First mm-vlbi Observations between the TRAO 14-m and the NRO 45-m Telescopes: Observations of 86 GHz SiO Masers in VY Canis Majoris

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1 PASJ: Publ. Astron. Soc. Japan 56, , 2004 June 25 c Astronomical Society of Japan. First mm-vlbi Observations between the TRAO 14-m and the NRO 45-m Telescopes: Observations of 86 GHz SiO Masers in VY Canis Majoris Katsunori M. SHIBATA, 1 Hyung-Soo CHUNG, 2 Seiji KAMENO, 1 Duk-Gyoo ROH, 2 Tomofumi UMEMOTO, 1 Kwang-Dong KIM, 2 Keiichi ASADA, 1 Seog-Tae HAN, 2 Nanako MOCHIZUKI, 3 Se-Hyung CHO, 2 Satoko SAWADA-SATOH, 1 Hyun-Goo KIM, 2 Takeshi BUSHIMATA, 1 Young Chol MINH, 2 Takeshi MIYAJI, 1 Nario KUNO, 4 Hiroshi MIKOSHIBA, 4 Kazuyoshi SUNADA, 4 Makoto INOUE, 1 and Hideyuki KOBAYASHI 1 1 National Astronomical Observatory, Osawa, Mitaka, Tokyo shibata@hotaka.mtk.nao.ac.jp 2 Taeduk Radio Astronomy Observatory, Korea Astronomy Observatory, 36-1 San, Whaam, Yusong, Taejon , Korea 3 The Institute of Space and Astronautical Science, Yoshinodai, Sagamihara, Kanagawa Nobeyama Radio Observatory, National Astronomical Observatory, Nobeyama, Minamimaki, Minamisaku, Nagano (Received 2004 February 16; accepted 2004 April 4) Abstract We have made VLBI observations at 86 GHz using a 1000-km baseline between Korea and Japan with successful detections of SiO v =1,J = 2 1 maser emissions from VY CMa and Orion KL in 2001 June. This was the first VLBI result for this baseline and the first astronomical VLBI observation for the Korean telescope. Since then, we observed SiO v =1,J = 2 1 maser emission in VY CMa in 2002 January and 2003 February and derived the distributions of the maser emissions. Our results show that the maser emissions extend over 2 4 stellar radii, and were within the inner radius of the dust shell. We observed other SiO maser sources and continuum sources, and 86-GHz continuum emissions were detected from three continuum sources. It was verified that this baseline has a performance comparable to the most sensitive baseline in the VLBA and the CMVA, and is capable of investigating the proper motions of maser features in circumstellar envelopes using monitoring observations. Key words: radio lines: stars stars: individual (VY Canis Majoris) stars: supergiants techniques: interferometric 1. Introduction Among circumstellar masers, SiO masers are known to locate closer to the star, about 2 6 stellar radii, than H 2 O and OH. Many rotational transitions of vibrationally excited states of SiO masers have been found in evolved stars (Cho et al. 1998; Pardo et al. 1998; and references therein). From these characteristics, SiO masers are considered to be good probes for investigating the physical conditions and kinematics in the circumstellar envelope close to the stars. Hence, VLBI observations of SiO J = 1 0 maser emission in evolved stars have been made since the early 1990 s (Diamond et al. 1994; Miyoshi et al. 1994; Greenhill et al. 1995; Kemball, Diamond 1997; Boboltz et al. 1997; Boboltz, Marvel 2000). Recently, a higher rotational transition of SiO v =1,J = 2 1 maser emissions has been observed in VX Sgr, o Cet, and R Cas by Doeleman et al. (1998), Phillips and Boboltz (2000), and Phillips et al. (2003), respectively, with the Coordinated Millimeter VLBI Array (CMVA) and the VLBA. Doeleman et al. (1998) imaged SiO maser emissions in the supergiant VX Sgr from a single VLBI baseline, and showed that the maser spots were arranged in a ring-like configuration with a diameter of about 30 mas. They found a systematic velocity gradient among the maser features on the southern portion of the ring. Phillips and Boboltz (2000) observed o Cet with the Present address: ASSIA, P.O.Box , Taipei 106, Taiwan. CMVA, and detected the maser emissions on two independent baselines. The maser image of o Cet showed that the maser features extended in the north south direction with a size of about 100 mas. Phillips et al. (2003) made simultaneous observations of the SiO v =1,J = 1 0 and J = 2 1 in R Cas at the stellar phase of 0.54 with the VLBA. They compared their images and the relative strengths of the J = 1 0 and J = 2 1 masers with the models of Humphreys et al. (2002) and suggested that the observed weaker emissions from J =2 1 than those from J = 