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BOZZA BOZZA BOZZZAAAA!

LBT FACILITY INSTRUMENTATION PROPOSAL

Title: Optical and NEar IR Interferometric or Combined Imager and
Spectrograph (ONEIRIC I/S)


PI A. Richichi*, Co-PI M. Gai**
Co-authors:, C. Baffa*, G. Brusa*, A. Cimatti*, C. Del Vecchio*, S.
Esposito*, L. Fini*, S. Gennari*, F. Lisi, R. Maiolino*, F. Mannucci, R.
Ragazzoni***, A. Riccardi*, P. Salinari* (all others to be added in
alphabetical order)
*) Osservatorio Astrofisico di Arcetri
**) Osservatorio Astronomico di Torino
***)Osservatorio Astronomico di Padova

Introduction

More than a single and specific instrument we are here proposing to develop
a family of instruments devoted to exploit the full potential of LBT for
high angular resolution work at high sensitivity. In order to obtain the
ultimate performances a considerable development of the adaptive optics
system is necessary.
We will discuss here only the conceptual grounds on which we base our
proposal. The work to evaluate in greater detail the instrumental options
and the potential cost of the proposed interferometer and of the associated
advanced Adaptive Optics system is still to be done, and will certainly
require a number of years and the involvement of many collaborating
Institutes.
In the following we will try to discuss various options for this ambitious
program, balancing scientific wishes with our present understanding of
technical possibilities, but we are fully aware of the tentative nature of
our present preliminary discussion.
The reasons of presenting this proposal at the present time are the
following:
We believe that in the first decade of next century instruments similar to
the ones we describe here will be able to provide a unique and essential
complement, both in imaging and in spectroscopy, to space born astronomy,
including not only HST but also NGST.
Consequently we think that a considerable part of the LBT instrumentation
budget has to be reserved to interferometry at short wavelength, and the
present Call for Proposal is therefore the appropriate occasion for
discussing a preliminary budget allocation.
A large fraction of the investment, both in terms of money and of work,
required by the high angular resolution instruments consists in developing
an adaptive optics system considerably more advanced than the one already
included in the telescope budget. Various other instruments can take
advantage of this investment on AO, while some of options that can be
implemented at the combined focus could also be part of instrument at
different foci. This is therefore an appropriate time for optimizing the
instrument complement of LBT.
This is not a "private" proposal. We look for the involvement of many of
the groups participating in LBT in a coordinated program aiming to make of
LBT the powerful instrument it can be. The work we propose is challenging
and vast, only a large participation of LBT partners can accomplish it in a
reasonable time. Of course we desire to contribute work in this program in
areas that will have to be agreed.
As mentioned above, the key point of this proposal is that we believe that
a further dramatic step can be done in adaptive optics, pushing this
technique to very high quality correction over a large fraction of the sky
even at optical wavelength. In order to achieve this ambitious objective a
number of techniques have to work together. The first section of this
proposal will therefore deal with the conceptual approach to the advanced
adaptive optics we need. The second section will discuss instrumental
options and priorities, the third will attempt to indicate a possible
configuration for the combined focus and for the instruments. Section 4
reports preliminary ideas on an implementation sequence, section 5 attempts
a census of required expertise areas, and section 6 reports very rough cost
estimates for at least a part of the necessary developments. Section
briefly 7 reports scientific priorities and possible areas of involvement
of the proponents.

