Issue |
A&A
Volume 508, Number 1, December II 2009
|
|
---|---|---|
Page(s) | 401 - 408 | |
Section | Stellar atmospheres | |
DOI | https://doi.org/10.1051/0004-6361/200912518 | |
Published online | 15 October 2009 |
A&A 508, 401-408 (2009)
New identified (3H)4d-(3H)4f transitions of Fe II from UVES spectra of HR 6000 and
46 Aquilae
,![[*]](/icons/foot_motif.png)
F. Castelli1 - R. L. Kurucz2 - S. Hubrig3
1 - Istituto Nazionale di Astrofisica-
Osservatorio Astronomico di Trieste, via Tiepolo 11,
34131 Trieste, Italy
2 -
Harvard-Smithsonian Center for Astrophysics, 60 Garden Street,
Cambridge, MA 02138, USA
3 -
Astrophysical Institute Potsdam, An der Sternwarte 16, 14482 Potsdam, Germany
Received 18 May 2009 / Accepted 4 September 2009
Abstract
Aims. The analysis of the high-resolution UVES spectra of
the CP stars HR 6000 and 46 Aql has revealed the presence of
an impressive number of unidentified lines, in particular in the
5000-5400 Å region. Because numerous 4d-4f transitions of Fe II lie
in this spectral range, and because both stars are iron overabundant,
we investigate whether the unidentified lines are Fe II.
Methods. ATLAS12 model atmospheres with parameters
= 13 450 K,
= 4.3 and
= 12 560 K,
= 3.8
were computed for the individual abundances of the stars HR 6000
and 46 Aql, respectively, to use them as spectroscopic sources to
identify Fe II lines and determine Fe II gf-values. After identifying several unknown lines in the stellar spectra as (3H)4d-(3H)4f transitions of Fe II, we derived astrophysical
-values
for them. The energies of the upper levels were assigned on the basis
of both laboratory iron spectra and predicted energy levels.
Results. We determined 21 new levels of Fe II with energies between 122 910.9 cm-1 and 123 441.1 cm-1. They allowed us to add 1700 new lines to the Fe II
linelist in the wavelength range 810-15 011 Å. Many of these
lines are sufficiently strong to contribute to the spectra of
Population I late B-type stars, even when their iron abundance is
subsolar. In the 5000-6000 Å region discussed in this paper, the
astrophysical and computed -values
show good general agreement and greatly improve the synthetic spectrum
of both HR 6000 and 46 Aql. However, many features remain
unidentified indicating that further work to classify Fe II high energy levels has still to be done
Key words: line: identification - atomic data - stars: atmospheres - stars: chemically peculiar - stars: individual: HR6000 - stars: individual: 46 Aquilae
1 Introduction
The analysis of the UVES spectrum of the chemically peculiar star HR 6000
performed by Castelli & Hubrig (2007) identified a large number of
unidentified lines across the entire observed range from
3050 Å to 9460 Å. Most impressive regions were those at 4404-4411 Å and 5000-5400 Å.
An attempt to identify unknown lines in the 5130-5136 Å interval using
the available lists of predicted Fe II lines (version 2003)
was laborious and unsuccessful.
An analysis of both the list of unidentified lines and of the plot of the
observed and computed spectra available at the Castelli
web-site
led Johansson to identify about half of them as transitions
between high excitation levels of Fe II. These identifications were
made
using unpublished line lists that he had obtained from laboratory spectra.
Johansson (2009) was able to identify a group of 13 lines
concentrated in the 4404-4411 Å interval as belonging to
the multiplet 4s(7S)4d 8D-4s(7S)4f 8F.
The upper terms have energies of between 132 145 cm-1 and
132 158 cm-1, and are therefore well above the Fe II ionization
limit of
130 563 cm-1. Another example of new identifications
can be found in Castelli et al. (2008),
where some unknown lines in the 5175-5180 Å interval of HR 6000
were identified with (3H)4d-(3H)4f transitions of
Fe II.
In this case, the energy of the upper levels is of the order of
123 000 cm-1, and therefore just below the ionization limit.
Laboratory spectroscopic sources produce emission lines and must populate upper energy levels to produce such a spectrum. Most stars have spectra with lines in absorption, so that the line strength is determined from the lower energy level populations; for this reason, stellar lines are stronger than laboratory lines. HR 6000 and 46 Aql are bright and can be observed at high resolution and high signal-to-noise ratio. They have low projected rotation velocity, 1.5 km s-1 for HR 6000, and 1.0 km s-1for 46 Aql. Thus, blending is minimal and wavelengths and line strengths can be determined well by fitting the spectrum.
It is likely that most of the unidentified lines in these stars are those of Fe II. Some could be due to other overabundant elements, in particular P II, Mn II, and Xe II. However, as far as Mn II is concerned, the unidentified lines are either weaker or do not appear at all in the HgMn star HD 175640, which is iron weak ([Fe/H] = -0.25) and manganese overabundant ([Mn/H] = +2.4) (Castelli & Hubrig 2004; Castelli et al. 2008). It is also improbable that such a large number of unidentified lines, mostly concentrated in the 5000-6000 Å region, are not also due to Fe II. Since these lines are present at subsolar iron abundances, they must be present in all Population I late B-type stars, or in any object with strong Fe II lines, but they are normally smeared out by rotation and difficult to see.