1 0 were not consistent with the model unless the stellar phase of the model can be shifted. VY CMa is a luminous supergiant, L = L,ata distance of 1.5 kpc (Lada, Reid 1978), which has a high massloss rate of about M yr 1 (Danchi et al. 1994). VY CMa is also an irregular variable star, whose period is unclear but is roughly estimated as being about 1300 days in the optical region (Marvel 1996). Strong OH, H 2 O, and SiO masers have been detected in VY CMa. Richards et al. (1998) imaged the H 2 O maser emissions in VY CMa with MERLIN, and showed the maser features extended over 700 mas east west and 400 mas north south. Compared with an image from 1985 (Bowers et al. 1993), they inferred that the H 2 O maser features in VY CMa had motions directed away from the assumed stellar position. The SiO v =1,J = 2 1 masers in VY CMa were discovered by Kaifu et al. (1975), and other masers in many rotational transitions of vibrationally excited states were found by Cernicharo et al. (1993). Miyoshi (2003)

2 476 K. M. Shibata et al. [Vol. 56, Table 1. Characteristics of telescopes at 86 GHz. Telescope Diameter Beam width Aperture efficiency Gain (m) ( ) (%) (JyK 1 ) TRAO NRO Table 2. Observational parameters. Epoch Date Polarization T sys (K) TRAO NRO TRAO NRO Epoch Jun. 02 Linear LHCP Epoch Jan. 06 LHCP LHCP Epoch Feb. 27 LHCP LHCP made a VLBA observation of SiO v =1and2,J = 1 0 in VY CMa in 1998 October. He suggested that the observed positional coincidence between the v =1andv =2masers within 0.2 mas (0.3 AU for the distance of 1.5 kpc) preferred the collisional pumping model in VY CMa. An image of SiO J = 2 1 emissions in VY CMa with high angular resolution has not yet been presented. In addition to the VLBA and CMVA telescopes, the Taeduk Radio Astronomy Observatory (TRAO) of the Korea Astronomy Observatory in Korea and the Nobeyama Radio Observatory (NRO) of the National Astronomical Observatory in Japan both have millimeter-wave telescopes. TRAO has a radome-enclosed 14-m telescope that operates from 86 to 175 GHz. This telescope had not previously joined in VLBI observations. The NRO 45-m telescope operates at the 22, 43, 80, and 100-GHz bands. This telescope has participated in some VLBI observations, such as global mm-vlbi observations, and is now leading the domestic VLBI network in Japan (J-Net). The large collecting area and the high efficiency of the NRO 45-m telescope in the mm-wavelength region helps to form a powerful baseline having a similar detection sensitivity to the most sensitive baselines in the VLBA or the CMVA. We have therefore promoted mm-vlbi observations using these two millimeter telescopes in east Asia. In this paper, we describe our observing systems and the results of our mm- VLBI observations. 2. Observations We have made observations at 5 epochs: 2001 June, 2002 January, March, and May, and 2003 February. However, the observations in 2002 March and May failed because of a problem in the standard reference system at the NRO. Hence, we will focus on our description on the observations in 2001 June (hereafter, epoch 0), 2002 January (epoch 1), and 2003 February (epoch 2) Telescopes and Receivers We made VLBI observations using the TRAO 14-m telescope and the NRO 45-m telescope. The characteristics of each telescope are listed in table 1. We used 100-GHz SIS receivers and H-maser frequency standards at both sites. At the NRO, the output signal in a 5 7 GHz IF band from the receiver was converted to frequencies in the MHz range, and then transferred to the VSOP type VLBI backend (see subsection 2.2). In order to observe Left Hand Circular Polarization (LHCP), we put a circular polarizer in front of the feed horn. At the TRAO, we converted GHz IF signals into the MHz frequency range, and observed Linear Polarization in epoch 0. From epoch 1, we used the same type of circular polarizer as for the NRO, setting it in front of the horn in the TRAO. We summarize our observational parameters in table 2. The system temperatures in table 2 were the values that we measured near to the source direction. At the TRAO, the weather conditions were clear during three epochs. At the NRO, it was clear for epochs 0 and 1, and cloudy for epoch 2. The very high system temperature for the NRO 45 m at epoch 2 was