Free Access
Issue
A&A
Volume 525, January 2011
Article Number A44
Number of page(s) 11
Section Atomic, molecular, and nuclear data
DOI https://doi.org/10.1051/0004-6361/201015006
Published online 30 November 2010

© ESO, 2010

1. Introduction

Manganese is an iron-peak element and the nucleosynthesis path that leads to its formation is relatively well understood. However, it remains unclear which objects are the main donors of manganese to the Galaxy at different times of its evolution. Nevertheless, manganese is widely used for the investigation of the chemical evolution of the disk and halo of our Galaxy (see Sobeck et al. 2006, and references therein).

Accurate atomic oscillator strengths (f-values and log (gf)) are required for the correct interpretation of the physical properties and processes in stellar and sub-stellar objects. In particular, oscillator strengths for atomic lines in the infrared (IR) spectral region are needed for the study of the spectra of late-type stars, ultracool dwarfs, and dust obscured objects such as young stars and the centre of galaxies. Accurate oscillator strengths are of particular importance to the study of late type stars and ultracool dwarfs where the spectral energy distribution peaks in the IR and is dominated by spectral lines from neutral atoms and molecules (see Lyubchik et al. 2004; Jones et al. 2005).

Over the past ten years there has been an increase in IR spectral observations of astrophysical objects with the advent of new IR spectrographs on ground based and satellite borne telescopes. However, there are only a relatively small number of experimentally measured oscillator strengths for IR spectral lines compared to the number of measured oscillator strengths for visible spectral lines available in the literature. The status of oscillator strengths in the atomic database has been discussed by Wahlgren & Johansson (2003), Johansson (2005), Brickhouse et al. (2006) and Blackwell-Whitehead et al. (2008), and there have been calls for more IR measurements including Lyubchik et al. (2004), and Bigot & Thévenin (2006).

The current laboratory atomic database for Mn i oscillator strengths is dominated by transitions in the UV and visible. Our previous oscillator strength measurements for Mn i, Blackwell-Whitehead et al. (2005a), include 44 transitions from 2090 to 27 800 Å of which six transitions are in the IR. Only eleven experimentally measured oscillator strengths for Mn i have been published for λ > 6520 Å, see Blackwell-Whitehead et al. (2005a). Booth et al. (1984) reported 58 Mn i f-values including the UV and visible resonance transitions from the 3d5(6S)4s4p z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$} and 3d5(6S)4s4p z 6PJ\hbox{$z~^6{\rm P}^{\circ}_J$} levels. However, no f-values have been published for the IR transitions from the z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$} and z 6PJ\hbox{$z~^6{\rm P}^{\circ}_J$} levels despite these transitions being amongst the strongest Mn i transitions for 10   000 < λ < 15   000 Å. In our current work we measured Mn i oscillator strengths for transitions from the z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$} and z 6PJ\hbox{$z~^6{\rm P}^{\circ}_J$} upper levels by combining branching fractions (BFs) measured by high resolution Fourier transform spectrometry with known level lifetimes (τ). These include the 12 899 Å and 12 875 Å spectral lines that are observed as strong features in the spectra of ultracool dwarf stars, see Lyubchik et al. (2007).

2. Laboratory measurements

The oscillator strengths have been determined by combining BFs with radiative lifetimes in the same manner as described in Blackwell-Whitehead et al. (2005a). The spectrum of Mn i was recorded in the UV to visible spectral range (1600 Å to 8000 Å) using the Imperial College high resolution Fourier transform spectrometer (FTS) (Pickering 2002), and in the visible to IR spectral range (3500 Å to 55 000 Å) using the 2m FTS at the National Institute of Standards and Technology (NIST) (Nave et al. 1997). The light source for the NIST and Imperial College measurements was a hollow cathode lamp (HCL) with a manganese cathode using either an argon or neon buffer gas. Two manganese cathodes were used in these measurements. The NIST cathode was an alloy of 95% manganese and 5% copper, the cathode used at Imperial College was an alloy of 88% manganese and 12% nickel. The optimal running conditions for the HCL were found to be a pressure of 340 Pa of neon with currents of 200 mA to 500 mA. Further details of the measurements and their analysis are given in Blackwell-Whitehead et al. (2005a). The spectra were recorded at a range of currents to determine an intensity versus current curve of growth for each line to determine if any lines were self absorbed, which would lead to erroneous relative line intensities. It was found that a HCL current of 200 mA was used to observe absorption free spectra of the UV and visible resonance lines and a higher HCL current of 500 mA was used for the measurement of the relatively weaker IR transitions. In addition, the line profile of each transition was fitted using published hyperfine structure (HFS) constants for the upper and lower level of the transition (Handrich et al. 1969; Dembczyński et al. 1979; Brodzinski et al. 1987; Blackwell-Whitehead et al. 2005b). The residual value of the fit was found to be the same as the background noise level indicating that no self absorption was present.