Because every new energy level accounts for hundreds of new lines throughout the
spectrum from the UV through to the IR wavelength regions, we extended to a wider
range previous work for the 5175-5181 Å interval
with the aim of increasing the number of the classified Fe II levels.
We used identifications both based on laboratory spectra and
derived from predicted energy levels.
Furthermore, because numerous high-excitation lines from the
3d6(5D)4d-3d6(5D)4f transitions
are observed in the 5000-5400 Å interval, we examined the
computed -values used in the synthetic spectra computations.
We compared them with both experimental (Johansson 2002) and astrophysical
values, which we derived from the lines of both HR 6000 and 46 Aql.
We used
two stars to check the consistency and estimate the reliability of
the astrophysical oscillator strengths. We then derived
astrophysical
-values
for the new identified (3H)4d-(3H)4f lines.
Finally, we show an example of synthetic spectrum computed both with old and
new line lists.
To use HR 6000 and 46 Aql as spectroscopic sources for
identifying Fe II lines, to determine Fe II -values, and to
compute synthetic spectra, we fixed at first the model atmosphere and
the abundances for each star. Special care was devoted to deriving the iron abundance.
2 The stars HR 6000 and 46 Aql
Both HR 6000 (HD 144667) and 46 Aql (HD 186122)
are B-type CP stars.
HR 6000 was extensively studied by Castelli & Hubrig (2007) in the
3050-9460 Å region, and
46 Aql, by Sadakane et al. (2001) in the 5100-6400 Å interval.
Because model atmospheres are needed to determine astrophysical
-values, we revised
the model atmosphere and the abundances of HR 6000 and determined
the model atmosphere and the abundances for 46 Aql on the basis
of our observations.
2.1 Observations
The observations of HR 6000 are described in Castelli & Hubrig (2007). Those for 46 Aql were obtained in the framework of the same observational program (ESO prg. 076.D-0169(A)). They are of the same quality and were reduced with the same procedures used for HR 6000.
For this paper, we revised our previous analysis of HR 6000
(Castelli & Hubrig 2007) because further investigation
indicated that the Balmer profiles used by us to derive the
model parameters are too affected in the UVES spectra
by imperfections related to the echelle orders. As we show in
Appendix A, the spectral distortions for H,
H
,
and, to a less extent, for H
are
so significant that it is impossible to determine a reliable
continuum level for the observed profiles. Only H
does
not appear to be affected by these problems, so it could be used for
the analysis with some confidence.
2.2 Model parameters and abundances
In Castelli & Hubrig (2007), the model parameters of HR 6000
(
= 12 850 K,
= 4.10,
km s-1) were
derived from both Balmer profiles and
Fe I and Fe II ionization equilibrium.
In this paper, the model parameters for both HR 6000 and
46 Aql were obtained from
the Strömgren photometry, from the requirement that there was no
correlation between the Fe II abundances derived from high and
low excitation lines, and from the constraint of
Fe I-Fe II ionization equilibrium.
This type of determination led to a revision of the model parameters
for HR 6000.
Table 1: Observed and dereddened Strömgren indices for HR 6000 and 46 Aql.
The parameters of the two stars obtained from the Strömgren photometry
are given in Table 1. They were derived with the method described in
Castelli & Hubrig (2007), where the reddening for HR 6000 is also discussed.
The observed indices were taken from the Hauck & Mermilliod (1998)
catalogue.
ATLAS9 models with parameters [13 800 K, 4.3] for HR 6000 and [12 750 K, 3.8]
for 46 Aql were computed for solar abundances and zero microturbulent
velocity.
The iron abundance was derived from the equivalent
widths of a selected sample of high-excitation Fe II lines
with experimental
-values available from Johansson (2002).
They
are due to (5D)4d-(5D)4f transitions and are listed
in Table 2.
The equivalent widths of these lines, as well as of the other lines
discussed in this paper, were measured by
integrating the residual fluxes over the profiles.
Abundances were obtained
with a Linux version (Castelli 2005) of the WIDTH code (Kurucz 1993).
We note that there are no measured equivalent widths
for the line at 5100.734 Å because
it is part of a strong blend formed by four Fe II lines.
The line was included in Table 2 to show the whole set
of (5D)4d-(5D)4f transitions for which
experimental
-values are available.
The lines from Table 2 are
particularly well suited to providing the iron abundance because they
are rather insensitive to the model parameters. In fact, for differences
in
of 500 K, the abundance difference is less than 0.05 dex;
for differences in
of 0.2 dex, the abundance difference is
of the order of 0.01 dex.
For example, for HR 6000, the Fe II abundance from ATLAS9
models computed for
= 4.3, and
= 13 300 K, 13 800 K, and 14 300 K,
is -3.66 dex , -3.65 dex (Table 2), and -3.61 dex, respectively.
For models computed for
= 13 800 K, and
= 4.1, 4.3, and 4.5, the iron
abundance is -3.65 dex, -3.65 dex (Table 2), and -3.64 dex.
Similar results were obtained for 46 Aql.
