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2 Abstract The regions adjacent to super massive black holes (SMBHs) and protostars are ideal laboratories in which study physical phenomena in extreme environments far removed from terrestrial conditions. It is now believed that at the center of nearly every galaxy is a SMBH which ranges in mass from solar masses which dramatically distorts the inner galactic potential. VSOP-2 can explore this huge gravitational potential well within light days to light months from the central engine depending on the distance to the active galactic nucleus (AGN). Furthermore, while it is known that stars are formed during the gravitational collapse of dense molecular cloud cores, much less certain is the detailed kinematics and morphological changes which take place during star formation which occurs on sub-solar diameter scales. Consequently astronomers would dearly like to have a powerful telescope to probe these fundamental astrophysical phenomena. The emergence of Very Long Baseline Interferometry (VLBI) technique with its direct imaging capability allows astronomers to study both these phenomena with the highest angular resolution of any astronomical technique. Moreover, Space-VLBI using an orbiting radio telescope operating in conjunction with ground telescope arrays allows a synthetic aperture considerably more than Earth s diameter in size to be created which enables new higher resolution observations to be undertaken such as imaging of the accretion disks surrounding the SMBHs in the nearest AGN. In 1997 the Japanese Institute of Space and Astronautical Science (ISAS) launched the world s first dedicated VLBI spacecraft HALCA (Highly Advanced Laboratory for Communication and Astronomy) which undertook VLBI observations in conjunction with ground-bases radio telescopes until November 2005 and was the critical space-based element of the Japaneseled VLBI Space Observatory Programme (VSOP). Recently ISAS/JAXA approved the next generation Space-VLBI mission VSOP-2 for development with a planned launch date in early The VSOP-2 spacecraft (ASTRO-G) will employ a 9-m off-axis paraboloid antenna with cryogenically cooled receivers operating at 8, 22, and 43 GHz in dual-polarization. An apogee height of 25,000 km will yield an angular resolution of 38 micro-arcseconds at 43 GHz. A phase-referencing capability is being considered which would not only increase the number of sources that can be observed with Space-VLBI but also enable forefront astrometrical observations to be undertaken. With these unique technical capabilities in our hands, we are aiming to

3 use ASTRO-G to provide answers to some of the fundamental problems in modern astrophysics This document is divided into three chapters: Chapter 1 is a brief executive summary of the VSOP-2 project, Chapter 2 presents the science goals of the mission. and finally Chapter 3 provides an overview of the satellite and mission. The text in this document will be continually revised and updated until the launch of ASTRO-G. The latest version can be found at:

4 Contents 1 Executive summary Background Key Science Proposals Other Science Proposals Overview of VSOP-2 Mission Overview of VSOP-2 Mission International collaboration Brief Current Status of VSOP Approval of VSOP-2 (ASTRO-G) Mission timeline VSOP-2 Science goals Introduction Active Galactic Nuclei (AGNs), the most powerful engines in the Universe Astrophysical Jets - Internal structure of relativistic flows Scientific results from ground VLBI and VSOP observations Accreting matter onto AGN Radio Lobe Expansion AGN jets and VSOP D magnetic structure Black holes the heart of the beast Imaging of black hole environments Accretion disks the power plants in AGN Accretion disk models Approaching to accretion disks with VSOP Mass accretion process onto AGNs Plasma distribution Extragalactic H 2 O maser Introduction VSOP-2 megamaser observations H 2 O masers in the Local Group Formation and evolution of stars A standard picture of young stellar objects A standard picture of dying stars Protostar s magnetospheres Rotation of protostars Binary systems H 2 O maser excitation and early stellar evolution The root of outflow ii

5 2.7.8 Proto-stellar disks Study on final stellar evolution Stellar SiO masers Stellar H 2 O masers Exploring magnetic fields in jets of YSOs and evolved stars Other Science Topics for VSOP Gamma-ray bursts Supernovae Mass mapping through microlensing VSOP-2 Mission Introduction From HALCA to ASTRO-G Other space-vlbi projects Requirements for VSOP Angular Resolution Sensitivity VLBI Polarimetry Outline of the satellite Improvements over VSOP Technical challenges for the ASTRO-G satellite Outline of the ground system International cooperation Responsibilities

6 Chapter 1 Executive summary 1.1 Background Countless generations of human beings have observed the night sky with awe and tried to explain celestial events. Our generation is now engaged in trying to understand exotic astrophysical phenomena such as black holes and the birth of stars which requires sophisticated forefront technology. Astrophysics itself is a well established discipline of modern physics and is the only domain in which certain fundamental physical questions can be addressed. Several different complementary instruments and techniques are used to explore the fine-scale structures of a wide range of astronomical objects from protostars to supermassive black hole accretion discs in order to study fundamental physics in extreme conditions to obtain an overall multi-waveband picture. In current astronomy the highest angular resolution is achieved using the technique of Very Long Baseline Interferometry (VLBI) in which ground-based radio telescopes typically located hundreds to thousands of km apart simultaneously observe the same source. Angular resolution is defined as the smallest angle two objects can be apart and yet still be individually distinguished. For any telescope angular resolution is given by the ratio of the observing wavelength to the telescope diameter. For VLBI in which ground-based radio telescope arrays are used this diameter is the maximum separation between the array elements. Space VLBI greatly enhances ground-based VLBI by adding an orbiting radio telescope to the array and consequently the maximum telescope separation is no longer limited to an Earth diameter and higher angular resolution can be achieved than just using the ground-based array. In 1997 the full potential of Space VLBI was realized when the Japanese Institute of Space and Astronautical Science (ISAS) launched the first dedicated VLBI radio telescope satellite (HALCA) which observed radio sources in conjunction with ground-based radio telescopes operated by the world-wide VLBI community in a project known as the Very Long Baseline Interferometry Programme (VSOP). In general, Space VLBI achieves extremely high angular resolutions up to 40 microarcseconds which is compared with other astronomical instruments in Table 3.1. The VSOP-2 mission consisting of the orbiting ASTRO-G radio telescopes operating with groundbased VLBI arrays will be able to image the innermost parts of both protostars and AGN accretions discs (whose linear sizes are thought to be less than 1 AU and light days to light months) as well as image the base of relativistic jets which emanate from the central region of radio-loud AGN such as radio-galaxies and quasars at a linear resolution of a few light-years. The VSOP-2 images will have a higher angular resolution by a factor of 2-3 compared to only ground-only VLBI images (see Figure 1.1). Images made by the VSOP mission demonstrated the imaging potential of Space VLBI very well. For example, Figure 1.2 shows that the inner jet of the quasar 3C345 (redshift 0.595) which is resolved into several components when observed by VSOP at 5 GHz. VSOP-2 observations of the 3C345 jet will certainly enable better constrain ts to be placed on the theoretical AGN jet production models. Thus, the VSOP mission quite clearly demonstrated capability of Space VLBI and can be thought 1

7 Table 1.1: Comparison of Angular Resolutions Telescope CHANDRA 1 HST 2 Ground VLBI VSOP-2 VLBA GMVA 3,4 Highest Resolution (arcsecond ) ( ) 5 Wave bands X-ray Optical/NIR ,(86) GHz 86 GHz 8,22,43 GHz 1. Space X-ray Observatory 2. Hubble Space Telescope 3. Global Millimeter VLBI Array 4. The longest baseline of 10,900 km is adopted. 5. The current highest resolution (NL-MK baseline) at 86 GHz (NRAO VLBA Status Report). Figure 1.1: Comparison of the spatial resolutions of ground-vlbis and VSOP-2 of as a precursor or pathfinder mission for VSOP-2. Indeed it showed that Space VLBI is a highly promising technique designed to address a wide range of fundamental questions in astrophysics. It will be able to reveal previously unexplored physical and astrophysical phenomena occurring in the vicinity of the central engine. However, due to the relatively low sensitivity of VSOP observations, sources such as X-ray binary systems, stellar masers in the Milky Way, and many of other radio-quiet objects could not be observed. The significantly improved sensitivity of VSOP-2 will allow observations of such weaker sources, and consequently the number of potential target sources for VSOP-2 able will increase greatly increased compared to VSOP. In the coming decade VSOP-2 will be the only Space VLBI mission. In summary, VSOP-2 is going to enable direct imaging of protostars and accretion disks around supermassive black holes (SMBHs) at the highest angular resolution ever attained in astronomy. 2

8 Figure 1.2: Upper: The 6-cm radio continuum image of the quasar 3C345 imaged using only the ground VLBI facility. The core (brightest part) is spatially unresolved. Lower: The zooming-up of the 3C345 core obtained by the VSOP observations. The highest angular resolution by addition of the data from the HALCA (The space-vlbi telescope) baselines enabled the inner jet structure to be more resolved into several components. The image courtesy of J.Klare and T.Krichbaum. 1.2 Key Science Proposals Attainment of the VSOP-2 science goals requires observing in in the millimeter wave-band, enabling imaging on the scale of the accretion disk surrounding the supermassive black holes in the heart of active galactic nuclei (AGN) as well as the AGN jet acceleration region. Furthermore it allows the structures of protostellar magnetospheres to be clarified. As a result, VSOP-2 would allow studies of regions where extreme physical conditions are encountered. Consequently the high-resolution imaging capability of VSOP-2 will enable new science in fundamental astrophysics to be undertaken. The design of the VSOP-2 instruments is intended to realize the following science goals: The structures and magnetic field configuration of accretion disks in active galactic nuclei (AGNs) The mechanism of jet acceleration and collimation The motion of H 2 O masers in galactic star forming regions Proto-stellar magnetospheres The Structures and Magnetic fields of Accretion Disks in Active Galactic Nuclei The size of the accretion disk around a supermassive black hole is tens of Schwarzschild radii (R s ). An angular resolution of tens micro-arcseconds is necessary to resolve the structure of such accretion 3

9 disks as they are expected to be about R s in extent. For instance, the mass of the black hole in M87 is M, and 1 R s corresponds to an angular size of of three micro-arcseconds. The VSOP-2 beam-size at 43 GHz is 38 micro-arcseconds, corresponding to 13 R s for M87, and so it is expected that the size of the accretion disk can be determined. Within 20 Mpc of our Galaxy, there are three sources for which the VSOP-2 beam size at 43 GHz is less than 3 R s, and 13 galaxies for which the beam size is less than 200 R s. Figure 1.3: The nearby elliptical galaxy M87 hosts the bright active galactic nucleus in its center, offering the possibility of studying an AGN at the highest angular resolution. VLA images (upper) show the large-scale radio jet, while the ground-vlbi image (bottom) indicates the inner part of the jet in the vicinity of a super massive black hole (SMBH). These images are shown at scales ranging from 3,3000 light years (10 kpc) to 0.4 light months (0.01 pc) (Junor et al. 1999). The M87 jet is thus thought to be powered by the SMBH of about 3 billion solar masses. Image courtesy of National Radio Astronomy Observatory. The Mechanism of Jet Acceleration and Collimation Ground-VLBI observations probe regions pc from the center of active galactic nuclei (AGN). However, VSOP-2 will allow regions interior to this to be observed and thus allow jet generation, acceleration and collimation models to be further constrained. Furthermore determination of magnetic field structures as can be done by VSOP-2 is is essential to understanding the collimation and propagation of jets in astrophysical environments. Figure 1.4 shows the VSOP image of the radio jet in the nearby giant elliptical galaxy M87. VSOP-2 observations of M87 will allow astronomers to obtain more detailed structural information on the M87 jet emanating from its central engine. The Motion of H 2 O Masers in Galactic Star forming Regions Water maser emission in the 22 GHz band is detected in star forming regions, late-type stars, and extra-galactic sources typically in compact features called maser spots. Combining phase referenc- 4

10 Figure 1.4: Nuclear jet of M87, imaged by the Space-VLBI, VSOP. The image reveals details of the inner structure of the nuclear jet; the jet is limb-brightened to within 2 milliarcsecond of the central engine. Image courtesy of M.Reid. ing VLBI or Space-VLBI observations with radial velocity information derived from the maser spot Doppler shift it is possible to reconstruct the the 3-D motions of the maser spots. Such 3-D motions are required to for determining the gravitational field in which gas is moving and for the study of other environmental conditions. In star forming regions, water maser motions appear as an outflow from the protostellar disk which allows important constraints to be placed on star formation mechanisms. Compared to ground-vlbi observations, VSOP-2 observations have the distinct advantage of being able to measure proper motions over a shorter period of time, and, as a result, will enable shortlifetime maser spots motions to be determined. Fig.1.5 shows an montage of both radio and optical images of the S106 FIR which includes a ground-vlbi image of the microjet from this star-forming region. The VLBI observations at a 0.4 milliarcsecond resolution allowed the microjet proper motion field as probed by the 22 GHz water masers to be determined. These masers proper motions reveal a microbowshock created by the impact of the microjet on a dense water-containing gas medium. Millimeter-wave interferometer molecular-line observations and continuum observations on the on the AU scale and VLBI maser observation on the AU scale have enabled research into protostar formation by mass accretion and its associated outflows to be undertaken. VSOP-2 star formation studies will be roughly divided into two subjects, 1) investigating the three dimensional motions of the accretion disk and outflow (jets) on the scale of 10 AU or less in the H 2 O maser emission, and 2) investigating protostar magnetosphere inside an accretion disk and interaction with an accretion disk. Figure 1.5: An optical and radio images of a low-mass star-forming site in the S106 HII region (upper left). Microbowshock probed by water maser, obtained from the four-epochs VLBI observations (upper right). Dense gas outflowing from the central star encounters the ambient molecular medium with a shock front, giving a rise to maser excitation. The image courtesy of R. S. Furuya. 5

