A survey of high-frequency SiO masers in supergiants and AGB stars

Transcription

1 Mon. Not. R. Astron. Soc. 304, 906±924 (1999) A survey of high-frequency SiO masers in supergiants and AGB stars M. D. Gray, 1 * E. M. L. Humphreys 1;2 and J. A. Yates 3 1 Department of Physics, University of Bristol, Tyndall Avenue, Bristol BS8 1TL 2 School of Chemistry, University of Bristol, Cantock's Close, Bristol BS8 1TS 3 Department of Physical Sciences, University of Hertfordshire, College Lane, Hat eld AL10 9AB Accepted 1998 December 9. Received 1998 September 10; in original form 1998 April 27 1 INTRODUCTION The present paper is based on observations and interpretation of high-frequency maser emission from SiO molecules in the circumstellar envelopes of supergiant and AGB stars. The AGB stars are subdivided, by variability type, into Mira/SRa, SRb, SRc and OH/ IR objects. The term `high-frequency' means, for the purposes of this work, greater than 200 GHz. Maser emission from circumstellar SiO is a widespread phenomenon: for details of the transitions, star types, isotopomers and possible pumping schemes involved, see for example Bujarrabal et al. (1986), Jewell et al. (1987), Cernicharo & Bujarrabal (1992), Cernicharo, Bujarrabal & SantareÂn (1993), Gray et al. (1995), GonzaÂlez-Alfonso & Cernicharo (1997), Humphreys et al. (1997) and Pardo et al. (1998). This work deals only with the major isotopomer, the vibrational states v ˆ 1 and v ˆ 2, and within these the rotational transitions from J ˆ 5 4 up to J ˆ 8 7, with the heaviest emphasis on J ˆ 7 6. Long-term monitoring, over more one period, of high-frequency SiO maser emission (J ˆ 7 6 in v ˆ 1 and v ˆ 2) has been performed for only one star, R Aqr (Gray et al. 1998). The highfrequency emission appears to be more vulnerable than the 43- and 86-GHz emission. The monitoring work is consistent with the idea that the high-frequency masers form in clumps, like those found in VLBI images for low-frequency SiO masers (Diamond et al. 1994; Boboltz, Diamond & Kemball 1997), but does not prove this. Clumps which support SiO maser emission may be produced by thermal instabilities (Cuntz & Muchmore 1994), by the Parker * Present address: Department of Physics and Astronomy, University of Wales, Cardiff, PO Box 913, Cardiff CF2 3YB. ABSTRACT We report the results of a survey of high-frequency SiO maser emission from a sample of 34 late-type supergiant and AGB stars. Four new sources were detected, including the rst detection of 301-GHz emission from an S-type star. Variability of the maser emission with stellar phase is discussed in detail. We nd that high-frequency SiO maser emission appears to be exceptionally weak or absent in Mira variables over an optical phase range of approximately 0.4 to 0.7. No signi cant correlation is found between maser photon luminosity and spectral type for non-supergiant types. Centroid velocities of spectra appear to be redder than the stellar velocity over a phase range that is wider than the band where emission is weak on average. However, the conclusion that only the blue features appearing in spectra after an optical phase of 0.7±0.8 correspond to new masing objects should be viewed with caution. Key words: masers ± surveys ± stars: AGB and post-agb ± circumstellar matter ± supergiants ± radio lines: stars. instability if the magnetic eld is suf ciently strong (Hartquist & Dyson 1997), or by a combination of these processes. The very limited amount of high-frequency data must be compared with trends which are derived from low-frequency surveys. Martinez et al. (1988) found that the peak maser emission in the SiO 43-GHz lines in v ˆ 1 and v ˆ 2 appears to lag behind the optical maximum by typically 0.1±0.2 of a stellar period, but lags of up to 0.3 are not exceptional. In the case of R Aqr, the phase lag in Martinez, Bujarrabal & Alcolea (1988) was 0.35, and it was 0.25 at 86 GHz (v ˆ 1; J ˆ 2 1) following the optical maximum of 1996 March 21 (Gray et al. 1998). Cho, Kaifu & Ukita (1996a) nd a good correlation between the peak ux in the v ˆ 1 and v ˆ 2 SiO lines at 43 GHz and the infrared continuum ux, in this case at 4 mm. This correlation echoes earlier surveys, which also found that SiO maser brightnesses at 43 GHz followed near-infrared continua (e.g. Hjalmarson & Olofsson 1979; Nyman & Olofsson 1986), which suggests that the maser pumping scheme has a radiative component. In the follow-up work (Cho, Kaifu & Ukita 1996b) several additional conclusions are drawn for low-frequency (43 GHz) masers. The likelihood of detecting maser emission appears to increase as the spectral type becomes later from M4 to M10, but ux ratios between v ˆ 1 and v ˆ 2, from the same rotational transition, appear to be independent of spectral type. The highest vibrationally excited line considered v ˆ 3; J ˆ 1 0 was detected only in stars which had an optical phase close to 0.2, and the corresponding v ˆ 1 and v ˆ 2 lines also peaked near this same phase. In addition to the already mentioned correlation between the uxes of SiO masers and near-infrared continua, sources which were detected in v ˆ 3 tended to have brighter v ˆ 2 masers than those which were not detected (for phases q 1999 RAS