due to not only the weather condition, but also the receiver condition VLBI System The data were recorded using the VSOP terminal (Kawaguchi et al. 1994) with a recording rate of 128 Mbps (one channel of 32-MHz bandwidth and 2-bit sampling) at both telescopes. The recorded data were correlated with the Mitaka FX correlator (Chikada et al. 1991) and the output had 1024 spectral channels resulting in a velocity resolution of 0.109kms 1. Because the TRAO had not been made any VLBI observations until our observations, we used its geographical coordinates, which were measured by a GPS system on 2001 June in the first correlation. We then added corrections that were calculated from the residual fringe rates of the radio source The geocentric coordinates that were used in the final correlation for the TRAO telescope were X = m, Y = m, and Z = m Phase Stabilities Before the VLBI observations, phase-stability tests were made by injecting an 86-GHz tone signal from the receiver horn at both telescopes. This 86-GHz tone signal was generated by a multiplier (doubler) and a synthesizer which was

3 No. 3] First mm-vlbi between TRAO and NRO 477 Table 3. List of detected sources. Name Correlated flux SNR SNR Baseline length Epoch (Jy) (Obs) (Calc) (Mλ: projected) SiO v =1J = 2 1 maser VY CMa 22 ± Epoch 2 Orion KL 31 ± Epoch 1 86 GHz continuum ± Epoch ± Epoch 2 3C 273B 2.5 ± Epoch 2 3C ± Epoch 2 At peak component for maser sources. Observed SNR is determined by the integration time of 30 s. SNR calculated from equation (1). phase-locked to a standard reference. The receiver at the TRAO was located at the Cassegrain focus, and tilted with the telescope. Hence, we measured the phase stability in both cases with the telescope pointing towards the horizon and when it was tracking stellar objects during observations. On the other hand, we measured the phase stability only in the case of pointing towards the zenith at the NRO because the receiver was stationary (located in the lower cabin of the telescope building). At the TRAO, the difference between the phase of the injected tone signal, which was converted to the MHz range, and that of the reference tone at the same frequency, which was generated from the phase-locked synthesizer, was monitored using a vector voltmeter; the difference was confirmed to be within (peak-to-peak) after some improvements. At the NRO, we used a phase-calibration tone detector that was installed in a sampler to monitor the phase stability, and confirmed that the phase noise was similar to that at the TRAO. We also have recorded the injected tone signal through the VLBI backend on tape in order to examine the phase stabilities with a higher time resolution than 1 second, which was the time resolution of the measurements by the vector voltmeter and the phase-calibration tone detector. The recorded data of the tone signal were transferred into the disk of a computer and were analyzed with a time resolution of 64 microseconds. The results showed that the phase noise was about 40 (peakto-peak) with modulations of about 30 Hz and 0.7 Hz seen in the TRAO data; the phase noise was about 30 (peak-to-peak) with modulations of about 60 Hz and 1 Hz in the NRO data. These modulations may result from the AC power supply and the receiver cryogenics. 3. Results 3.1. Detection Limits We successfully obtained fringes from SiO v =1,J = 2 1 maser emission at GHz in VY CMa and Orion KL, and from 86-GHz continuum emission in for the observations at epoch 0. These were the first detections for the TRAO NRO baseline, and first VLBI fringes ever for the TRAO. In addition to the above sources, we have observed SiO 86-GHz maser emissions in 12 evolved stars and 86-GHz continuum emissions in two QSOs. We present the correlated flux density and the SNR of the detected sources in table 3. The error in the correlated flux density was mainly caused by the variation of the system temperature during the observations. Undetected maser sources were TX Cam, IK Tau, GX Mon, RLeo,UOri,RLMi,oCet,RCrt,WXSer,UHer,WHya, and IW Hya. The maser spots in these evolved stars might be resolved out by our 1000-km baseline, because VY CMa is the farthest star, at a distance of 1.5 kpc (Lada, Reid 1978), among the evolved stars that we observed. We present