11 Proto-stellar Magnetospheres Magnetospheres in protostars controls both the angular momentum of the accretion disk and the out-flow of material in these systems. It is thought that magnetic reconnection in the the protostar magnetosphere causes huge flares which result in the observed X-ray emission from these objects. It is also known that there is a correlation between X-ray emission and radio emission, and it is important to determine the time variation and the brightness temperature of this radio emission, which in turn requires sensitive high angular resolution observations. At 150 pc, the VSOP-2 angular resolution of 38 µas corresponds to AU (1.2 R ) in the star forming region, and detailed imaging of the area in the magnetosphere where the flare occurs can be undertaken. The brightness temperature of weak line T Tauri stars can be measured with VSOP-2 using phase referencing. Non-thermal radio emission from the magnetosphere of the protostar is believed to be produced by gyro-synchrotron radiation. Furthermore, the structure of the magnetic field in the protostar can be clarified for the first time with dual polarization imaging. 1.3 Other Science Proposals Extragalactic H 2 O Masers Extragalactic water masers are found within a few parsecs of several active galactic nuclei. Figure 2.10 shows an image of the 22 GHz masers in the spiral galaxy NGC 4258 which are located 40,000 Schwarzschild radii from the central engine in a rapidly rotating accretion disk. Monitoring observations to determine the three dimensional motions of these maser spots were used to prove the existence of a supermassive black hole in center of NGC The motions of such spots allows not only the supermassive black hole mass in the AGN center to be obtained but the temperature and density of the emitting gas can also be probed revealing the physical conditions of the accretion disk. At present about 50 extragalactic water megamaser sources are known. An advantage of VSOP-2 observations for the megamasers is, again, the ability to measure the proper motions of the maser spots in a short time than required for ground-based observations. Figure 1.6: The image shows the results of ground VLBI observations of water maser emission in the nearby spiral galaxy NGC The geometry of the maser distribution on the disc and the radio jet in the center. The dots denote the locations of each maser spot. Image courtesy of J.Herrnstein. 6

12 Relativistic Jets from Microquasars A microquasar is a galactic binary system in which material is ejected from the more compact object, such as a neutron star or a stellar-mass black hole, as two oppositely directed radio-emitting relativistic jets. Also due to angular momentum conservation a hot X-ray emitting accretion disk forms around the compact object as it is accreting material from its larger companion. Since the mass of the microquasars is very small about a 1-10 solar mass, compared with those of quasars of a few billion solar masses, microquasar jet evolution is very rapid on timescale ranging from days to weeks. Thus these object are ideal targets for VSOP-2 which will be able to follow the jet time-evolution at high angular resolution with full polarimetry. Radio Quiet Quasars At least 90% of active galaxies are radio-quiet. Many Seyfert galaxies and weak radio quasar are low brightness AGN. Most of these sources were not able to be observed with VSOP due to HALCA s limited sensitivity. However, high resolution VSOP-2 observations will be possible for a number of radio-quiet AGN which will allow the core and parsec-scale structures to be studied for these low radio luminosity sources. Final Evolution of Stars H 2 O and SiO masers associated with evolved stars or asymptotic giant branch (AGB) stars probe the important stage in the stellar evolution. H 2 O and several SiO maser transitions fall in the 43 GHz observing band of VSOP-2. Recently discovered 43 GHz SiO maser spots are known to be bright and extremely compact. Multi-transitional observations of these SiO masers excited in different inversion schemes will give new information on the local physical conditions of the maser emission, thereby tracing the mass loss process. Radio jets are also found in evolved stars, before or during the the central star forming process in planetary nebulae. Some of them are probed by H 2 O masers and probes for elucidating the jet formation mechanism in young stellar objects (YSOs). More detailed information on these science goals can be found in the next section. 1.4 Overview of VSOP-2 Mission Overview of VSOP-2 Mission The VSOP-2 (VLBI Space Observatory Programme 2) mission having been selected by the Japanese Institute of Space and Astronautical Science (ISAS)/JAXA became a new project at the start of the 2007 fiscal year. One of the reasons for it being selected was the success of the ISAS/JAXA VSOP mission which latest from VSOP observations were conducted at 1.6 and 5 GHz using an 8-m diameter deployable space radio antenna operating in conjunction with a global network of 40 ground-based radio telescopes with the space data being collected by 5 different tracking stations and with all the data be processed at three correlation facilities. To support the ambitious science goals, the next generation space-vlbi spacecraft requires improvements in both sensitivity and angular resolution when compared with the HALCA spacecraft used for the VSOP mission (see Table 1.2). The VSOP-2 spacecraft (ASTRO-G) will have a deployable 9-m off-axis paraboloid antenna with cryogenically cooled receivers operating at 8, 22, and 43 GHz in dual-polarization. With an apogee height of 25,000 km 43 GHz observations will have an angular resolution of 38 micro-arcseconds achievable. Furthermore, a phase-referencing capability is being actively considered which will not only increase the number of observable sources but will also allow state of the art astrometric measurements to be undertaken. Using these impressive technical capabilities, we aim to tackle several fundamental problems in modern astrophysics with ASTRO-G. 7

13 Table 1.2: Technical evolution of VSOP-2 from VSOP Mission VSOP VSOP-2 Antenna Diameter 8m 9m Mission Period Frequency bands 1.6, 5 GHz 8, 22, 43 GHz Highest Angular Resolution 0.4 mas (5 GHz) 0.04 mas (43 GHz) Polarization LCP LCP/RCP (dual) Data Recording Rate 128 Mbps 1 Gbps Phase-referencing No Yes ASTRO-G will be the first space radio telescope operating at millimeter wavelengths using cryogenicallycooled systems, and will achieve the highest angular resolution ever attained in the history of astronomy. The telescope is also capable in observing at 22 GHz, in order to make movies of both molecular outflow in galactic water maser sources and rotating disks in the center of AGN (also traced by water maser emission) that are believed to occur about 10 4 Schwarzschild radii from AGN central engines. At the lowest band of 8 GHz, we will aim to image continuum VLBI sources, such as relativistic AGN jets, microquasars, gravitational lensing objects, and other exotic astrophysical phenomena. To achieve our science goals, the space satellite will incorporate the most advanced forefront technologies. The most important of these technologies to be used on the ASTRO-G spacecraft are the deployable 9-m antenna which has to operate at 43 GHz, space-qualified low-noise amplifiers in cryogenically cooled refrigerators, and finally, the technologies required for determining the satellite orbit down to a few centimeters, which requires using a combination of GPS/Galileo signal receiver, accelerometers, and laser pointing systems International collaboration Space-VLBI projects are scientifically and technically complementary to other ground VLBI facilities, both inside and outside Japan. A large international collaboration which included ground telescope arrays, correlators, tracking stations, and orbit determination was required for the success of the VSOP mission and a similar degree of collaboration will be required if VSOP-2 is to achieve its full potential (see Figure 1.7). Astronomers in Japan, China and South-Korea have recently established a VLBI consortium, aiming to combine East-Asian VLBI facilities to form a regional VLBI Network. The first meetings of this new consortium were held in China and South-Korea in In South Korea, the KVN (Korean VLBI Network) project has commenced. When completed the KVN will consist of three 21 m telescopes which are designed to be operational at all the 3 VSOP- 2 frequency bands. The construction of these telescopes should be completed by the end of Participation of these KVN antennas operating in conjunction with other East Asia telescopes will enable VSOP- 2 observations to be undertaken with a sensitive ground array (Figure 1.8). Also, VLBI activity in Japan is rapidly growing one outcome of which is the ability to undertake a series of Space VLBI missions. The current very healthy state of Japanese VLBI is illustrated the existence of the 10-element Japanese VLBI Network (JVN)( Figure 1.8). Looking beyond East Asia, one notices that VLBI antennas have world-wide geographical distribution including antennas in the United States, Europe, Canada and Australia. These VLBI antennas and their associated VLBI correlators greatly contributed to the success of the VSOP mission. Indeed many of the most sensitive VLBI antennas and their correlators are located in these countries and participation of these facilities will also be critical to the success of the VSOP-2 mission as well. 8

14 Figure 1.7: VLBI network in the world. Image courtesy of C.Walker. Furthermore the international collaborations that were originally developed for VSOP can now be extended for VSOP-2 where it is expected that an even larger collaboration will be developed which will include participation by East Asia, Australia, New Zealand, Europe, US, and the rest of the world. It is also expected that the Square Kilometer Array (SKA) will be developed during the next decade and the SKA should be operational during the half of the VSOP-2 mission. VSOP-2 observations with the SKA would provide the highest sensitivity Space-VLBI observations and will enable new classes of faint objects to be studied on space baselines. 1.5 Brief Current Status of VSOP Approval of VSOP-2 (ASTRO-G) The VSOP-2 mission has been approved as the 25th scientific mission of Institute of Space and Astronautical Science (ISAS) and the Japan Aerospace Exploration Agency (JAXA). The mission received its final formal approval within ISAS at a meeting of the Board of Councillors in April VSOP(HALCA) was an ISAS engineering test mission, but VSOP-2 has been selected as a full ISAS/JAXA science mission. The ASTRO-G spacecraft differs from the HALCA spacecraft in many respects. However, it will require an equally challenging combination of advances in the same diverse areas of space engineering and operations as HALCA. Among the diverse technical challenges in satellite developments required to achieve the science mission goals, the deployment of the 9-m antenna seems to be the most challenging. The deployment scheme of the 9m antenna is prototyped from the ETS-VIII satellite. In December 2006, two reflectors (19m 17m each, operating at 2.3 GHz band) of the ETS-VIII (KIKU-8) were successfully deployed. This success gives high confidence that the ASTRO-G 9-m antenna can be deployed on orbit and will achieve the 0.4 mm RMS surface accuracy required for 43 GHz observations. VSOP-2 received a new start at the beginning of Japanese fiscal year 2007 and funding is currently in place for satellite development from The VSOP-2 satellite will be called ASTRO-G until its launch, which is currently projected to be in January or February With the VSOP-2 mission now approved it time to start international collaboration discussions. 9

15 Figure 1.8: Lef t:the locations of Japanese VLBI Network (JVN) telescopes operating at 8.4 GHz. Right: Overview of East-Asia VLBI Network (EAVN). Growing-up VLBI activities in East Asia will support ground-based VLBI network for VSOP Mission timeline The VSOP-2 mission timeline is briefly presented in Table

16 Table 1.3: VSOP-2 overview Satellite: overview Main institution ISAS/JAXA Main reflector diameter 9 m Antenna design Offset Cassegrain Mass 910 kg Lifetime 5 years Satellite: nominal orbit Apogee/Perigee height 25,000/1,000 km Orbital period 7.5 hours Orbital inclination 31 Precession of AOP( ω) / +258 yr 1 / 167 yr 1 LAN( Ω) Precession of argument of +258 yr 1 perigee ( ω) Precession of longitude of ascending 167 yr 1 node ( Ω) Satellite: observing system Polarization Dual Polarization Reception (LCP/RCP) Phase-referencing Capability Fast-Switching Method Bandwidth per channel 128/256 MHz Observing band 8 GHz 22 GHz 43 GHz Angular resolution µas 75 µas 38 µas SEFD 4080 Jy 2200 Jy Jy 2 Sensitivity 8 GHz 22 GHz 43 GHz 7-σ detection sensitivity 3 23 mjy 50 mjy 107 mjy... to large telescope 4 5 mjy 12 mjy 22 mjy Image noise level mjy/beam 0.45 mjy/beam 0.71 mjy/beam 1. The angular resolutions given here are the beam FWHM for a source at the orbit normal direction. 2. Goal values 3. To a single VLBA (25 m) antenna, assuming a coherence time of 6, 2, and 1 minutes at 8, 22 and 43 GHz. 4. To the phased-vla (27 25 m) 5. For an observation lasting one orbit (7.5/,hr) with the full VLBA. Table 1.4: VSOP-2 Mission Timeline 2005 Formal approval of the project proposal 2006 Budget request Prototype Model Development Flight Model Development 2012 Launch Mission end (after >5 years operation) 11

17 Chapter 2 VSOP-2 Science goals 2.1 Introduction VLBI (Very Long Baseline Interferometry) enables the highest spatial resolution images in all of astronomy. This resolution is determined by the ratio of the observing wavelength to the baseline length. By extending the baseline into the space, the resolution can be further improved. Space-VLBI uses a radio telescope in space in conjunction with telescopes on the ground to realize interferometer baselines longer than the diameter of the Earth. The Japanese VLBI Space Observatory Programme (VSOP), in operation from 1997 to 2005, was the first dedicated space VLBI mission achieving 0.35 milliarcsecond (mas) resolution at 5 GHz. This mission revealed the fine structures of jets emanating from active galactic nuclei (AGNs), AGN cores with extremely high brightness temperatures, and, in several AGN, the plasma torus surrounding the nucleus. Following the successes of the VSOP, VSOP-2 has now been selected as the next ISAS/JAXA science mission. This next generation space-vlbi mission will achieve a resolution 10 times higher than VSOP, directly observe AGN accretion disks, determine AGN jet formation and acceleration mechanisms, and elucidate water maser properties in both distant galaxies and in Galactic star forming regions. In this Chapter, first we review scientific results from ground-vlbi and VSOP observations, and then introduce the science goals for the VSOP-2 mission Active Galactic Nuclei (AGNs), the most powerful engines in the Universe AGNs are the most powerful engines in the universe, emitting W ( L ). Their power source is believed to be the gravitational energy of matter accreting onto a black hole. A massive black hole at the heart of galaxies creates a deep gravitational potential well. When gas in the galactic disk falls into this well, it forms an accretion disk, gets heated, and radiates at very high temperatures. Energy from accretion is converted to radiation, advection into the black hole, and jet kinetic energy. The jets are a bipolar flow of relativistic plasma emanating from the central engine and can extend as far as the Mpc scale. This flow ultimately terminates at hot spots and spreads out into radio lobes. This model has become the standard picture of AGNs. Significant progress on understanding the nature of AGNs has been brought by high resolution observations over the last decade. Hubble Space Telescope (HST) observations of the radio galaxy M87 revealed a rotating gas disk with an enclosed mass of M, implying the existence of a super massive black hole (Harms et al. 1994). VLBA observations of H 2 O maser emission in the Seyfert galaxy NGC 4258 disclosed a Keplerian disk rotating at speeds of up to 1000 km s 1 around the center, indicating the presence of a black hole weighing M (Miyoshi et al. 1994; Herrnstein et al. 1999). The discovery of extremely broad iron line emission from the active galaxy MCG with the X-ray satellite ASCA (Tanaka et al. 1995) has been interpreted as gravitational redshift of 12