2 High-frequency SiO masers in supergiants and AGB stars 907 Table 1. SiO rest frequencies. Transition Frequency Short /GHz v ˆ 1; J ˆ v ˆ 2; J ˆ v ˆ 1; J ˆ v ˆ 2; J ˆ v ˆ 1; J ˆ v ˆ 2; J ˆ v ˆ 1; J ˆ v ˆ 2; J ˆ close to 0.2), and the average ratio of the SiO maser to 4-mm ux is also higher for those sources detected in v ˆ 3 for both v ˆ 1 and v ˆ 2 masers. Another interesting relationship is the ratio of v ˆ 1 to v ˆ 2 maser ux. Cho et al. (1996b) show that this tends towards unity as the ratio of maser to infrared ux rises, and take this as an indication of the effectiveness of infrared pumping of the SiO masers. 2 OBSERVATIONS All observations used in this work were carried out on the 15-m James Clerk Maxwell Telescope (JCMT) in Hawaii with either the A2 receiver (210±280 GHz) or the B3i receiver (298±380 GHz). Eight SiO frequencies were used, and these are listed in Table 1, together with a three-digit short notation which will be used for identi cation in the rest of the text, and the corresponding rovibrational transition. The SiO frequencies were derived from the theory by Dunham (1932). Several observing epochs were used, but for most of these the A2 receiver was not available and the majority used only the B3i receiver. In detail, the observing epochs were 1994 December 08, 1995 January 14, 1995 March 03, 1995 April 18, 1995 June 05, 06, 07 and 08, 1995 December 28, 29 and 30, 1996 February 29/March 01, and 1996 July 27, 28 and 29. Receiver A2 was used only during the fourth epoch (1995 April) and the sixth epoch (1995 December). Although very rapid variability has been observed in some SiO masers (Balister et al. 1977; Pijpers, Pardo & Bujarrabal 1994), data taken for the same object and frequency have been co-added if the observations were only a day or two apart, since there is no conclusive evidence of signi cant short-term spectral change, at the 5j level or above, in the present data. Five objects were detected at epochs separated by 48 h or less, three of which were Miras and two of which were supergiants. In none of these cases did the residual spectrum deviate signi cantly from a at baseline. The J ˆ 8 7 and J ˆ 7 6 lines of SiO were observed with a system based on the liquid-helium-cooled single-channel SIS mixer receiver B3i with the Digital Autocorrelation Spectrometer (DAS) as the backend, providing a bandwidth of 250 MHz and a channel spacing of MHz. Observations of the SiO v ˆ 2; J ˆ 7 6 transition at 299 GHz were all tuned in the lower sideband: this was usual for the v ˆ 1; J ˆ 7 6 transition (301 GHz) also, but the upper sideband was used in 1995 April and June. The SiO J ˆ 8 7 lines were also usually tuned in the lower sideband with the exception of 1995 January and March when the upper sideband was used for v ˆ 2; J ˆ 8 7 at 342 GHz. The lower frequency SiO lines, corresponding to J ˆ 6 5 and J ˆ 5 4 transitions, were observed with a system based on the liquid-helium-cooled single-channel mixer receiver A2. The backend, bandwidth and channel spacing were as used with B3i. All lines observed with the A2 receiver at epoch 1995 March used the upper sideband, whilst the lower sideband was employed for the 1995 December observations. All observations were made in beam-switching mode, with the secondary mirror chopped by 2 arcmin at a frequency of 1 Hz. Onsource times, inclusive of overheads for sky subtraction, varied between 10 and 70 min for any observed object, at any one frequency, at any one epoch. Pointing calibrations were carried out as necessary, using a variety of convenient thermal sources and quasars from the JCMT pointing source catalogue. Focusing was carried out at the start of each observation, and again after sunset or sunrise if either of these events happened before the end of the observing run. The JCMT feedhorn recorded single polarization from a linear feed system. There is therefore a potential loss of ux, dependent on the fraction which is polarized and the plane angle of the polarized component relative to the feedhorn. Moreover, the plane angle varied during each observation due to the parallactic angle effect, since the JCMT is an altitude-azimuth instrument. However, we have no option here but to proceed assuming that any ux variations introduced by polarization are less important than those related to the overall ux of the source (Martinez et al. 1988). The parallactic angle effect, within any given observation, is certainly small, given their short duration. The possible consequences of the parallactic angle effect on the interpretation of observations at different epochs is discussed in detail in Section 5.1. The observations at the J ˆ 5 4 and J ˆ 6 5 frequencies of SiO taken in 1995 December have already been reported in Humphreys et al. (1997, tables 3a and b), and these observations only appear in the present work as part of the statistical study in Section 5. The only additional data at these frequencies are from the Miras R Leo and S CrB. R Leo was observed on 1995 April 18, and was not detected above a 3j noise level of 3.2 Jy at 258 GHz; it was also not detected at 215 GHz, where the 3j level was 4.1 Jy. S CrB was observed only at 258 GHz at this same epoch, and was not detected down to a 3j level of 3.2 Jy. The sources observed are listed in Table 2. Columns of Table 2, from left to right, are: (1) the name of the star, (2) the Hipparcos Catalogue number (HC) of the star, if applicable, (3) and (4) the right ascension and declination respectively, (5) the LSR velocity used in the observations, (6) the stellar velocity derived from infrared CO, or other molecular observations (Cho et al. 1996b), (7) the spectral type, (8) the stellar distance, (9) the stellar period, (10) the type designation (Mira, SRb, supergiant, etc.), and (11) the Julian day number of an epoch of maximum light, as close as possible to the observations, for those objects with a single wellde ned period. 3 DATA REDUCTION All data were reduced using SPECX V6.4 (Padman 1993). The basic operations involved merging two independent DAS quadrants and subtraction of a quadratic baseline, chosen from regions of the spectrum well separated from the LSR velocity. The same regions of the spectrum were used to estimate rms noise levels. Conversion factors used to convert the TA scale to Jy were as used in Humphreys et al. (1997) for all the SiO lines. These conversion factors were derived from the JCMT User Guide. To improve signal-to-noise ratios, the MHz spectral resolution was halved to MHz for strong detections, but for weaker detections, spectral resolution was degraded by factors of 3, 4 or 6. Limits