the calculated SNR, which was derived from the parameters given in the tables using the following equation, where F is the observed correlated flux in Jy K 1 is the telescope gain in (Jy K 1 ), T sys is the system temperature, ν is the bandwidth in Hz, τ is the integration time, and η s is the correlation loss (we use η s =0.88), in table 3: SNR = η s F 2τ ν K i T sysi K j T sysj. (1) Although the observed and calculated SNR for 3C 279 show good agreement, the observed SNRs for other sources are slightly higher, 10 40%, than that of the calculated SNRs. These discrepancies suggest that either the derived correlated flux densities were underestimated, or the system temperatures applied to equation (1) were overestimated. It is more likely that the variation of the system temperature during the observation caused an underestimation of the correlated flux density because we had measured the system temperature only two or three times during observations. The visibility amplitudes were calibrated based on the system temperatures, which were obtained by interpolation of the measured temperatures listed in table 2 in our analysis. If the system temperature at the time when we had observed a source was higher than the system temperature that was used in the amplitude calibration, the correlated flux density calibrated by the interpolated system temperatures would be underestimated. We consider that our observations with the baseline between the TRAO and the NRO exhibit almost the best performance that could be expected from the characteristics of both telescopes. The estimated detection limits for the SNR of 7 for a velocity resolution of

4 478 K. M. Shibata et al. [Vol. 56, 0.5kms 1 and a bandwidth of 32 MHz are 11 Jy and 0.8 Jy, respectively, assuming a system temperature of 300 K in both telescopes Distribution of SiO v = 1, J = 2 1 Maser Emissions in VY CMa The red supergiant VY CMa was observed for five 7-min scans at epoch 1 and five 40-min scans at epoch 2. Also, was observed as a delay and bandpass calibrator at epoch 1, and and 3C 273B were observed as calibrators at epoch 2. Figure 1 shows the (u, v) coverage for VY CMa. The total power spectra derived from autocorrelation data of the NRO telescope are shown in figure 2. The data were analyzed by a standard procedure (e.g., Doeleman et al. 1998) using the NRAO AIPS package. Maser peaks at the radial velocity with respect to the LSR (hereafter we use the radial velocity with respect to the LSR) of 10.2 and 25.5kms 1 in epoch 1 and epoch 2, respectively, were selected as the reference channel, which showed fringes during the whole observing run and has no indication of multi-spot structure. We fringe-fit these components and applied the results to the other channels. Then, maser images were obtained by CLEAN with five frequency-channels average, resulting in a velocity resolution of 0.54 km s 1 ; the location and flux density of maser spots were derived by a 2-dimensional onecomponent Gaussian fitting. The resultant maser distributions are shown in figures 3a and b for epoch 1 and epoch 2, respectively, and the cross-power spectra in figure 2. The beam size was 1.70mas 0.26mas (P.A. = 179 )and1.73mas 0.25mas (P.A. = 4. 3) in epoch 1 and epoch 2, respectively. The origin of each map was originally arbitrary, and the position of the reference maser spot was different in each map. Hence, we shifted the origin of the maps to the mean position of the features, which have radial velocities of 0 5 km s 1, because these features are relatively concentrated in both maps and appear to be located on the same part of the whole distribution. Due to the poor (u,v) coverage and/or high sidelobe level of about 90% in both observations, images occasionally have fake spots in these distributions, especially in the velocity channel which has a weak component and/or multi-component structure. Figures 3a and b show that the maser emissions distribute over the same extent at both epochs except for the lack of features at the eastern portion in epoch 2, and the velocity structures have similar trends at both epochs. These facts indicate that most of the derived maser spots in figure 3 are reliable. The total-power spectra in figure 2 show the variation of the SiO v =1,J = 2 1 emission