Gemini Near-Infrared Integral-Field Spectrograph (NIFS)
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Science Drivers
NIFS science is adaptive optics science, so much of the Gemini core science that has been described in the science case for ALTAIR will be realized using NIFS in combination with ALTAIR. NIFS is a fast-tracked, limited capability spectrograph and so will not reproduce all of the GNIRS capabilities. The key features of NIFS are its high spatial resolution IFU and its moderate spectral resolution in the 1-2.5 mm wavelength range; imaging spectroscopy at the maximum spatial resolution attainable with Gemini will be the primary role of NIFS. The application of this technique to the study of the demographics of massive black holes in nearby galactic nuclei, and the related study of the excitation and dynamics of the inner narrow-line regions of nearby Seyfert galaxies are defined to be NIFS core science. NIFS will excel in all observations requiring moderate spectral resolution data of spatially complex regions having high surface brightness in either a spatial or spectral sense; spectroscopy of compact, high surface brightness continuum sources and imaging of extended, narrow emission-line regions are examples of these two extremes. These themes feature strongly in the broader science programs described below.
The science requirements of NIFS cannot be treated in isolation from the requirements imposed by ALTAIR. We therefore attempt to estimate AO natural guide star and performance requirements where possible.
3.1 NIFS Core Science
3.1.1 Massive Black Holes in Nearby Galactic Nuclei
One of the most profound results from the Hubble Space Telescope (HST) is the evidence for the existence of massive (107 to 109 Mʘ) black holes in the nuclei of many nearby early-type galaxies (e.g., Kormendy & Richstone 1995; Lauer et al. 1995; Faber et al. 1997). An apparent relationship exists between the black hole mass and bulge mass which suggests that either central black holes grew by accreting inner bulge stars, or else that the central black hole and the bulge formed coevally in major merger events. However, this correlation suffers from strong observational selection effects (Ford 1997), and it is yet to be determined whether elliptical galaxies and spiral galaxies follow the same or different correlations. Given the potential close link between black hole formation and galaxy evolution, it is important to define the mass distribution and frequency of occurrence of central black holes in both classes of galaxies. These are still poorly known; most particularly in late-type spiral galaxies in which the nuclear regions are obscured by dust and the bulge masses are smaller so the black hole masses may also be smaller. Observations with NIFS will help determine the demographics of massive black holes in galactic nuclei. Spatially resolved, high resolution dynamical studies of the innermost nuclear stellar populations (Figure 1) and LINER-like gaseous accretion disks (Figure 2) at near-infrared wavelengths are necessary to do this. Observations of both surface brightness distributions, mean rotation, and radial velocity dispersion profiles with spatial resolutions of a few parsecs and spectral resolutions of 3000-5000 are required to model the stellar dynamics or gaseous accretion disk dynamics and infer properties of the central black hole. The high spatial resolution required dictates the use of adaptive optics correction. This, and the presence of obscuring dust in the central regions of many galaxies which complicates the interpretation of optical data from HST, dictates the use of near-infrared observations. The CO (2-0) absorption bandhead at 2.3 mm is ideal for measuring velocity dispersions of cool stellar populations in the nuclei of low redshift galaxies (e.g., Gaffney, Lester, & Doppmann 1995; Shier, Rieke, & Rieke 1996; Böker, van der Marel, & Vacca 1999). The presence of a mass concentration is indicated by a rising stellar velocity dispersion profile near the nucleus. The emission-lines of H I Pb 1.282 mm, H I Brg 2.166 mm, [Fe II] 1.257 mm, and [Fe II] 1.644 mm are expected from shock-excited gas in circumnuclear LINER-like accretion disk. The enclosed mass is inferred from the rotational velocity, assuming that the gas follows Keplerian orbits about the mass concentration.
Figure 1: HST/WFPC2 F606W images of the central 18˛´18˛ regions of spiral galaxies (Carollo, Stiavelli, & Mack 1998). Many have cuspy cores, and many are obscured optically.
Figure 2: Dusty circumnuclear LINER-like disk in the black hole galaxy NGC 7052 (van der Marel, & van den Bosch 1998).
The velocity dispersion profile of M32 rises from sV ~ 60 km s-1 at 1.0˛ radius to sV ~ 95 km s-1 at 0.1˛ radius (Bender, Kormendy, & Dehnen 1996; Figure 3). Stellar velocity dispersions are therefore expected to be ~ 50 km s-1 in the outer parts and ~ 100 km s-1 at the centers of galaxies containing low mass black holes. These velocity dispersions correspond to Gaussian FWHMs of ~ 118 km s‑1 and ~ 235 km s-1, respectively. A FWHM velocity resolution of ~ 100 km s-1 will therefore suffice to measure stellar velocity dispersions from the CO (2-0) absorption bandheads in the K band. The LINER gas disk in the elliptical galaxy M84 has peak rotational velocities of ±400 km s‑1 (Bower et al. 1998), and the gas disk in NGC 4261 has peak rotational velocities of ±200 km s‑1 (Ferrarese, Ford, & Jaffe 1996). A velocity resolution of < 100 km s-1 will be required to measure disk rotational velocities in a range of lower mass objects. Similarly high spectral resolving powers of R ~ 4000-5000 are required to significantly separate individual OH airglow lines in the J and H bands in order to perform sensitive measurements of the emission-lines from circumnuclear LINER-like disks.
Figure 3: Velocity dispersion and rotational velocity profiles for M32 which contains a 3´106 Mʘ black hole (Bender, Kormendy, & Dehnen 1996).
High Strehl ratios are required for these observations. Galaxy nuclei are in general too faint and diffuse to be efficiently used as guide objects for ALTAIR. Approximately 10% of nearby galaxies in the Shapley-Ames catalog have suitably bright AO guide stars within ~ 30˛ of the nucleus. However, the laser guide star facility will allow this sample to be greatly extended and will deliver near optimal Strehl ratios on objects with suitable OIWFS guide stars. It will be necessary to accurately determine the PSF by frequent measurement of a PSF star, by reconstructing the PSF from OIWFS frames, or by modeling based on the AO control loop output (Véran et al. 1997).
The performance of NIFS for measurements of stellar velocity dispersions in galactic nuclei has been modeled in NIFSSIM using empirical K band surface brightness and velocity dispersion distributions for our Galactic center. Saha, Bicknell, & McGregor (1996) fit the Galactic center cool star and hot star surface brightness distributions with Reynolds-Hubble law profiles and the stellar rotational velocity and velocity dispersion profiles with quartic polynomials. These profiles have been adopted for our model galactic nuclei after shifting them to distances ranging up to 20 Mpc and convolving with the instrumental point spread function, P(q). Each spatial position in the nuclear profile is assumed to have an M1III type spectrum in the K band, characterized by the Kleinmann & Hall (1986) spectrum of 75 Cyg, before convolution with a Gaussian velocity dispersion profile.