3 908 M. D. Gray, E. M. L. Humphreys and J. A. Yates Table 2. Source list. Source HC RA (1950) DEC (1950) v(lsr) v Spectral Type Distance Period Type JD max hh h mm m ss: s s 6dd8mm9ss0 /km s 1 /km s 1 /pc /d SV And h 01 m 46: s M5-M M o Ceti h 16 m 49: s M5-M M S Per h 19 m 15: s M3-M SRc R Cet h 23 m 29: s M4-M M U Cet h 31 m 20: s M2-M M OH h 29 m 23: s OH/IR IK Tau 03 h 50 m 43: s M6-M M S Tau 04 h 26 m 27: s M6.5-M M TX Cam 04 h 56 m 43: s M8-M M NV Aur 05 h 07 m 19: s M M S Ori h 26 m 32: s M6.5-M M U Ori h 52 m 50: s M6-M M VY CMa h 20 m 54: s M SG DU Pup 07 h 32 m 54: s M M OH h 39 m 59: s M6I OH/IR IW Hya 09 h 42 m 58: s M M R Leo h 44 m 52: s M6-M M R Crt h 58 m 06: s M SRb RT Vir h 00 m 05: s M SRb R Hya h 26 m 58: s M6-M M W Hya h 46 m 12: s M7.5-M SRa RX Boo h 21 m 57: s M6.5-M SRb S CrB h 19 m 21: s M6-M M H1-36 Ara 17 h 46 m 24: s M8 +w M VX Sgr h 05 m 03: s M4-M SRc OH h 34 m 52: s VM OH/IR R Aql h 03 m 57: s M5-M M x Cyg h 48 m 38: s S6-S M NML Cyg 1 20 h 44 m 33: s M4.5-M ? SG SV Peg h 03 m 31: s M SRb R Cas h 55 m 51: s M6-M M R Peg h 04 m 08: s M6-M M V Cas h 09 m 31: s M5-M M R Aqr h 41 m 14: s M5-M8.5 +w M Notes: Column 1: 1 =V1489 Cyg. Many of the stars in the table also have IRAS and/or IRC numbers which appear in the SIMBAD data base. Column 2: from the Hipparcos Catalogue (Perryman et al. 1997). Column 3, 4 and 5: 1 (Yates, Cohen & Hills 1995); 2 (Benson et al. 1990) with additional detail from the SIMBAD data base; 3 (Jewell et al. 1991); 4 (Ivison, Seaquist & Hall 1994). Column 6: non-blank stellar velocities (Cho, Kaifu & Ukita 1996b); others taken to be equal to v(lsr) (column 5). Column 7: unmarked spectral types from the GCVS (Kholopov et al. 1985); 1 (Benson et al. 1990). The addition of +w implies that the star has a white dwarf companion. Column 8: unmarked distances with errors from trigonometrical parallax (Perryman et al. 1997); 1 SiO maser moving cluster parallax (Marvel 1997); others: 2 (Yates, Cohen & Hills 1995); 3 (Cho, Kaifu & Ukita 1996b); 4 (Engels et al. 1983); 5 (Ivison, Seaquist & Hall 1994); 6 (Svignanam et al. 1989); 7 (Svignanam, Le Squeren & Foy 1988); 8 (Heske et al. 1990); 9 estimated from the Mira period±luminosity relation. Column 9: unmarked periods are from Kholopov et al. (1985), others from 1 Benson et al. (1990); 2 Allen et al. (1989). Column 10: most variability classes are from Kholopov et al. (1985); see also Section 3. Column 11: unmarked epochs of maximum light from AAVSO and VSNet; 1 from Kholopov et al. (1985); 2 from Hall et al. (1990); DU Pup at phase 0.4 on 1988 August 27; 3 from Allen et al. (1989), Feast et al. (1983). For full Julian day add to gure shown. on non-detection are based on the noise level in a bandwidth of MHz, corresponding to one-eighth of the original spectral resolution. Spectral resolutions used for individual spectra presented in this work are given in the gure captions in Section 4. Where possible, objects were assigned a stellar velocity, spectral type, distance, period and variability type; see columns (6) to (10) in Table 2. Stellar velocities were assigned only where there was reason to believe that they represented a more accurate value for the true stellar velocity than the LSR velocity used in the observation. Spectral types were assigned, keeping only the colour part of the code: luminosity class information has been discarded. Distances were obtained for all the objects, but the reliability varies considerably. Where possible, distances have been obtained from parallax measurements, either trigonometric (Perryman et al. 1997) or moving cluster (Marvel 1997). For objects in the Hipparcos Catalogue, the parallaxes have been used if they are at least twice the stated error; the parallax of R Aqr has also been used, as it is quite large (5.1 mas) in spite of the high uncertainty. For other

4 High-frequency SiO masers in supergiants and AGB stars 909 sources, where errors are not quoted in column (8) of Table 2, an error of 30 per cent in the distance will be assumed. Stellar periods were used, in conjunction with the epochs of maximum light in column (11) of Table 2, to obtain the optical phase of the object at each epoch of observation. In general, this operation was possible only for the Mira and SRa variability types: most of the SRb and SRc types were too irregular to obtain reliable phases, and it was not possible to nd a period and/or an epoch of maximum light for some of the OH/IR objects. Classi cation of variability closely follows Kholopov et al. (1985) and Benson et al. (1990), but any extremely red object which has no spectral type, no period, or a period >1000 d is classi ed here as an OH/IR star. The designation `SG', supergiant, is reserved for those bright optical stars which cannot be reliably assigned a period and tted into the SRc type. In order to study phase-dependent effects, epochs for maximum light were obtained as close as possible to the range of observing dates (see Section 4). Julian dates were converted to calendar dates where necessary, using JD as the equivalent of 2000 Jan 01. In most cases the epoch of maximum light was extracted from AAVSO data, the archives of the Variable Star Network (VSNet) in Japan, or additional AAVSO data in the Hipparcos Catalogue (Perryman et al. 1997). For objects where AAVSO data were not available, the source of the information is cited in the notes to Table 2. 4 RESULTS In this section we present spectra of lines which were previously undetected in the objects concerned, some comparisons demonstrating the variability of spectra when compared with Humphreys et al. (1997), and tables which list the full set of observations, including the more marginal detections and upper limits on sources which were undetected. Variability is quanti ed by the introduction of a variability index, v, equal to js 2 S 1 j= S 2 S 1, where S 2 and S 1 are the integrated uxes recorded at times 2 and 1 respectively (Yates & Cohen 1996). The possible consequences of maser polarization, and the effect of variation in the parallactic angle between observations at different epochs, are discussed in Section 5.1. Table 3. Observations at 299 GHz. Source JD Phase S n R Sn dn /Jy /Jy km s 1 o Ceti < 4.1 o Ceti < 2.1 S Per < 4.0 IK Tau < 4.2 IK Tau < 5.3 TX Cam TX Cam < 3.8 TX Cam < 3.4 VY CMa R Leo R Leo W Hya H1-36 Ara < 10.4 VX Sgr < 6.0 x Cyg NML Cyg < 3.2 Note: If two very close epochs have been co-added, column (2) relates to the rst day. Table 4. Observations at 301 GHz. Source JD Phase S n R Sn dn /Jy /Jy km s 1 SV And <2.8 o Cet o Cet S Per S Per <2.9 R Cet <4.8 U Cet <4.0 OH <3.8 IK Tau <2.1 IK Tau <3.4 S Tau <4.0 TX Cam <3.3 TX Cam <9.8 NV Aur NV Aur <6.6 S Ori <3.9 S Ori <2.4 U Ori VY CMa DU Pup <3.1 OH <3.6 IW Hya <3.9 R Leo <1.8 R Leo R Crt RT Vir R Hya W Hya H1-36 Ara <3.4 H1-36 Ara <5.0 VX Sgr <5.9 OH <10.7 R Aql <5.2 x Cyg NML Cyg SV Peg R Cas <2.6 R Peg <2.9 V Cas <3.6 Note: If two very close epochs have been co-added, column (2) relates to the rst day. Previously unreported observations for SiO at frequencies of 299 GHz and above appear in Tables 3 to 6. Data taken during 1995 December appear in table 3(c) of Humphreys et al. (1997), and observations of R Aqr in Gray et al. (1998). Table 3 contains the data for 299 GHz, Table 4 the data for 301 GHz, whilst the J ˆ 8 7 data are divided between Table 5 (342 GHz) and Table 6 (345 GHz). In each table the rst column is the source name, the second column is the observing epoch as a Julian day, the third column is the phase of the star, calculated with respect to the maximum in Table 2, the fourth column is the peak ux observed, or the 3j upper limit and the fth column is the integrated line ux. In all the Tables 3 to 6, the observations are listed by object, in the same order as in Table 2, and where there are multiple observations of an object, these are presented chronologically. Errors on the phase are taken to be 2 d, for the estimation of the reference maximum, plus one unit in the last digit quoted for the period per whole period elapsed between the reference and the epoch of observation.