Table 2:
Iron abundance derived from a selected sample of high excitation (4d-4f) Fe II lines
with experimental -values and ATLAS9 models.
After having determined the iron abundance, we checked
the model parameters by inquiring whether the adopted model gives
the same abundance for Fe I lines and for a large sample
of Fe II lines including low-excitation transitions.
To this purpose, we measured the
equivalent widths of the Fe I and Fe II lines
listed in Table B.1 and derived the corresponding abundances.
They are -3.68 0.06 dex and -3.68
0.15 dex,
respectively, for HR 6000 and
-3.93
0.06 dex and -3.91
0.11 dex,
respectively, for 46 Aql.
For both stars, the ionization equilibrium condition is fulfilled
within the error limits, and the Fe II abundance derived
from the large
sample of lines listed in Appendix B agrees, within the error limits,
with the abundance yielded
by the small sample of high-excitation Fe II lines listed in Table 2.
We conclude that the [13 800, 4.3] ATLAS9 model for HR 6000 and the [12 750 K, 3.8]
ATLAS9 model for 46 Aql not only reproduce the respective Strömgren colors,
but also satisfy the constraints of both the Fe I-Fe II
ionization equilibria and
Fe II abundance, independent of the excitation potential.
The H
profile is also fairly well reproduced by the models in
both stars.
The ATLAS9 models were used to obtain abundances for elements other than iron. For He I, we adopted the lines listed in Castelli & Hubrig (2004) and analyzed them as described in that paper. For the other elements, the lines listed in Table B.1 were used. For most lines, equivalent widths were measured. For weak lines or lines that are blends of transitions belonging to the same multiplet, such as Mg II 4481 Å, and most O I profiles, we derived the abundance from the line profiles with the synthetic spectrum method. Synthetic spectra were computed with a Linux version (Sbordone et al. 2004) of the SYNTHE code (Kurucz 2005). When no lines were observed for a given element an upper abundance limit was determined by reducing the intensity of the computed line at the noise level. To compute synthetic spectra, rotational velocities equal to 1.5 km s-1 and 1.0 km s-1 were adopted for HR 6000 and 46 Aql, respectively. They were derived by comparing the observed and computed Mg II profiles at 4481 Å.
Because the abundances found for most elements were far from solar,
we computed ATLAS12 models (Kurucz 2005) for the individual abundances
with the same parameters as determined for the ATLAS9 models.
The structure of the ATLAS12 models is heavily affected by the large iron
overabundance,
while the helium underabundance, although large, has a negligible effect,
in contrast to what was wrongly stated in Castelli & Hubrig (2007).
As a consequence, the Fe I-Fe II ionization equilibrium is
no longer achieved by the ATLAS12 models unless the parameters are changed.
Because the gravity affects the wings of the Balmer profile more than the temperature does,
we kept fixed the gravity which reproduces the H
profile rather well and modified
until the Fe I abundance
from the lines listed in Appendix B agrees with the Fe II abundance obtained from the lines shown in Table 2.
The final ATLAS12 parameters are
= 13 450 K,
= 4.3, and
= 0.0 km s-1for HR 6000 and
= 12 560,
= 3.8, and
km s-1 for 46 Aql.
The computed indices (b-y), m, and c for HR 6000 are -0.064, 0.126,
and 0.535, respectively. For 46 Aql, they are -0.052, 0.116, and 0.662.
For both stars, the observed (b-y) is reproduced by the models
within the observational uncertainties, while the c index is not.
This means
that the models are able to predict the optical spectrum but not
the spectrum shortward of the Balmer discontinuity.
We note that we could simply have used the ATLAS9 models computed for solar abundances with the parameters given in Table 2, which reproduce both the Strömgren colors and the criteria for Fe I-Fe II ionization equilibrium. However in this case, the number densities of the input model are different from those computed by the final synthetic spectrum on the basis of non solar abundances for several elements, in particular for helium which heavily affects the state equation results. We preferred to use consistent computations in model and synthetic spectra rather than use different sets of abundances in the two cases, even though predictions more close to obsevations may be achievable when different abundances are indeed used for model atmosphere and synthetic spectrum.
Table 3 summarizes the final
stellar abundances used to compute ATLAS12 models and synthetic
spectra. For HR 6000, abundances that differ from previous
determinations (Castelli & Hubrig 2007) by 0.2 dex or more are those for
He I (+0.20), Ca II (+0.2),
Ti II (+0.30),
Cr II (+0.27),
Mn II (+0.42), Fe I and Fe II (+0.21),
and Y II (+0.2).
For 46 Aql, abundances differing by more than 0.2 dex
from those derived by Sadakane et al. (2001) are those
for C II (-0.33), S II (+0.77), Ti II (-0.44),
and Fe I (-0.32).
The number in parenthesis is the difference
between the abundance determined in this paper and that from the other analyses.
Sadakane et al. (2001) adopted an ATLAS9 model atmosphere with parameters
= 13 000 K,
= 3.65, and
km s-1.
The abundances of 46 Aql have approximately the same pattern as
in HR 6000, but the deviations from solar values are generally lower.
The most remarkable differences between the two stars are
the overabundance of copper, zinc, and arsenic in 46 Aql, while
no lines of these elements were observed in HR 6000.