Figure 4 shows a simulated 3600 s K grating exposure on the nucleus of the model spiral galaxy shifted to 10 Mpc with no interstellar extinction. The simulated spatial image of the galaxy (compressed in the spectral direction) is shown in Figure 5. The peak K band surface brightness in the central 0.1˛´0.1˛ region is 10.53 mag arcsec‑2. Spectra with continuum signal-to-noise-ratios of ~ 30 are generally required to measure stellar velocity dispersions for dynamical studies (Kormendy, priv. comm.). The simulated observation has a signal-to-noise ratio of ~ 100 in the central 0.1˛´0.1˛ region (excluding systematic effects). The K band surface brightness at a fiducial radius of 1.0˛ (~50 pc at 10 Mpc) is 12.8 mag arcsec‑2. Figure 6 shows the K grating spectrum extracted from the central 0.1˛´0.1˛ region of the simulated exposure in Figure 4 after subtraction of a 3600 s sky exposure, transformation for 2D wavelength calibration, and division by a smooth spectrum star. The stellar velocity dispersion has been estimated from this frame using a least-square fit to a Gaussian-smoothed template spectrum which is identical to the simulated galaxy spectrum. Various subtleties complicate the analysis of real data including interstellar extinction, smearing due to the PSF, incomplete sky subtraction and correction for terrestrial atmospheric absorption, and systematic differences between the galaxy and template spectra. None of these is considered here. However, this simple analysis is sufficient to demonstrate the potential of this type of observation. The stellar velocity dispersions measured for each 0.1˛´0.05˛ spatial pixel are shown in Figure 7. The lighter dots are azimuthal averages of the individual velocity dispersion measurements, and the solid curve is the velocity dispersion profile from which the data were simulated. The formal velocity dispersion errors are more than adequate for the purpose of detecting a central massive black hole, but these are likely to be under-estimates of the uncertainties in real data. The K band surface brightness profile in the Galactic center declines as r-0.8 so the signal-to-noise ratios for the annular azimuthal averages should be proportional to ~ r0.2, thus accounting for the approximate constancy of the signal-to-noise ratios with radius in the simulation. This highlights the benefits of using an integral field unit for observations of this type.
Figure 4: Simulated 3600 s K grating exposure of the nucleus of a spiral galaxy at 10 Mpc. The CO Dv=2 absorption bands in the galaxy spectrum are apparent at right.
Figure 5: Spatial image of the galaxy nucleus shown in Figure 4. The K band surface brightness in the central 0.1˛´0.1˛ region is 10.53 mag arcsec-2. Each vertical slitlet is 0.1˛ wide. The field-of-view is 3.1˛´3.3˛.
Figure 6: Extracted spectrum of the central 0.1˛´0.1˛ region of the simulated galaxy at 10 Mpc shown in Figure 4.
Figure 7: Radial stellar velocity dispersion profile derived from the spectra in Figure 4. The black points indicate the velocity dispersion measurements for each 0.1˛´0.05˛ spatial pixel. The lighter dots are azimuthal averages of these data. The solid curve shows the velocity dispersion profile from which the simulated data were derived.
The radial velocity dispersion profile inferred from a simulated 3600 s exposure with the K grating of the same simulated galaxy now shifted to 20 Mpc is shown in Figure 8. Two effects are apparent here; the linear spatial resolution is degraded, and the signal-to-noise ratios of the measurements are worse. The poorer spatial resolution will be the limiting factor in detecting massive black holes at distances beyond ~ 20 Mpc. The velocity dispersion uncertainties are expected to limit attempts to infer total stellar masses, and hence mass-to-light ratios, in objects such as ultra-luminous IRAS galaxies which are typically at distances between 20 and 100 Mpc. The K band surface brightness (with no extinction) in the 20 Mpc simulation at a fiducial radius of 1˛ (~ 100 pc at 20 Mpc) is 13.4 mag arcsec‑2. A signal-to-noise ratio of 30 in 0.1˛´0.1˛ will be reached in ~ 4.5 hr of on-source integration with NIFS (§2.8.4) at this surface brightness. Clearly, integration times longer than 3600 s will be required to reduce the velocity dispersion uncertainties.
Figure 8: Radial stellar velocity dispersion profile for a simulated 3600 s observation of a spiral galaxy at 20 Mpc using the K grating. The black points indicate the velocity dispersion measurements for each 0.1˛´0.05˛ spatial pixel. The lighter dots are azimuthal averages of these data. The solid curve shows the velocity dispersion profile from which the simulated data were derived.
The finite spatial resolution achieved with NIFS is the main factor limiting our ability to detect massive nuclear black holes. The enclosed stellar mass at any radius is derived from the velocity dispersion profile and the light distribution using the collision-less Boltzmann equation. Under the reasonable assumptions that the velocity dispersion is approximately constant in the inner region of the galaxy and that the volume mass density is proportional to ~ r-2, the mass enclosed within radius r is given very approximately by
where D is the distance to the galaxy, s is the velocity dispersion at radius r, and Menc is in units of Mʘ. Figure 7 indicates that a galaxy like the Milky Way at a distance of 10 Mpc will have a velocity dispersion of ~ 70 km s-1 at r ~ 0.1˛ (~ 5 pc at 10 Mpc), the innermost measurement. This corresponds to an enclosed mass of » 1.2 ´ 107 Mʘ according to the above equation. The K band luminosity of the Milky Way within a radius of 5 pc is ~ 3.5 ´ 107 Lʘ, so the mass-to-light ratio of our simulated galaxy within the region probed by NIFS is M/LK » 0.34. Normal stellar populations can have M/LK in the range 0.0-2.5 (Figure 9), so we will require mass-to-light ratios in excess of at least ~ 5.0 to unambiguously indicate the presence of a massive black hole. This will be the case if the mass enclosed within r ~ 0.1˛ is > 1.8 ´ 108 Mʘ due to the presence of a > 1.7 ´ 108 Mʘ black hole. If this was the case, the velocity dispersion at r = 0.1˛ would be ~ 270 km s-1. We therefore expect that NIFS will unambiguously detect black holes with masses > 2 ´ 108 Mʘ in spiral galaxies out to distances of ~ 10 Mpc. The 2´106 Mʘ black hole in the Milky Way nucleus is uncharacteristically small; the best estimate of the mean relation between black hole mass and bulge mass (Magorrian et al. 1998) predicts that the 2´1010 Mʘ Milky Way bulge would typically be associated with a 1.2´108 Mʘ nuclear black hole mass. If this mean relation proves to be appropriate for late-type spiral galaxies, our simulations suggest that it will be possible to detect only the more massive of the black holes that may be present in late-type spiral galaxies to distances of at least 10 Mpc with NIFS.
Figure 9: M/LK versus cluster age for a exponentially declining star formation rate with 109 yr time constant (Thatte et al. 1997).