5 910 M. D. Gray, E. M. L. Humphreys and J. A. Yates Table 5. Observations at 342 GHz. Source JD Phase S n R Sn dn /Jy /Jy km s 1 TX Cam < 4.7 VY CMa < 5.0 VY CMa R Leo < 5.4 Note: If two very close epochs have been co-added, column (2) relates to the rst day. Table 6. Observations at 345 GHz. Source JD Phase S n R Sn dn /Jy /Jy km s 1 o Cet o Cet VY CMa Note: If two very close epochs have been co-added, column (2) relates to the rst day. The detection statistics for each line are six detections in 16 observations ( ve sources out of 11) at 299 GHz, and 14 detections from 39 observations (13 sources out of 31) at 301 GHz. Too few observations were carried out at the other frequencies to be statistically signi cant, but for the record there was one detection from four observations at 342 GHz, and three detections from three observations at 345 GHz. No objects were detected in the three observations at the lower frequencies (see Section 2). Apart from the marginal detections, individual objects are discussed below. 4.1 o Ceti (Mira) The archetypal Mira variable was observed at two phases on either side of a cycle boundary in The two 301-GHz spectra are shown in Fig. 1, with dates of observation. At the time of the rst observation, the bright maser was detected at an optical phase of This had decayed to a smaller emission peak by an optical phase of 0.37, with a signi cantly redder centroid of 46.7 km s 1, compared with 45.1 km s 1 at the earlier epoch. The variability index was 0.05 over 149 d, which is not signi cant, the broadening of the spectrum making up for the lowering of the peak. Unfortunately, no observations were made during the passage through optical maximum, so the phase where the maser emission peaked is unknown. There is no signi cant shift in the velocity of the peak, 45.1 km s 1, between the observations. Although the 299-GHz line was sought at both epochs, it was not detected. It is very important that o Ceti was also detected in the 345-GHz (J ˆ 8 7) line on two occasions, as it is the only Mira variable which has so far been observed to emit in either of these lines. Both detections, however, were only at the 3j level; see Table R Leo This nearby Mira variable was detected as a strong 301-GHz maser source at the beginning of 1996 March. The most interesting feature of the R Leo observations in this work is that it was observed at almost exactly the same phase, one period apart. Although the rst Figure GHz spectra of o Cet (Mira). In the February 29 spectrum, the spectral resolution is MHz (0.31 km s 1 ), degraded by a factor of 2 from the observations. In the July 27 spectrum, degradation by a factor of 3 gives the spectral resolution shown, equal to MHz (0.47 km s 1 ).

6 High-frequency SiO masers in supergiants and AGB stars 911 Figure and 301-GHz spectra of R Leo. The uppermost gure has spectral resolution of MHz (0.31 km s 1 ), degraded by a factor of 2 from the observations. The middle and bottom gures show the growth of 299-GHz ux density. The 1994 December observation, a 3j detection, used a degradation factor of 6 to give a spectral resolution of MHz (0.94 km s 1 ), whilst a spectral resolution of MHz (0.47 km s 1 ), degraded by a factor of 3 from the observations, was used for the 1995 January gure. observation resulted in a 3j upper limit of less than 1.8 Jy (at phase 0.05), the later observation, at phase 0.07, was a strong detection with a peak of 22.7 Jy. The lower limit on the variability index at 301 GHz is 0.97 over 318 d. Unfortunately, it is not possible to conclude that the 301-GHz masers switched on in R Leo between phases 0.05 and 0.07, because the two observations were not in the same cycle and the `switch-on' could easily be accounted for by variation between cycles. Previous observations failed to detect R Leo at 301 GHz, when the star was at phase 0.87 (Humphreys et al. 1997). Evidence for a signi cant increase in ux density, within a single stellar cycle and before optical maximum, comes from observations of R Leo in 1994 December and 1995 January at 299 GHz. At the earlier date, R Leo was marginally detected at a phase of 0.63, but by phase 0.75 the peak had reached 11.1 Jy. The bulk of the growing peak is considerably redshifted with respect to the marginal detection, with a centroid at 5.0 km s 1, compared to a stellar velocity of 0.5 km s 1. All three detections of R Leo are shown in Fig. 2. The variability index at 299 GHz is 0.87 over 37 d. This is the shortest time-scale over which de nite variability was observed in this work. 4.3 R Hya This Mira variable was detected by Humphreys et al. (1997) at an optical phase of 0.75 at several SiO frequencies, and is therefore useful for observing changes in the spectrum with phase. The new observation, at 301 GHz only, was in the same cycle, with the phase having advanced to As for the 299-GHz spectrum in R Leo, there is evidence of a general growth in ux density, before the optical maxmium, with the appearance of new spectral features. The variability index is 0.36 over 62 d. In this case, the spectrum has become redder with phase, the centroid moving from ±9.1 to ±8.0 km s 1 and the peak from ±8.9 to ±8.6 km s 1. Fig. 3 shows the R Hya spectrum from Humphreys et al. (1997)