The arsenic overabundance cannot be quantified owing to the
lack of
data for As II. Lines of As II observed
in the spectrum are listed in Table B.1. To each line, we assigned
the corresponding transition on the basis of the
two separate lists of lines and energy levels taken from the NIST
database
,
.
Furthermore, Cr II is slightly overabundant
in HR 6000 ([0.3]) and underabundant in 46 Aql [-1.1], Ca II
is solar in HR 6000 and slightly underabundant in 46 Aql [-0.3],
and Y II and Hg II are less overabundant in
HR 6000 than in 46 Aql.
Table 3:
Abundances (N(elem)/Ntot) for HR 6000 and 46 Aql
from ATLAS12 models.
In both HR 6000 and 46 Aql, the He I profiles cannot be
reproduced by the same abundance. We adopted the abundance that reproduced
the wings of the lines at
3867, 4026, and 4471 Å the most closely. The cores of
all He I lines would require a lower abundance than that
reproducing the wings.
In both stars, the average abundances of phosphorous and manganese have
deviations larger than 0.2 dex. For phosphorous, these are caused
by a difference of the order of 0.5 dex between the abundance
from lines with < 5000 Å
and the abundance from lines with
Å.
For manganese, the large deviation
is due to a difference of 0.6 dex in HR 6000 and 0.4 dex
in 46 Aql between the abundances
from lines lying shortward and longward of the Balmer discontinuity.
The Mn II abundance was derived from both equivalent widths and
line profiles. In the first case,
no hyperfine structure was considered in the computations,
while in the second case hyperfine components were taken into account
for all the lines except for 3917.318 Å, for which no hyperfine constants
are available for its upper energy level.
The hyperfine components
were taken from Kurucz
.
The differences between the average abundances obtained by the two methods
are of the order of 0.01 dex.
A plausible explanation of the discrepancy in the Mn II abundance is that the model structure is inadequate to reproduce the ultraviolet spectrum, as already deduced by comparing the computed and observed c indices. Vertical abundance stratification, which is a consequence of radiative diffusion acting in CP stars (Michaud 1970), can be invoked to explain the anomalous He I line profiles as well as the phosphorus and manganese inhomogeneous abundances. A comprehensive discussion of observational evidence for the abundance stratification in CP stars is given by Ryabchikova et al. (2003). It describes both the impossibility of fitting the wings and the core of strong spectral lines with the same abundance, and the differences in the abundances measured from the lines of the same ion that form at different optical depths, as, for instance, in the case of lines lying longward or shortward of the Balmer discontinuity. Finally, for a reliable discussion on phosphorus, more oscillator strengths data of P II and P III lines observed in the optical region are needed.
In both HR 6000 and 46 Aql, the Hg II line at 3983.890 Å is due mostly to the heaviest isotope of Hg. The lines of the Ca II infrared triplet at 8498, 8542, and 8662 Å are red-shifted by 0.14 Å in HR 6000 and by 0.13 Å in 46 Aql, so indicating a non solar Ca isotopic composition. In HR 6000, emission lines of Cr II, Mn II, and Fe II were observed. In the spectrum of 46 Aql, instead, there are emission lines of Ti II and Mn II.
3 The (5D)4f and (3H)4f states of Fe II
Energy levels of a given atom are described most often by the LS coupling
in which the total orbital angular momentum L of the atom is
coupled with the
total spin angular momentum S to produce the total angular momentum
J = L + S. Some high levels, such as the (5D)4f and (3H)4f
states of Fe II,
tend to appear in pairs, so they are more accurately described by the jK coupling with the
notation [K]J, where j
is the total angular momentum of
the core and K = j
+ l is the coupling of
j
with the orbital angular momentum l
of the active electron. The level pairs correspond to the two
separate values of the total angular momentum J obtained
when the spin s =
1/2 of the active electron is added to K.
While the energy levels of the 3d6(5D)4f states are known and available for instance in the NIST database, this is not the case for the levels with the higher parent term 3d6(3H)4f. Most of these levels have not been observed in the laboratory and remain unclassified.
4 The 3d6(5D)4d-3d6(5D)4f transitions of Fe II
The 4d-4f transitions discussed in this paper appear in the
optical region, mostly between 4800 and 6000 Å.
Their presence in stellar spectra was found a long time
ago by Johansson & Cowley (1984). They are also present
in the spectra of the iron-rich peculiar stars HR 6000
and 46 Aql. We used UVES spectra of these stars to derive
values of astrophysical
that we compared with experimental and computed
values.
4.1 Experimental log gf-values
All the experimental data described in this section were made available
to Castelli by Johansson and are briefly described in
Johansson (2002).
Radiative lifetime measurements of five 3d6(5D)4f levels,
i.e., 4[6]13/2, 4[7]13/2, 4[7]15/2, 3[6]13/2,
3[5]11/2, and branching fraction measurements for 13 transitions 4d-4f with wavelengths in the 4800-5800 Å region
were performed at Lund.
Einstein coefficients A, derived by combining the measured branching fractions
with the lifetime measurements, were converted into experimental
oscillator strengths.
The 4d-4f transitions together with the experimental
-values are given in Tables 2 and 4.