3.1.2 Nearby Active Galactic Nuclei
Many nearby galaxies possess active nuclei which are characterized by broad (FWHM ~ 500 km s-1) emission-lines originating in their central regions over size scales of 100 pc up to ~ 2 kpc. This is the so-called narrow-line region (NLR). The ultimate energy source is believed to be accretion onto a massive black hole in most objects, although intense starbursts in dense regions may be responsible for some LINER-like activity (Terlevich & Melnick 1985). Emission from the immediate vicinity of the accretion disk produces the broad-line region (BLR) which remains unresolved with existing telescopes. Understanding the nature of the central energy source, its interaction with the host galaxy, and the global implications for the evolution of galaxies are continuing themes in the study of Active Galactic Nuclei (AGN). High spatial resolution optical studies of AGN with HST (e.g., Capetti et al. 1996; Winge et al. 1997; Axon et al. 1998; Figure 10; Figure 11) have revealed a wealth of information about the structure and excitation of the inner NLR. While it has traditionally been believed that the NLR clouds are photoionized by the central source (Ferland & Netzer 1983; Wilson & Tsvetanov 1994), these recent high spatial resolution imaging and dynamical studies have demonstrated that NLR clouds may instead be predominantly shock-excited by energetic thermal and non-thermal mass outflows from the central object. Strong dynamical interactions between the emission-line gas and radio-emitting ejecta can be explained if the NLR is formed from shells of ambient interstellar medium swept up and compressed by the supersonic expansion of hot-gas heated by interactions with the advancing radio jet (Pedlar, Dyson, & Unger 1985; Taylor, Dyson, & Axon 1992; Steffen et al. 1997). The nuclear regions of Seyfert galaxies are invariably obscured by dust clouds making near-infrared observations of the inner NLR desirable. The near-infrared region also offers the best ground-based spatial resolution using adaptive optics correction. [Fe II] 1.257 mm, [Fe II] 1.644 mm, H I Pb 1.282 mm, and H I Brg at 2.166 mm emission-lines are well-suited to excitation and dynamical studies of the high-excitation precursor zones associated with fully radiative shocked regions (Figure 12). Strong coronal emission-lines are the primary initial coolants of hot gas in partially radiative shocks. With NIFS, the [Si VI] 1.961 mm coronal line will become accessible at modest redshift. The mechanical energy flux from the jet can be estimated from the [Fe II] and H I Pb lines in less obscured regions, and from H I Brg in more obscured regions. H2 1-0 S(1) 2.122 mm emission in Seyfert galaxies is also collisionally-excited, but generally has a smaller velocity width of ~300 km s‑1 suggesting that it may arise in a different emission region (Veilleux, Goodrich, & Hill 1996). X-ray-heating from the AGN core, shock-heating by the interaction of the radio jets with the interstellar medium, and shock-excitation in outflows from star formation regions may all contribute to the H2 emission from Seyfert galaxies. High spatial resolution dynamical studies may provide a means of distinguishing between these alternatives.
Figure 10: HST/WFPC2 image in [O III] l5007 of the NLR clouds near the nucleus of the Seyfert 1 galaxy NGC 4151 (Hutchings et al. 1998).
Figure 11: HST/WFPC2 image of NGC 1068 in [O III] l5007 light (grayscale) and 6 cm radio emission (contours) from Capetti, Macchetto, & Lattanzi (1997). Note how the radio jet penetrates between [O III] clouds. The 3˛´3˛ field-of-view is similar to that of NIFS.
Figure 12: Near-infrared K band spectra of the Circinus galaxy obtained with the 3D integral-field spectrograph in 0.5-0.6˛ seeing and with spectral resolving powers of 1000 (top) and 2000 (bottom) (Maiolino et al. 1998).
Seyfert activity is frequently associated with circumnuclear starbursts (Figure 13), often in rings, but the role these play in fueling or refueling the active nucleus is still unclear. The presence of circumnuclear starburst rings demonstrates that large quantities of gas have been channeled into the region close to the nucleus. This gas may accrete directly onto the black hole, but it must lose its remaining angular momentum to do this. It may be the stars, or their remnants, formed in the starburst ring or the central star cluster that feed the central black hole. High angular resolution spectral imaging of Seyfert galaxy cores will reveal structure interior to the starburst ring. The morphology and dynamics of the emission-line regions will permit new insight into how gas is funneled into the core and the role possibly played by stellar bars in driving such gas flows. Measurement of stellar mass-to-light ratios, M/LK, will probe the star formation histories of the regions. Detailed comparison of spatially-resolved spectra with starburst models (e.g., Leitherer et al. 1999) will provide estimates of the starburst ages, masses, and star formation histories. These can then be compared to particular models for AGN fueling (e.g., Norman & Scoville 1988). The potential of AO-corrected imaging of Seyfert galaxy cores in addressing these issues is beginning to be explored (Marco, Alloin, & Beuzit 1997; Chapman, Walker, & Morris 1998; Rouan et al. 1998; Marco & Alloin 1998, 1999; Figure 14).
Figure 13: K band images of NGC 7469 obtained with the AOB on CFHT (left) and without AO correction (right). The FWHM with AO correction is ~ 0.13˛. The field is 10˛´10˛. Note that the bright central Seyfert core has been subtracted from the AO corrected image (ALTAIR PDR).
Figure 14: Raw AO-corrected K band image of the core of NGC 1068 obtained with CFHT (Rouan et al. 1998). The field-of-view of 2.2˛´2.2˛ and pixel size of 0.0344 are similar to those of NIFS. The image is plotted on a logarithmic scale.
Measurement of the central black hole masses in Seyfert galaxies is also highly desirable. Although, these can be estimated using reverbation-mapping techniques for a few objects (e.g., Wandel, Peterson, & Malkan 1999), stellar velocity dispersions provide a more direct determination. Stellar velocity dispersions can be measured using the CO (2-0) absorption bandheads and interpreted in the same way as for normal galaxies (§3.1.1) to place limits on the enclosed mass and hence detect or constrain black hole masses. The analysis is complicated in the case of common Seyfert galaxies by the intense Seyfert core emission (Figure 14). Minimizing this contamination will depend on achieving high Strehl ratios in the telescope, ALTAIR, and NIFS. However, approximately 43% of nearby galaxies show a detectable level of nuclear activity (Ho, Filippenko, & Sargent 1997). These galaxies may contain either lower mass black holes or massive black holes that are currently accreting at well below their Eddington limit. These galaxies will be prime candidates for observation with NIFS. The core of the nearest Seyfert galaxy, Circinus, has a radius of < 1.5 kpc at K (Maiolino et al. 1998), corresponding to < 0.08˛. The stellar velocity dispersion within ~ 40 pc (~ 2˛) of the nucleus is ~ 75 km s-1, corresponding to a Gaussian FWHM of ~ 180 km s-1. We require a velocity resolution of ~ 100 km s-1 to confidently measure such velocity dispersions.
The AO
requirements on Seyfert galaxy programs are less severe than for normal
galaxies. Many Seyfert nuclei are
bright enough and sufficiently compact to use as AO guide objects for ALTAIR. A large number of nearby Seyfert galaxies in
the Shapley-Ames catalog were checked for nearby stars. Approximately 20% have suitably bright guide
stars within ~ 30˛ of the nucleus. Many of the 486 low luminosity Seyfert
galaxies of Ho, Filippenko, & Sargent (1997) also have suitable AO guide
stars and are at distances < 20 Mpc.
High Strehl ratios are needed to measure black hole masses, and are
desirable when studying emission from NLR clouds.