1. An advanced AO system

The LBT baseline AO system is already fairly advanced and includes
technology that is still under development, such as Sodium lasers and
adaptive secondary mirrors. D. Sandler et Al. [1] have shown that an
adaptive system such as the one envisaged for LBT, with a single artificial
star and a single corrector, can achieve excellent correction (Strehl ratio
of about 0.5 in the K band) over a large fraction of the sky.
Both, quality of correction and sky coverage drop unfortunately going to
shorter wavelength if only a single artificial star is used for correction,
due to focus anisoplanatism. Although the use of bright natural stars
allows excellent correction even at optical wavelength, the resulting sky
coverage is extremely low. Removing focus anisoplanatism is therefore one
of the key factors for extending AO to shorter wavelength with a
significant sky coverage.
In a paper of several years ago Tallon and Foy [2, TF] have shown that by
using a modest number of artificial stars the focus anisoplanatism effect
can be totally removed. Moreover the Tallon and Foy multi-laser technique
makes also possible to obtain turbulence tomography, separating the
contributions to the instantaneous wave front phase error that are due to
different layers of the atmosphere. Although not yet tested in practice,
for the obvious reason that even single laser AO is still in its infancy,
the TF method is essentially based on geometry and is therefore fairly
solid if laser stars work. Extensive simulations of this method are
presently in preparation in the frame of a European TMR program.
With the TF technique it becomes therefore possible not only to correct
high orders of the wave front error, but also to apply separate corrections
for different atmospheric layers. The consequence is that one can use
"multi-conjugate AO", a technique described by J. Beckers [3], that extends
considerably the isoplanatic angle. The combination of the two techniques
therefore provides the possibility of correcting the wave front not only
without the limit posed by focus anisoplanatism, but also over a much
larger field of view.
The extension of the corrected field is crucial for assuring sufficient sky
coverage at short wavelength. Like a single laser AO system, also multi-
laser systems cannot measure the global wave front inclination, which is
the largest term of the phase error. Although sophisticated techniques have
been proposed to solve this problem, the only reliable way of measuring the
wave front inclination proven until now is that of using a natural star. As
pointed out in [1], the image of a faint star within the field where the
high orders are corrected using the artificial stars is essentially
diffraction limited (most photons fall in a circle of ~ ( /D), except that
the entire image moves around because of the non-compensated wave front
inclination. In this conditions even a modest signal to noise ratio on the
star signal allows to correct the wave front inclination error to the
required level of a fraction of ~ ( /D. A wider corrected field increases
therefore the probability of finding a suitable natural star.
We have outlined only very qualitatively the basic features of the advanced
AO system we aim to, but it is possible to give a very preliminary
evaluation of what we expect from such a system using current technology
and reasonable assumptions on turbulence parameters. If we assume that we
will use nights with Fried coherence length ro ( 20 cm (in the V band),
that the turbulence is evenly distributed between ground and about 10 Km
height and that both correctors have actuator separations corresponding to
~30 cm on the beam, we obtain a Strehl ratios S ~ 0.7 and a corrected field
for finding a natural guide star of ~30 arcsec in the R band. If the
turbulence is mainly concentrated in two layers, close to the conjugates of
the two correctors, the isoplanatic field becomes even larger.
A star of spectral type K or later (about 2/3 of the stars are like that at
faint magnitudes) with mV ( 21 in this isoplanatic field provides enough
photons in about 30 ms to correct tip-tilt to (/5*D or better in the R or I
bands. As the average density of such stars at high galactic latitude is
about 0.6 per sq. arcmin, the sky coverage would be ~ 15 % even at high
galactic latitude, while it would be essentially 100% for the average sky
at lower galactic latitudes.
Table 1 shows the Strehl ratio of the dominant residual error, the "fitting
error", in the above assumptions as calculated from a model of the adaptive
secondary mirror. Fig. 1 shows a section of the corresponding PSF. Although
a significant fraction of the energy is missing from the central peak, the
residual wave front error is on small spatial scales, therefore diffracts
energy on large angular scales so that the contrast between the peak and
the scattered light is extremely high, >103. This is an important feature
for interferometry, where the reduction of the fringe contrast caused by
scattered light is negligible.
Table 1: Strehl ratios corresponding to the dominant error term in a multi-
laser double corrector system, the "fitting error". The assumption is that
each one of the two correctors, both with the same resolution, is used to
correct turbulence with the same r0 (30 cm at (=0.5 (m), corresponding to
a total r0 of 20 cm. The fitting error Strehl ratio of a single corrector
is reported in column 2, while column 3, due to the simplified assumptions,
is simply the square of column 2.
|lambda |Fitting error SR |Fitting error SR |
|[(m] |of one corrector |of both correctors |
| |r0 =30 cm |r0=20 cm |
|0.5 |0.723 |0.523 |
|0.7 |0.846 |0.715 |
|0.9 |0.903 |0.815 |
|1.2 |0.944 |0.891 |
|1.6 |0.968 |0.937 |
|2.2 |0.983 |0.966 |



Figure 1: A cut through the calculated PSF in the I band . The assumed
total r0 is 20 cm at 0.5 (m.