7 912 M. D. Gray, E. M. L. Humphreys and J. A. Yates Figure GHz spectra of R Hya. The spectral resolution in both cases is MHz (0.31 km s 1 ), corresponding to a degradation by a factor of 2 from the observations. (upper spectrum) and the new detection from observations two months later on the same scale (lower spectrum). In spite of the small overall reddening, it is obvious from Fig. 3 that major new features have appeared both to the blue and to the red of the 1995 December peak. 4.4 W Hya This star, classed as a Mira-like SRa, was detected by Humphreys et al. (1997) at an optical phase of 0.57 in 1995 December at several SiO frequencies, including both 299 and 301 GHz. W Hya was observed again in 1996 July, when the phase had shifted across the next cycle boundary and reached Increases in peak ux density were recorded in both 299- and 301-GHz lines (see Fig. 4). In Fig. 4, lines of the same frequency have been plotted on the same vertical scale. At 301 GHz the increase was very dramatic, with the peak ux density rising from 12.8 to Jy, yielding a variability index of 0.89 and forming the brightest 301- GHz maser observed to date. A signi cant shift to the red in the 301- GHz spectrum was observed, with the peak moving from 39.4 to 45.2 km s 1 and the centroid from 41.1 to 43.6 km s 1. By contrast, the 299-GHz spectrum moved to the blue between the two observations, the peak moving from 47 to 42 km s 1 and the centroid moving through a very similar range. The variability index for 299 GHz was The earlier observation shows a detection at a phase which, for low-frequency SiO masers, is associated with low maser output. It is possible that the maser emission is more robust to phase effects in an SRa variable than in the Miras, where the amplitude of oscillation is greater. 4.5 TX Cam The Mira variable TX Cam was observed at three well-separated epochs in the 299-GHz line, and a strong detection was obtained only on the rst occasion, when the optical phase was 0.23 in 1995 March. Subsequent observations at phases of 0.87 (1996 February) and 0.14 in the next cycle (1996 July) failed to detect any maser emission (see Fig. 5). The variability index for the rst two epochs was 0.91 over 354 d. Attempts to detect 301-GHz emission at the latter two epochs were also unsuccessful. The fact that a strong maser of peak 47 Jy was replaced, in the next cycle, by a nondetection at a very similar phase is additional evidence for poor repeatability of high-frequency maser structure between adjacent pulsational cycles. The peak of the 1995 March detection was at 10.6 km s 1, bluer than the estimate of the stellar velocity given in Table 2 by 2.7 km s VY CMa The supergiant VY CMa is aperiodic, but over the time-range of these observations it appears to have been generally increasing in optical output (AAVSO archives). For each of the four lines observed in this work, Figs 6 and 7 show the previously unpublished spectrum and the residual, formed by subtracting the spectrum of the same frequency from Humphreys et al. (1997) to show changes over time. Fig. 6 shows the results for the J ˆ 8 7 lines, and Fig. 7 the results for J ˆ 7 6. In the case of 299 GHz, the order of subtraction is reversed, because the new spectrum in fact dates from an earlier epoch than the data in Humphreys et al. (1997). In all cases, the residual spectra have been tted to a new baseline, based

8 High-frequency SiO masers in supergiants and AGB stars 913 Figure and 299-GHz spectra of W Hya. The spectral resolution in the 301-GHz spectra is MHz (0.31 km s 1 ), degraded by a factor of 2 from the observations, whilst in the 299-GHz spectra, degradation by a factor of 3 has been applied to produce a spectral resolution of MHz (0.47 km s 1 ). Spectra at the same frequency are shown to the same scale. on the 30 km s 1 ranges at the extreme ends of the spectra shown. In the case of 345 GHz, there is evidence, at the 3j level, that the ux density increased over the two months between observations, with a variability index of The increase is not uniform over the whole spectrum, however, and the spectrum appears to have shifted slightly to the red. There is no statistically signi cant change in line strength at 342 GHz. At 301 GHz the shape and strength of the spectrum changed little between the end of 1995 December and the end of 1996 February, as can be seen from the 301-GHz residual in Fig. 7, and a variability index of only 0.02 was recorded. The double peaks which form the maximum of the spectrum appear to have grown very slightly and become better de ned, but any change is small. Over a considerably longer interval the 299-GHz spectrum reddened by a small but