The feasibility of measuring emission-line profiles for Seyfert galaxy NLR clouds at high spatial resolution has been assessed using NIFSSIM. Near-infrared emission-lines from the NLR are seen against stellar emission from the galaxy and thermal dust emission from the Seyfert core. In a similar way to §2.8.5, we use the signal-to-noise ratio predictions for point sources observed with a Strehl ratio of 1.0 (§2.8.3) to infer the signal-to-noise ratios per 0.1˛´0.1˛ aperture obtained in 3600 s for different continuum surface brightnesses, and from these estimate 3s detection limits for emission-lines seen against these continua. We adopt an emission-line FWHM of 500 km s-1 typical of [Fe II] emission in Seyfert galaxies (Fig. 14 of Veilleux, Goodrich, & Hill 1997). H2 1-0 S(1) emission in Seyfert galaxies typically has a FWHM of 300 km s-1, so our surface brightness detection limits should be decreased by 3/5 for this line. The 3s emission-line detection limit predictions for a 3600 s integration time are listed in Table 1 for each of the NIFS gratings at their full resolving power and for a smoothed two-pixel resolving power of R = 1000 matched to the adopted emission-line width. The latter is better suited to detecting NLR emission-lines, but would provide little profile information.
Table 1: 3s emission-line detection limits in 3600 s integrations for 500 km s-1 line width.
|
J1
grating |
J2
grating |
||||
|
mJ |
R =
5090 |
R =
1000 |
mJ |
R =
6100 |
R =
1000 |
|
(mag arcsec-2) |
(W cm‑2 arcsec‑2) |
(W cm‑2 arcsec‑2) |
(mag arcsec-2) |
(W cm‑2 arcsec‑2) |
(W cm‑2 arcsec‑2) |
|
9.0 |
1.7´10-21 |
7.4´10-22 |
9.0 |
1.4´10-21 |
5.7´10-22 |
|
10.0 |
1.2´10-21 |
5.2´10-22 |
10.0 |
8.9´10-22 |
3.6´10-22 |
|
11.0 |
7.4´10-22 |
3.3´10-22 |
11.0 |
5.8´10-22 |
2.4´10-22 |
|
12.0 |
4.8´10-22 |
2.1´10-22 |
12.0 |
3.7´10-22 |
1.5´10-22 |
|
13.0 |
3.1´10-22 |
1.4´10-22 |
13.0 |
2.4´10-22 |
9.9´10-23 |
|
14.0 |
2.2´10-22 |
9.7´10-23 |
14.0 |
1.7´10-22 |
7.0´10-23 |
|
15.0 |
1.5´10-22 |
6.8´10-23 |
15.0 |
1.3´10-22 |
5.5´10-23 |
|
|
|
||||
|
H
grating |
K
grating |
||||
|
mH |
R =
5340 |
R =
1000 |
mK |
R =
5340 |
R =
1000 |
|
(mag arcsec-2) |
(W cm‑2 arcsec‑2) |
(W cm‑2 arcsec‑2) |
(mag arcsec-2) |
(W cm‑2 arcsec‑2) |
(W cm‑2 arcsec‑2) |
|
9.0 |
8.8´10-22 |
3.8´10-22 |
9.0 |
5.2´10-22 |
2.3´10-22 |
|
10.0 |
5.4´10-22 |
2.4´10-22 |
10.0 |
3.6´10-22 |
1.5´10-22 |
|
11.0 |
3.4´10-22 |
1.5´10-22 |
11.0 |
2.3´10-22 |
1.0´10-22 |
|
12.0 |
2.2´10-22 |
9.6´10-23 |
12.0 |
1.9´10-22 |
8.1´10-23 |
|
13.0 |
1.4´10-22 |
6.1´10-23 |
13.0 |
1.5´10-22 |
6.3´10-23 |
|
14.0 |
1.0´10-22 |
4.3´10-23 |
14.0 |
1.4´10-22 |
6.0´10-23 |
|
15.0 |
7.9´10-23 |
3.4´10-23 |
15.0 |
1.3´10-22 |
5.7´10-23 |
|
|
|
||||
|
J
grating |
HK
grating |
||||
|
mJ |
R =
2790 |
R =
1000 |
mK |
R =
2530 |
R =
1000 |
|
(mag arcsec-2) |
(W cm‑2 arcsec‑2) |
(W cm‑2 arcsec‑2) |
(mag arcsec-2) |
(W cm‑2 arcsec‑2) |
(W cm‑2 arcsec‑2) |
|
9.0 |
1.2´10-21 |
7.2´10-22 |
9.0 |
3.6´10-22 |
2.3´10-22 |
|
10.0 |
7.4´10-22 |
4.4´10-22 |
10.0 |
2.4´10-22 |
1.5´10-22 |
|
11.0 |
4.7´10-22 |
2.8´10-22 |
11.0 |
1.7´10-22 |
1.0´10-22 |
|
12.0 |
3.2´10-22 |
1.9´10-22 |
12.0 |
1.3´10-22 |
8.3´10-23 |
|
13.0 |
2.1´10-22 |
1.3´10-22 |
13.0 |
1.2´10-22 |
7.4´10-23 |
|
14.0 |
1.6´10-22 |
9.3´10-23 |
14.0 |
1.1´10-22 |
6.9´10-23 |
|
15.0 |
1.3´10-22 |
7.7´10-23 |
15.0 |
9.9´10-23 |
6.2´10-23 |
The underlying continuum surface brightness in AO observations of Seyfert galaxies can be estimated in several ways. Veilleux et al. have measured near-infrared spectra of 33 Seyfert galaxies through 3˛´3˛ or 1.5˛´1.5˛ apertures. The average emission-line and continuum surface brightnesses from Veilleux et al. are listed in Table 2. The J and K average surface brightnesses provide lower limits to the values that will be detected from the Seyfert core at AO resolution. NICMOS images suggest central H band surface brightnesses of < 10 mag arcsec-2 in typical Seyfert galaxies (Quillen et al. 1999), and AO images of NGC 3227 obtained with the Adaptive Optics Bonnette on CFHT have surface brightnesses at 0.1˛ radius of J ~ 7.5 mag arcsec-2, H ~ 6.5 mag arcsec-2, and K ~ 6.2 mag arcsec-2 (Chapman, Morris, & Walker 2000). Peletier et al. (1999) have compiled a database of ground-based subarcsecond resolution J, H, and K surface photometry of a sample of CfA Seyfert galaxies[1]. The K band surface brightnesses of this representative sample of Seyfert galaxies at a fiducial radius of 1˛ (which largely avoids the non-stellar core emission) are shown in Figure 15 as a function of redshift. Most of this sample have K surface brightnesses at 1˛ radius in the range 12-16 mag arcsec-2. Clearly, the underlying core surface brightness depends on the galaxy distance and on the Strehl ratio achieved in the observations, and may vary from < 10 mag arcsec-2 to > 15 mag arcsec-2 over the NIFS field-of-view.