1.1 The Sodium lasers

We will now go through the various components of the multi-laser, multi-
conjugate AO system just outlined to understand qualitatively what we need.
Concerning the laser system we refer to [2] for a description of the
geometry, of the algorithms and of the capabilities of the multi-laser
technique, and we only answer a few simple practical questions:
How many lasers do we need per pupil?
Based on the analysis of Tallon and Foy the minimum number of lasers that
accommodates our preliminary field requirements is four per pupil. The
maximum field angle over which a full correction is possible in this case
is about 50 arcsec, while it would be too small (~7 arcsec) with 3 lasers.
The four lasers would have to be projected at an angle of 83 arcsec from
the telescope axis or from the telescope side. Depending on how the lasers
are projected and on the efficiency of stray light rejection the number of
lasers could become 5. We will assume the conservative number of 5 lasers
in the following.

How powerful lasers do we need?
The requirements on power are somewhat less than for a single laser
systems. Lasers currently under development for single laser systems are
therefore more than adequate.
How will we project the lasers?
The laser beams can be projected exactly like in the single laser system we
have foreseen, from the back of the secondary or from the side of the
pupil. This second option has not been foreseen in the telescope design
and could therefore be more expensive to implement, although it has the
advantage of minimizing the number of lasers and the total amount of light
back-scattered by the atmosphere toward the telescope. Projecting the
lasers from the back of the secondary is simpler in LBT, but may require an
odd number of lasers, therefore not less than 5 in our case, to avoid
superposition of the Rayleigh beacon of one of the lasers with the Sodium
star of the one on the opposite side.
How will we sense the laser stars?
Each artificial star will be analyzed by a "normal" wave front sensor, for
instance a Shack Hartman sensor, with a slightly lower spatial resolution
than the one needed for single-laser systems for the same atmospheric
conditions and correction requirements.
What new technology is needed?
None, if we assume that the laser technology will be available in the next
few years. Certainly the reliability and cost requirements are different
from those of a single laser. The current development at University of
Arizona seems to fit all requirements including moderate unitary cost (~250
k$ per laser).


1.2 The correctors

The proposed scheme of double-conjugated adaptive correction requires two
correctors for each beam, conjugated to suitably different heights in the
atmosphere.

1.2.1 The adaptive secondary
The first corrector is obviously the Gregorian adaptive secondary, whose
conjugated plain lays about 100 m above the primary. The secondary is
undersized to be the pupil, determining an effective telescope aperture of
8.25 m. The actuator spacing will be between 25 and 30 mm, corresponding
to 25 to 30 cm on the wave front, as required for use as single corrector
during initial phases and for other focal stations where double-conjugate
AO will not be implemented. It must be noted, in passing, that the
implementation of a multiple laser star system is of advantage also for the
foci where only a single corrector is used. Overcoming the focus
anisoplanatism problem allows better correction at all wavelengths and
opens the possibility of observing at short wavelength, although with small
field.
The basic technology of the adaptive secondary has been already developed
and is currently adopted to construct the first telescope unit, the MMT
F/15 adaptive secondary. An important feature of this device is that an
internal position loop controls the position of each actuator. It is
therefore possible to operate the corrector by updating "absolute" actuator
positions rather than zeroing the residual error on a star. This is
essential when using multiple laser stars and multi-conjugated correction,
because in this case the control loop cannot be closed directly on a star
in the usual way.
In addition to correct the high orders for low atmospheric layers, the
adaptive secondary will correct the wave front inclination produced by the
entire atmosphere and in the optical path inside the telescope and the
instrument. Using a single corrector for tip-tilt causes negligible de-
correlation of the higher orders. In our configuration, where the
inclination is corrected at the low conjugate, the de-correlation on the
highest layers (H~10 km) is <1 cm, while the value of ro for a high layer
in reasonably good seeing conditions is likely to be more than 30 cm in V.