Table 1

New and remeasured laboratory oscillator strengths for Mn i.

The manganese spectra were intensity calibrated using tungsten intensity standard lamps. The Imperial College tungsten intensity standard lamp was calibrated by the National Physical Laboratory, UK, and the NIST tungsten intensity standard lamp was calibrated by Optronics Laboratories1. Both lamps have a minimum radiance uncertainty of 3 percent in the spectral region used for the intensity calibration. The tungsten spectra were recorded before and after the manganese spectra and compared to determine if the instrumental response had changed during measurement of the manganese spectrum. In each case, the instrument response did not vary by more than the uncertainty in the relative radiance of the tungsten lamp. The measured tungsten spectra were used to determine the instrument response, and this was used to calibrate the relative intensity of the Mn i lines. The Mn i line profiles were fitted by employing a centre of gravity fit using the XGremlin software by Nave et al. (1997). To observe all lines from the z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$} and z 6PJ\hbox{$z~^6{\rm P}^{\circ}_J$} upper levels, spectra were recorded in two overlapping spectral regions. The intensity calibration of the visible to IR region was placed on the same intensity scale as the UV to visible region using intermediate transitions in the 18 000 cm-1 to 18 500 cm-1 region (3d6(5D)4s a   6DJ−3d5(6S)4s4p y 6PJ\hbox{$y~^6{\rm P}^{\circ}_J$}). The intermediate transitions are from upper levels with the same configuration and comparable level energy, which indicates that the z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$}, z 6PJ\hbox{$z~^6{\rm P}^{\circ}_J$} and y 6PJ\hbox{$y~^6{\rm P}^{\circ}_J$} levels have comparable level populations. Curves of growth for the intermediate lines indicated that no self absorption was present. Furthermore, an estimate of the change in level population between the z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$}, z 6PJ\hbox{$z~^6{\rm P}^{\circ}_J$} and y 6PJ\hbox{$y~^6{\rm P}^{\circ}_J$} levels was measured by comparing the intensity ratios between the a   6DJy 6PJ\hbox{$_{J}{-}y~^6{\rm P}^{\circ}_J$} transitions and transitions from the z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$} and z 6PJ\hbox{$z~^6{\rm P}^{\circ}_J$} upper levels under different HCL conditions. The intensity ratios did not vary by more than the uncertainty in the measured relative line intensities, indicating that the level populations remained constant to within a few percent.

3. Laboratory oscillator strengths

Table 1 presents the BFs, transition probabilities and oscillator strengths. The oscillator strengths were obtained by combining the BFs with the published radiative lifetimes of Kronfeldt et al. (1985) and Schnabel et al. (1995). The Mn i Ritz wavenumbers in Table 1 are determined from the upper and lower energy level values from the NIST atomic Spectra Database (Ralchenko et al. 2009), which are taken from the term analysis of Catalán et al. (1964). The air wavelengths in Table 1 have been determined with the Edlén (1966) equation, and include the more recent update for the refractive index of air, Eq. (3) in Birch & Downs (1994). All lines measured in this work have a peak signal to noise ratio of more than 100, and the BF uncertainties are dominated by the tungsten lamp calibration uncertainty and the uncertainty in the intensity calibration “cross-over” between the two spectral regions. The uncertainty in the oscillator strength is determined from the BF and lifetime uncertainty using the criteria discussed by Sikström et al. (2002) and follows the NIST guidelines for evaluating and expressing uncertainty (Taylor & Kuyatt 1994).