Table 2: Average emission-line surface brightnesses (from Veilleux, Goodrich, & Hill 1997).
|
Object |
Aperture (˛) |
J Continuum (mag arcsec-2) |
K Continuum (mag arcsec-2) |
Pb 1.282 mm (W cm‑2 arcsec‑2) |
[Fe II] 1.257 mm (W cm‑2 arcsec‑2) |
H2 1-0 S(1) 2.122 mm (W cm‑2 arcsec‑2) |
H I Brg 2.166 mm (W cm‑2 arcsec‑2) |
|
|
|
|
|
|
|
|
|
|
MCG-05-23-16 |
3´3 |
13.2 |
12.0 |
2.2´10-22 |
1.4´10-22 |
… |
… |
|
Mrk 1 |
1.5´1.5 |
… |
14.3 |
… |
… |
2.0´10-22 |
1.8´10-22 |
|
Mrk 3 |
3´3 |
14.6 |
… |
1.3´10-21 |
1.8´10-21 |
… |
… |
|
Mrk 78 |
3´3 |
15.6 |
… |
1.5´10-22 |
2.2´10-22 |
… |
… |
|
Mrk 176 |
3´3 |
14.3 |
14.6 |
4.2´10-23 |
1.8´10-22 |
2.4´10-23 |
6.2´10-24 |
|
Mrk 266SW |
3´3 |
14.6 |
15.2 |
6.1´10-22 |
2.9´10-22 |
3.1´10-23 |
4.1´10-23 |
|
Mrk 266NE |
3´3 |
14.8 |
… |
4.5´10-22 |
5.7´10-22 |
… |
… |
|
Mrk 348 |
3´3(J) 1.5´1.5(K) |
14.9 |
13.9 |
1.5´10-22 |
2.6´10-22 |
5.2´10-23 |
3.9´10-23 |
|
Mrk 403 |
3´3 |
15.3 |
… |
6.7´10-23 |
… |
… |
… |
|
Mrk 463E |
3´3 |
14.9 |
12.6 |
3.3´10-22 |
2.3´10-22 |
5.8´10-23 |
3.0´10-23 |
|
Mrk 477 |
3´3 |
15.5 |
14.5 |
1.2´10-21 |
8.6´10-22 |
1.0´10-22 |
1.1´10-22 |
|
Mrk 533 |
3´3(J) 1.5´1.5(K) |
14.5 |
12.5 |
7.3´10-22 |
4.1´10-22 |
1.5´10-22 |
4.6´10-22 |
|
Mrk 573 |
3´3(J) 1.5´1.5(K) |
14.3 |
13.0 |
2.9´10-22 |
2.7´10-22 |
5.6´10-23 |
1.2´10-22 |
|
Mrk 622 |
3´3 |
14.4 |
… |
1.1´10-22 |
1.7´10-22 |
… |
… |
|
Mrk 917 |
1.5´1.5 |
… |
13.5 |
… |
… |
1.6´10-22 |
1.3´10-22 |
|
Mrk 1066 |
3´3(J) 1.5´1.5(K) |
13.8 |
12.4 |
1.3´10-21 |
1.3´10-21 |
6.0´10-22 |
6.0´10-22 |
|
Mrk 1073 |
3´3(J) 1.5´1.5(K) |
14.2 |
13.1 |
8.1´10-22 |
5.8´10-22 |
1.7´10-22 |
2.1´10-22 |
|
Mrk 1157 |
1.5´1.5 |
… |
13.8 |
… |
… |
1.9´10-22 |
1.0´10-22 |
|
Mrk 1210 |
6´3(J) 3´3(K) |
14.8 |
14.0 |
2.1´10-22 |
4.3´10-22 |
4.5´10-23 |
7.9´10-23 |
|
Mrk 1388 |
3´3 |
15.4 |
… |
2.2´10-22 |
8.7´10-23 |
… |
… |
|
NGC 1068 |
1.5´1.5 |
… |
8.2 |
… |
… |
1.3´10-21 |
5.6´10-21 |
|
NGC 2110 |
3´3(J) 1.5´1.5(K) |
13.5 |
12.0 |
… |
1.7´10-21 |
3.3´10-22 |
>8.3´10-23 |
|
NGC 2992 |
3´3 |
13.2 |
13.1 |
3.2´10-22 |
9.6´10-22 |
1.3´10-22 |
3.4´10-23 |
|
NGC 3081 |
3´3 |
13.8 |
… |
6.1´10-22 |
3.5´10-22 |
… |
… |
|
NGC 4388 |
3´3 |
14.3 |
13.5 |
9.0´10-22 |
8.2´10-22 |
2.1´10-22 |
1.3´10-22 |
|
NGC 5252 |
3´3 |
14.3 |
15.1 |
1.9´10-22 |
3.4´10-22 |
… |
7.3´10-24 |
|
NGC 5506 |
3´3 |
12.9 |
11.1 |
9.5´10-21 |
3.2´10-21 |
… |
1.3´10-21 |
|
NGC 5728 |
3´3 |
15.0 |
13.9 |
2.4´10-22 |
3.3´10-22 |
7.8´10-23 |
2.5´10-23 |
|
NGC 7212 |
1.5´1.5 |
… |
14.3 |
… |
… |
8.9´10-23 |
1.2´10-22 |
|
ESO 428-G14 |
3´3 |
13.4 |
… |
9.2´10-22 |
1.0´10-21 |
… |
… |
|
IC 5135 |
1.5´1.5 |
… |
13.2 |
… |
… |
3.5´10-22 |
2.8´10-22 |
Figure 15: Measured K surface brightness (mag arcsec-2) at a radius of 1˛ for CfA Seyfert galaxies (Peletier et al. 1999).
We expect the NLRs of Seyfert galaxies to resolve into individual near-infrared emitting clouds with sizes of ~0.2-0.5˛, as has been achieved with HST for the [O III] l5007 emitting clouds (Figure 10 and Figure 11). We estimate surface brightnesses for these clouds in two ways. First, the average emission-line surface brightnesses measured by Veilleux et al. through 3˛´3˛ or 1.5˛´1.5˛ apertures (Table 2) are broadly indicative of the 3s emission-line surface brightness detection limits NIFS must achieve in order to study individual NLR clouds efficiently. The second way of estimating emission-line surface brightnessess for NLR clouds is based on [O III] l5007 surface brightnesses measured with HST. Schmitt & Kinney (1996) tabulate 3s detection limit surface brightnesses for this line based on archival HST images (Figure 16). We use the near-infrared emission-line data of Veilleux et al. and the [O III] flux data of Whittle (1992) to estimate empirical flux ratios relative to the [O III] line of Pb/[O III] ~ 0.049, [Fe II] l1.257/[O III] ~ 0.049, H2 1-0 S(1)/[O III] ~ 0.009, and Brg/[O III] ~ 0.006. The average Seyfert 1 and Seyfert 2 galaxy [O III] detection limits then convert to the 3s detection limits listed in Table 3 for the near-infrared emission-lines. [Fe II] l1.257/[Fe II] l1.644 = 1.36 in the absence of reddening. The predictions and data in Table 1, Table 2, and Table 3 suggest that we should be able to detect NLR clouds in many objects away from the central Seyfert core at moderate spectral resolution, but that these clouds will only be detectable against the bright core by smoothing the spectra to R ~ 1000 or worse. NIFS is unlikely to reach the faint limits of HST [O III] images for some objects, but this is not necessary in order to study the brightest NLR clouds.