9 914 M. D. Gray, E. M. L. Humphreys and J. A. Yates Figure GHz spectra of TX Cam. A spectral resolution of MHz (0.312 km s 1 ) has been used for the 1995 March (top) spectrum, whilst the nondetections are based on a spectral resolution of 1.25 MHz (1.25 km s 1 ), degraded by a factor of 8 from the observations. signi cant amount between 1995 March 12 and the end of December in the same year. In particular, the small blue wing, peaking near 7kms 1 in the 299-GHz spectrum in Fig. 7 had entirely disappeared (see g. 1d of Humphreys et al. 1997). The variability index for the 299-GHz line was 0.02, the same as at 301 GHz, but over a different, and longer, period. 4.7 R Crt The SRb variable R Crt was observed only at 301 GHz in 1996 February, and the spectrum is shown in Fig. 8, together with the residual formed by subtracting off the analogous spectrum from Humphreys et al. (1997). There is perhaps a suggestion of a slight reddening of the spectrum, but this is of low signi cance. As with the other detected SRb variables and supergiants, the peak of the 301-GHz spectrum lies signi cantly to the red of the stellar velocity, in this case by 4.7 km s 1. The variability index over 63 d was RT Vir The SRb variable RT Vir was observed at 301 GHz only on 1996 February 29, two months after the observation reported in Humphreys et al. (1997). The residual (Fig. 9) shows that there was no signi cant change in the spectrum over the interval of 62 d separating these epochs, supported by a low variability index of 0.21 with an uncertainty of The group comprising the SRb variables and VY CMa appear to have 301-GHz spectra that are remarkably constant, over time-scales of months, in both shape and total ux. This peak of this spectrum lies to the red of the stellar velocity by 4.9 km s Other detected stars In this section we show the spectra of all the other objects which were new strong detections in this work, but where only a single observing epoch was available. The new objects (see Fig. 10) comprise the two Mira variables NV Aur and x Cyg, together

10 High-frequency SiO masers in supergiants and AGB stars 915 Figure and 342-GHz spectra of VY CMa. The spectral resolution at 345 GHz is MHz (0.814 km s 1 ), degraded from the observations by a factor of 6. At 342-GHz degradation from the observations by a factor of 8 has been applied to yield a spectral resolution of MHz (1.09 km s 1 ). Both residuals have been formed by subtracting the data from Humphreys et al. (1997) from the 1996 March spectra shown here. with the SRb SV Peg and the supergiant NML Cyg. Additional objects, detected at the 3j level at 301 GHz and not shown in Fig. 10, are S Per and U Ori. Note that NV Aur was observed at a second epoch in 1996 July, but the result is not shown here, as the system temperature was very high on this latter occasion and the spectrum could have been swamped by noise with only a small decrease in peak ux density. The detection of x Cyg is signi cant because it is an S-type star, the rst such object to be detected as an SiO maser source in the J ˆ 7 6 lines. Both the supergiant NML Cyg and the SRb variable SV Peg follow the pattern of having a 301-GHz

11 916 M. D. Gray, E. M. L. Humphreys and J. A. Yates Figure and 301-GHz spectra of VY CMa. The spectral resolution in all cases is MHz (0.312 km s 1 at 299 GHz and km s 1 at 301 GHz), resulting from degradation of the observational resolution by a factor of 2. At 301 GHz the residual is the result of subtracting the spectrum in Humphreys et al. (1997) from the 1996 February spectrum shown here. At 299 GHz the residual results from subtracting the spectrum shown here from the spectrum shown in g. 1(d) of Humphreys et al. (1997). spectrum that peaks well to the red of the stellar velocity. Separations are 13.5 km s 1 for NML Cyg and 8.4 km s 1 for SV Peg. 5 ANALYSIS OF VARIABILITY AND PHASE The rst test to be carried out in this section is to plot the calculated photon luminosities, or upper limits, for all the available observations as a function of stellar phase. Calculation of photon luminosities assumes isotropic beaming in all cases. In order to estimate upper limits for the non-detections, the maximum integrated ux is assumed to be the quoted 3j limit over the spectral resolution of MHz. The observations discussed in Section 4 are augmented

12 High-frequency SiO masers in supergiants and AGB stars 917 Figure GHz spectrum of R Crt with the residual formed by subtracting the analogous spectrum from Humphreys et al. (1997) from the spectrum shown here. The spectral resolution, obtained by degrading the observational value by a factor of 3, is MHz (0.465 km s 1 ). Figure GHz spectrum of RT Vir, together with the residual resulting from subtraction of the analogous spectrum in Humphreys et al. (1997) from the spectrum shown here. The spectral resolution, degraded from the observational value by a factor of 3, is MHz (0.465 km s 1 ).

13 918 M. D. Gray, E. M. L. Humphreys and J. A. Yates Figure 10. New detections of 301- and 299-GHz emission. For the Mira NV Aur and the 301-GHz spectrum of x Cyg, the spectral resolution is MHz (0.465 km s 1 ). The 299-GHz spectrum of x Cyg has a spectral resolution of MHz (0.624 km s 1 ). For NML Cyg the spectral resolution is MHz (0.620 km s 1 ) and for SV Peg MHz (0.930 km s 1 ).