While the 4f levels with
J > 11/2 decay only to 4d levels,
the 4f 3[5]11/2 level may decay to 3d
levels as well. As a consequence, the
-value
for the transitions involving the
3[5]11/2 level may be less accurate than those
related to levels with J>11/2, which have an estimated error of 0.05 dex.
4.2 Computed log gf-values
The computed -values were taken from both
Kurucz's line lists (K09, January, 2009
version)
and
Raassen & Uylings (1998) (RU98) data. We note that
Fuhr & Wiese (2006) (FW06) adopted
the RU98 data for the few (5D)4d-(5D)4f transitions that
they listed in their critical compilation.
Both K09 and RU98 results were obtained with semi-empirical methods,
although different ones. The K09 results were obtained with the use of the
Cowan (1981) atomic structure code, while the RU98 results were obtained by the
orthogonal operators method.
The computed -values are listed in Table 4.
4.3 Astrophysical log gf-values
To derive astrophysical -values we computed
synthetic profiles for the Fe II lines listed in Table 4.
We used the ATLAS12 models discussed in Sect. 2, the SYNTHE
code (Kurucz 2005),
and line lists based on the Kurucz database that
we continually update with new data when available (Castelli & Hubrig 2004).
By fixing the iron abundance to -3.65 dex for HR 6000 and
to -3.91 dex for 46 Aql (Table 3), we adjusted
the
-values in the calculated
spectrum, for the lines listed in Table 4, until observed and
computed profiles were in optimal agreement.
All the lines can be fitted well except for the very strong ones
with
higher than 0.9. Their observed cores are
stronger than those computed and can never be reproduced by the computed
spectrum because increasing the
-value broadens the wings instead of
increasing the core.
An example is the line at 5260.254 Å, which we decided to exclude from
our comparisons. The strongest Fe II lines are possibly affected by
iron abundance vertical stratification, which does not affect
the medium-strong and weak lines.
![]() |
Figure 1:
Comparison of astrophysical |
Open with DEXTER |
![]() |
Figure 2:
Calculated |
Open with DEXTER |
Figure 1 shows the difference between the astrophysical -values from
both HR 6000 and 46 Aql as function of the astrophysical
of
HR 6000. The average difference, shown by the dashed line, is
-0.019
0.042 dex, but for
single lines the difference increases with increasing
from 0.00
to 0.15 dex, in the sense that the values from 46 Aql become
larger than those from HR 6000.
The largest difference of -0.15 dex occurs for
5961.705 Å.
Because each astrophysical value reproduces the observed line
in each star well we do not have an explanation of this discrepancy.
As astrophysical -values we assumed the average of the values
obtained
from HR 6000 and 46 Aql. Astrophysical
-values for both the
two stars and the average are listed in Table 4.
4.4 Comparison of values of log gf
The difference between experimental and computed -values
versus the experimental
is shown
Figs. 2 and 3, where the computed data are those from K09 and
RU98, respectively. The average difference of 0.028
0.061 dex
yielded by the RU98 values is lower than the average difference of
0.043
0.071 dex given by the K09 values, but they agree within the
error limits. The largest difference of +0.130 in Fig. 2 is
due to the transition at 4883.292 Å. However,
the J value of the 4f upper level of this line is 11/2, which is not
high enough to ensure that there are no additional decays
to 3d levels (Johansson 2002).
This fact could affect the experimental
-value and the closer
agreement
with the RU98 data may be fortuitous. In fact, Table 4 shows that
the astrophysical
-value agrees more closely with the K09 than
the RU98 value.
![]() |
Figure 3:
Calculated |
Open with DEXTER |
![]() |
Figure 4:
Calculated |
Open with DEXTER |
Astrophysical -values are compared with
-values
from K09 and RU98 in Figs. 4 and 5, respectively.
The mean difference of 0.006
0.116 dex given by the K09 data
is fully comparable with the average difference of
-0.007
0.144 dex yielded by the RU98
-values.
In both cases, the dispersion around the mean value is
rather large. The lines giving the largest discrepancies are
different for K09 and RU98. In K09, the lines are those
at wavelengths of
5257.119 (-0.47),
5359.237 (-0.314), 5358.872 (+0.286), 5355.421 (+0.285),
5366.210 (-0.267), 5062.927 (-0.260) Å.
In RU98 they are those with wavelength
5070.583 (-0.760), 5140.689 (-0.543),
5093.783 (-0.399),
5200.798 (+0.287), 5081.898 (-0.279) Å.
The values in parentheses are the difference
(computed) -
(astrophysical).
For all of these transitions,
the difference between the astrophysical
-values
from HR 6000 and 46 Aql is less than 0.06 dex, so that
the cause of the disagreements is probably due to the computed values.
We note that Kurucz updates his calculation whenever new Fe II levels become available. The January 2009 version of the Fe II linelist used for the (5D)4d-(5D)4f transitions discussed in this paper includes only a few of the new (3H)4d-(3H)4f levels presented in Sect. 5. In the near future, a new Fe II linelist with all the new levels given in Tables 5 and 6 will be made available at the Kurucz web-site.