profile between epoch 1 and epoch 2, especially in the velocity ranges of 5 12 and km s 1. A similar trend can also be seen in the crosspower spectra. This variation of the maser profile caused a lack of maser features at the eastern portion in the radial velocity of km s 1 in epoch 2, and weakened the intensity of maser spots of the corresponding velocities. We could detect no maser spots whose radial velocity was smaller than 0kms 1 or greater than 35kms 1 in epoch 2. Similarly to other circumstellar masers, such as OH, H 2 O, and SiO J =1 0,theSiOv =1,J = 2 1 masers in VX Sgr and R Cas had ringlike configurations of maser features, which probably have a star at the center (Doeleman et al. Fig. 1. The (u,v) coverage of VY CMa. The locus for the observation on 2002 January is plotted with 90% reduced scale in order to avoid confusion. Fig. 2. Total and cross-power spectra of SiO v =1,J =2 1maserin VY CMa. The date of the observation is indicated in each panel. The total-power spectrum is derived from auto-correlation data of the NRO 45m and is shown by the thin line. The cross-power spectrum is the sum of the flux density in CLEANed maser image, having a five-times coarser velocity resolution of 0.54 km s 1, and is shown by the thick line. 1998; Colomer et al. 1996; Phillips et al. 2003). In the case of VY CMa, it seems that the maser emissions are also distributed along a ringlike configuration, although we see only the southern half, but it is not clearer than those in VX Sgr and R Cas. Monnier et al. (2000) have estimated the stellar radius of VY CMa as being 10 mas, and the inner radius of the dust shell as 50 mas from mid-infrared observations. In figure 4, we put a circle of radius 10 mas, representing the star, at the possible center of the hypothetical ring, which was fitted to

5 No. 3] First mm-vlbi between TRAO and NRO 479 Fig. 3. SiO v =1,J = 2 1 maser distribution in VY CMa; (a) on 2002 January and (b) on 2003 February. The line color indicates the LSR radial velocity of maser emission. The integrated flux of maser spot is indicated by the filled circle, small open circle, middle open circle, and large open circle for the fluxes of F < 2Jy, 2 F < 10Jy, 10 F < 30Jy, and 30Jy F, respectively. The origin of map, (0, 0), is the mean position of masers, whose radial velocity with respect to the LSR is between 0 and 5km s 1. Fig. 4. SiO v =1,J = 2 1 maser distribution in VY CMa; (a) on 2002 January and (b) on 2003 February. These are representations of figure 3 without the flux scale and the radial velocity indication. In the lower left-hand corner, an inset shows the beam size of each map. The broken circle shows the stellar size derived by Monnier et al. (2000), but its position is decided by eye, assuming that the maser emissions have a ringlike configuration centered on the star. the maser distributions by eye. Assuming that the masers in VY CMa have a ring configuration and that the star position is correct, the maser emissions are distributed over radii of 2 4r. This is comparable to the maser emissions in R Cas (Phillips et al. 2003), but farther than that in VX Sgr (<1.5r ;Doeleman et al. 1998) and o Cet (1.5 2r ; Phillips, Boboltz 2000). Most of the masers are emitted within the inner radius of the dust shell (50 mas). This is consistent with the theoretical argument that the silicon would be incorporated into dust and the formation of SiO will be suppressed in the dust shell (Lockett, Elitzur 1992; Diamond et al. 1994; Greenhill et al. 1995). Although the existence of systematic velocity structures in SiO maser emissions was reported for NML Cyg and R Cas (Boboltz, Marvel 2000; Phillips et al. 2003), we did not find a systematic velocity structure in the SiO v =1, J = 2 1 maser emissions in VY CMa. We have not investigated the proper motions of maser features because the (u, v) coverage for epoch 1 was very poor compared with that of epoch 2, and the interval of one year between epoch 1 and epoch 2 was considered to be too long to discuss the variation of the maser features. We would like to continue the observations for the SiO masers in VY CMa more frequently with shorter intervals in order to derive the proper motions and to investigate the kinematics in the circumstellar envelope of VY CMa.