14 High-frequency SiO masers in supergiants and AGB stars 919 Figure 11. Mira variables: photon luminosities are shown as a function of phase for the Miras and the SRa variable W Hya (marked). Star symbols denote detections at 345 GHz, open circles are detections at 301 GHz, lled circles are those at 299 GHz, and open squares are those at 258 GHz. All non-detections shown are marked as down-arrows representing the upper limit of the photon luminosity. here with additional data from Humphreys et al. (1997) and from the study of R Aqr (Gray et al. 1998). Objects such as supergiants, SRb variables and OH/IR objects have been excluded, since a phase cannot be assigned to them. Variables of type SRc have also been excluded, even if a phase can be assigned, since these objects (S Per and VX Sgr) are very distant compared to most of the Miras in the sample and their derived photon luminosities, or upper limits, are too high to make a meaningful comparison with the Mira variables. Because of the strong dependence of the photon luminosity on the distance of the object, it was also found necessary to exclude nondetections of all Miras over 500 pc, since the upper limit on nondetection of H1-36 Ara, say, greatly exceeds the photon luminosity of the brightest detected object. The sample for the photon luminosity versus phase plot therefore consists of the detected Miras and the SRa, W Hya, together with non-detections towards the closer Miras. The largest upper limits shown under the stated restrictions correspond to photon luminosities which are all well below s 1. A plot of the photon luminosities versus phase appears in Fig. 11. The intrinsically brightest masers belong to the Mira variables TX Cam at 299 GHz and NVAur at 301 GHz, followed by the SRa variable W Hya. Fig. 11 appears to show that there is a band of phase, between approximately 0.4 and 0.7 in which high-frequency maser emission is very weak or absent. If the null hypothesis is assumed, namely that the graph represents random scatter about a horizontal line at the mean value of 41.2, then the hypothesis should be rejected at the 95 per cent con dence level on the basis of a chisquared test. This conclusion is reinforced if the only masers above s 1 in this phase range are excluded on the grounds that W Hya, as an SRa rather than a Mira, might have signi cantly different properties. Even in the case of W Hya, however, the 301-GHz maser at phase 0.16 is much brighter than the detection in the `forbidden' range at 0.57 (see also Fig. 4). The re-appearance of high-frequency masers near an optical phase of 0.7 is consistent with the idea that these masers result from a new outward-moving shockwave driving into the photosphere following this phase, the equivalent of model phase 0.0 in Humphreys et al. (1999). The second main test is to plot the stellar photon luminosities as a function of spectral type. This test is searching for a correlation between maser output and the stellar temperature. In the case of those variables which change spectral type signi cantly during a cycle, the median spectral type has been used. All objects that appear in Table 2, whether periodic or not, have been plotted in Fig. 12 unless they have a very high upper limit on non-detection due to a large stellar distance, or unless they are OH/IR objects for which a spectral type cannot be assigned. In Fig. 12 the photon luminosity of objects has been plotted as a function of median M-spectral type. All classes of object are included. A restriction has, however, been placed on non-detections: these are not shown if the upper limit is greater than s 1. In general, the supergiant/src objects appear to occupy a region in the upper left part of the diagram, with photon luminosities typically greater than s 1. The only object in this grouping not detected was VX Sgr. There seems to be little difference between the output of SRb stars and Miras in the case of detected objects. The main difference between these two classes is that the SRb objects were detected without exception, whilst the Miras produce maser emission which is highly variable over their cycles. Only R Crt exceeds s 1 in both observations at 301 GHz, which is weaker than the highest measured outputs from Miras (though at different spectral type). Cho et al. (1996b) suggested that the maser output increased as spectra became later for low-frequency SiO masers, although the correlation was not good. In this work the brightest Mira detections do come from the stars with the latest spectral types (TX Cam and NV Aur) but, as in Cho et al. (1996b), the correlation is not signi cant: in particular, more data are required for spectral types later than M8 and earlier than M6. It is also not possible to say whether, on average, the 301-GHz emission is signi cantly stronger or weaker than that at 299 GHz. To investigate the relationship between high-frequency SiO

15 920 M. D. Gray, E. M. L. Humphreys and J. A. Yates Figure 12. Photon luminosities are plotted against M spectral type for objects of all classes. Where non-mira objects exceed s 1, the object is specially marked. For VYCMa, the 301-GHz observations only are marked: all detections at type 5 in other lines are also VY CMa. Symbols identifying lines are as in Fig. 11. masers and infrared emission from the host stars, we plot the maser photon luminosities against an IRAS colour, de ned by C ˆ F 12 F 25 = F 12 F 25, where F 12 and F 25 are the stellar uxes at 12 and 25 mm respectively. Points for all objects in all lines and all observing epochs are plotted in Fig. 13, including in this case the higher limits on non-detection. As for the plot of photon luminosity against spectral type (Fig. 12 above), there are regions of the graph that correspond roughly to the stars of different classes observed. Detected supergiants and SRb variables tend to have higher photon luminosities and lower values of C (,0.2±0.35) than the Miras which have typical values of C from,0.30 to 0.6. There are several possible trends, but more data are required for the redder objects, particularly those with negative colours, before anything signi cant can be deduced. First, there are no detections at 299 GHz in any object redder than 0.44 except VY CMa, suggesting that this line requires a stronger near-infrared continuum than at 301 GHz. This would support a stronger radiative component in the pumping mechanism of the former line, as would be expected from its higher Figure 13. All objects: photon luminosities are shown as a function of the IRAS colour, C (see text for de nition). Symbols have the same meaning as in Fig. 11.

16 High-frequency SiO masers in supergiants and AGB stars 921 Figure 14. For all detections of strongly periodic stars, the velocity shifts of the spectral peak (a) and the velocity centroid (b) are shown with respect to the stellar velocity. Non-Mira variables (S Per and W Hya) are specially marked. Symbols are as de ned in Fig. 11. vibrational excitation. The second possible trend is a general fading of low-luminosity detections at 301 GHz with decreasing (reddening) colour. If true, this would suggest that only stars with very large envelopes could continue to support maser emission as the envelopes become redder. However, before any strong conclusion can be drawn, more data are needed with much stronger upper limits for OH/IR objects and very late-type Miras. Finally, we look at the shift of the spectral peaks and centroids relative to the nominal stellar velocities from Table 2. We note that for all objects detected, except R Aqr, the more accurate stellar velocities from column (6) of Table 2 are available. Vertical error bars in Fig. 14 assume these stellar velocities are accurate to 0.5 km s 1, whilst the stellar velocity of R Aqr has been assumed to be accurate to only 1.0 km s 1. The overall pattern of the plot is similar in style to gs 5(a) and (b) of Cho et al. (1996b), so that a comparison may be made between the behaviour of low- and highfrequency emission. Spectral-shifting shows that the maser emission is strongly coupled to the motions within the circumstellar envelope of the host star. If the intensity of the maser spectra are coupled to a particular shift pattern in the envelope, and this pattern is out of phase with the infrared light curve, this is evidence that envelope dynamics are important in forming and pumping the masers. Tracing the envelope pulsation therefore yields useful information about the effectiveness of both the collisional and radiative elements of the pumping scheme. In Fig. 14(a) we plot the difference between the velocity corresponding to the spectral peak and the stellar velocity against phase for all detections where a phase can be reliably assigned. Fig. 14(b) is similar, but the centroid (mean) velocities of the emission replace those of the peak. As for the previous gures in this section, the data in Tables 3 to 6 have been augmented with data from Humphreys et al. (1997) and Gray et al. (1998). On average, the spectra lie to the red of the stellar velocity, which follows the