![]() |
Figure 5:
Calculated |
Open with DEXTER |
5 The new identifications in HR 6000 and 46 Aql
The spectrum of HR 6000 contains an enormous number of unidentified lines mostly concentrated in the 5000-5400 Å region (Castelli & Hubrig 2007). The same unidentified lines can also be observed in the spectrum of 46 Aql. Johansson (2006) remarked that a great number of unidentified lines in the plots of HR 6000 available at the Castelli web-site (see footnote 2), are also present in laboratory iron spectra. He therefore identified several unknown features in the 4000-5500 Å interval of HR 6000 as beign produced by iron. These identifications are at different levels of completeness. In a few cases, the transition is identified only as Fe , in most cases, as Fe II, and in some cases, as Fe II with both levels of the transition being classified.
5.1 The 3d6(3H)4d-3d6(3H)4f transitions of Fe II
Most of the Fe II lines classified by Johansson are due to the (3H)4d-(3H)4f transitions. Their lower excitation potential is higher than 103 800 cm-1 (12.87 eV) and their upper excitation potential is of the order of 123 000 cm-1 (15.25 eV), and therefore close to the ionization limit of 130 563 cm-1(16.19 eV).
In preliminary work, Castelli et al. (2008) provide an example in both HR 6000 and 46 Aql of four lines at 5176.711, 5177.3896, 5177.7762, and 5179.536 Å identified for the first time with Fe II (3H)4d-(3H)4f transitions. Only for two lines (at 5177.3896 and 5179.539 Å) were energies and terms known for both levels, while for the other two lines (at 5176.711 and 5179.536 Å) the term of the upper level was unknown, except for the J quantum number. Further lines and energies indicated by Johansson are those marked with a ``J'' in Table 5.
To complete and extend the number of the identifications, we proceeded
as follows. We started from the ``J'' lines.
Knowing the term, the energy of the lower level of a ``J'' line was determined
using either the NIST database (see footnote 4) or the Kurucz
linelists (see footnote 1). A given (3H)4d-(3H)4f transition was
then
searched for among the Fe II predicted lines available
in the Kurucz database (version 2007). For this search,
both the lower energy level and the J quantum number of the upper level
were used as key parameters. Once the predicted line corresponding to the new identified
transition was fixed, the predicted energy was replaced by the
energy assigned by Johansson to the level. This substitution was made
not only for the given line but also for all lines with that predicted
level as their upper level. We then compared the pattern of
the computed -values with the pattern of astrophysical
-values
for those lines. If they were similar, we accepted the identification.
If not, we tried another match. In this way, because of the J lines, we
fixed 11
new energy levels and all the transitions with a new level as an upper level.
In addition, we determined another 10 new levels from predicted lines.
For a given predicted upper level,
we searched in the spectrum for unidentified lines with similar intensity
and the same wavelength difference as the predicted lines. We selected
the lines with positive
-values, or even negative, but
close to zero. The observed
wavelength and the known lower energy level were then used to fix the upper
level of the transitions.
Table 5 lists both the (3H)4d-(3H)4f transitions originally
identified by S. Johansson and those that we derived from the above
described procedure. The letter ``J'' is associated with the first group of lines,
the letter ``K'' with the second group.
Not all the lines listed in Table 5 are observable in
the spectra, so that values of astrophysical
cannot be
assigned to all
lines. The last column of Table 5 lists lines observed in HR 6000 that were
listed as unidentified lines by Castelli & Hubrig (2007).
Figure 6 compares astrophysical -values from HR 6000 and 46 Aql.
It is the analogous to Fig. 1, but for the
(3H)4d-(3H)4f transitions. The average difference, shown
by the dashed line, is negligible, but the scatter is larger than
that obtained for the (5D)4d-(5D)4f transitions.
The reason is the very low intensity of some of the
(3H)4d-(3H)4f transitions, which makes the fitting of the
profile problematic.
As for the (5D)4d-(5D)4f transitions, we assumed
astrophysical
-values to be the average of values
obtained for HR 6000 and 46 Aql, but we excluded the lines for which
astrophysical
-values
differed by more than 0.1 dex. Astrophysical
-values
for the two stars and
the average values are listed in Table 5.
![]() |
Figure 6:
Comparison of astrophysical |
Open with DEXTER |
The astrophysical -values are compared with the calculated
-values in Fig. 7. From this comparison, we excluded
the lines at
5177.777 (-1.45),
5250.632 (-2.61), 5346.098 (-1.21), and 5420.234 (-1.76) Å,
for which the
difference between computed and
astrophysical
-values, given in parenthesis,
is larger than 1.0 dex. The upper energies of the lines are 122 952.73,
123 026.35, 123 015.40, and 123 251.47 cm-1, respectively.
The most probable explanation of these large discrepancies
is the presence of some additional unidentified component
contributing to the absorption, so that the astrophysical
-value
is largely overestimated.
The mean difference between the two sets of data is
-0.07
0.22 dex, indicating that the
astrophysical
-values are,
on average, higher than the computed
-values.
The differences for the individual lines increase with decreasing
,
namely with decreasing line intensity.
Weak lines are more difficult to be fitted by the synthetic
spectrum owing to the contribution of the noise and the
non-negligible effect of the position for the continuum.