18 Figure 2.1: (Left) Illustration of the standard model for AGN (Urry & Padovani 1995) showing the accretion disk, jet, and and obscuring torus. (Right) HST image of of the nearby AGN M87 (Harms et al. 1994) the spectrum at only several Schwarzschild radii from the central black hole. Recently, studies of the orbit motions of stars around Sgr A, the central object of our galaxy, have shown it contains a M black hole (Ghez et al. 2003) The inferred existence of a massive black hole supports the standard picture of AGN. However, to date, no direct image of either a black hole or an accretion disk have been made. Accretion disks are thought to be times larger than the Schwarzschild radius of the black hole, and their outer radius corresponds to µas in nearby AGNs. They are too small to be imaged, even with the best telescopes currently available. VSOP-2 will be the first telescope to resolve the structure of accretion disks. This will provide a deeper understanding of AGN structure and will allow comparisons with simulations to test the validity of theoretical models. AGN jets are another important target of VSOP-2 observations. It is still unclear how the jet plasma flow is initiated, collimated, and accelerated, although several well-developed physical models have been proposed. It is important to investigate, with the appropriate resolution, the relationship between accretion disks and jets in sources where these outflows are produced and accelerated. Because magnetic fields are expected to play an important role in these mechanisms, polarization observations are crucial. 2.2 Astrophysical Jets - Internal structure of relativistic flows Scientific results from ground VLBI and VSOP observations Almost 800 observations were undertaken by the VSOP mission with spatial resolutions 3 times higher than those with ground VLBI arrays at the same frequency (Hirabayashi et al. 1998). Here we briefly review the major scientific results from the VSOP mission to illustrate the importance of space-vlbi. Jet Motion Statistics AGN jets are relativistic plasma flows that emit synchrotron radiation at radio wavelengths. Variabilitytimescale estimates of brightness temperature T b of the radiation indicate that in some cases this temperature should exceed the inverse Compton limit of K but this prediction cannot be directly tested using only ground-vlbi with its limited baseline lengths but requires greater than Earthdiameter space-vlbi baselines. 13

19 The first results from the VSOP survey showed that 53 objects from a sample of 98 had source frame brightness temperatures of T b > K (Scott et al. 2004, Figure 2.3). This indicates that Doppler boosting is important in these sources. A Doppler factor of δ = 1 β 2 /(1 β cos θ) will boost the intrinsic brightness temperature by δ/(1 + z), where β is the jet speed relative to the speed of light, θ is the viewing angle of the jet, and z is the redshift (Piner et al. 2000). Consequently a high relativistic speed are required to amplify the observed brightness temperatures to > K. Lister et al. (2001) performed VSOP observations of 27 sources selected from the Pearson Readhead survey. The measured apparent jet speeds, β app, are related to the intrinsic speed in units of the speed of light, β, and the viewing angle, θ by β app = β sin θ/(1 β cos θ). Statistics of the Lorentz factor Γ = 1/ 1 β 2 can be derived from the measured β app by Γ = (β app + δ 2 + 1)/2δ. The distribution of the Lorentz factors indicates that the largest observed values are Γ 30. This high value is a challenge for some current theoretical jet acceleration models Frequency log10[tb(k)] Figure 2.2: Results from VSOP Survey Program observations. The source frame core brightness temperatures for 98 Survey sources with identified cores and known redshifts. The open green histogram has been used for those sources with resolved cores and the open red portion of the histogram represents brightness temperature lower limits for unresolved cores (Scott et al. 2004). Morphology and motion of jets VSOP observations provide a resolution of 0.35 mas at 5 GHz and have revealed the extent to which pc-scale jets can appear to bend (in projection). Figure 2.3 shows VSOP images of the quasar with a distinct kink in the jet being observed to propagate ballistically from the core at constant speed (Murphy et al. 2003). Such bends and kinks are also seen in VSOP observations of NGC 6251 (Sudou et al. 2000), M87 (Biretta et al. 2002), and Mrk 501 (Edwards et al. 2000). Inner structures of jets and Kelvin Helmholtz instabilities Why do some AGN jets show multiple bends or wiggles in their structure? The three major possibilities are thought to be the Kelvin-Helmholtz instability (caused by an interaction between the jet plasma 14

20 Figure 2.3: Seven epoch VSOP observations of the low-redshift quasar made over a 4.04 year period showing the evolution of the parsec-scale jet. The vertical extent of the image is 13 mas, with the spacing between images corresponding to the time interval between epochs (Murphy et al. 2003). flow and ambient matter), precession of the jet nozzle, and the current-driven kink instability (a relative of the magnetic pinch). Lobanov & Zensus (2001) observed the quasar 3C 273 with VSOP and resolved the inner structures of the jet. They discover that twin ridges of brightness peaks across the jet weave into a double helix along the jet out to 300 pc from the nucleus. They claim that the double helix can be explained produced by superposition of Kelvin Helmholtz modes up to 5th order in the plasma flow whose Lorentz factor and Mach number are Γ = 2.1 ± 0.4 and M = 3.5 ± 1.4 (see figure 2.4). Ballistic ejection of precessing jets In contrast, the quasar 3C 380 shows evidence for a precessing jet nozzle. This quasar at z = has relativistic jets with apparent superluminal speeds of β app = 7 to 13. The motion has been monitored with ground VLBI (Polatidis & Wilkinson 1998). Kameno et al. (2000) observed this quasar with the VSOP. It turned out that every knot component shows a ballistic motion, radially from the core, at a constant speed. The apparent speed and ejection direction of components change as a sinusoidal function of ejection epoch, whose amplitude and period is 5.1 and 47 yr, respectively. The ejection direction is thought to reflect the plane of the accretion disk. Precession of the disk plane can be caused by a disk instability, orbital motion of the black hole, or torques from a companion object. Higher resolution is essential to clarify the cause of the precession. Magnetic structures of jets If jets have a significant toroidal magnetic field (which is expected in the case of magnetic jet flow acceleration), then they may be subject to the m = 1 current-driven (or kink) instability. This instability is closely related to the well-known m = 0 pinch current-driven mode. (Nakamura & Meier 2004) studied the kink instability in AGN jets in detail and found that it causes a fundamental physical change in the jet flow itself, causing it to travel in a corkscrew pattern. The jet s strong magnetic field backbone keeps the flow from becoming disrupted and, in fact, can suppress Kelvin- Helmholtz instabilities. This new method of producing jet wiggles has not yet been applied to any observed sources. However, time-dependent, high-resolution, polarization observations of AGN jets by VSOP-2 will be able to distinguish sources whose wiggles are caused by current-driven kinks from those caused by Kelvin-Helmholtz instabilities. 15

21 Figure 2.4: VSOP 5 GHz image of the parsec-scale radio jet in 3C273. The dot-dashed white lines denote the locations of the four flux density profiles shown in the inset. A total of 240 such profiles have been measured along the jet. Each of these profiles is centered on the smoothed ridge line (the dashed black line) and oriented orthogonally to it. Each of the measured profiles is fitted by two Gaussian components designated as P1 and P2. The locations of the peaks of P1 and P2 are marked in the image by the red and blue lines. The double helical pattern formed by P1 and P2 suggests that they result from K-H instability developing in the jet (Lobanov & Zensus 2001). Sub-parsec scale acceleration High resolution VSOP and VLBA observations have enabled the precise study of jet motions. It has been found that jet speeds on the sub-pc scale are slower than those on pc-scale in nearby AGN such as Cygnus A (Krichbaum et al. 1998), 3C 84 (Dhawan et al. 2000; Asada et al. 2006), M87 (Junor et al. 1999; Biretta et al. 2002), NGC 1052 (Vermeulen et al. 2003). In the case of 3C 84, the jet travels at β app = on sub-pc scales, but at β app on the pc-scale. Jet motion in M87 on sub-pc scales has not been detected (although recent reports suggest that 43 GHz observations may have yielded evidence for motion close to the core), indicating it s not faster than β app = 0.01, while HST observations showed β app = at 1 kpc (Biretta et al. 1999). There are no sub-pc scale superluminal motions found to date. This fact may imply acceleration in sub-pc scale, or that the apparent motion is a pattern speed rather than the bulk speed Accreting matter onto AGN VSOP and ground VLBI observations have provided some clues as to how matter is distributed in the vicinity of AGN. Dense plasma tori were found on the pc-scale in several AGN (Jones et al. 2001; Walker et al. 2000; Kameno et al. 2001; Vermeulen et al. 2003; Kameno et al. 2003). VSOP and VLBA observations of the radio galaxy NGC 4261 (Jones et al. 2001) showed that a plasma disk with a radius of 0.2 pc casts a shadow, produced by free free absorption, near the base of the counter-jet. 16

22 In the active galaxy NGC 1052, multi-frequency VLBI observations were performed to determine the spatial distribution of the free free absorber (Kameno et al. 2001; Vermeulen et al. 2003; Kameno et al. 2003). These constrained the electron density ( cm 3 < n e < cm 3 ) and the size ( 1 pc) of the plasma torus. These plasma disks (or tori) may be the fuel reservoir that supplies accreting matter to the central engine of AGN Radio Lobe Expansion The jets in AGN interact with the ambient medium to form radio lobes, which are often found on the kpc to Mpc scale. A series of VSOP observations detected expansion of the inner radio lobe in 3C 84 (Asada et al. 2006). The angular distance from the nucleus to the end of the southern lobe increased from 18.4 mas (7.4 pc) in 1998 to 19.5 mas (7.8 pc) in The expansion corresponds to a subluminal velocity of β adv 0.5. During its expansion the lobe dimmed, from a flux density of 17.4 Jy to 13.4 Jy. The rates of expansion and dimming could be explained by adiabatic expansion from the results of X-ray observations of radio lobes (Tashiro et al. 1998). It may be that VSOP-2 observations at higher frequencies will be able to trace the lobe expansion more accurately in optically thin plasma above the synchrotron peak frequency AGN jets and VSOP-2 Introduction As described in the previous sections, relativistic flows are the most prominent factors contributing to the overall feedback from AGN into intergalactic and intracluster medium. Understanding the efficiency and magnitude of this feedback is crucial for building a detailed account of the galaxy evolution and large-structure formation in the Universe. This understanding can only be built on an extended knowledge about physical mechanisms governing generation and propagation of relativistic flows and their interaction with the ambient medium surrounding the jets. VSOP-2 observations will provide unparalleled observational capabilities for addressing this problem. VSOP-2 will image, with high- fidelity, the total and polarized emission on angular scales sufficient for resolving transversely a substantial fraction of extragalactic jets and analyzing their interior. This will make possible addressing, in unprecedented detail, the dynamics and evolution of extragalactic relativistic flows. Direct imaging of extragalactic jets with VSOP-2 enables exploration of the following three major topics. Relativistic jet formation Morphological studies will provide information about the flow stability and enable estimates of the fundamental flow parameters (bulk Lorentz factor, Mach number, jet/ambient medium density ratio) to be made. These estimates will be essential for studying the physics and propagation of relativistic jets with numerical relativistic magneto-hydrodynamic (RMHD) simulations and advanced analytical models. High-fidelity imaging will be used for monitoring the evolution of the jet ridge line and obtaining direct assessments of the pattern speeds in jets. In the most prominent jets, it would also be used for determining a two-dimensional velocity field of the flow and answering the questions about the flow stratification and the intensity of the flow/ambient medium interaction. Jet magnetic fields Full polarimetry of transversely resolved flows, using imaging of the linearly polarized emission and mapping the rotation measure along and across the jets, will be employed to determine the threedimensional structure of the magnetic field. This will provide a unique perspective on the role played 17

23 by the magnetic field in the collimation, large- scale stability, and dynamics of relativistic flows. The linear polarization information will be also used for tracing the evolution of relativistic shocks in extragalactic outflows. A linear polarization image reflects the projected magnetic fields on the plane of sky, in the case of optically thin synchrotron emission. The component along the line of sight can be inferred from mapping the distribution of the Faraday rotation measure (FRM) in the jet. It has been shown (Asada et al. 2002) that combining linear polarization and FRM imaging yields an information about the magnetic field configuration in relativistic jets. VSOP-2 will employ this method to make such measurements in a large number of objects and uncover the fundamental physical properties of the magnetic field in relativistic flows. Jet/ambient interaction Combining VSOP-2 observations with ground VLBI observations at matched angular resolution will enable detailed spectral index maps and obtain the distribution of the synchrotron turnover frequency in relativistic flows. The spectral information will be essential for understanding the process of shock dissipation and development of Kelvin-Helmholtz instability in extragalactic jets. The turnover (or peak) frequency of the synchrotron emission is sensitive to changes of physical conditions in the jet such as velocity, particle density, and magnetic field strength. It can be used as an excellent tool for probing the physics of the jet in more detail than is allowed by analysis of the flux and spectral index properties of the jet (Lobanov 1998, Dodson et al. 2006). Jet astrophysics Space VLBI observations of relativistic jets with the VSOP-2 will offer two major advantages compared to ground VLBI. (1) VSOP-2 will have much better sensitivity for studying extended regions of the jets with optically thin synchrotron spectrum. For such regions, ground VLBI observations, albeit providing a matching resolution, will be highly inefficient in reaching similar sensitivity to optically thin emission. The VSOP-2 observations at 8, 22, and 43 GHz will be, 40, 60, and 500 times, respectively, more efficient for recovering the internal structure of relativistic flows than ground VLBI observations at frequencies providing roughly the same resolution (22, 43, and 86 GHz, respectively). (2) VSOP-2 will be a better tool for studies of the magnetic field structure. First, a higher resolution reduces beam depolarization resulting from magnetic fields tangled on smaller scales than the resolution. Second, the higher observing frequency is more sensitive to the innermost jets, where radiation is optically thick at lower frequencies, in which case the polarization tends to vanish. Combined together, all these factors will make VSOP-2 a superb astrophysical tool for studying the relativistic flows and their connection to the general physical processes in AGN D magnetic structure A linear polarization image reflects the projected magnetic fields on the plane of sky, in the case of optically thin synchrotron emission. The component along the line of sight can be inferred from the Faraday rotation measure (FRM). Faraday rotation is an effect of the propagation of radiation through magnetized plasma in which the plane of polarization is rotated. The polarization angle, χ, is rotated from the intrinsic polarization angle, χ 0, to χ = χ 0 + FRMλ 2, (2.1) 18