It can be seen that our oscillator strengths for the transitions from the z 6PJ\hbox{$z~^6{\rm P}^{\circ}_J$} levels agree, to within the uncertainty, with the previous UV and visible measurements by Booth et al. (1984). There is also a good agreement with our laboratory log (gf) values for the transitions from the z 6PJ\hbox{$z~^6{\rm P}^{\circ}_J$} levels and the semi-empirical calculated log (gf) values of Kurucz & Bell (1995). The measured log (gf) values for UV transitions from the z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$} levels agree with both Booth et al. (1984) and Kurucz & Bell (1995) to within the uncertainties. However, the semi-empirical log (gf) values of Kurucz & Bell (1995) for IR transitions from the z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$} levels are approximately 30% (0.12 dex where the unit dex is log 10 of the ratio of the two values) stronger than our values. The difference between the measured log (gf)s and the semi-empirical log (gf)s is larger than the uncertainty in the measured values. It is possible that the semi-empirical calculations predict more level mixing than is present in the actual system. If the z 4PJ\hbox{$z~^4{\rm P}^{\circ}_J$} levels have less mixing than predicted by the semi-empirical calculations of Kurucz & Bell (1995), then the measured log (gf)s will be weaker.

The effect of hyperfine splitting on the fine structure levels and on the line profiles of Mn i transitions in the IR can be seen in an example shown in Fig. 1. The hyperfine splitting increases the width and decreases the peak intensity of the line profile. Prochaska & McWilliam (2000) discuss the importance of including hyperfine splitting in the analysis of elemental abundances in stars. In particular, if HFS is not taken into account, the chemical elemental abundance may be underestimated or it may erroneously be assumed that the line is blended with some unknown feature. To assist in the correct interpretation of the oscillator strengths in Table 1 we provide wavenumbers, wavelengths and oscillator strengths in Table 2 for the HFS component lines in the IR transitions 3d6(5D)4s a   6DJ − 3d5(6S)4s4p z 6PJ\hbox{$z~^{6}{\rm P}^{\circ}_{J}$} and 3d6(5D)4s a   4DJ − 3d5(6S)4s4p z 4PJ\hbox{$z~^{4}{\rm P}^{\circ}_{J}$}.

thumbnail Fig. 1

The upper plot shows the hyperfine splitting of the fine structure levels 3d6(5D)4s a 4D7/2 − 3d5(6S)4s4p(3P) z 4P5/2\hbox{$z^{~4}{\rm P}^{\circ}_{5/2}$} with allowed hyperfine transitions. The lower plot shows the hyperfine split profile of the transition observed in the uncalibrated laboratory spectrum at 12 975 Å, together with an indication of the positions and relative line strengths of the individual HFS transitions. A complete list of the Ritz wavelength and log (gf) for each HFS transition is available in Table 2.

The wavenumber of each transition from the upper hyperfine structure level to the lower hyperfine structure level, σHFS   trans, in Table 2 is determined from the wavenumber for the fine-structure transition, σFS, using: σHFS trans=σFSKlAl2+KuAu2\begin{equation} \label{hfs} \sigma_{\rm HFS~trans}= \sigma_{\rm FS} - \frac{K_{\rm l} A_{\rm l}}{2} + \frac{K_{\rm u} A_{\rm u}}{2} \end{equation}(1)where Au and Al are the magnetic dipole hyperfine interaction constants for the upper and lower fine structure levels; and K is defined as: K=F(F+1)J(J+1)I(I+1)\begin{equation} K=F(F+1)-J(J+1)-I(I+1) \end{equation}(2)where F is the quantum number associated with the total angular momentum of the electrons, J, and the nuclear spin, I. For manganese the spin of the nucleus I = 5/2. Equation (1) excludes the contribution from the electric quadrupole hyperfine interaction constant B which is relatively small when compared to the magnetic dipole hyperfine interaction constant A for the levels considered in this paper, see Kuhn (1964).

Several of the IR transitions included in our work (12 899.8, 12 975.9, 13 281.5, 13 319.0, and 13 642.9 Å) have been studied by Meléndez (1999) who has determined the wavelength of individual hyperfine structure components by analysing Mn i transitions in solar photospheric spectra. However, Meléndez notes that several of these lines have blended features and he does not provide HFS constants for the upper and lower levels of the transitions.