Table 3: 3s detection limits based on [O III] l5007 flux.
|
Type |
[O III] l5007 limit (W cm-2 arcsec-2) |
Pb l1.282 limit (W cm-2 arcsec-2) |
[Fe II] l1.257 limit (W cm-2 arcsec-2) |
H2 1-0 S(1) l2.122 Limit (W cm-2 arcsec-2) |
Brg l2.166 Limit (W cm-2 arcsec-2) |
|
|
|
|
|
|
|
|
Seyfert 1 |
4.1´10-21 |
2.0´10-22 |
2.0´10-22 |
3.7´10-23 |
2.5´10-23 |
|
Seyfert 2 |
1.4´10-20 |
6.9´10-22 |
6.9´10-22 |
1.3´10-22 |
8.5´10-23 |
Figure 16: [O III] images of Seyfert 1 (top) and Seyfert 2 (bottom) galaxies from HST (Schmitt & Kinney 1996). The horizontal bars correspond to 200 pc.
The feasibility of measuring the stellar velocity dispersions needed to measure black hole masses depends in the K band surface brightness due to the near-nuclear stellar light in Seyfert galaxies, as well as the concentration of AGN core light achieved with the AO-corrected PSF. We again base our analysis on the ground-based surface photometry of CfA Seyfert galaxies at a fiducial radius of 1˛ in Figure 15 (Peletier et al. 1999). The brightest of these Seyfert galaxies have surface brightnesses at 1˛ radius brighter than the value of 13.4 mag arcsec-2 simulated for the 20 Mpc galaxy in Figure 8. However, most of the sample have K surface brightnesses at 1˛ extending to ~ 15.0 mag arcsec-2. This surface brightness is reached at 1˛ radius if our simulated galaxy is moved to a distance of 120 Mpc. Reliable velocity dispersions cannot be inferred for individual spatial pixels at a surface brightness of 15 mag arcsec-2. Figure 17 shows the velocity dispersion profile inferred from a simulated 3600 s exposure using the K grating with our simulated galaxy at a distance of 120 Mpc. These dispersions were obtained by forming annular averages of the spectra before deriving the velocity dispersion. The signal-to-noise ratios of these averaged spectra are sufficient to derive velocity dispersions. However, velocity dispersion cannot be separated from rotational velocity in these annular averages so it remains unclear how scientifically useful such data will prove to be.
Figure 17: Radial stellar velocity dispersion profile for a simulated 3600 s observation with the K grating of a spiral galaxy at 120 Mpc having a K band surface brightness at 1˛ radius of 15 mag arcsec-2. The light dots are the values inferred from azimuthally averaged spectra. The solid curve shows the velocity dispersion profile from which the simulated data were derived.
3.2 Gemini Core Science
3.2.1 Brown Dwarfs and Low Mass Stars
Recent near-infrared sky surveys have succeeded in identifying large numbers of low mass stars and brown dwarfs. However, objects in binary systems still offer the only means of empirically determining precise masses and absolute magnitudes for this class of object. These data are basic to an understanding of the substellar mass function, and ultimately the transfer of angular momentum during star formation and the universal proportion of matter bound up in sub-stellar companions. The distances to objects in binary systems can be determined, so they provide empirical calibration of the color versus absolute magnitude relation that can be applied to field and cluster brown dwarf candidates. NIFS with occulting disks will record moderate resolution near-infrared spectra of the close companions that will provide effective temperatures and other physical parameters for the companions. The imaging capability of NIFS will be invaluable in removing the complex residual “speckle” pattern of the bright primary star.
Spectra in the J, H, and K bands with a resolving power of R ~ 1000 are sufficient to determine molecular absorption band strengths for temperature determination (e.g., Gl 229B; Geballe et al. 1996; Figure 18; K ~ 14.8 mag). The central star can be used as the AO guide star in all conceivable cases. Finding a nearby OIWFS guide star will be subject to random field statistics (§2.9).
Figure 18: Near-infrared spectrum of the methane absorption brown dwarf companion Gl 229B (Geballe et al. 1996).
NIFSSIM has been used to assess the performance of NIFS in detecting binary companions by determining the signal-to-noise ratio that will be achieved in spectra obtained with the K grating and smoothed to a two-pixel resolving power of R ~ 1000. We consider a faint point source at a radius of 0.5˛ from a bright star (Figure 19 and Figure 20). The extent to which the halo of the bright star can be subtracted will be a limiting factor in detecting faint companions. Conventional PSF fitting routines, such as currently in IRAF, cannot cope with the AO-corrected PSFs typical of NIFS images. There will also be residual speckle structure which is not well-fit by analytic PSF profiles. These complex issues have not been addressed in the present analysis; we simply subtract a spectrum extracted from the diametrically opposite position in the primary star image.
Figure 19: Simulated compressed NIFS image showing a K=18 mag companion star 0.5˛ to the left of a K=12 mag primary star in 0.4˛ FWHM seeing with a Strehl ratio of 0.6. The exposure time with the K grating is 1800 s which just saturates the K=12 mag star.
Figure 20: Spectrum of the K=18 mag M1 III star companion in Figure 19 smoothed to a two-pixel resolving power R ~ 1000. Atmospheric absorption features have been removed by division with a smooth spectra star.
NIFSSIM predicts that the NIFS detector will just saturate using the K grating in 1800 s on a K = 12 mag primary star in 0.4˛ FWHM seeing with a Strehl ratio of 0.6. The R=1000 signal-to-noise ratios achieved on a 0.5˛ offset companion star in this case are listed in Table 4. The minimum full frame integration time will be limited by the detector readout speed to ~ 10 s. A K = 6.4 mag star will saturate in this time under the above seeing and Strehl ratio conditions. Consequently, it will be necessary to occult stars brighter than K ~ 6.5 in order to search for their companions. Table 4 also lists R=1000 signal-to-noise ratios for companions to a K = 7 mag primary recorded with an integration time of 15 s. These estimates are scaled to a total integration time of 1800 s assuming SNR µ Öt. It can be seen from Table 4 that there will be a 1.5-2.0 mag penalty in detecting companions around the brighter star.
Table 4: Binary Companion SNR, R = 1000.
|
K = 12 mag Primary Star |
K = 7 mag Primary Star |
||||
|
Secondary Mag |
SNR in 1800 s |
Secondary Mag |
SNR in 15 s |
SNR in 1800 s |
|
|
|
|
|
|
|
|
|
14.0 |
342.3 |
14.0 |
13.5 |
147.5 |
|
|
15.0 |
180.6 |
15.0 |
6.0 |
65.2 |
|
|
16.0 |
107.1 |
16.0 |
2.3 |
25.2 |
|
|
17.0 |
43.4 |
17.0 |
1.0 |
10.5 |
|
|
18.0 |
17.6 |
18.0 |
0.4 |
3.9 |
|
|
19.0 |
8.6 |
|
|
|
|
|
20.0 |
3.3 |
|
|
|
|
3.2.2 Young Star Clusters
A knowledge of the initial stellar mass function over the full range of masses from the Eddington limit to below the hydrogen-burning limit, and its dependence on environment, are fundamental to an understanding of the star formation process. The upper stellar mass cut-off needs to be explored in nearby regions of massive star formations in order to better understand the nature of massive star formation occurring in more extreme regions such as starburst galaxies. Concentrations of high mass stars are found in young Galactic star clusters. These are often obscured by dust due to their youth or their large distance from Earth (Figure 21). High spatial resolution observations are needed to probe the cores of dense star clusters associated with Galactic giant H II regions (e.g., Blum, Damineli, & Conti 1999) and clusters in the vicinity of the Galactic center (e.g., Cotera et al. 1996). The physical parameters of embedded massive stars can be derived from high (~ 70) signal-to-noise ratio, moderate resolution (R > 1000) spectra in the H and K bands (Blum, Damineli, & Conti 1999; Hanson, Howarth, & Conti 1997; Hanson, Conti, & Rieke 1996).