24 Figure 2.5: Distribution of Faraday rotation measure (FRM) in the jet of the quasar 3C 273 (Asada et al. 2002). Contours and color indicate total intensity and FRM, respectively. Systematic gradient of FRM across the jet can be explained by a helical magnetic field winding around the jet. where λ is the observing wavelength. Thus, multi-frequency polarization observations are necessary to determine the intrinsic polarization angle and FRM. The FRM depends on the magnetic field along the line of sight, B [G], and the electron density, n e [cm 3 ] as FRM = n eb d l[rad m 2 ]. (2.2) LOS Hence, the combination of linear polarization and FRM images provides clues to the 3-D magnetic structure. This method has been established by using ground-based VLBI. VLBA observations of the quasar 3C 273 revealed helical magnetic fields in the pc-scale jet (Asada et al. 2002). VSOP-2 offers two advantages for studies of magnetic field structure. First, a higher resolution reduces beam depolarization resulting from magnetic fields tangled on smaller scales than the resolution. Second, the higher observing frequency is more sensitive to the innermost jets, where radiation is optically thick at lower frequencies, in which case the polarization tends to vanish. The higher frequency is also important to determine the FRM. Because of 180 ambiguities of the polarization angle, a suitable range of observing wavelengths is required to cover the wide possible FRM range. 2.3 Black holes the heart of the beast The size of a black hole is characterized by the Schwarzschild radius, r s, which is given by r s = 2GM BH c 2, (2.3) 19

25 where c is the speed of light, M BH is the black hole mass, and G is the gravitational constant. For a massive AGN with M BH = 10 9 M, for instance, the Schwarzschild radius is 20 AU and its apparent size will be 2 µas if the distance is 10 Mpc. Thus, µas resolution is required to image black holes directly. Although it will be challenging to image black holes (or, rather, detect their shadows) with VSOP-2, observations of the closest and most massive AGN will be a high priority. In the case of stellar-mass black holes in our galaxy, e.g., 10 M at a distance of 1 kpc, the apparent size will be only nanoarcsec and so direct imaging will not be feasible Imaging of black hole environments AGNs shine with a luminosity of L = L. This power is believed to be generated by the release of gravitational energy through mass accretion toward the massive black hole at the center. Imaging of a black hole and accreting matter is essential to understanding the details of power generation in AGN. Although the black hole itself doesn t radiate (except for Hawking radiaton which can be ignored for AGN-size black holes), it should be possible to image the shadow it casts on the shining accretion disk. If it is a Schwarzschild (non-rotating) black hole, the shadow radius will be the radius of the marginally stable orbit, i.e., 3 r s. Although the central object in the galactic center, Sgr A, is expected to casts the largest shadow of known AGN, 45 µas, the image will be scattered by interstellar plasma to a size of 7 mas at 43 GHz. To see the black hole shadow, observations at frequencies higher than 500 GHz are required to avoid interstellar scattering. Thus, VSOP-2 observations will not be able to resolve the shadow of the black hole in Sgr A. M87 is the best alternative candidate for approaching to direct black hole imaging. Its apparent shadow diameter is expected to be 20 µas, which is comparable to the VSOP-2 resolution, as shown in figure 2.6 (Takahashi 2004; Takahashi et al. 2007). It is still unclear how bright the background is and how much plasma will scatter, or defocus, the image, nevertheless, VSOP-2 is the best telescope to approach the direct image of a black hole. Figure 2.6: The supermassive black hole in M87 for a variety of inclinations and spin parameters (a/m) shown with unlimited resolution and with the angular resolution expected for VSOP-2 (Takahashi et al. 2007) 2.4 Accretion disks the power plants in AGN Accretion disks are considered to be the power plant where the gravitational energy of the accreting matter is converted to radiation and where jets are produced and accelerated. No telescope has ever 20

26 succeeded in imaging an accretion disk directly. Imaging studies of accretion disks in AGNs could be the promising subject for Space-VLBI. The priority of VSOP-2 observations is to test the AGN paradigm a massive black hole and accretion disk system by imaging the most inner part of an AGN or an accretion disk. Imaging studies also allow us to discriminate disk models by means of the size of the accretion disks, distributions of temperature and surface density, and the pattern of the magnetic fields. Since the size of an accretion disk is proportional to the Schwarzschild radius, nearby AGNs are the best candidates to resolve, as shown in section Accretion disk models There are three major models for accretion disks; the standard model (Shakura & Sunyaev 1973), the ADAF (Advection-Dominated Accretion Flow) (Narayan & Yi 1994), and the Slim disk model (Abramowicz et al. 1988). The physical properties of the equilibrium conditions in these models are summarized in Abramowicz et al. (1995). Standard disk Standard disks are characterized by a balance between radiative cooling and viscous heating. The effective temperature, T eff is given by T eff [ 3GM BH Ṁ 8πσ(r/r s ) 3 ] 1/ η 1/4 ( L L edd ) 1/2 ( ) L 1/ r 3/4, (2.4) W where Ṁ is the mass accretion rate, σ is the Stefan Boltzmann constant, and L edd = ( M M ) W is the Eddington luminosity. For a typical quasar near the Eddington limit, with M BH = 10 8 M and L W and efficiency η 10%, the effective temperature will be T eff 10 5 K at a radius of r 100r s. The big blue bump, commonly seen in the ultraviolet spectrum of quasars, is thought to be produced by multi-temperature black body radiation from the standard disk. Such black body radiation is too faint to be detected by VSOP-2 whose brightness temperature sensitivity is K. Nevertheless, there remains a chance to detect standard disks with the VSOP-2. The EUV to hard X-ray emission from quasars cannot be produced by the black body radiation of the standard disks. Haardt & Maraschi (1991, 1993) proposed the disk corona model, which proposes that a diffuse hot corona with kt kev (T K) sandwiches the cold standard disk to produce the hard X-ray radiation via inverse Compton scattering of UV photons from the disk. Such a hot corona can emit at high brightness temperatures detectable with the VSOP-2 at radio wavelengths. ADAF and RIAF Advection-Dominated Accretion Flows are radiatively inefficient, so that most of the heat produced through accretion is preserved as the flow is compressed toward the black hole. Since the accreting gas is not cooled efficiently, it is heated to the local virial temperature of T ion m p c 2 /k(r/r s ) = (r/r s ) 1 K for ions and T e m e c 2 /k = K for electrons. The condition is far from equilibrium, because the gas is collisionless (Begelman & Chiueh 1988). Similarity of the sound speed with the free-fall and rotation speeds in this hot gas results in a geometrically thick (H r) torus shape. Hence, the gas density must be optically thin. This is self-consistent with radiative inefficiency. The primary radiation from ADAF disks is thought to be thermal synchrotron or bremsstrahlung, with Compton enhancement. The spectral peak frequency, ν m, of synchrotron radiation is determined by the brightness temperature T b and the magnetic field B as 21

27 ν m = Hz ( Tb ) 2 ( B ) 10 9 K 10 4 G (Pacholzcyk 1970). If we assume that the magnetic pressure is half of the gas pressure, then B (r/r s ) 1/2 G. And, with T b T e, we find ν m Hz at r = r s, with the spectral peak in the millimeter band. The primary emission falls off at higher radio frequencies as a power law or faster. However, the inverse Compton process generates a second peak at Hz in the optical and a third in the hard X-ray and γ-ray wavebands. The ADAF model predicts that the mass accretion rate is a factor of 10 4 of the Eddington limit and that the electron temperature is above 10 9 K for r < 100r s (Manmoto et al. 1997). Thus, ADAF disks can be detected and imaged with VSOP-2. Most low luminosity AGNs (LLAGNs) show spectral energy distributions (SEDs) which are consistent with the ADAF model. The standard disk model cannot explain their low luminosity, hard spectrum above UV wavelengths, and the absence of a big blue bump. LLAGNs are common: about 40% of nearby galaxies host an LLAGN (Ho et al. 1997). Thus an ADAF disk may exist in the majority of galaxies. The RIAF (Radiatively Inefficient Accretion Flow) model has been proposed as a variation of the ADAF models to explain the SED of Sgr A which is relatively dim in the optical and X-ray bands (Yuan et al. 2003). This model predicts a higher electron temperature (T e K) than for an ADAF because of radiative inefficiency. A higher fractional polarization can be also expected. Thus the existence of RIAFs is even more encouraging for VSOP-2 observations. Slim disk The slim disk model, with a extremely high mass accretion rate above the Eddington limit, was proposed by Abramowicz et al. (1988). The disk is hotter than the standard disk because of the high accretion rate. It is geometrically thick and supported by radiation pressure. It is also optically thick, radiatively efficient and luminous. Heat production via mass accretion is balanced with radiative and advective cooling. Narrow-line Seyfert 1 galaxies are thought to host a slim disk. Their Balmer line widths from the BLR (Broad Line Region) of FWHM < 2000 km s 1 are narrower than those of usual Seyfert 1s, and their black hole masses are in the range M. Their luminosities are close to the Eddington luminosity. Their spectra extend to soft X-ray energies, indicating thermal emission of kt 100 ev (T 10 6 K), 100 times hotter than a standard disk. These properties are consistent with those predicted from the slim disk model. Collin & Kawaguchi (2004) suggest that narrow-line Seyfert 1 galaxies host a growing black hole with an age estimated to be 10 7 yr from their low mass and high accretion rate. These objects are important for understanding the evolution of massive black holes. The brightness of the blackbody radiation from slim disks is 10 6 K, which is below the detection limit of VSOP-2. However, a hot thin corona above the disk layer can be a bright source of radiation, as in the case of standard disks, by means of inverse Compton scattering of UV photons from the disk. Narrow-line Seyfert 1 galaxies show a power-law spectrum at hard X-ray energies, which is probably produced in the hot corona layer. The temperature of the corona is estimated to be kt kev (T K), so that VSOP-2 will be able to detect synchrotron or bremsstrahlung radiation from it Approaching to accretion disks with VSOP-2 Each of the three above major disk models offers the possibility of imaging the accretion disk with VSOP-2. The ADAF and RIAF models are the most promising models for detection, as they predict 22

28 the highest brightness temperatures. Here we discuss the detectability in terms of the disk size and the flux density. The resolution of the VSOP-2 is comparable to the apparent sizes of accretion disks in nearby AGNs. The synthesized beam of 38 µas at 43 GHz corresponds to 13 r s for M87, so that the image of the accretion disk can be resolved. Table 2.1 lists AGN within 20 Mpc with accretion disks that can be resolved by the VSOP-2. There are at least three objects for which the VSOP-2 beam size is less than 20 r s, and at least 13 objects with for the beam size is less than 200 r s, within 20 Mpc of the earth. Table 2.1: AGN disk candidates Name Other Name D M BH θ g S 15GHz Remarks [Mpc] [10 8 M ] [µas] [mjy] NGC 3031 M NGC NGC S at 5 GHz NGC NGC S at 8.4 GHz NGC NGC 4374 M NGC 4486 M NGC NGC 4594 M S at 8.4 GHz NGC 5128 Cen A S at 8.4 GHz IC 1459 PKS S at 8.4 GHz Sgr A S at 8.4 GHz LLAGNs are good targets for the imaging of accretion disks. They are thought to have an ADAF disk and are less affected by radio emission from powerful jets. Most LLAGNs are weak radio sources with flux densities of mjy in the nearby universe (within 20 Mpc) (Nagar et al. 2000), nevertheless, there are more than 50 sources that are detectable with phase-referencing VSOP-2 observations. Ground-based VLBI observations of nearby LLAGNs have shown that they have no jet, or at most weak jets (Ulvestad & Ho 2001), so that accretion disks probably can be imaged in the absence of jets. The broad-band SEDs of LLAGNs can be modeled by ADAF or RIAF models (Quataert & Narayan 1999; Yuan et al. 2002), with disk temperatures of K. Thus, they are good candidates for the imaging of accretion disks with VSOP Mass accretion process onto AGNs The sub-pc region of AGNs, corresponding to Schwarzschild radii, is a mass reservoir where accreting matter is accumulated. The total mass and velocity fields in this region are tightly related to the mass accretion rate which is an important indicator of AGN activity. The physical phase of matter in the sub-pc region is complex. Some type-2 Seyfert galaxies, LINERs (Low-Ionization Narrow Emission line Regions), and radio galaxies show H 2 O or OH maser emission from their sub-pc regions (Moran et al. 1999). This indicates that dense (n H > 10 9 cm 3 ) warm (T > 400 K) molecular gas forms a rotating disk in this region. HI absorption is detected in some galaxies. VLBI observations of HI absorption have revealed rotating neutral gas disk with the radius of 10 pc and a rotation speed of several 100 km s 1 (Conway & Blanco 1995; Mundell et al. 1995; Taylor 1996; Mundell et al. 2003). Furthermore, a plasma torus or disk is found on pc scales 23

29 via free free absorption (Walker et al. 2000; Jones et al. 2001; Kameno et al. 2001, 2003). Besides radio observations, X-ray spectroscopy has revealed dense cool matter, named warm absorber, in the vicinity of the central engine via soft X-ray absorption and OVII and OVIII edges (Krolik & Kriss 2001). Tanaka et al. (1995) detected a broad iron fluorescent line in the Seyfert galaxy MCG The line at 6.4 kev in the rest frame is spread and redshifted, indicating that cool matter does exist within 10 r s of the black hole where the gravitational redshift is significant. These facts suggest that a mixture of gas in multiple phases, such as molecular, neutral, and plasma, exists in the sub-pc region (see figure 2.7). Figure 2.7: AGN Gas phases as functions of temperature and density. How does the gas remain in multiple phases in the sub-pc region? Wada (2001) made a threedimensional fluid simulation in the central 100 pc region of an active galaxy and found that supernovae in cold (T < 100 K) gas produce thin hot (T e > 10 5 K) plasma spheres which emit ultraviolet radiation resulting in the generation of warm (T = K) regions. Their simulations, with a resolution of 0.5 pc, showed a mottled picture of the multiple phase gas. VSOP-2 offers the potential to illustrate the mixture of multiple-phase gas in these regions with 0.01 pc resolution observations Plasma distribution Ionized gas can be probed via free free absorption at radio wavelengths. Its optical depth, τ ff, is an indicator of cold dense plasma, as τ ff n 2 ete 1.5 dl, (2.5) LOS where LOS dl is the integration along the line of sight. Kameno et al. (2000) established a method to measure the spatial distribution of free free opacity from multi-frequency VLBI images. Plasma distributions with resolutions of 0.1 pc were obtained in some AGNs using VLBA and VSOP observations in this way. Denser plasma at the heart of AGNs absorbs radio emission at higher frequencies, so that higher resolution at higher frequencies is essential. VSOP-2 is the best telescope for such observations as it works at millimeter wavelengths and offers a resolution better than 100 µas, corresponding to 24