Table 2

Wavenumber and wavelength of the HFS component lines in the IR transitions 3d6(5D)4s a   6DJ−3d5(6S)4s4p z 6PJ\hbox{$z~^{6}{\rm P}^{\circ}_{J}$} and 3d6(5D)4s a   4DJ−3d5(6S)4s4p z 4PJ\hbox{$z~^{4}{\rm P}^{\circ}_{J}$}.

thumbnail Fig. 2

The best fit to the observed solar spectrum feature found from the minima of Eq. (3) for the Mn i line at 12 899 Å.

thumbnail Fig. 3

The best fits to the observed solar spectrum features found from the minima of Eq. (3) for Mn i lines at 12 899 Å, 12 975 Å, 13 281 Å (left column) and 13 293 Å, 13 318 Å, 13 415 Å (right column). The wavelength scale is taken from the observed spectrum of the Sun.

4. Modelling Mn i lines in the spectra of the Sun, Arcturus and ultracool dwarfs

We have computed the synthetic spectra of late-type stars and brown dwarfs using the WITA6 programme (Pavlenko 2000). The model atmospheres for the Sun (spectral classification = G2V, Teff/log (g) = 5770/4.44, and abundances from Gurtovenko & Kostik 1989), and Arcturus (K2III, 4300/1.5, and abundances from Peterson et al. 1993) have been computed with ATLAS12 (Kurucz 1993; Pavlenko 2003). The calculated spectra of the brown dwarf 2MASSW 0140026+270150 were computed with the DUSTY 2000/4.5/0 model atmosphere (Allard et al. 2001). The atomic line data for species other than Mn i were taken from the Vienna Atomic Line Database (VALD) (Kupka et al. 1999). We used molecular line lists from different sources: TiO (Plez 1998), FeH (Dulick et al. 2003), CrH (Burrows et al. 2002) as well as H2O line list BT2 (Barber et al. 2006). For Arcturus we used the abundances of Peterson et al. (1993) and the spectra of the Sun and brown dwarfs were computed with the solar abundances reported by Anders & Grevesse (1989). The absorption lines are hyperfine split and each hyperfine component line (Table 2) has been fitted using the Voigt function H(a,v) and the formulae of Unsöld (1955) to calculate the damping constants. Theoretical spectra were computed with a wavelength step 0.01Å and convolved with Gaussians to match the instrumental broadening. For Arcturus, a rotational broadening was added corresponding to a projected equatorial radial velocity vsini = 7 km s-1, where v is the equatorial velocity and i is the inclination of the stellar rotation axis to the line of sight to the Earth, by following the Gray (1976) formulae.

We have fitted our synthetic spectra to the observed spectra of the Sun and Arcturus atlases by Kurucz (1991) and Hinkle et al. (1995) respectively. To obtain the best fit to the observed spectra we followed the minimisation procedure described by Pavlenko & Jones (2002) and Jones et al. (2002). In summary, we find the minima of the 3D function S(fs,fh,fg)=ν(FνFνx)2,\begin{equation} S\left(f_{\rm s}, f_{\rm h}, f_{\rm g}\right) = \sum_\nu \left(F_{\nu} - F_{\nu}^x \right)^2, \end{equation}(3)where Fν and Fνx\hbox{$F_{\nu}^x$} are the observed and computed spectra respectively, and fs, fh, fg are the wavelength shift, the normalisation factor, and the profile broadening parameter, respectively. To estimate the uncertainty of the best fit we use the parameter ΔS=SN(N1)\hbox{$\Delta S = \sqrt{S\over N(N-1)}$}, where N is the number of points in the observed spectrum. We provide an example of the profile fit for the Mn i at 12 899 Å line in the spectrum of the Sun in Fig. 2 and the fitted profiles of the other IR transitions (12 975, 13 281, 13 293, 13 318 and 13 415 Å) are shown in Fig. 3.

4.1. Mn i lines in spectra of the Sun

Fits of our synthetic spectra to the observed spectrum of the Sun are shown in Fig. 3. The manganese abundances obtained from the fits are given in Table 3. The log N(Mn) values in Table 3 are determined from a grid of manganese abundances, using the solar value log N(Mn) = −6.64 of Gurtovenko & Kostik (1989) and 0.05 and 0.1 dex as the abundance steps for the log N(Mn) measurements.