1.2.2 The Correcting Beam Combiner option
We have preliminarily explored a few optical layouts for the second stage
of correction:
1. a spherical corrector illuminated by the telescope F/15 beam (two or
three added reflections)
2. a double pass off-axis parabolic collimator illuminating a flat
corrector (four added reflections).
3. a corrector coincident with the beam combiner flat mirror (no extra
optics).
Although configuration a) and b) give more freedom in choosing the
conjugate position, the optical performances where unacceptable at a first
attempt, due to the large off-axis angles required in order to avoid
vignetting of the beams. The optical efficiency of the first two
configurations is also not ideal at visible wavelength due to the extra
reflections. Further work could identify better solutions, what counts for
the moment is that at least configuration c) seems to be optically viable,
and, in fact, very interesting. We will therefore continue the present
qualitative discussion of the corrector requirement with reference to the
"correcting combiner" (CC).
Of course placing the second corrector at the beam combining flats
introduces some constraints, but fortunately these don't seem to limit
seriously the performances. If we don't want to introduce extra
reflections, the position, and therefore the size, of the second corrector
are determined by the available space (less than 2 m at the central
combined focus). Preliminary optical design shows that in this case a
simple doublet of ~14 cm diameter can form on the beam combiner flat an
image of a layer located about 6 km above the primary (~9 km above sea
level). This is an excellent conjugate position for observations within
about 40 degrees from zenith, where the height of the conjugate would
become ~4.6 Km (~7.6 Km above sea level).
The on-axis image of the layer formed on the flat combiner mirror has a
diameter of about 80 mm. The doublet has excellent Strehl ratio over about
2x2 arcmin. Of course the beam combiner must be at 45 degrees inclination,
therefore the conjugates of opposite edges of the corrector are at a
different height. The conjugate surface is inclined, going from about 5.5
to about 6.5 km at the opposite edges of the beam. This means that even if
we had only a single turbulence layer exactly at 6 km, there would be a
finite isoplanatic angle due to the change in conjugation height across the
beam. In most practical circumstances this effect is likely to be
negligible.
The CC needs to correct not only on axis but also over some field. Assuming
we want to be able to correct a field of about one arc-minute, provided by
the 4 laser system and certainly useful at least for finding natural guide
stars in favorable atmospheric circumstances, the corrector minor axis
would be about 80+36=116 mm. The actuator density (on sky) should be
similar to the one of the secondary or slightly less dense, say about 30 to
40 cm on the primary, corresponding to about 3-4 mm physical separation
along the minor axis. The separation could be 1.4 times wider along the
major axis. These separations are close to present high density correctors
using piezo actuators, while the total number of actuators, ~1000, is at
the high end of what has already been done. The correction range can be (
10 (m PtV, because the tip-tilt term is corrected at the secondary, and is
again compatible with current piezo actuated correctors. An attractive
alternative to piezo actuated correctors seems to be that of electrostatic
correctors, although currently available only in much smaller formats,
because of the potential for use in vacuum and cryogenical environment and
of the potentially much lower cost.
The only requirement that makes these correctors different from most
currently available devices is that we want to close an internal position
loop for each actuator with nanometer accuracy as done for the adaptive
secondary. The reduced gap (~5 (m) between deformable mirror and back plate
partially compensates the reduced area per actuator, so that using the same
type of capacitive sensors adopted in the adaptive secondary seems to be
possible, in particular for the electrostatic devices. An alternative could
be that of fast optical sensing with a suitable device (fast interferometer
or wave front sensor), using normal reflection directly on the corrector
surface. A slow interferometer is needed in any case for periodical
calibration even if the position sensors are internal to the device.
Although the basic technology for building the CC is (nearly) available,
there is no doubt that developing this component is one of the challenging
aspects of our proposal. The CC has about the same complexity of the two
adaptive secondary mirrors and is likely to be about as expensive. On the
other hand it can provide the best possible optical efficiency, because no
extra optics for the second stage of correction is needed in the beam
combiner.
There is a serious drawback in adopting double-conjugated AO: it is more
difficult to cool down the entire instrument to cryogenical temperature, in
particular because of the second corrector. This is not a real problem for
R, I, J, and H, but limits the achievable sensitivity in K to about 1/2
magnitude less than the theoretical limit for a cold system. The
possibility of cooling the CC to approximately -60 degrees C, sufficient
for recovering the maximum sensitivity in K, has to be examined carefully
in deciding on CC technology.
The CC is not only the second stage of adaptive correction, but must also
compensate optical path differences and pupil geometry variations of the
two interfering telescopes. The first compensation can be performed by
translating rigidly the entire unit, the second by small changes of the
mirror spacing.
Once the system is initially aligned, the optical path difference
fluctuations are mainly due to atmospheric piston changes between the two
telescopes. The typical piston coherence time is about b/3*V, where V is a
typical wind speed, and is therefore of the order of 1 s, while the OPD
amplitude is of the order of ~ 10 (m in good seeing conditions. Translating
even a relatively massive unit at low frequency and with small amplitudes
is not a major problem.
Pupil geometry must be preserved with a relative accuracy of the order of a
few PpM if we want to preserve astrometric accuracy over the entire
interferometric field. The primary mirror spacing is subject to changes of
about 10 PpM per 0C, due to thermal expansion of the steel structure and to
changes due to gravity deformation of the telescope which are of a fraction
of a mm between zenith and horizon pointing. Both can be corrected open
loop.