![]() |
Figure 7:
The calculated |
Open with DEXTER |
![]() |
Figure 8:
Comparison of the UVES spectrum of HR 6000 with a synthetic
spectrum computed after this study (upper plot) and before this
study (lower plot). The black line is the observed spectrum, and
the red line is the synthetic spectrum. Nine Fe II lines corresponding
to (3H)4d-(3H)4f newly identified transitions are marked in the
upper plot. Their calculated |
Open with DEXTER |
5.2 More new Fe II identified lines
There are a few other lines that do not belong to the
(3H)4d-(3H)4f transitions, that were identified by
S. Johansson in HR 6000. They are indicated
with a ``J'' in Table 6.
The lines marked with a ``K'' were then obtained from
predicted
lines,
because of the coincidence of the term of the upper predicted level with
the term of a ``J'' line.
Table 6 shows that we fixed two new (5D)6d levels.
Some computed
-values are much weaker than
the astrophysical
-values.
Some other unknown transition is probably the main component of the
observed line so that the astrophysical
-value is unreliable.
6 Conclusions
Figure 8 compares the synthetic spectra for HR 6000 computed before and after the study presented in this paper. The interval plotted is an example of the quality of the improvement that we obtained for the whole 5100-5400 Å region. It shows that as many as nine new Fe II lines have been identified within this 10 Å range. Nevertheless, several absorption lines remain unidentified. They are probably Fe II lines whose levels have still to be fixed. A very similar plot was obtained for 46 Aql. We can infer that a large part of the unidentified lines observed in the spectra of B-type stars are due to unknown high-excitation Fe II transitions.
The new lines identified in this paper correspond to high excitation transitions of Fe II with an upper level just below the ionization limit. We fixed 21 new levels of Fe II whose energies range from 122 910.9 cm-1 to 123 441.1 cm-1, and added 1700 lines to the Fe II linelist for the range 810-15011 Å. Furthermore, Johansson (2009) identified in the spectrum of HR 6000 the Fe II lines of the multiplet 4s4d8D-4s4f8F at 4410 Å. Their lower energy level is near the ionization limit and their upper energy level is above it.
Among the newly identified high-excitation Fe II lines, several have residual flux in HR 6000 and 46 Aql of the order of 0.7 and numerous others are observable as weak absorption lines or parts of blends. The two stars are iron-overabundant stars, but these lines are also present with lower intensity in the UVES spectrum of HD 175640, a B-type peculiar star with an iron underabundance of -0.25 dex with respect to the Sun (Castelli & Hubrig 2004). This implies that Fe II lines from the new high excitation levels contribute to the spectrum of all Population I late B-type stars, even when their abundance is less than solar. The lines are clearly observable in high resolution, high signal-to-noise spectra of slowly rotating stars, while they contribute to the broad observed features in B-type stars with high rotational velocities. In general, they would appear in any object with strong Fe II lines.
We conclude that we have clarified the nature of several unidentified lines observed in the optical spectra of B-type stars and concentrated mostly in the 5000-5400 Å region (see also Wahlgren et al. 2000), but that a large amount of work remains to be done to reproduce stellar observations well. More than 1000 energy levels of Fe II are known, but we have seen that they are not enough. Ignorance of them and of the transitions involved still remain an outstanding shortcoming affecting the model atmosphere and synthetic spectra computations.
References
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- Younger, S. M., Fuhr, Y. R., Martin, G. A., & Wiese, W. L. 1978, J. Phys. Chem. Ref. Data, 7, 495 [NASA ADS]
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Online Material
Appendix A: The Balmer lines of HR 6000 observed on the UVES spectra
Figure A.1 shows the UVES spectra of HR 6000 reduced
by the UVES pipeline
Data Reduction Software (version 2.5,
Ballester et al. 2000) that were used by Castelli & Hubrig (2007) and
also in this paper. All spectra are FLUXCAL-SCIENCE products. Those at
3290-4520 Å and 4780-5650 Å are flux-calibrated
spectra in 10-16 erg s-1 cm-2 A-1 corrected for terrestrial
extinction. The red spectrum at 5730-7560 Å is in non-physical units
``quasi-ADU'' because the flux calibration
procedure is not implemented in the reduction software for the REDL and REDU
data taken with the red mosaic CCD's.
The sizable distortions in UVES spectra make
the difficulty in drawing a true
continuum over H
and H
evident.
The use of H
also
causes problems because of the position of this line at the left end
of the spectrum order. Only H
does not have
significant problems.
Computed spectra from the final ATLAS12 model ([13450, 4.3], Sect. 2.2) are also plotted in Fig. A.1 to show the different slopes of the observed and computed continua. The computed fluxes were scaled by a given arbitrary quantity to be roughly overimposed on the UVES spectra.
![]() |
Figure A.1:
The observed UVES spectra of HR 6000 (black line)
are plotted together the computed spectra (red line) in order to show
the different slopes of the observed and computed continua. The computed
fluxes are scaled by a given arbitrary quantity to be roughly overimposed
on the UVES spectra. The ATLAS12
final model with parameters
|
Open with DEXTER |
Appendix B: Lines used for the abundance analysis
Table B.1 lists the lines examined in the spectra
of HR 6000 and 46 Aql to derive the stellar abundances.