Table 3

The Mn i abundances determined from the best fits to the solar spectrum.

Table 4

The Mn i abundances determined from the best fits to the Arcturus spectrum.

thumbnail Fig. 4

The best fits to the observed Arcturus spectrum features found from the minima of Eq. (3) for Mn i lines at 12 899 Å, 12 975 Å, 13 281 Å (left column) and 13 293 Å, 13 318 Å, 13 415 Å (right column). The green line shows the section of the observed profile used to determine manganese abundance. The wavelength scale is taken from the observed spectrum of Arcturus.

4.2. Mn i lines in spectra of Arcturus

Manganese lines in the Arcturus spectrum are considerably more blended. Our fits to the observed profiles of Mn i lines are shown in Fig. 4, and the results of manganese abundance determination are shown in Table 4. The increase in the dispersion of the abundance results for Arcturus in comparison to the Sun is predominantly caused by the increase in blended features in the Arcturus spectrum. To minimise the effect of blending we exclude spectral regions with a relatively high number of unidentified blended features when fitting our calculated spectra to the observed spectrum. The corresponding wavelength range for each line is shown in Table 4 and the unblended spectral region used for our fit is marked by a green line in Fig. 4. In addition, in Fig. 5 we provide plots of the dependance of S with log N(Mn) for the Mn i lines (12 899 Å, 12 975 Å, 13 281 Å, 13 293 Å, 13 318 Å, 13 415 Å) observed in the spectrum of Arcturus.

thumbnail Fig. 5

Dependence of S on log N(Mn) for the Mn i lines at 12 899 Å, 12 975 Å, 13 281 Å (left column, top to bottom) and 13 293 Å, 13 318 Å, 13 415 Å (right column, top to bottom) observed in spectrum of Arcturus, see Fig. 4.

4.3. Mn i lines in spectra of ultracool dwarfs

The blending of Mn i lines in the IR spectrum of late-type objects increases significantly for effective temperatures lower than 3000 K. Numerous water lines form a pseudo background in the spectral region of many of the IR Mn i lines. However, for M9 dwarfs it is possible to fit the observed profiles of some Mn i lines. We provide an example of the profile fit for the Mn i line at 12 899 Å in the spectrum of 2MASSW 0140026+270150. The observed spectrum is described in detail in Lyubchik et al. (2007). The computation was performed with an initial assumption of solar like abundances for manganese, and other elements, and we have determined a value of log N(Mn) = −6.7 ± 0.2 in the atmosphere of 2MASSW 0140026+270150, which agrees to within the joint uncertainties with our derived solar abundance for manganese.

Both the line intensity and line profile can be fitted with the solar value of the manganese abundance. However, the spectra of cooler objects are dominated by the water bands in the near-IR, see the computed spectrum of LP944-20 (M 9.5, 2000/4.5 from Pavlenko et al. 2007). For these cooler objects only the Mn i line at 12 899 Å can be used for the analysis because the other lines are too blended with the H2O bands.

thumbnail Fig. 6

Top: the manganese line at 12 899 Å in the spectrum of M 9.5 dwarf 2MASSW 0140026+270150, where the red line is the observed spectrum and the green line is the model spectrum using our oscillator strengths and hyperfine component line positions (Teff/log (g)/ [Fe/H]  = 2500/5.0/0.0). Bottom: a comparison of the calculated spectrum of LP944-20 with only absorption features from Mn i in red and all other atomic and molecular species including the water vapour bands in blue.

5. Summary

Branching fractions for 20 Mn i transitions have been measured using high resolution Fourier transform spectroscopy and placed on an absolute scale using radiative lifetimes. Fifteen of these transitions have no previously published experimentally measured oscillator strengths. The remaining five transitions agree with previous published oscillator strengths to within the uncertainty of the measurements.

Using our new experimental log (gf) values we have determined the manganese abundance in the atmosphere of several late type stars. Our solar manganese abundance, log N(Mn) = −6.60 ± 0.05, agrees well with the manganese abundances of Anders & Grevesse (1989) log N(Mn) = −6.65, Gurtovenko & Kostik (1989) log N(Mn) = −6.64, and Biemont (1975) log N(Mn) = 6.67. Blending affects Mn i lines in the IR spectra of stars cooler than the Sun. As a result, our manganese abundances obtained from the fits to different lines in the spectrum of Arcturus are in the range −6.75 < log N(Mn) < −7.15, with a mean log N(Mn) = −6.95 ± 0.20, which agrees to within the uncertainty with the log N(Mn) = −6.97 determined by McWilliam et al. (2003) for Arcturus. Furthermore, our new laboratory measured log (gf)s for IR Mn i spectral lines can be used in the study of dust obscured objects and as a secondary abundance check to abundance studies using Mn i spectral lines in the visible.