17 922 M. D. Gray, E. M. L. Humphreys and J. A. Yates pattern of the aperiodic variables. However, if the detections of the SRa variable W Hya and the SRc variable S Per, specially marked in Figs 14(a) and (b), are removed, it can be argued that there is a phase regime where there is no maser emission from objects with a blueshift relative to the stellar velocity. This conclusion is stronger for the velocity centroid (Fig. 14b) than for the spectral peak (Fig. 14a). In this respect, it can be argued that the high-frequency data shown here follow the behaviour of the low-frequency spectra in Cho et al. (1996b). The fact that there does appear to be a correlation of spectral shift with phase, and that the phase where blueshifted emission reappears coincides with the phase where bright emission returns (,0.7), indicates that the high-frequency masers appear to be correlated with the dynamics of the circumstellar shell. Fig. 11 suggests that most masing objects fade away by an optical phase of,0.4, so when the bright emission returns, it is likely that all of it, both blueshifted and redshifted, is new. This is far less clear in low-frequency work, where it is rare for the maser spectra to be completely depleted in the 0.4±0.7 phase band. The new masing objects are probably driven by shock waves with a strong collisional element to their pumping scheme, since the behaviour of these objects seems to be more in phase with the envelope pulsation than the infrared light curve. Pulsation may also have important effects on the radiative pumping scheme via velocity-shift-induced line overlap in the infrared pumping lines. This scheme operates well for the rare isotopomers of SiO and for vibrational states v $ 3 (GonzaÂlez-Alfonso & Cernicharo 1997), and may in uence the transitions studied in this work. The statistics on the data are not good enough to say clearly whether, at J ˆ 7 6, there is any difference between the velocity-shift behaviour of 299 and 301 GHz. The phase variation detected in this work is in good agreement with the low-frequency behaviour discussed in Cho et al. (1996b), where the velocity pattern of the masers also appears to be related to a dynamical model of the circumstellar envelope which includes pulsation: there appears to be a band of optical phase, from approximately 0.3 to 0.8 where no stars have a mean blueshift in their spectra. This in turn agrees well with work by Hinkle (1978), Hinkle & Barnes (1979), Hinkle, Hall & Ridgeway (1982), Hinkle, Scharlach & Hall (1984) and Hinkle, Lebzelter & Scharlach (1997), which uses 1.6-mm CO emission to derive the velocity of a layer very close to the stellar photosphere as a function of phase. The reappearance of stars with blueshifted maser spectra following an optical phase of,0.8 ts well with the start of a new blueshifted portion of the radial velocity curve, at approximately the same phase, as derived from the CO observations: however, the conclusions drawn by Cho et al. (1996b) regarding the interpretation of the blueshifted maser emission as out ow and the redshifted as in ow are probably too simplistic, since from VLBA observations (Boboltz et al. 1997) we know that the 43-GHz maser zone in R Aqr was contracting, on average, until an optical phase of 0.05, well beyond the gure of 0.8 suggested by Cho et al. The interpretation of the dynamics from the CO observations agrees well with modelling of pulsating circumstellar envelopes by Bowen (1988), with SiO maser information added in Humphreys et al. (1996, 1999) in which the model phase 0.0, corresponding to the maximum outward velocity of the subphotospheric `piston' is set at an optical phase of approximately 0.7, to correspond roughly with the line-doubling phase (Hinkle et al. 1997) in the CO data. These models site the SiO masers in the pulsating inner regions of the circumstellar envelope and not in the steady wind zone. Further indirect evidence for pulsation, as a vital ingredient for the formation of SiO maser features, is that more evolved objects which appear to have lost their pulsational structure, protoplanetary nebulae (PPN) for example, appear also to have lost the ability to generate SiO masers (Nyman, Hall & Olofsson 1998). Finally, Cho et al. (1996b) suggest that the bright peak in the maser output near optical phase 0.2 could be due to a shock wave impacting on the maser zone, evidence for a strong collisional element in the maser pumping scheme at that phase. 5.1 Polarization effects The discussion in Sections 4 and 5 ignores various effects which may be important if the high-frequency masers have a high degree of linear polarization. The rst effect is that the single polarization feed and detector system would lose half the unpolarized emission, and between zero and 100 per cent of the linearly polarized maser emission, depending on the orientation of the source to the telescope feed (the parallactic angle) and the plane of the polarization characteristic of the particular source. If the de nition of variability is broadened to include both the variability due to a change in the intensity and the apparent change in intensity due to a change in the plane of polarization within the source, as a function of phase or time, then the results and conclusions of this work remain largely unaffected, provided that the parallactic angles are similar at different epochs of observation for a particular source. The only problem is that the `real' intensity variation and that due to rotation of the plane are convolved in some unknown manner, and a deconvolution is not possible with the current set of data, since the degree of polarization is unknown. Is it possible to ignore the variation of the parallactic angle between two observing epochs for the same object? The parallactic angle is de ned (Lang 1974) as cos l sin h q ˆ arcsin ; 1 sin z where l is the observer's latitude, h is the hour-angle of the source, and z is the source zenith distance. For the JCMT on Mauna Kea, cos l ˆ 0:941. The following observations had a spread of parallactic angles smaller than 15 per cent at all epochs and therefore have an error of < 7 per cent, even if the source were 100 per cent linearly polarized: the 299-GHz observations of R Leo, o Cet and x Cyg, the 301-GHz observations of W Hya, the 345-GHz observations of o Cet, and the 342-GHz observations of VY CMa. The huge variation of the 301- GHz maser in W Hya (Fig. 4) is therefore entirely due to phase effects. In the case of the stars with no well-de ned periodicity, there were only small changes in observed intensities and lineshapes, although all these objects were observed at a variety of parallactic angles. In the case of VX Sgr and NML Cyg, spectra taken less than 48 h apart, but at very different parallactic angles, were not signi cantly different. There were also only small differences in the spectra of R Crt and RT Vir and in the 301- and 299- GHz spectra of VY CMa, taken in observations separated by months, which also had large parallactic angle shifts. In general, we can therefore conclude that the spectral features of the supergiant and SRb stars are largely unpolarized. Only the variation in the VY CMa 345-GHz spectrum, and the disappearance of the wing feature at 299 GHz (Fig. 7) might perhaps be due to parallactic angle change. We note that wing features are more likely to be polarized than the core spectrum (Cernicharo et al. 1997). For the remaining observations of periodic variables, most observations fall into the category of random parallactic angles at random phases, and little can be said about these observations. A few observations were, however, made at very similar phases: the