. Blended with O II 4649.14 & 4650.84. |
6. Blended with O II 4705.35. |
7. Also includes C II 3920.69. |
8. Blended with O II 4277.40. |
9. Blended with C II 4374.27. |
10. Blended with N II 4530.40. |
11. Blended with N II 4621.29. |
12. Blended with S II 4815.52. |
Species and Line | LS5127 | ![]() |
O II 3911.96 | 50c | 12a |
C II 3918.98 | 340b | 72a |
N II 3919.01 | - | |
O II 3919.28 | 8a | |
C II 3920.69 | 88a | |
O II 3954.37 | 44c | 20a |
He I 3964.73 | 210c | 132a |
N II 3995.00 | 108a | 40a |
N II 4041.31 | 57b | 19a |
N II 4043.53 | 59b | 9a |
O II 4069.62 | 182a | 16a |
O II 4069.89 | 19a | |
O II 4072.16 | 27a | |
C II 4074.48 | 185a | 15a |
C II 4074.85 | 9a | |
O II 4075.86 | 32a | |
C II 4076.53 | - | |
N II 4076.91 | - | |
O II 4119.22 | 324a | - |
O II 4120.28 | ||
He I 4120.81 | ||
He I 4120.99 | ||
O II 4121.46 | ||
Fe III 4122.06 | ||
Fe III 4122.98 | ||
He I 4168.97 | 70b | 81a |
N II 4241.78 | 67a | 16a |
S III 4253.59 | 69a | 13a |
O II 4253.90 | 4a | |
O II 4254.13 | 2a | |
C II 4267.02 | 338a | 184a |
C II 4267.27 | ||
O II 4267.71 | ||
Si III 4267.80 | 26a | |
S II 4294.43 | 43a | 24a |
O II 4294.79 | 3a | |
O II 4303.84 | 39a | 9a |
Species and Line | LS5127 | ![]() |
O II 4325.76 | 55c | 6a |
C II 4325.83 | - | |
C II 4326.16 | - | |
S III 4361.53 | 39c | 4a |
O II 4366.89 | 42c | 16a |
C II 4372.49 | 48c | - |
Fe III 4395.78 | 43b | 15a |
O II 4414.90 | 101b | 28a |
O II 4416.97 | 23a | |
Fe III 4419.59 | 53a | 23a |
He I 4437.55 | 101a | 102a |
N II 4447.03 | 65b | 18a |
S II 4463.58 | 36b | 14a |
Mg II 4481.13 | 423a | 205a |
Mg II 4481.33 | ||
Al III 4512.54 | 60a | 20a |
S II 4524.95 | 47b | 22a |
Al III 4528.91 | 100b | 6a |
Al III 4529.20 | 28a | |
N II 4530.40 | 7a | |
Fe II 4549.47 | 49b | - |
S II 4552.38 | 191a | 71a |
Si III 4552.62 | ||
Si III 4567.82 | 156a | 59a |
Si III 4574.76 | 93a | 27a |
Fe II 4583.83 | 39b | 12a |
O II 4590.97 | 39c | 17a |
O II 4596.18 | 32c | 12a |
Species and Line | LS5127 | ![]() |
N II 4601.48 | 49c | 15a |
N II 4607.16 | 47c | 14a |
C II 4618.40 | 114c | - |
C II 4619.23 | ||
N II 4621.29 | ||
Si II 4621.42 | ||
Si II 4621.72 | ||
Fe II 4629.34 | 107a | - |
N II 4630.54 | ||
O II 4638.86 | 43a | 21a |
O II 4641.82 | 108a | 24a |
N II 4643.09 | 17a | |
S II 4648.17 | 170a | 5a |
O II 4649.14 | 34a | |
O II 4650.84 | 17a | |
S II 4656.74 | 26b | 10a |
O II 4661.63 | 65b | 15a |
Al III 4663.05 | 13a | |
He I 4713.14 | 306a | 239a |
He I 4713.37 | ||
Si III 4813.30 | 81b | 8a |
S II 4815.52 | 40a | |
Si III 4819.72 | 47c | - |
S II 5014.03 |
522a | 38a |
He I 5015.68 | 233a | |
S II 5032.41 | 128a | 47a |
Si II 5041.03 | 85a | 42a |
N II 5045.09 | 190a | |
S II 5047.28 | ||
He I 5047.74 | ||
Si II 5055.98 | 182a | 82a |
Si II 5056.31 |