1.3 Wave front sensing

Although in principle no special developments are necessary for sensing
high and low orders when using a Tallon and Foy multiple star system, we
will briefly discuss here some aspects connected with our proposed scheme.
The separation of the laser beams from the object beam is conveniently
feasible in proximity of the bent Gregorian focus. The laser star spots
could be easily reflected off the main beam by small reflectors, as they
are more than 80 arcsec off axis. Depending on efficiency in the science
bands of Sodium light suppression it might be convenient to reflect off the
main beam all the Sodium light coming from the artificial stars and from
lower atmospheric back scatter.
The unshielded out of focus Sodium light is in fact several orders of
magnitude more intense than the broad band natural sky background in the V
band, and has to be suppressed very efficiently when working with detectors
sensitive at 589 nm and on extremely faint sources. Sodium stray light can
even reduce the sensitivity of the high order wave front sensors, that also
must be shielded, although they can tolerate a much higher stray light
flux.
A convenient option for baffling most of the stray sodium light is offered
by the Gregorian configuration of LBT. A suitable field stop, with
apertures for the science beams and for the laser stars at the primary
focus can remove most of the out of focus Sodium light produced in the low
atmosphere and entering the instrument.
The consequence of separating the laser beams in front of the instrument is
of course that we cannot measure high order wave front error produced
within the instrument with the lasers. We therefore have to take care of
reducing as much as possible "internal seeing" in the instrument, avoiding
the formation of turbulence on scales comparable or smaller than the beam
cross-section (~100 mm). An ideal system would be that of evacuating the
entire instrument. This would certainly be the preferred option if even a
moderate cooling of the CC is possible. A simpler alternative is to enclose
the instrument in a sealed thermal protection and fill it with Helium at
atmospheric pressure. The ~ 12 cm (minor axis) laser beam splitter would
act in both cases as instrument windows, but without differential pressure
if the system is filled with Helium. Within the Helium filled instrument a
temperature uniformity of a fraction of degree C has to be achieved to
effectively prevent "instrument seeing".
The low order sensing must be done after the second stage of correction to
profit of the wider corrected field. It is difficult to imagine a metrology
system with the same level of completeness and reliability as the beam from
a natural star that goes through most of the optical system with a full
beam. Sensing the low orders as close as possible to the final images is
therefore a good way of correcting errors due to mechanical flexures,
thermal expansion and refraction index variations during observations. Of
course a "system test mode" with higher order sensing on brighter stars can
be used several times, if needed between observations.
We need to sense at least five quantities, tip and tilt for the individual
beams and differential piston. Whether we need to sense other low order
errors is a crucial question that will have a clear answer only after a
sufficient experience with the use of single laser AO systems. The main
question concerns the focus term. The artificial stars are produced by
resonant scattering in a layer that is > 10 km thick, therefore any change
in the Sodium height distribution in the layer affects the focus of the
laser stars. In a binocular telescope like LBT the Sodium vertical
distribution can in principle be monitored continuously without the need of
extra stars or extra telescopes, because the laser stars projected from one
of the telescopes are significantly elongated (several arcsec) when seen
from the other. As the laser spot are separated by typically 40 m on the
Sodium layer, even small variations of Sodium equivalent height on smaller
spatial scales might force to use a natural star for focus correction. This
of course has a cost in photons and therefore in sky coverage.
We want to cover a wide range of wavelengths (0.6 to 2.5 (m) where we are
likely to use two different types of detectors (Silicon CCD and, for
instance, HgCdTe detectors) to optimize the detection efficiency, dark
current and readout noise. In each detector band we only have two bands
that we can use for sensing (R and I for silicon detectors, J and H for the
neat IR because the K band has a somewhat reduced sensitivity due to the
(probably) warm optics. We therefore should arrange the low order sensors
to share one of the two available bands in each range while the other one
is used for the astronomical measure. We expect that only 5 parameters will
have to be measured, but even if the number is 7 or even more, this can be
done in various ways using modified versions of curvature, Shack Hartman
and Foucault sensors.
The four degrees of freedom associated with tip-tilt are more difficult to
sense than differential piston because they change at higher frequency
(roughly by a factor of two) and because they only use the photons from one
of the beams. The focus term, if not provided by the laser stars for the
reasons discussed above, is somewhat less demanding than tip-tilt because
it is only one degree of freedom and it has less power, even if its rate of
change could be somewhat higher than that of tip tilt. The previous
considerations on sky coverage wouldn't therefore be severely affected by
the need of measuring focus.