30 0.01 pc in nearby AGNs. The detailed distributions of plasma will be a key to solve the question of the ionizing source: is it supernovae in the molecular disk, photoionization caused by secondary UV radiation, or direct X-ray radiation from the central engine? Figure 2.8 displays ground-vlbi map of the active nuclear region of the nearby elliptical galaxy NGC 1052 (Kameno et al. 2001). The radio map of the galaxy shows a core-jet structure projected nearly on the plane of the sky. The galaxy shows the radio free-free absorption due to the edge-on torus-like structure probed by plasma, which might trace an outer part of an accretion disc. The galaxy also contains strong water maser emission in the nuclear region. Figure 2.8: The distribution of the radio free-free absorption opacity (upper) in the nuclear region of the elliptical galaxy NGC 1052 and the schematic view of the plasma torus and a twin-jet system emerging from the central engine of AGN (lower). Image courtesy of S.Kameno. 2.6 Extragalactic H 2 O maser Introduction The extragalactic H 2 O masers with high apparent luminosity (L H2 O > 100 solar luminosity) are apt to be located in or nearby an active nucleus, or to be associated with jet-activity in active galactic nucleus (AGN). These H 2 O masers which associate with the AGN-activity is called a megamaser since the apparent luminosity is more than a million times higher than that of masers found in the Milky Way. This large luminosity accounts for the excitation of megamasers relating to the nuclear activity within the central few parsecs of a central engine. When the maser in a molecular disc is observed close to edge-on in line of sight with well-ordered velocity field, the radiation of the maser is beamed and enhanced in line of sight. This results in the strong maser emission at 22 GHz. VLBI observations of the megamasers revealed the presence of a Keplerian masering disc surrounding super massive black hole and other structures of inner parsecs of AGN, from which some important physical parameters have been derived. VLBI observations of 22GHz H 2 O maser emission in the active galaxy NGC 4258 revealed the 25

31 Figure 2.9: Spectral profiles of well-known H 2 0 megamaser in the nearby active galaxy NGC 3079, monitored by the single-dish telescope at Effelsberg in Germany. Intensity of the maser shows the strong variability of each spectral component. (The image courtesy of W.Sherwood.) presence of a sub-parsec-scale H 2 O maser disc with a Keplerian rotation around a central massive object (Figure 2.10; Miyoshi et al. 1995). After the discovery of the masing torus surrounding a black hole in NGC 4258, very few other such sources suitable for studying the central sub-parsecs of AGN have been discovered, but astronomers need more candidates to probe the full range of the kinematics and dynamical structures of sub-parsec-scale circumnuclear regions in AGN. The megamaser distributes from parsecs from the center of AGN, so that one needs millarcsecond angular resolution of VLBI. Intensity variabilities on time scales of weeks to months are commonly observed in these megamasers, which has posed difficulties in determining the optimum timing of interferometric experiments (Figure 2.9). Figure 2.11 presents the schematic view of the central region of AGN, in which the obscuring media consisting of dusty neutral materials forms a circumnuclear disk/torus surrounding a central engine. When observers see this circumnuclear disk /torus from (nearly-)edge-on view, the torus/disk blocks our direct view from a central engine. When we see this system from face-on view, the central engine is visible from observers. It has been considered that nuclear H 2 O maser arises from the cold and neutral circumnuclear media, which may be an extention of the warmer inner disk. VLBI observations have thus revealed the inner few parsecs of the circumnuclear disk/torus by measuring the distribution of the H 2 O maser spots in active galaxies. Previous VLBI observations show that most AGN H 2 O maser spots are located at distances of pc from the central engine, which corresponds to Schwarzschild radii. Within 0.1 pc from the central engine, molecules diassociate and consequently there are no H 2 O masers closer than this to the AGN core. Currently, about 80 extragalactic nuclear H 2 O masers, excluding the low luminosity masers from extragalactic star-forming sites, are known (Figure 2.12). However, only about ten of them have been so far imaged at milliarcsecond resolution using VLBI. With 80 micro arcseconds angular resolution of VSOP-2 at 22 GHz, we will investiagate further H 2 O megamasers to derive fundamental physical parameters around or in accretion disks and super massive black holes. 26

32 Figure 2.10: The VLBI observations of H 2 O maser emission in the edge-on active galaxy NGC The geometry of the maser distribution and velocity structures of the maser were first revealed by VLBI observations. The image courtesy of L.Greenhill VSOP-2 megamaser observations Direct imaging of the innermost regions of active galactic nuclei with H 2 O megamasers at 80 microarcsec angular resolution will enable to address the fundamental astrophysical phenomena in the followings. Probing accretion discs The successful VLBI studies of NGC 4258 clearly proved the existence of a sub-pc-scale molecular disk obscuring the very center of the galaxy. The H 2 O maser emission that is symmetrically offset from the velocity of the galaxy ± 1000 km/s arise along each part of the rotating disk seen tangentially. It is important to identify other similar systems as many as possible towards AGN, and to find out how typical the sub-pc-scale maser disc in NGC 4258 really is. Kinematics of inner sub-parsecs of active galactic nuclei We need more investigations of megamasers in AGN, in order to study the kinematics of the innermost region of AGN. Currently, the measurement of the megamasers coupled with VLBI is an only method to probe outer part of accretion disc. From direct imaging of distributions and velocity structures of megamasers, we will be able to look at inner few sub-parsecs by the highest angular resolution attained by VSOP-2. 27

33 Figure 2.11: Schematic view of the central region of an active galactic nucleus. For the detection of maser emission, observers are to be aligned to the plane of the masering disk or torus in their line of sight. Distance measurement Distances to galaxies can be measured towards rotating masering disk systems, with a technique established in the study of NGC 4258 (Herrnstein et al. 1999). In the rotating disk model where the inner disk radius and velocity drift of the maser are measured, the distance to a galaxy hosting the maser can be estimated by the detection of angular proper motion of the maser. Ground-VLBI enabled proper motions of the maser spots in NGC4258 at 30 microarcsec per year (Herrnstein et al. 1999) to be measured, and showed possibility that this method can be applied to more distant objects with space-vlbi. Among about 80 nuclear H 2 O maser sources that are presently known, there are several sources whose proper motions are measurable at < 100 Mpc. It is also expected that on-going extragalactic H 2 O maser survey will double the number of distant nuclear maser up to about 200 Mpc in the next five years. Thus, using VSOP-2 combined with sensitive ground radio telescopes at higher angular resolution, we will be able to calibrate Hubble constant at Mpc at unprecedented accuracy better than 5 %. The 22GHz-band receivers of the VSOP-2 is currently designed to observe highly redshifted nuclear masers further to z=0.0793, which corresponds to cz = km/s. The most distant nuclear maser known to date is found in a type 2 QSO at z=0.66 (Barvainis & Antonucci 2005), while the second most distant nuclear maser is found in radio galaxy 3C403 (Tarchi et al. 2003) at z= The VSOP-2 22GHz receivers will be able to the maser in 3C403, if the flux density of the maser is sufficiently strong for VLBI observations H 2 O masers in the Local Group The VSOP-2 phase-referencing capabilities will enable proper motions of H 2 O masers also in the Large/Small Magellanic Clouds (LMC/SMC) to be determined. Measurement of the proper motions within a galaxy in the Local Group should indicate galactic rotations, peculiar motions reflecting the detailed galactic dynamics or star formation process, or secular motions of the galaxies. Proper motions of galaxies in the Local Group would indicate the total mass and angular momentum of the Local Group. These parameters will tell us how the Local Group has been developed during the 28

34 Figure 2.12: The number of known H2 O megamaser. The image courtesy of J.Braatz (National Radio Astronomy Observatory ). cosmological timescale (Brunthaler et al. 2005). Star formation process in LMC/SMC is one of the most important issues to elucidate how star formation is processed in metal-poor environments in LMC/SMC, similar to those seen in the earlier stage of the universe evolution. Accurate locations and three-dimensional spatio-kinematics of the H2 O masers in LMC/SMC should tell us the star-formation activity in detail and evolutionary status of the maser sources. These results might help in better understanding of how easily star could be created under the metal-rich environments as seen in the Galactic disk. M33/19 IC133 Figure 2.13: H2 O maser proper motions seen M33/19 and IC133 in M33 (Brunthaler et al. 2005). Left: Proper motion vectors of the two maser sources. Right: Position drifts in R.A. and decl. directions. 29

35 2.7 Formation and evolution of stars A standard picture of young stellar objects Star formation begins with the gravitational contraction of a molecular cloud and continues until a main sequence star has been formed. Most stars are believed to be formed in star clusters, implying that fragmentation as well as gravitational contraction of an interstellar cloud occurs at the same time. Since the lifetime of a star depends on its mass (t M 3 M 2 ), for a better understanding of the evolution and destiny of individual stars it is important to investigate accretion process during the star formation period when the final mass is determined. Star formation starts when dense molecular cores are formed. These molecular cores are supported by magnetic pressure and collapse quasi-statically. Formation of the molecular core occurs on a timescales of 10 7 yr. As the center of the molecular core gets denser, the infall of gas contributes to the formation of the protostellar core. The mass of a protostar increases to solar-mass scale within yr. After the gas accretion phase, the protostar becomes a Classical T Tauri Star (CTTS) which is visible at optical wavelengths. Accreting matter forms a proto-planetary disk around the central CTTS. This CTTS phase continues for yr. The protostar evolves into a Weak Line T Tauri Star (WTTS) and then a main sequence star within several tens of million years. Figure 2.16 is the standard picture of a protostar. A protostar is surrounded by a molecular core consisting of gas and dust accreting while in free fall. Such a accreting gas disk was discovered in the T Tauri HL Tau in 1.4 mm observations with the BIMA array and found to have quite sharply defined edges (Welch et al. 2003). The scenario mentioned above has been relatively well examined for low and intermediate mass stars. However, a standard picture of massive stars (M 10 M ) is still quite obscure. There are two competing scenarios for massive star formation, mass accretion onto a single massive young stellar object through a gas disk/torus with an extremely high mass accretion rate (Ṁ 10 3 M year 1 ) and a sequence of merger of lower mass stars. Because massive star formation sites are located relatively far away ( 500 pc) and the time scale for star formation is much shorter (t 10 5 years), so far only a small number of suitable massive star formation sites have been studied and even these have only been examined on scales larger than 100 AU A standard picture of dying stars A main sequence star evolves into a red giant or a red supergiant, depending on its initial main sequence mass, after hydrogen burning at its center has terminated. At this stage hydrogen burning occurs in a shell located outside of either a helium or carbon core, while the outer stellar layer extends up to as 100 times ruther than a Solar radius in order to maintain hydrostatic equilibrium. At the stellar surface, the gravity is much smaller and the temperature is much lower ( 3000 K) than a main sequence star. Because the stellar luminosity is as 10,000 times high as a Solar luminosity, strong radiative pressure leads the stellar surface dynamically unstable, drives energetic mass loss flow, and forms a thick circumstellar envelope. At a later stage of stellar evolution, a star as heavy as or heavier than the Sun reaches the asymptotic giant branch (AGB) stage, during the final phase of which the stellar luminosity changes with a period longer than 100 days (Mira variable stars) and mass loss occurs at the highest rate ( M year 1 ). Because the stellar luminosity and physical condition of the circumstellar envelope change periodically and a flow velocity is faster than a sonic velocity, shock waves in the envelope are formed in the vicinity of the stellar surface and transferred outwards through the circumstellar envelope. Finally such energetic mass loss ends and photo-evaporation and photo-ionization of the circumstellar envelope are processed by the central star that is becoming a white dwarf to form a planetary nebula or a Wolf-Rayet star following super nova explosion. Note that a large fraction of planetary nebulae are significantly asymmetric while the parental stars are spherically symmetric. Some of them 30

36 exhibit highly collimated bipolar jets as seen in AGNs and YSOs. Figure 2.14: Schematic picture of protostar magnetosphere (Hartmann 1998). Figure 2.15: Left: X-ray flares from the protostar YLW15 observed with the X-ray satellite ASCA (Tsuboi et al. 2000). Right: The star-disk magnetic reconnection model of YLW15 (Montmerle et al. 2000). Magnetic field lines are twisted by the differential velocity of the central star and the inner edge of its accretion disk Protostar s magnetospheres Recent theoretical and observational researches have revealed that extended magnetic structures, magnetospheres, exists around Young Stellar Objects (YSOs) (André 1996). Theoretically, they will play an important role in the accetion disk angular momentum. (Bouvier et al. 1993) and acceleration of the outflow (Shu et al. 1994, Pudritz & Norman 1986). There is evidence of large-scale magnetic structures in YSOs from X-ray observations of T Tauri stars and protostars, with X-ray flares frequently observed (Tsuboi et al. 2000). 31

37 Figure 2.16: The physical scales associated with a protostar. Matter accretes via an accretion disk, with the innermost material being directed toward the star via the star s magnetic field lines. Outflow from the star takes the forms of jets, stellar winds and bipolar outflow. Moreover, the detection of a non-thermal radio continuum by VLBI and circular/linear polarization provide direct evidence of magnetospheres around YSOs. The non-thermal radio emission is most likely gyrosynchrotron radiation from mildly relativistic electrons in the magnetospheres (André et al. 1992). There is a correlation between the X-ray and the radio luminosities as shown in figure 2.17 (Güdel & Benz 1993, Furuya et al. 2003), indicating that there it a close relationship in the radiation mechanisms and locations. The stellar flares have similarities with solar flares, as shown in figure 2.17 (Shibata & Yokoyama 1999). It is thought that the X-ray flare and the non-thermal radio continuum are caused by a large scale flare due to magnetic reconnection. Non-thermal radio continuum from YSOs VLBI observations of evolved Weak line T Tauri stars (WTTSs) have revealed magnetic structures with sizes of up to 5 25 times the stellar radius and with brightness temperature of K or more (André 1996). However, there are to date only 14 objects from which non-thermal radio continuum emission has been detected(see table 2.2), and there are only several images of the magnetic structures from global VLBI observations, for example, V773 Tau (HD ) (Phillips et al. 1996). The large-scale strong magnetic structures in young protostars are expected from the detection of circular polarization (Feigelson et al. 1998) and hard X-ray emission (Tsuboi et al. 2000). However nonthermal radio continuum emission from protostars has not been detected to date. Because previous observations were limited to lower frequency bands (<8.4 GHz), free-free obscuration or interstellar scattering of opaque ionized winds/jets associated with the protostars may explain the non-detections (André et al. 1992, Girart et al. 2004). Therefore, in order to detect non-thermal radio emission from protostars, VLBI observations at high frequency such as 22 and 43 GHz, in the optically thin regime, are highly desirable. Observations by VSOP-2 As mentioned above, the size of the magnetic structure around YSOs is, from X-ray and radio observations, tens of solar radii in extent. It is, therefore, expected that imaging of the magnetosphere will be able to be carried out in detail because the VSOP-2 angular resolution of 38 µas at 43 GHz corresponds to 1.2 R in a star forming region at a distance of 150 pc. Figure 2.18 shows the time variation of the spectrum of a giant flare from GMR-A, a WTTS in the Orion star-forming region. Spectra are shown 2, 4, 18 and 69 days after the detection of the flare with the BIMA array at mm wavelengths. During the flare the flux density increased by more than a factor of 5 on a timescale 32