The wording ``not obs'' is given for
lines not present in the spectra, while the wordings ``profile''
and ``blend'' are given for lines well observed in the spectra
that do not have measurable equivalent widths either because they
are too weak to be measurable or because other components affects
the line. These wordings also indicate lines for which
adequate equivalent widths cannot be computed, as in the cases of
Mg II at 4481 Å and most O I lines, which are
blends of transitions belonging to the same multiplet.
The abundances from the final ATLAS12 models derived from the equivalent
widths or from the profiles are given in the table, as well as
upper abundance limits for lines not observed,
but predicted at solar abundance by the synthetic spectrum.
For Fe I and Fe II, -values were taken
from Fuhr & Wiese (2006) (FW06) when available. Otherwise Kurucz's
last determination was adopted (Kurucz 2009), except for Fe II
at 5257.119 Å. In this case, the previous values (Kurucz 2007)
produce synthetic profiles in closer agreement with
the observations.
Table B.1: Analyzed lines in the stellar spectra, measured equivalent widths in mÅ, and relative abundances.He I is not included.
Table 4:
Astrophysical -values for a sample of (5D)4d-(5D)4f lines of Fe II observed in HR 6000 and 46 Aql.
The values of
from the two stars are averaged and compared with experimental
-values from Johansson (2002)
and calculated
-values
from Kurucz (2009) (K09) (footnote 7) and from Raassen & Uylings (1998) (RU98).
Table 5: Lines due to (3H)4d-(3H)4f transitions of Fe II.
Table 6: More new Fe II identified lines.
Footnotes
- ...
46 Aquilae
- This study is the result of a collaboration with Sveneric Johansson, who unfortunately died before this paper started to be written.
- ...
- Tables 4-6, and Appendices A and B are only available in electronic form at http://www.aanda.org
- ... II lines
- http://kurucz.harvard.edu/atoms/2601/gf2601.lines0600
- ...
web-site
- http://wwwuser.oat.ts.astro.it/castelli/hr6000/hr6000.html
- ...
catalogue
- http://obswww.unige.ch/gcpd/gcpd.html
- ...
database
- http://physics.nist.gov/PhysRefData/ADS/lines_form.html
- ...
- http://physics.nist.gov/PhysRefData/ADS/levels_form.html
- ... Kurucz
- http://kurucz.harvard.edu/atoms/2501/hyper250155.srt
- ...
version)
- http://kurucz/harvard/edu/atoms/2601/gf2601kjan09.pos
- ...
lines
- http://kurucz.harvard.edu/atoms/2601/gf2601.lines0500
- ...
- http://kurucz.harvard.edu/atoms/2601/gf2601.lines0600
- ... pipeline
- http://www.eso.org/observing/dfo/quality/UVES/pipeline/pipe_reduc.html
All Tables
Table 1: Observed and dereddened Strömgren indices for HR 6000 and 46 Aql.
Table 2:
Iron abundance derived from a selected sample of high excitation (4d-4f) Fe II lines
with experimental -values and ATLAS9 models.
Table 3:
Abundances (N(elem)/Ntot) for HR 6000 and 46 Aql
from ATLAS12 models.
Table B.1: Analyzed lines in the stellar spectra, measured equivalent widths in mÅ, and relative abundances.He I is not included.
Table 4:
Astrophysical -values for a sample of (5D)4d-(5D)4f lines of Fe II observed in HR 6000 and 46 Aql.
The values of
from the two stars are averaged and compared with experimental
-values from Johansson (2002)
and calculated
-values
from Kurucz (2009) (K09) (footnote 7) and from Raassen & Uylings (1998) (RU98).
Table 5: Lines due to (3H)4d-(3H)4f transitions of Fe II.
Table 6: More new Fe II identified lines.
All Figures
![]() |
Figure 1:
Comparison of astrophysical |
Open with DEXTER | |
In the text |
![]() |
Figure 2:
Calculated |
Open with DEXTER | |
In the text |
![]() |
Figure 3:
Calculated |
Open with DEXTER | |
In the text |
![]() |
Figure 4:
Calculated |
Open with DEXTER | |
In the text |
![]() |
Figure 5:
Calculated |
Open with DEXTER | |
In the text |
![]() |
Figure 6:
Comparison of astrophysical |
Open with DEXTER | |
In the text |
![]() |
Figure 7:
The calculated |
Open with DEXTER | |
In the text |
![]() |
Figure 8:
Comparison of the UVES spectrum of HR 6000 with a synthetic
spectrum computed after this study (upper plot) and before this
study (lower plot). The black line is the observed spectrum, and
the red line is the synthetic spectrum. Nine Fe II lines corresponding
to (3H)4d-(3H)4f newly identified transitions are marked in the
upper plot. Their calculated |
Open with DEXTER | |
In the text |
![]() |
Figure A.1:
The observed UVES spectra of HR 6000 (black line)
are plotted together the computed spectra (red line) in order to show
the different slopes of the observed and computed continua. The computed
fluxes are scaled by a given arbitrary quantity to be roughly overimposed
on the UVES spectra. The ATLAS12
final model with parameters
|
Open with DEXTER | |
In the text |
Copyright ESO 2009
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