6 480 K. M. Shibata et al. 4. Summary We conducted first VLBI observations at 86 GHz on the 1000-km baseline between the TRAO and the NRO, and detected SiO v =1,J = 2 1 maser sources and continuum sources. We also made mapping observations of SiO 86 GHz masers in VY CMa with an interval of about one year, and derived the maser distributions at the two epochs. The maser emissions in VY CMa extend over 2 4r and within the inner radius of the dust shell. From our results, it has been confirmed that this baseline has a good performance for the mm-vlbi. However, the long interval between epochs and the poor (u,v) coverage of the observation in 2002 January prevented us from any detailed discussion about the variation of the maser features and/or maser motions. It is necessary to observe more frequently with shorter time intervals. We thank the staff of TRAO and NRO for their kind support during our test measurements and observations. We specially thank the NMA group of NRO for providing the polarizer. KMS also thanks Dr. Y. Fukuzaki for his advice about coordinate calculations. A part of this work was supported by a Grant-in-Aid for Scientific Research from the Ministry of Education, Culture, Sports, Science and Technology (No ) and by KAO s Main Program from Korea Research Council of Fundamental Science & Technology (No , No ). References Boboltz, D. A., Diamond, P. J., & Kemball, A. J. 1997, ApJ, 487, L147 Boboltz, D. A., & Marvel, K. B. 2000, ApJ, 545, L149 Bowers, P. F., Claussen, M. J., & Johnston, K. J. 1993, AJ, 105, 284 Cernicharo, J., Bujarrabal, V., & Santarén, J. L. 1993, ApJ, 407, L33 Chikada, Y., et al. 1991, in Frontiers of VLBI, ed. H. Hirabayashi, M. Inoue, & H. Kobayashi (Tokyo: Universal Academy Press), 79 Cho, S.-H., Chung, H.-S., Kim, H.-R., Oh, B.-Y., Lee, C.-H., & Han, S.-T. 1998, ApJS, 115, 277 Colomer, F., Baudry, A., Graham, D. A., Booth, R. S., de Vicente, P., Krichbaum, T. P., Gómez-González, J., & Schalinski, C. 1996, A&A, 312, 950 Danchi, W. C., Bester, M., Degiacomi, C. G., Greenhill, L. J., & Townes, C. H. 1994, AJ, 107, 1469 Diamond, P. J., Kemball, A. J., Junor, W., Zensus, A., Benson, J., & Dhawan, V. 1994, ApJ, 430, L61 Doeleman, S. S., Lonsdale, C. J., & Greenhill, L. J. 1998, ApJ, 494, 400 Greenhill, L. J., Colomer, F., Moran, J. M., Backer, D. C., Danchi, W. C., & Bester, M. 1995, ApJ, 449, 365 Humphreys, E. M. L., Gray, M. D., Yates, J. A., Field, D., Bowen, G. H., & Diamond, P. J. 2002, A&A, 386, 256 Kaifu, N., Buhl, D., & Snyder, L. E. 1975, ApJ, 195, 359 Kawaguchi, N., Kobayashi, H., Miyaji, T., Mikoshiba, H., Tojo, A., Yamamoto, Z., & Hirosawa, H. 1994, in VLBI Technology Progress and Future Observational Possibilities, ed. T. Sasao, S. Manabe, O. Kameya, & M. Inoue (Tokyo: Terra Scientific Publishing Company), 26 Kemball, A. J., & Diamond, P. J. 1997, ApJ, 481, L111 Lada, C. J., & Reid, M. J. 1978, ApJ, 219, 95 Lockett, P., & Elitzur, M. 1992, ApJ, 399, 704 Marvel, K. B. 1996, PhD Thesis, New Mexico State University Miyoshi, M. 2003, in Mass-Losing Pulsating Stars and their Circumstellar Matter, ed. Y. Nakada, M. Honma, & M. Seki (Dordrecht: Kluwer Academic Publishers), 303 Miyoshi, M., Matsumoto, K., Kameno, S., Takaba, H., & Iwata, T. 1994, Nature, 371, 395 Monnier, J. D., Danchi, W. C., Hale, D. S., Lipman, E. A., Tuthill, P. G., & Townes, C. H. 2000, ApJ, 543, 861 Pardo, J. R., Cernicharo, J., González-Alfonso, E., & Bujarrabal, V. 1998, A&A, 329, 219 Phillips, R. B., & Boboltz, D. A. 2000, AJ, 119, 3015 Phillips, R. B., Straughn, A. H., Doeleman, S. S., & Lonsdale, C. J. 2003, ApJ, 588, L105 Richards, A. M. S., Yates, J. A., & Cohen, R. J. 1998, MNRAS, 299, 319