38 Table 2.2: Objects from which non-thermal radio continuum emission has been detected by VLBI. No CTTS or protostar has been detected to date. Spectral Brightness Object Region type log L 5GHz Radius temperature Reference (erg s 1 Hz 1 ) (R 3R ) (10 8 K) V773 Tau Taurus K < Phillips et al. (1991) HDE Taurus G < Phillips et al. (1991) Hubble 4 Taurus K < Phillips et al. (1991) HP Tau/G2 Taurus K < 3.5 > 6 Phillips et al. (1991) T Tau Sb Taurus M 16.9 (8.4 GHz) 14.5 R Smith et al. (2003) DoAr 21 ρ Oph K > André et al. (1992) ROXs 39 ρ Oph K < 10 > 0.1 André et al. (1992) WL 5 ρ Oph F7(IR) 17.2 < André et al. (1992) VSG 11 ρ Oph M3(IR) 16.6 < 10 > 0.1 André et al. (1992) ROC 16 ρ Oph cool(ir) 16.9 < 10 > 0.2 André et al. (1992) VSSG 14 ρ Oph A < 10 > 0.6 André et al. (1992) S1 ρ Oph B André et al. (1992) Θ 1 Ori A Orion B3/TTS Felli et al. (1989) GMR-A Orion K (86 GHz) 36 R > 0.5 Bower et al. (2003) Figure 2.17: Left: Correlation of X-ray and radio luminosities of stars (Furuya et al. 2003). Right: The similarity between the stellar flares and solar flares (Shibata & Yokoyama 1999). 33

39 of hours, to a peak of 160 mjy at 86 GHz, making it one of the most luminous stellar radio flares ever observed. Remarkably, the Chandra X-Ray Observatory was in the midst of a deep integration of the Orion nebula at the time of the BIMA discovery; the source s X-ray flux increased by a factor of 10 approximately 2 days before the radio detection (Bower et al. 2003). Just after the mm-flare, the source has a rising spectrum which is weak at frequencies as low as 8 GHz. The detection of nonthermal radio emission from protostars is therefore more likely at 22 or 43 GHz, where the emission appears to be optically thin. In Figure 2.18, the upper triangles correspond to the detection limits of VSOP-2, and the lower inverse triangles correspond those by using the phase-referencing observing mode. We can conclude that it is possible to detect the radio emission at higher frequencies soon after a flare using phase-referencing. Figure 2.18: Time variation of spectra of GMR-A (Bower et al. 2003) Spectra are shown 2 days (22 Jan 2003), 4 days (24 Jan), 18 days (7 Feb) and 69 days (29 Mar) after the flare. Detection limits of VSOP-2 (upper inverse triangles) and those by using phase-referencing mode (lower inverse triangles). The expected flux densities (mjy) at 22 and 43 GHz for T Tauri stars and protostars which have been detected at X-ray energies by satellites such as ASCA and Chandra are shown in table 2.3, where the flux density is inferred from the known correlation between X-ray and radio luminosities. These YSOs that undergo bright X-ray flares will be good candidates for VSOP-2 observations. It is thought that non-thermal radio emission from the magnetospheres of YSOs is strongly circular polarized due to its origin as gyrosynchrotron radiation. The ability to undertake imaging observations at both left and right circular polarization simultaneously with VSOP-2 will enable clarification of the structure of the magnetic field around YSOs Rotation of protostars The rotation period of a protostar is a key parameter for the accretion onto the protostar because it is a direct indicator of the angular momentum accumulated by the star. However, it is very difficult to determine the rotation period of protostars as they are deeply embedded and optically undetectable. The detection of X-ray emission from protostars has opened a new window on the rotation of protostars. The X-ray light curves of two protostars in the ρ Oph region were observed by ASCA. The sinusoidal light curve of WL6 (figure 2.19) was interpreted to be the rotation modulation of the protostar (Kamata et al. 1997). The light curve of YLW15 shows quasi-periodic X-ray flares at 20 hour intervals (Tsuboi et al. 2000) (figure2.15). The flares can be interpreted as star disk magnetic reconnection in which 34

40 Table 2.3: Candidates YSOs for VSOP-2 observations which have been detected at X-ray energies. The expected flux densities (mjy) at 22 and 43 GHz inferred from the X-ray luminosities are shown. Object Distance Log L x Flux (22 GHz) Flux (43 GHz) Class Ref. (pc) (erg/s) (mjy) (mjy) DoAr III Imanishi et al. (2003) S II Imanishi et al. (2003) GY II Imanishi et al. (2003) YLW15A I Tsuboi et al. (2000) YLW16A I Imanishi et al. (2003) V773 Tau III Skinner et al. (1997) (quiescent) GMR-A III Bower et al. (2003) the central star itself is rotating rapidly with a 1 day spin period close to the break-up rotation velocity (Montmerle et al. 2000). These results suggest that the protostar is a fast rotator, however confirmation from observations at other wavelengths in which the emission is optically thin is essential Effesberg (2.8 cm) VLA (2 cm) Flux Density (mjy) Julian Date [ ] Figure 2.19: Left: The sinusoidal light curve of the protostar WL6 as observed by the X-ray satellite ASCA (Kamata et al. 1997). Right: Periodic radio flares of V773 Tau close to 3.4 day rotational period of the star spots (Massi et al. 2002). If periodic flares at radio wavelengths from X-ray protostars are detected, then the radio continuum emission will be the another new window in which to investigate the rotation of the protostar. Indeed, the periodic radio flares from the WTTS V773 Tau (HD283447) are observed with a modulation in agreement with the 3.4 day rotational period of the star spots (Massi et al. 2002) (figure 2.19). This Classical T Tauri Star (CTTS) observed at optical wavelengths is a slow rotator with a velocity of v sin i km s 1 (e.g., Bouvier et al. 1993) whereas protostars is expected to rotate faster. To reduce the angular momentum, momentum loss by such mechanisms as the stellar wind (Shu et al. 1994) and momentum transfer by magnetic interaction between the star and its accretion disk have been proposed. The mass of the central star can be estimated dynamically from the size of the magnetic structure of the protostar imaged by VSOP-2 and its rotation period if assume that the star co-rotates with the inner edge of its accretion disk due to the disk locking mechanism (Ghosh & Lamb 1979). 35

41 2.7.5 Binary systems Multi-epoch, 0.1 angular resolution Very Large Array (VLA) observations for years have revealed orbital motions and masses in several proto-binary systems in star forming regions (Rodríguez 2004). For instance, from 8.4 GHz observations spanning 1990 to 2002, Curiel et al. (2003) reported the proper motions of the two components of the proto-binary YLW 15, which shows periodic X-ray flares. The orbital proper motions suggest a lower limit to the total mass and upper limits to period of the binary system are 1.7 M and 360 yr, respectively. Figure 2.20: Predicted proper motions of YLW15 binary system (Curiel et al. 2003). The solid lines assume a mass ratio of 2.4 and show the predicted orbits for a timespan of 40 (left) and 500 (right) years. The absolute proper motions can be obtained by multi-epoch VSOP-2 phase referencing observations. If the non-thermal radio continuum can be detected from the binary system, the highly accurate astrometry achieved by VSOP-2 over a short time span will allow the orbital motions of close binary and spectroscopic binary systems on the plane of the sky to be determined. The double radio source V773 Tau is one of the best targets (Phillips et al. 1996) H 2 O maser excitation and early stellar evolution The non-thermal emission in a protostar may indicate that the protostar forms a convective zone at the stellar surface, which creates a dynamo action to amplify the stellar magnetic fields and drive magnetohydrodynamical process on the stellar surface. Alternatively, in the case of a massive young stellar object, the non-thermal emission may indicate that some parts of a strong wind are relativistically accelerated (e.g. Reid et al. 1995). In the both case, mass accretion has occurred at a much earlier stage than these activity phases. Such an earlier phase of star formation is identified in a molecular cloud core with gas density n H cm 3. Following the mass accetion, an outflow is also created for releasing angular momentum contained in the accreting material. In some of them, H 2 O maser emission, a target of VLBI observations, is also detected. Most of H 2 O masers are associated with energetic outflows from YSOs. A typical dynamical age of the outflows traced by H 2 O masers is 1000 years, much shorter than those derived from CO emission observations. This implies either that maser excitation occurs at the later stage of outflow formation or that the dynamical ages derived from H 2 O maser observations do not indicate true ages of the outflows. However, some H 2 O maser outflows shows extremely short dynamical ages, for which no extended CO outflow is detected (e.g. S106 FIR, Furuya et al. 2000). On the other hand, some H 2 O masers are likely associated with accretion disks/torii of YSOs as mentioned later. These indicate that H 2 O masers are associated with 36

42 the earliest phase of YSO s growth. A H 2 O maser source usually consists of many compact maser clumps, or maser features, which are separately detected with high angular resolution VLBI and whose proper motions can be determined on short time scales (longer than one month). Most exciting H 2 O masers (outflow or inflow) sources are still obscure because of the small number of measured maser feature proper motions in each YSO, which are necessary for elucidating true spatio-kinematics of gas around them. More than half of the individual maser features have lifetimes shorter than 2 3 months and their proper motions could not be measured because they were too slow. Most such short-lived maser features in the whole our our Galaxy should be targets for VSOP-2 proper motion measurements. The next two subsections describe studies on outflows and proto-stellar disk with H 2 O maser observations The root of outflow YSO outflows are extended to up to a few parsecs and are highly collimated. Such highly collimated outflows is an important YSO property and is very different from stellar winds driven by massive stars which are fundamentally less collimated. Recent observational and theoretical studies have concluded that such collimated outflows are jets that are driven magnetohydrodynamically very close to the YSOs. Note that colder molecular components in the outflows are dragged by hotter components in them and still be highly collimated in many observed outflows (Bachiller 1996). Most YSO H 2 O masers are associated with such molecular outflows. They are located in very vicinity of YSOs or roots of the outflows, therefore dynamical time scales of the outflows, which are derived from maser spatio-kinematics, are quite short (<1000 years) or much shorter than true ages of the outflows. However, some of the H 2 O maser dynamical time scales are really equal to true ages of the outflows and extremely short ( 100 years, e.g. S106FIR, Furuya et al. 2000). The maser spatio-kinematical structures are less collimated than those seen on larger scale in CO emission because the maser excitation regions are very close to YSOs, where YSO kinematical structures are complicated, including helical motions and equatorial flows. Some VLBI observations have revealed such complicated motions traced by H 2 O masers. Measurements of helical motions may enable us to estimate where the YSO jet is driven because a helical rotation velocity is comparable to a Keplerian rotation velocity of the jet s root. However, it is difficult to divide the complicated motions into individual kinematical components (e.g., poloidal flow and helical motion). Because previous VLBI observations determined only a small numbers of maser feature proper motions per maser source, it was impossible to construct any detailed maser spatio-kinematical models. Measutement of proper motions and its time derivate (proper acceleration), therefore, should provides us a great progress in constructing detailed maser spatio-kinematical models including helical motions. The first success in such a proper motion measurement is shown in figure 2.21(Moscadelli et al. 2006). Multi-epoch observations with VSOP-2 within a few months should be a great opportunity for such measurements described above Proto-stellar disks It is believed that a dusty and gaseous proto-stellar disk is present on scales of AU from the center of a protostar, and that this is where the mass accretion to the central star and its fluctuations play an important role in planet formation in the following stage. Sites of the dynamical accretion accompanying rotation have been imaged and studied on AU-scales using millimeter-wave arrays (e.g. Hayashi et al. 1993; Momose et al. 1998). On scales of 1 10 AU, the gas kinematics can be probed by H 2 O masers using VLBI. To date, 22 GHz H 2 O masers have been found in locations ranging from low-mass stars to massive stars in starforming sites. The physical conditions necessary for excitation of the H 2 O masers are T = K and densities of cm 3, which can occur in molecular outflows or at shock fronts in the disk. 37

43 R.A. Relative offset from the feature 8(mas) Decl. Figure 2.21: H 2 O maser features detected with the VLBA (Moscadelli et al. 2006). Left: Spatial maser distribution. Right: Proper motions of the feature 6 with respect to the feature 8. The motions are well fit to constant acceleration motions along the jet axis. Many of these masers are associated with molecular outflows, while some of them, suggesting rotation or accretion, have been found in protostars. It is not sufficient to distinguish rotation and accretion from outflows just by observing the velocity fields of H 2 O masers. It is crucial to measure the proper motions of the masers at mas resolution using VLBI. VLBI observations will be able to reveal the 3-D dynamics of the disk, measuring both the proper motion and the line-of-sight velocity of the masers. The first success in such a proper motion measurement is shown in figure 2.24 (Imai et al. 2006). VLBI mapping of maser features enables us to estimate an accurate mass for the protostar, assuming they are gravitationally bounded to the central star. The accretion rate can be estimated using the mass and luminosity of the protostar. Obtaining these disk parameters directly from VSOP-2 observations will have a great impact on the study of the star- and planet-formation. Table 2.4 displays a list of all H 2 O masers that have observed in proto-stellar gas disks to date. The VSOP-2 angular resolution at 22 GHz is 80 µas, corresponding to 0.04 AU at distance of 500 pc. At this resolution, the maser proper motion at a velocity of 10 km s 1 can be detected in seven days. On the other hand, the intensity variability of H 2 O masers in low-mass stars is very short and typically less than a month so that the interval of VLBI observations must be as short as possible. The high angular resolution of VSOP-2 will make such observations easier. It is also interesting to note that those masers that are located at the far side of the Galactic center or outer edge of our Galaxy can be measured on comparable scales at observing intervals of 2 3 months. With VSOP-2, we can resolve maser features located at as distant of several kpc with comparable linear resolution with those at 100 pc, and image the initial stage of the star-forming process and emergence of H 2 O maser emission more clearly. VLBI imaging of non-thermal continuum emission from protostars is a very powerful tool for determining the location of the central star with respect to each maser feature. The location of the star measured by VLBI can be a reference point to measure the relative proper motion of each maser feature. Thus, the detection of the non-thermal radio emission from protostars would give us the location of the dynamical center in a star, which can help in clarifying intrinsic proper motions from 38