Figure 21: K band image of the central W43 star cluster (Blum Damineli, & Conti 1999) showing three stars identified as early-type supergiants and a Wolf-Rayet star.
Knowledge of the low mass end of the stellar initial mass function in different environments is needed to determine the amount of Galactic mass locked up in low mass stars, to understand chemical enrichment and recycling in galaxies, and to determine the impact of starbursts on galaxy evolution. Most low mass stars currently forming in the Galaxy appear to be forming in star clusters associated with giant molecular clouds (Lada et al. 1991). Stellar masses for lower mass pre-main-sequence stars (M < 5 Mʘ) cannot be determined unambiguously from broadband near-infrared photometry alone due to the indeterminate effects of interstellar extinction and the nature of their evolutionary tracks. Moderate resolution K band spectra of these obscured, low mass, cluster pre-main-sequence stars are required to assign them spectroscopic temperatures (Hodapp & Deane 1993; Luhman & Rieke 1998), and hence infer their masses based on evolutionary tracks.
Source confusion
and the irregular backgrounds from complex reflection and emission nebulosity
associated with dense young star clusters make slit spectroscopy of faint,
embedded, young stars difficult. NIFS
with its IFU will allow more accurate removal of these irregular
backgrounds. The W43 cluster (Figure 21) contains stars measured photometrically to K ~ 16 mag, corresponding to a spectral
type of ~ A0 on the main sequence. NIFS
will be capable of measuring K band
spectra of these stars with R ~ 5000
and signal-to-noise ratios of ~ 40 in 1800 s.
3.2.3 YSO Jet Driving Mechanism
The driving mechanism for outflows from young stellar objects (YSOs) has not been observationally identified. Shocked, collimated jets are seen at large distances from the star (Figure 22), but the properties of the winds at their origins, and even the mass loss rates, remain uncertain and model dependent. High spatial resolution spectral imaging in emission-lines probing shocked gas, such as H2 1-0 S(1) 2.122 mm and [Fe II] 1.644 mm, will allow observation of the energetic, highly collimated jets as they emerge from the inner regions of the accretion disks. The high spectral resolution of NIFS, relative to narrow-band line filters, will enable better discrimination against continuum emission making NIFS the preferred Gemini instrument for near-infrared spectral imaging of faint, narrow emission-line sources. High resolution spectral imaging of YSO jets with NIFS will provide simultaneous morphological, excitation, and kinematic data which, over time, will allow the evolution of features in these stellar jets to be traced as they progress along the jet and interact with the surrounding material. Such observations are crucial to understanding the role played by high energy outflows in terminating infall and determining the final stellar mass. For example, a “Herbig-Haro” emission knot located in a nearby dark cloud ~ 150 pc from Earth and moving at 100 km s-1 traverses 0.13˛ in one year. Proper motions of such Herbig-Haro knots could be followed over a 2-3 yr period, allowing the acceleration mechanism to be probed as well as the interaction of these knots with the ambient cloud. Temporal variations in an extremely young Herbig-Haro flow ejected from XZ Tau have been seen in the optical with HST (Krist et al. 1999; Figure 23). Emission-line spectroscopy of such features with NIFS will reveal details of how the flows expand and evolve.
Figure 22: Jets and circumstellar disks from young stellar objects recorded by HST/WFPC2 in optical line and continuum emission.
Figure 23: HST/WFPC2 F675W images of XZ Tau in 1995 and 1998 (Krist et al. 1999). The size and appearance of the bubble has changed between these epochs.
The targeted emission-lines will be the H2 lines in the K band and [Fe II] 1.644 mm in the H band. The ratio of H2 1-0 S(1) 2.122 mm to H2 1-0 Q(3) at 2.424 mm can be used to derive the interstellar extinction correction. This is one motivation for extending the NIFS K grating response to the edge of the atmospheric window near 2.5 mm. These observations require the highest possible spatial resolution and velocity resolutions in the H and K bands of 50-100 km s-1. Velocity centroids can be determined to Dv ~ FWHM/SNR which should be ~ 5-10 km s‑1 with typical signal-to-noise ratios. Visible T Tauri stars can be used as natural guide stars for ALTAIR, when available, but a different star will be required for the OIWFS. The near-infrared responsivity of the OIWFS will greatly increase the availability of suitable OIWFS stars in star formation regions. Many YSOs are either not visible objects or are resolved in the optical. Laser guide stars will be essential for AO-corrected observations of these stars.
3.2.4 YSO Jet-Cloud Interactions
YSO outflows remove excess angular momentum from the protostar system, they contribute to the turbulent support of molecular clouds, and they may be responsible for disrupting molecular clouds and ultimately terminating star formation within them. YSO outflows generally consist of a highly collimated, high velocity bipolar jet embedded in a less well collimated low velocity bipolar molecular outflow that is detected at millimeter wavelengths The physical parameters of the highly collimated YSO jets are still incompletely understood. They emit most strongly in shock excited transitions of H2 and [Fe II] in the near-infrared and low excitation emission-lines typical of Herbig-Haro objects in the optical in regions where the jet material impacts the surrounding medium. A bow shock forms where shocked gas impacts quiescent material in front of a Mach disk which forms where the jet impacts previously shocked material. Observed emission-line strengths can be modeled either as J-shocks, C-shocks, or a combination of the two (e.g., Buckle, Hatchell, & Fuller 1999). J-shocks occur where the magnetic field is weak and the gas properties change suddenly. C-shocks occur in the presence of a strong magnetic field. A C-shock can form at the bow shock and a J-shock can form at the Mach disk when the magnetic field is slightly weaker. Which type of shock applies in YSO jet-cloud interactions is still controversial. The relative emission-line strengths give an indication of the type and speed of the shock (Smith 1995). The bow shock and Mach disk are expected to be separated by ~ 500 AU (~ 3.4˛ at 150 pc) in most cases (Hartigan 1989), but the curved structure of the bow shock complicates identification of this feature. High spatial resolution spectral imaging with NIFS may succeed in separating these components.
YSO jets often have a knotty appearance, possibly due to jet instabilities or episodic ejection. Emission arising from between the knots may be due to the jet being partially molecular, due to entrainment of ambient material in a mixing layer, or simply due to the existence of unresolved emission knots. YSO jets have many similarities (and differences) with relativistic jets emanating from radio galaxies. Understanding the physical processes occurring in YSO jets may also help in understanding the nature of these extragalactic jets.
Spatially resolved NIFS spectra in the K band are needed to determine H2 emission-line fluxes, flux ratios, and radial velocities. Proper motions are also needed for knots within the jet to test jet models quantitatively. Typical jet velocities are ~ 200 km s-1. Projected velocities are correspondingly lower, so velocity resolutions of ~ 50 km s-1 are required. Optical AO guide stars may be scarce in dark clouds, making laser guide star observations highly beneficial. It will be possible to use nearby embedded stars as the near-infrared OIWFS guide star in many cases.