1

Certain trade names and products are mentioned in the text in order to adequately identify the apparatus used to obtain the measurements. In no case does such identification imply recommendation or endorsement by NIST or any of the coauthor institutes.

Acknowledgments

We thank Drs. L. Prato and I. S. McLean for providing the 2MASSW 0140026+270150 spectrum. R.B.W. gratefully acknowledges the European Commission for a Marie Curie fellowship. J.C.P., Y.P. and H.R.A.J. gratefully acknowledge funding from the Leverhulme Trust and STFC, UK. The work of Y.P. and L.Y. was partially supported by the Microcosmophysics-2 program of the National Academy and Space Agency of Ukraine. H.N. acknowledges the support of the Linnaeus grant to the Lund Laser Centre from the Swedish Research Council.

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All Tables

Table 1

New and remeasured laboratory oscillator strengths for Mn i.

Table 2

Wavenumber and wavelength of the HFS component lines in the IR transitions 3d6(5D)4s a   6DJ−3d5(6S)4s4p z 6PJ\hbox{$z~^{6}{\rm P}^{\circ}_{J}$} and 3d6(5D)4s a   4DJ−3d5(6S)4s4p z 4PJ\hbox{$z~^{4}{\rm P}^{\circ}_{J}$}.

Table 3

The Mn i abundances determined from the best fits to the solar spectrum.

Table 4

The Mn i abundances determined from the best fits to the Arcturus spectrum.

All Figures

thumbnail Fig. 1

The upper plot shows the hyperfine splitting of the fine structure levels 3d6(5D)4s a 4D7/2 − 3d5(6S)4s4p(3P) z 4P5/2\hbox{$z^{~4}{\rm P}^{\circ}_{5/2}$} with allowed hyperfine transitions. The lower plot shows the hyperfine split profile of the transition observed in the uncalibrated laboratory spectrum at 12 975 Å, together with an indication of the positions and relative line strengths of the individual HFS transitions. A complete list of the Ritz wavelength and log (gf) for each HFS transition is available in Table 2.

In the text
thumbnail Fig. 2

The best fit to the observed solar spectrum feature found from the minima of Eq. (3) for the Mn i line at 12 899 Å.

In the text
thumbnail Fig. 3

The best fits to the observed solar spectrum features found from the minima of Eq. (3) for Mn i lines at 12 899 Å, 12 975 Å, 13 281 Å (left column) and 13 293 Å, 13 318 Å, 13 415 Å (right column). The wavelength scale is taken from the observed spectrum of the Sun.

In the text
thumbnail Fig. 4

The best fits to the observed Arcturus spectrum features found from the minima of Eq. (3) for Mn i lines at 12 899 Å, 12 975 Å, 13 281 Å (left column) and 13 293 Å, 13 318 Å, 13 415 Å (right column). The green line shows the section of the observed profile used to determine manganese abundance. The wavelength scale is taken from the observed spectrum of Arcturus.

In the text
thumbnail Fig. 5

Dependence of S on log N(Mn) for the Mn i lines at 12 899 Å, 12 975 Å, 13 281 Å (left column, top to bottom) and 13 293 Å, 13 318 Å, 13 415 Å (right column, top to bottom) observed in spectrum of Arcturus, see Fig. 4.

In the text
thumbnail Fig. 6

Top: the manganese line at 12 899 Å in the spectrum of M 9.5 dwarf 2MASSW 0140026+270150, where the red line is the observed spectrum and the green line is the model spectrum using our oscillator strengths and hyperfine component line positions (Teff/log (g)/ [Fe/H]  = 2500/5.0/0.0). Bottom: a comparison of the calculated spectrum of LP944-20 with only absorption features from Mn i in red and all other atomic and molecular species including the water vapour bands in blue.

In the text

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