2 Observations and data reduction

 

 

Table 1: The M-type Mira Sample

Star	P[days]	Sp.	$V_\ast$[km/s]
RR Sco	281	M6e-M9e	-37a
R Aql	284	M5e-M9e	29b
R Car	309	M4e-M8e	20c
R Leo	310	M6e-M9.5e	6b
S Scl	363	M3e-M9e(Tc)	(12)
R Hya	389	M6e-M9eS(Tc)	-11b

^a

$V_\ast$ from the midpoint of the CO J=3-2 transition (Young 1995).

^b

$V_\ast$ from the midpoint of the CO J=2-1 transition (Cernicharo et al. 1997).

^c

$V_\ast$ from the midpoint of the SiO J=3-2, v=0 transition (Groenewegen et al. 1999).

The observations presented here were taken from December 1998 to December 1999. The observed line spectra were obtained with the coudé echelle spectrograph and 81 cm camera of the Mount Stromlo Observatory 1.88 m telescope. A Site 2$\times$4k CCD was used as detector, allowing coverage of the spectral region 3600-5700Å in each exposure: the corresponding spectral range in each echelle order was 70-60Å. The echelle grating gave a dispersion of 0.0204 Å pixel^-1 at 4000Å, equivalent to 1.53 km s^-1 pixel^-1. Maximum resolution observations were done with a 300 $\mu$m (1.2 $\hbox{$^{\prime\prime}$ }$) slit, giving a 3-pixel resolution of 4.6 km s^-1. For spectophotometric observations, a 1500 $\mu$m (6 $\hbox{$^{\prime\prime}$ }$) slit was used, giving a resolution of 15.3 km s^-1. Because of the high slit losses experienced with the 300 $\mu$m slit, and the fact that the emission lines are mostly considerably broader than 15.3 km s^-1, a large fraction of all observations were in fact done using the wide slit.

The periods, spectral types and stellar center-of-mass velocities of the Miras in our sample are listed in Table 1. The periods and spectral types are taken from the index to variable stars of the Variable Star Network (VSNET) (Nogami et al. 1997). The heliocentric stellar center-of-mass velocity $v_\ast$ was determined from the midpoint of the circumstellar molecular emission of CO or SiO given in the papers cited below the table. This method of deriving stellar center-of-mass velocities has already been described by Reid & Dickinson (1976) and the velocities should be correct to 1-2 km s^-1. For SScl it was not possible to find a direct estimate of $v_\ast$ in the literature. We adopted the value given in parentheses: it was obtained by choosing $v_\ast$ to give the same velocity (relative to the stellar center-of-mass) for various emission lines as found for stars with known $v_\ast$. This estimate is probably correct to 4 km s^-1.

We observed the Balmer lines H$\gamma$ (4340.46Å), H$\delta$(4101.73Å), H$\zeta$ (3889.05Å) and H$\eta$ (3835.38Å) near maximum visible light when possible in order to get the shock velocity deep in the atmosphere. Between maximum and minimum light we searched for emission lines of the metals MgI, SiI, MnI, FeI and FeII. Table 2 lists the multiplet number M, wavelength $\lambda$, upper and lower energy levels of the transition $E_{\rm u}$ and $E_{\rm l}$, and the Einstein-coefficients $A_{\rm ul}$ for the metal emission lines observed. This data was obtained from the NIST database (Fuhr et al. 1988). Note that the detected lines include the forbidden lines of the multiplets [FeII]6F, [FeII]7F and [FeII]21F.

 

 

Table 2: Metal emission line properties

Ion	M	$\lambda$[Å]	$E_{\rm u}$[cm-1]	$E_{\rm l}$[cm-1]	$A_{\rm ul}$[s-1]
SiI	2	4102.95	39760	15394	9.6+4
MnI	2	4030.75	24802	0	1.7+7
MgI	1	4571.10	21870	0	2.2+2
	3	3829.32	47957	21850	8.9+7
	3	3832.35	47957	21870	6.7+7
	3	3838.29	47957	21911	4.5+6
FeI	2	4461.65	23111	704	2.9+4
	2	4375.93	22846	0	2.9+4
	3	4206.70$^\ast$	24181	416	7.2+3
	3	4216.18	23711	0	1.8+4
	3	4291.46$^\ast$	23711	416	4.1+3
	42	4307.90	35768	12560	3.4+7
	42	4202.03	35768	11976	8.2+6
	73	3852.57	32499	17550	2.9+6
	648	4374.49$^\ast$	49477	26624	4.9+5
	828	4427.30	22997	416	3.4+4
FeII	38	4583.84	44449	22637	3.8+5
	6F	4457.95	22810	385	2.9-1
	7F	4359.33	23318	385	1.1+0
	7F	4287.39	23318	0	1.5+0
	21F	4276.83	25805	2430	6.5-1
	21F	4243.97	25429	1872	9.0-1

$^\ast$

Searched for, but not observed in our spectra.

For data reduction, the Image Reduction and Analysis Facility (IRAF) of the National Optical Astronomy Observatories was used. The wavelength calibration of the spectra was obtained from thorium-argon arcs taken during each run. The standard stars HR 718, HR 3454 or HR 7596 were observed on photometric nights to allow calibration of the spectra to absolute fluxes. The calibrated fluxes from multiple observations of standard stars during each night were compared to estimate the errors in the fluxes tabulated in Tables 4-8: these comparisons indicate that the flux errors should be less than 5%. This error is much less than the typical variation of flux throughout the pulsation cycle (see Sect. 4).

In the case of non-photometic nights, we had to estimate the absolute fluxes from the spectra by assuming that the flux in the V band at 5100Å (a quasi-continuum point) varies in the same way as the light curves of The American Association Of Variable Star Observers (AAVSO) (Mattei et al. 1980). With this method we were able to estimate the continuum flux at 5100Å by scaling the non-photometric flux relative to the calibrated flux from a photometric night. We were then able to calibrate the flux at other wavelengths using the relative spectral response of the spectrograph. The line fluxes obtained in this way are marked with a colon in Tables 4-8.

In order to derive the emission line fluxes, it is necessary to have an estimate of the continuum level. This was estimated by eye: initially, the whole order containing the line was examined to get an idea of the overall continuum shape, then a wavelength interval on either side of the line approximately equal to the line width was used for the final estimate. In the case of P-Cygni line profiles, the associated absorption was ignored in the continuum estimate. For strong emission lines (ex. Balmer lines), the continuum level is not a significant contributor to line flux errors, but in the case of weak emission lines in these stars which exhibit many absorption lines (ex. FeI 4307.90Å at phase 0.63 in RLeo, Fig. 16), the continuum level can be a major contributor to the flux error. In the worst cases, we estimate that the flux could be wrong by a factor of 2.

The phases for each observation of each Mira variable were determined from the AAVSO maxima: phase zero corresponds to visible maximum. Table 3 lists the Julian dates and the corresponding phases for the variables observed.

   

Table 3: Julian dates and phases of observation

JDa	Phase
	RR Sco	R Aql	R Car	R Leo	S Scl	R Hya
1153	-	-	-0.16	-	-	-
1169	-	-	-	-	0.02	-
1206	-	-	-0.09	0.43	-	-
1268	-	0.01	0.11	0.63	-	-
1357	0.29	0.32	0.39	-	0.52	0.08
1375	-	0.39	0.45	-	0.58	0.12
1412	0.48	0.50	0.57	-	0.67	0.21
1440	0.58	-	0.66	-	0.75	-
1464	-	0.69	-	-	-	-
1542	-	-	-	1.49	-	0.53

^a Julian date - 2450000.