Figure 24: HST/NICMOS image of H2 1-0 S(1) emission from the IRc2/BN region of OMC-1 (Stolovy et al. 1998). The inner H2 “bullets” are identified.
Figure 25: Integrated velocity profiles for inner H2 1-0 S(1) emission-line “bullets” in OMC-1 (Stolovy et al. 1998). The line profiles were obtained through 0.6˛ diameter apertures with 1.5˛ spatial resolution (Chrysostomou et al. 1997).
A simulation of a 1800 s K grating spectrum of a 1.0˛ diameter circular “bullet” in OMC-1 (Stolovy et al. 1998; Figure 24; Figure 25) is shown in Figure 26. Molecular hydrogen line wavelengths and strengths relative to the H2 1-0 S(1) line were taken from Oliva & Moorwood (1988) and Brand et al. (1988) for OMC-1 and Oliva, Moorwood, & Danziger (1990) for the supernova remnant RCW 103. All H2 lines are assumed to have a Lorentzian profile with FWHM = 40 km s-1 typical of OMC-1 (Chrysostomou et al. 1997). Peak 1 in the Orion Molecular cloud is among the highest surface brightness line emitting regions with an H2 1-0 S(1) surface brightness of ~ 10-20 W cm-2 arcsec-2 (Stolovy et al. 1998). Analysis of the image in Figure 26 indicates that the 5s flux limit for detecting H2 1-0 S(1) line emission spread over 3 spectral pixels in 1800 s with 0.1˛´0.1˛ spatial resolution will be ~ 2´10‑23 W cm-2 arcsec-2; ~ 500 times fainter than OMC-1. Flux limits appropriate to other lines depend on their proximity to terrestrial OH emission-lines and to the thermal emission beyond 2.3 mm. Figure 27 shows the simulated H2 1-0 S(1) line spectrum of the 0.1˛´0.1˛ square region of a source having a surface brightness of 2´10-23 W cm-2 arcsec-2.
Figure 26: Molecular hydrogen spectrum in the K band for a uniform 1.0˛ diameter circular region with a H2 1-0 S(1) line surface brightness appropriate for OMC Peak 1.
Figure 27: Simulated spectrum of the H2 1-0 S(1) line with surface brightness 2´10-23 W cm-2 arcsec‑2 (at pixel 597) extracted from a 0.1˛´0.1˛ square aperture from an 1800 s exposure. Other apparently significant features are residual noise from subtracted OH lines.
3.2.5 Late Stages of Stellar Evolution
In recent years, HST and deep ground-based imaging of asymptotic giant branch (AGB) stars, proto-planetary nebulae, and young planetary nebulae have revealed remarkable but previously unknown structures. These show exquisite detail of central, point symmetric, often bipolar, cavities being carved out of the centers of spherical AGB star mass-loss envelopes (Sahai et al. 1998; Sahai et al. 1999a,b; Sahai & Trauger 1998). How such non-spherical cavities can be produced inside the recent spherical AGB star mass-loss wind is an unsolved mystery. Binarity, rotation, and magnetic fields have all been suggested. Most importantly, it is the immediate post-AGB phase where the asymmetries develop, so it is here that we should look for objects beginning the aspherical mass-loss process. Samples of candidate objects are currently being examined from the ground (e.g., van der Steene & Wood 1999, priv. comm.): they often show Ha emission with a central peak of width ~ 100 km s-1, a P Cygni type absorption on the blue edge of the emission-line, and broad wings with widths of order 1000 km s-1. High spatial and spectral resolution observations of these objects are required in order to examine the gas dynamical processes occurring in them. In particular, K band observations in the lines of H2 1-0 S(1) 2.122 mm and H I Brg 2.166 mm are required in order to study the beginning of cavity generation (as evidenced by H2 emission from shocked gas) and to determine the location of the H I emission region: is it a wind from the AGB star remnant, is it a jet from an accretion disk around a companion star, do the broad 1000 km s-1 and 100 km s-1 components of the H I lines come from the same place spatially? Velocity resolutions of a few km s-1 are required for this work in order to make models of the gas flows in these systems. Since the H2 emission has been detected by NICMOS, it should be measurable with NIFS. Many proto-planetary nebulae and most young planetary nebulae have central stars that can be used as guide stars for ALTAIR. OIWFS guide stars will be subject to random field statistics (§2.9).
3.2.6 Galactic Center
The Galactic center is a unique region of the Galaxy populated by old stars forming the inner Galactic bulge as well as young clusters of massive stars indicative of recent intense star formation activity. There is now strong evidence for the existence of a central massive black hole in the Galactic center (Eckart & Genzel 1996,1997; Genzel et al. 1997; Ghez, et al. 1998). High spatial and moderate spectral resolution observations of stars in the vicinity of the black hole are required to determine their radial velocity dispersion to complement available high spatial resolution proper motion data, and to study the nature of the stellar population. Stars close to the black hole should interact with the black hole and with each other frequently and may show evidence of these interactions in their spectral or morphological properties. Understanding the star formation history of the central region of the Galaxy will lead to a clearer understanding of the formation of the Galactic bulge, the nature of star formation in an environment of extreme gas temperature, pressure, velocity dispersion, magnetic field strength, and tidal shear, and of the processes fueling the central black hole in our Galaxy and perhaps in other more active galaxies.
The Galactic center is obscured at wavelengths shorter than ~ 1.5 mm. Most of the bright stars detected in the K band are young late-type supergiant and asymptotic giant branch stars or older luminous red giant stars. Main sequence stars are intrinsically fainter in the near-infrared, and so are more difficult to detect. Nevertheless, main sequence stars in the volume around the Galactic center are being found in deep photometric surveys using conventional techniques (e.g., Blum, Sellgren, & DePoy 1996) and AO image correction (Davidge et al. 1997b). Near-infrared spectra are needed for significant samples of faint stars in the crowded Galactic center region in order to estimate effective temperatures, extinctions, luminosities, and hence masses and ages. NIFS will be well-suited to this task (Figure 28); the spectral resolution, spatial resolution, sensitivity, and ability to accurately characterize complex background emission are all essential. Techniques for spectrally classifying early-type stars based on high signal-to-noise ratio R ~ 1000 H band and K band spectra are now in placed (e.g., Ali et al. 1995; Blum et al. 1997; Hanson, Rieke, & Luhman 1998). Higher spectral resolving powers of R > 3000 are required for radial velocity measurements of stars in the immediate vicinity of the central massive black hole. Present AO-corrected imaging (Davidge et al. 1997b) extends to K ~ 16 mag. Early-type main sequence stars are expected at K > 14 mag. High signal-to-noise ratio spectra of these stars can be measured with NIFS in ~ 1 hr per field. High Strehl ratios are essential to separate individual stars in this crowded region. A star with R = 13.9 mag (“star A”) is located 18.8˛ from Sgr A* and can be used as the AO guide star for observations of the central region. A range of near-infrared-bright OIWFS guide stars exist.
Figure 28: K band image of the Galactic center region (Eckart et al. 1995) centered on Sgr A*, the black hole candidate, with the NIFS field-of-view superposed.