44 Figure 2.22: Distribution of C 18 O molecular gas associated with the protostar L1551 IRS5, obtained by the Nobeyama Millimeter Array (Momose et al. 1998). Weak-line T Tauri stars are known to show high radio and X-ray luminosity. The velocity fields in this image show that there is rotation with a velocity of 0.24 km s 1 at r =700 AU from the center, as well as the accretion with a rate of 0.5 km s 1. internal motions of maser features. This is quite important when the maser emission is located very close to the dynamical center (within 2 3 AU) where a maser proper motion exhibit a centrifugal acceleration motion due to disk rotation or outflow helical motion. These provide great opportunities to directly estimate the stellar mass and dynamical process of the disk or outflow. ALMA will be able to resolve disk structures at spatial resolutions of 0.01 arcsec in sub-millimeter bands, which can reveal the mass distribution in a disk by the thermal continuum emitted from dust around the disk. ALMA is, however, unable to resolve the distribution and kinematics of molecular gas at such spatial resolutions and line-frequencies. On the other hand, VSOP-2 resolution allows the 3-D structures of the maser features to be determined although it is not able to observe the mass distribution. If the mass distribution from ALMA observations and infall motions from contemporaneous VSOP-2 observations can be obtained, the mass accretion rate of the star can be estimated. These studies of disks around protostars are complementary, utilizing both ALMA and VSOP-2 data in order to investigate the physics of protostars. Table 2.4: H 2 O masers in proto-stellar disks Object D r V rot Central mass References (pc) (AU) (km s 1 ) (M ) Cep A HW ±10 70±40 Torrelles et al. (1996) G Shepherd et al. (2004); Imai et al. (2006) L Fiebig et al. (1996); Fiebig (1997) NGC 2071 IRS1& Torrelles et al. (1998); Seth et al. (2002) AFGL5142 C Hunter et al. (1999) 39

45 Figure 2.23: Radial velocity distribution of water masers in the star-forming region L1287 and a rotating-infalling disk model best fit to the maser distribution (Fiebig 1997). A young stellar object (3 Sun mass) is located at the center of the disk. If maser proper motions, with curved trajectory, can be detected, this should be the most direct evidence for the masers associated with a circum-protostellar disk. Figure 2.24: An H 2 O maser disk in the protostar G ,imaged by VLBI (Imai et al. 2006). Analysis of the proper motion and absolute coordinates of maser features within 70 AU from the central star, which may be located close to the 2.6 mm continuum emission peak found by Shepherd & Kurtz (1999) rather than the 7 mm peak found by Shepherd et al. (2001) and has a mass of 8 M, provides direct evidence for rotation-infall motion around the central star. These observations suggest that a single star with at least M = 8 M can be formed through mass accretion from a thick gas torus/disk. The maser excitation may occur at the mid-plane of the disk/torus. 40

46 2.8 Study on final stellar evolution Stellar H 2 O and SiO maser emission is associated with energetic outflows driven by stellar mass loss at the highest rate of mass loss in the stellar evolution. Previously, H 2 O and SiO masers associated with evolved stars or asymptotic giant branch (AGB) stars were considered to be out of scientific targets of VSOP-2 observations because these maser spots are extended so that the VSOP-2 angular resolution completely resolves them out. However, Colomer et al. (1992) discovered extremely compact 43 GHz SiO maser spots. Although they are significantly resolved (detected flux fraction less than 50% with the VLBA beam, e.g. Yi et al. 2005), they are still detectable with inter-continental VLBI baselines and one of main targets of VLBA observations. Imai et al. (1997) discovered H 2 O maser spots detected in a VLBI baseline longer than 1000 km. Nowadays many stellar H 2 O masers have also been observed and detected on long VLBA baselines. Some of important stellar masers are located as far as 8 kpc from the Sun and many H 2 O and SiO masers have been detected towards the Galactic center, bulge, and the whole disk, higher angular resolution is essential for elucidating maser spatio-kinematics. Therefore stellar H 2 O and SiO maser emission also should be one of the VSOP-2 science targets Stellar SiO masers R.A. offset (mas) mas) Decl. off R.A. offset (mas) V (km s -1 LSR ) t (m Figure 2.25: SiO v=1 and v=2 J=1 0 maser emission mapped with the VERA in 2005 January 17 (Matsumoto et al., in prep). Dashed circles represent a stellar uniform disk with a diameter of 20.2 mas, obtained by near infrared observations. (Monnier et al. 2004). The location of the stellar disk is conjecture. Lef t : Radial velocity distribution of the SiO (v=2 J=1-0) maser emission. Right : Comparison of the v=1 and v=2 maser distribution. The two maser maps at different transitions are superposed so that the distribution patterns are spatially matched. It is important to register accurate distributions of these two maser transitions at the angular resolution of better than 1 milliarcsecond realized by VSOP-2 for tracing temporal variation of SiO maser excitation that may be indicated by position shifts of the two maser transitions. Recent studies of SiO masers have revealed that they are located very close to the stellar surface, at 2 3 stellar radii from the surface. The spatio-kinematics of 43 GHz SiO masers (J =1 0 transitions) have revealed complicated phenomena in the molecular atmosphere of AGB stars, where material released from the stellar surface is still in gas phase before forming circumstellar dust (e.g. Diamond 41

47 and Kemball 2003). The maser spatio-kinematical structure exhibits not only outflow due to radiative pressure of the star but also contraction towards the stellar surface, indicating that gas in the molecular atmosphere is still gravitationally bound by the star. They occur asymmetrically while the central star is spherically symmetric. Some SiO maser sources show the existence of rotating circumstellar envelopes (Imai et al. 1999b; Boboltz & Marvel 2000; Sánchez-Contreras et al. 2001), but they should be carefully examined because of drastic temporal variation of the spatio-kinematics of the SiO masers (e.g., Yi et al. 2005, see figure 2.25). Note that some maser features exhibit apparently extremely fast motions which are are probably not true bulk motions. Spatial and velocity structure of individual maser features and their temporal variation should be monitored in more detail to elucidate true temporal variation of physical conditions and dynamics of the molecular atmosphere. One of the most important issues of SiO masers is their excitation mechanism. SiO v = 1 and v = 2 J =1 0 lines are spatially coincident despite the vibrational levels of the two lines being quite different ( 1800 K). Investigation of the two lines at 43 GHz have been investigated in detail with the VLBA (e.g. Yi et al. 2005, see figure 2.26). The spatio-kinematical relation between the two lines seems to be dependent on phase of stellar pulsation, implying that the maser excitation mechanism may be mixture of radiative and collisional pumping schemes, in which the most dominant scheme is time dependent. On the other hand, the same relation also seems to be dependent on stellar type (O-rich or C-rich). The excitation scheme in term of line overlapping has been proposed on basis of VLBI observations of SiO v = 1 and v = 2 J =2 1 lines at 86 GHz as well as SiO v = 1 and v = 2 J =1 0 lines at 43 GHz (Soria-Ruiz et al. 2004). Thus VLBI observations of multi-transition maser lines enable us to examine these maser pumping schemes with time dependence and to elucidate the local conditions of the emission region. Precise tracing of such physical conditions in detail should lead us to understanding mass loss process through stellar pulsation. The ground-based VLBI observations at 86 GHz, especially with Global mm VLBI Array (GMVA), need space VLBI observations at 43 GHz to compare maser maps obtained in these two bands with similar spatial or angular resolution ( 100µas). Note that angular resolution of VSOP-2 at 43 GHz and GMVA at 86 GHz correspond to a segment of the surface of a Mira-type star. Mass ejection from gas cells floating from the stellar convective zone should be one of the origins of asymmetrical mass loss and may be detectable in SiO masers with the angular resolution mentioned above. A phase referencing capability for the VSOP-2 observations is essential for direct comparison of masing regions at the 43 GHz (J =1 0) transitions (v =1,2) and the similar 86 GHz (J =2 1) transitions. Astrometry accuracy necessary for such map comparison should be better than 1 mas. This capability also enable the data to be coherently integrated over a nominal coherence time and increases the probability of detecting the source across the entire VSOP-2 orbit. Even without comparative astrometry, comparisons of the angular scales of the emitting regions provides a measurement of the spot brightness temperatures or, in other words, of the percentage of the total flux that is emitted at each scale. Thereby we will study the maser saturation and consequentially the maser gains a direct function of the pumping mechanism. This will be potentially a further powerful discriminator between the radiative and collisional pumping mechanisms, as these have an order of magnitude difference in gain. The are large science returns to be had in being able to test the various models which have been proposed and it is critical to be able to compare the predictions of the models to observed physical conditions Stellar H 2 O masers Although H 2 O masers are located at much larger distances (5 50 AU) from a central pulsating AGB star than SiO masers, unique phenomena should be expected in these locations. At this location gas ejected from an AGB star cools and dust grains are formed. Pulsation in the molecular photosphere of an AGB star is generated by strong stellar radiative pressure. This pulsation may be enhanced at 42

48 km s May 11 May 31 June 19 July 13 August 1 (1998) pc Figure 2.26: H 2 O maser proper motions seen in the semiregular variable RT Virginis (Imai et al. 2003b). left: The three-dimensional maser feature motions. A location and a velocity vector are indicated by the root position and opening direction/length of a cone, respectively. right: A proper motion of a maser feature (a cluster of sequential velocity components of masers) indicating (constant) acceleration. Cyan arrows indicates intensity peak position of the accelerating maser feature in the individual observation epochs. the outer envelope where radiative pressure to dust grain plays more important role of the envelope dynamics. Shock waves produced by such stellar pulsation are transferred from the SiO to H 2 O maser region with acceleration. Such shock waves may be observed as acceleration/deceleration of H 2 O maser features (Imai et al. 2003b, see figure 2.26). Because the acceleration/deceleration driven by a shock wave may be seen on a time scale shorter than 2 3 month, multi-epoch VLBI observations should be made within this duration with sufficient angular resolution. Therefore, detection of such acceleration/deceleration in motions of H 2 O maser features should be one of scientific targets for VSOP-2. These observations enable us to precisely understand the dynamics of the circumstellar envelopes of AGB stars and to more accurately estimate stellar mass loss rates, an important parameter for tracing the final stellar evolution. Asymmetric stellar mass loss is also one of the most important issues for elucidating stellar evolution. At the late AGB and post-agb stages, highly collimated stellar jets are observed (e.g. Imai et al. 2002, see figure 2.27). The astronomical jets are now recognized not only in AGNs and YSOs but also in dying stars shown here. But driving mechanism of such stellar jets are still obscure, driven by either a single AGB star evolving to a compact white dwarf or a binary system consisting of an AGB star providing gas and a compact object (white dwarf) generating the jet. These stellar jets are identified as H 2 O fountains, H 2 O maser emission exhibiting high speed ( 100 km s 1 ) components with spatial and kinematical collimation. These H 2 O fountains show extremely short dynamical ages ( 100 years), but the number of the H 2 O fountain candidates is slowly increasing (8 sources by 2006 November). Most of the water fountains are located at far distances, they are also targets of the VSOP-2 science. 43

49 North-South Offset (milliarcsecond) H 2 O maser LSR velocity (October 1994) km s km s km s km s km s km s -1 Proper Motion 100 km s -1 East-West Offset (milliarcsecond) x OH maser LSR velocity (June - October 1994) km s km s km s AU at 2.6 kpc Figure 2.27: Highly collimated precessing jet in the OH/IR star W43A. Top: Collimated jet and spherically-expanding circumstellar envelope traced by H 2 O and 1612 MHz OH maser emission, respectively (Imai et al. 2002). W43A is located at the heliocentric distance of 2.6 kpc. The jet precession is suggested by the time variation of H 2 O maser distribution during last 10 years. The dynamical age of the jet is 50 years (Imai et al. 2005). Bottom: Magnetic field lines in W43A, which show equal magnetic field strength are elucidated by VLBI polarimetry (Vlemmings et al. 2006). The W43A magnetic field is widely well aligned and indicates toroidal magnetic fields along the jet, indicating that the W43A jet is driven by a magneto-hydrodynamical mechanism. Table 2.5: H 2 O/SiO/OH masers in the water fountain class. Object D l jet V exp t jet SiO OH References (kpc) (AU) (km s 1 ) (year) IRAS ? ? x Morris et al. (2003) OH ?? 23? 70 x Boboltz & Marvel (2005) IRAS ? 100??? x Deguchi et al. in prep. IRAS ? 150??? x Deguchi et al. in prep. W43A Imai et al. (2002, 2005) IRAS x x Imai et al. (2004, 2007) 44

50 2.8.3 Exploring magnetic fields in jets of YSOs and evolved stars In studying astronomical jets in YSOs and evolved stars mentioned above, as well as in AGNs, measurements of circular and linear polarization of radio emission are quite important for elucidating the magnetic field structures in the jets. VLBI polarimetry is possible when the VLBI observation is made with both left- and right-circular polarization detection. In both of H 2 O and SiO masers, both of circularly and linearly polarized emission are detected, which indicate magnetic field strength projected in the line-of-sight and a direction of the magnetic field, respectively. VLBI observations are necessary to measure the polarization for individual maser features, otherwise the observed polarization is depolarized. For example, the recent VLBI polarimetry of H 2 O masers revealed that a stellar collimated jet is driven by the magneto-hydrodynamical force in a toroidal magnetic field (Vlemmings et al. 2006, see figure 2.27). Although the degree of polarization is small in H 2 O masers ( 10%), by contrast in SiO masers it is quite high (over 40%). Some Mira variable stars have been revealed their magnetic field around the stars, which provide an important probe to investigate the stellar internal structure driving dynamo action and to examine shaping circumstellar envelopes and planetary nebulae by the magnetic fields. In star-forming regions, the interstellar magnetic fields play important roles not only for driving YSO jets but also for controlling mass accretion onto the YSOs. Empirically the magnetic field strength in star-forming regions is known to be expressed as a function of gas density B n 0.47 (Crutcher 1999), which is quite consistent with the magnetic field strength observed in the H 2 O maser clump density (Sarma et al. 2001). Thus the magnetic field will be enhanced in such maser clumps and well trace the magnetic fields around young stellar objects in more detail (e.g., Imai et al. 2003). Even if ASTRO-G does not have capability of VLBI polarimetry, polarimetry with ground-based VLBI baselines for compact maser features with VSOP-2 observations to understand the physical link between the maser feature dynamics and the magnetic field on the scale comparable to a maser feature size ( 1 AU). 2.9 Other Science Topics for VSOP Gamma-ray bursts Introduction Gamma-ray bursts (GRBs) were discovered in the late 1960 s. Isotropy of burst locations on the celestial sphere was established with BATSE (CGRO) in the 1990 s. The cosmological origin of GRBs was confirmed after BeppoSAX localizations of GRB positions made possible observations of afterglows in the late 1990 s. Radio observations Radio afterglows are faint, typically < 1 mjy (see Figure 2.28), so in general too faint for VSOP-2. However, radio observations have been very important. Observations of scintillation for the first few weeks after a burst provide evidence for the fireball model and imply superluminal motion in the fireball s expansion, as expected (see Figure 2.28). Proper motion studies also favour the fireball model over the competing cannonball model (see Figure 2.29). In the fireball model, the gamma-ray burst drives a relativistic blast wave into the circum-burst medium. In the relativistic fireball model a shift in the flux centroid is expected due to the spreading of the jet ejecta. For a jet viewed off the main axis the shift can be substantial. However, since gamma-rays were detected from GRB030329, it is likely that we are viewing jet pointing close to the line of sight. 45

51 In the cannonball model, plasmoids are ejected (like cannonballs) during the supernova explosion with relativistic speeds, and the emission comes from these compact knots. Proper motion in the cannonball model originates from the superluminal motion of the plasmoids, and a displacement of 0.55 mas over the 80 days of the VLBI observations was predicted. This is ruled out by the observations. Also, a more general problem for the cannonball model is the absence of rapid fluctuations in the radio light curves of the GRB. Strong diffractive scintillation is expected between 1 and 5 GHz with a modulation index of order unity and a timescale of a few hours, because the size of the plasmoids ( 0.01 micro-arcsec) always remain below the Fresnel scale ( 5 micro-arsec) of the turbulent ionized medium. While interstellar scintillation can be a powerful tool to test predictions of GRB afterglow models, such observations are rare. Most GRBs are expected to remain in the strong diffractive regime for only a few days, and hence there are practical difficulties obtaining well-sampled light curves. Nearby GRBs (z<0.2), although equally rare, enable us to image the afterglow directly, measuring the expansion or proper motion of the emitting region. VLBI observations of GRBs before GRB did not yield constraining limits on the angular size, as the GRBs were much more distant (z 1). GRB at z = was 50 times brighter than any previously studied event. This is a roughly once per 10 year event so if such an event occurred during the lifetime of VSOP-2, it should be observed with high priority. Number of Afterglows F ν [mjy] GHz Flux Density (µjy) t [day] Figure 2.28: Lef t: Histogram distribution of flux densities (or upper limits) at 8.5 GHz for a complete sample of bursts. The hatched histogram shows the distribution for the detections only. (From the review of Frail 2003.) Right: Light curve of the radio afterglow of GRB at 8.46GHz, compared with fireball model predictions of Waxman (1997) Supernovae Studies of extragalactic radio supernovae and supernova remnants produced when massive stars explode in the starburst galaxies have been undertaken. VLBI is a powerful tool for imaging the expansion of the supernovae shells in the early stage after their explosions (e.g., Marcaide et al. 1996; Bietenholz et al. 2003). A sequence of the supernova expansion of SN1993J in the nearby spiral of M81 is shown in Fig.2.30, in which images are giving important physical parameters of the expansion and deceleration rate and the distance to the galaxy. The peak flux densities of SN1993J at the brightest stage reached to about 50 mjy, which is well above the imaging sensitivity limit of VSOP-2 at 8.4 GHz. The size of the supernova shell at the earliest stage right after the explosion has not ever been resolved with ground-vlbi facilities, however, the VSOP-2 with improved angular resolution is expected to resolve the earliest stage ( 1 month after the explosion) of the explosion, revealing that 46

52 Figure 2.29: The positions derived from the observations of the radio counterpart to GRB in the first five epochs, relative to the first determination on April 1st at 8.4 GHz. A circle with a radius of 0.26 mas (2σ) is shown to encompass all measurements except for two observations at 5 GHz, which are known to suffer from systematic errors. These observations provide a constraint on the proper motion of 0.10 ± 0.14 mas over 80 days: evidence in favor of the fireball model and against the cannonball model. (From Taylor et al ) surrounding materials around the star are blown off and complete scenario of the explosion from its emergence to the end by monitoring with a certain interval Mass mapping through microlensing A significant fraction (23%) of total energy of the universe is believed to be in the form of dark matter, which, as its name suggests, has not been optically identified yet. Searches for dark matter objects have been conducted in the mass range M 10 7 M (Alcock et al. 2000) to M (Nemiroff et al. 2001). The integrated mass in the examined mass ranges is insufficient to satisfy the total mass of dark matter. The unexamined range of 10 1 < M < 10 4 M may make an important contribution to the total mass. Objects in the unsearched mass range can be identified as the darkmatter. One candidate is black holes of M, referred to as MASCOs (MASsive Compact Objects), and which are remnants of massive population III stars. Inoue & Chiba (2003) pointed out that microlensing by MASCOs can potentially be detected with VSOP-2. If a MASCO is located in front of a powerful radio source, it produce a dimple the size of the Einstein radius θ E given as θ E = ( M 10 2 M ) 1/2 ( ) DL D S /D 1/2 LS mas. (2.6) Gpc Here D L and D S are the distances to the MASCO and the radio source from the earth, and D LS is the distance between the MASCO and the radio source. The typical apparent size of 30 µas is detectable with the VSOP-2. Since the apparent size is proportional to mass of the MASCO, an imaging study of multiple MASCOs enables us to estimate the mass function. Even though such MASCOs may not be detected, the resulting upper limits will constrain the mass distribution of dark matter, a valuable contribution to cosmology. 47

53 Figure 2.30: Ground-VLBI images of SN 1993J in M81 at 8.4 GHz (Bietenholz et al. 2003). Lef t: Supernova evolution sequence, obtained from mid-1993 to late Right:1 mas scale, corresponding to 4000AU at d=3.8 Mpc is in the bottom. Brightness scale (color) ranges from 16 to 100 percent of the peak brightness value in each image. VSOP-2 observations at 8.4 GHz will be able to resolve compact sources in the last stage of evolution before the supernova explosion, and the structures and evolution processes of supernovae in distant galaxies. 48

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59 Chapter 3 VSOP-2 Mission 3.1 Introduction From HALCA to ASTRO-G VSOP-2 is built upon the success and experience gained with VSOP. A number of space VLBI missions were proposed in the 1980s, by the USA, Europe, Russia and Japan, but imaging Space-VLBI observations were never realized until the launch of the VSOP satellite of Highly Advanced Laboratory for Communications and Astronomy (HALCA) in The first launch of the M-V rocket enabled the development of the engineering satellite MUSES-B (HALCA) to test the technologies required of a Space-VLBI mission. Following the demonstration of these technologies, VSOP observations made in collaboration with the world s radio telescopes and with the cooperation of other space agencies also yielded many scientific results, taking full advantage of the improved resolution at 1.6 and 5 GHz. The success of VSOP has paved the way for future missions with improved sensitivity and resolution. Thus, VSOP has evolved to VSOP-2 with the fully scientific satellite ASTRO-G and improved angular resolution at higher observing frequency bands than the VSOP observations Other space-vlbi projects As of 2005, no next-generation mission has yet been approved. ARISE, and iarise, missions were proposed to NASA but were not selected. The main reason is the competition from X-ray and gamma-ray observations, which must be placed in orbit, and optical and infrared missions, for which the atmosphere imposes significant limitations. Space-VLBI observations, by their very nature, require the participation of ground-based telescopes, and in these the atmosphere impacts upon the sensitivity by limiting the coherence time, particularly at higher frequencies. The limitations imposed by atmospheric turbulence at high frequencies suggest that in the future, high frequency arrays of telescopes will be placed in orbit. VSOP-2 with the phase referencing capabilities is thus a logical, and significant, step in this direction, given its higher observing frequencies and improved sensitivity. 3.2 Requirements for VSOP-2 In this section, the technical requirements for VSOP-2 to achieve the science goals detailed in the previous chapters are briefly outlined Angular Resolution One of the main purposes of the VSOP-2 mission is to obtain the highest angular resolution images of any existing VLBI facility. It is expected that an angular resolution of less than 40 µ arcseconds at 43 GHz will enable the highest angular resolution images to be made. The disadvantage of higher 54

60 Table 3.1: Resolution on the maximum baseline (mas) 5 GHz 8 GHz 22 GHz 43 GHz VERA VLBA VSOP (0.081) VSOP Observing band 8 GHz 22 GHz 43 GHz Frequency range [GHz] SEFD [Jy] Coherence time [seconds] * Estimated ranges of the SEFDs: Jy (K-band), Jy (Q-band). Table 3.2: Basic properties of the ASTRO-G satellite observing bands. Typical coherence time for ground based VLBI observations are given. These depend highly on the location of the ground antenna and the weather. The coherence time of the satellite is determined by other factors, but here we assume these coherence times for space-ground baselines frequency VLBI (at >86 GHz) is that atmospheric effects are very large, coherence times are very short, telescopes are fewer and (u, v) coverage is poorer. At this time of 2006, there has been no VLBI array effectively operating at higher than 86 GHz. Space-VLBI offers solutions to all these problems. There are trade-offs between angular resolution and imaging fidelity for the VSOP-2 satellite s apogee height, but from the M-V class rocket launch capability and the nominal satellite mass, we assume an apogee height of 25,000 km. The angular resolutions for VSOP-2 observations in comparison to several other facilities are given in table Sensitivity There are some kinds of sensitivities, such as baseline (fringe detection) sensitivity, imaging sensitivity, and brightness temperature sensitivity. Baseline sensitivity can be improved with longer integration time, which is limited by the phase fluctuation by the atmosphere and the observing instruments. It is possible to correct the phase fluctuation using phase-referencing observations, which extend the integration time and enable higher sensitivities to be achieved. The rms thermal noise (σ ij ) for the baseline between antennas i and j is here: SEF σ ij = (ηs 1 Di SEF D j ) 2Bτ (3.1) η s = efficiency for the VLBI system: 0.64 for 1 bit sampling, 0.88 for 2 bit sampling SEF D i = System Equivalent Flux Density for antenna i, in Jy. (See tables 3.2) B = Bandwidth, in Hz. For VSOP-2, this will be 128 or 256 MHz. τ = Integration time in sec. Typically, the coherence times in table 3.2. are defined. 55

61 Table 3.3: 7-σ detection limit ( ν (MHz) = 256, η loss = 0.8) Telescope 1 Telescope 2 SEFD 1 (Jy) SEFD 2 (Jy) ν (GHz) τ (s) 7-σ (mjy) VERA VSOP VERA VSOP VERA VSOP VLBA VSOP VLBA VSOP VLBA VSOP VLA 27 VSOP VLA 27 VSOP VLA 27 VSOP * Adopting the goal values. The 7-σ baseline sensitivity between the VSOP-2 satellite and typical ground telescopes a VLBA 25 m, VERA 20 m and the phased 27-element VLA are given in table 3.3. With sensitive telescopes with a large aperture such as the Effelsberg 100 m, phased-up VLA (Y27) or GBT (Green Bank Telescope) 100 m, more sensitive observations are feasible (Table 3.3). An extremely high-sensitivity array, the SKA (Square Kilometre Array), is currently being planned by a large international collaboration with an observing capability in at least the 8 GHz band anticipated. The use of these antennas will improve the sensitivity by another order of magnitude VLBI Polarimetry ASTRO-G is equipped with dual polarizations (LHCP and RHCP), to get better quality, better resolution polarization image, and to increase sensitivity for the spectral line observations. This polarization capability was missing in HALCA. 3.3 Outline of the satellite Improvements over VSOP Ten times higher observing frequency The dense regions in the central parsecs of active galaxies become increasingly optically thin as the observing frequency, and resolution increase. VSOP-2 observations at 43 GHz are almost an order of magnitude higher than the 4.9 GHz observations of VSOP. Ten times better resolution The improved angular resolution of VSOP-2 is a result of the higher observing frequency and an increased apogee height of the satellite s orbit. The VSOP-2 angular resolution of 38 microarcseconds at 43 GHz is an order of magnitude better than the 0.4 milliarcsecond angular resolution of VSOP at 4.9 GHz. Ten times higher sensitivity 56

62 Figure 3.1: The schematic design of the VSOP-2 satellite. The effective aperture of the main reflector is 9 m, and that of the 40 GHz band link antenna is 80 cm. The VSOP-2 front-end will employ cryogenically cooled, low noise receivers and the increased downlink data rate, of 1 Gbps, enables wider observing bandwidths to be sampled. Additional gains in aperture efficiency are obtained from the off-set Cassegrain design. For continuum radio sources, this results in an order of magnitude improvement in sensitivity. Astrometric capability and further sensitivity gains from phase referencing Phase referencing observations can be made by nodding the satellite between the target source and a nearby calibrator (with a cycle time of about one minute. Centimeter level orbital accuracy will enable astrometric observations with higher resolution and better relative baseline errors than ground based observations. The phase referencing capability yields further gains in sensitivity as the integration time can be extended beyond the coherence time imposed for ground telescopes by the atmosphere. Measurement of magnetic fields through dual polarization observations HALCA was sensitive to only left-circular polarized radio emission, however the VSOP-2 satellite will enable both left and right circular polarizations to be detected, at all three frequencies. Magnetic field orientation, strength, and observations Faraday rotations will therefore become possible. Features of the VSOP-2 satellite To achieve the improvements listed above, the satellite design incorporates the following characteristics: 1. Large-scale deployable antenna of offset Cassegrain design. 2. Observation bands are at 8, 22, and 43 GHz (wavelengths of 4, 1.3, and 0.7 cm), with both LCP and RCP radiation being received. 3. Receivers at 22 and 43 GHz cryogenically cooled to 30 K. 4. High-speed data transmission from the satellite at 1 Gbps. 57