2. Instrumental options

The type of interferometry that can be done with LBT is unique both at
short and at long wavelengths, and the first question to answer is
therefore about the spectral coverage. We are convinced that both spectral
regions (NIR, with extension to I and R, and, if possible V, and the Middle
IR) must have high priority, but it is extremely difficult to combine the
entire range of wavelengths in a single instrument. The question is
therefore how to divide the spectral range in an optimum way, taking into
account efficiency, cost and risk.
There are a number of technical reasons for dividing the spectral range
such as cooling requirements, needs of adaptive correction of the wave
front, types of detectors. Fortunately these criteria give approximately
the same answer: the dividing line is near to 2 (m. We believe that the
best way of splitting the above wide range could be that of having an
instrument working at ( < 2.5 (m and another one working at ( > 2 (m. The
K band would then be covered by both interferometers, but only the
(necessarily) cryogenically cooled Long Wavelength Interferometer (LWI)
could be certainly optimized for K, because the optics of the short
wavelength instrument (SWI) is likely to operate at a temperature not too
far from ambient. The possibility of a moderate cooling of the SWI, say to
~220 K, that could bring the K band to its limiting sensitivity has to be
examined at a later stage of design.
The two instruments are highly complementary for many fields of
astronomical research and should both be available as common user
instruments in parallel, therefore will have to be designed for two
different focal stations. The natural division, dictated by the very
different sensitivity to differential retardation induced by the tertiary
and beam combiner flats, is to have the LWS at the rear beam combined focus
and the SWI at the central one. Both instruments need adaptive optics,
although at a different level of complexity, and both must have operational
modes that could produce significant science even with reduced or missing
adaptive correction. LWS could be forced in this mode by a delay of the
baseline adaptive optics system, while SWI should work, at the beginning,
without the advanced adaptive correction that will be discussed in the
following.
Assuming the above division of wavelength ranges between the two
Interferometers we now go through a list of options that could be
considered for the SWI

2.1 Imaging modes of the SWI

Interferometric imaging at the highest angular resolution and sensitivity
in each band is clearly the prime scope of the LBT Interferometers. This is
not a simple task, especially at short wavelength, but is one that could
make LBT unique not only in comparison with other ground based telescopes
of similar size but also in comparison with HST and NGST. In particular at
optical wavelengths the atmospheric background is not dramatically higher
than in space, and the possibility of competing with space borne instrument
is based essentially on the quality of the adaptive correction of the
effects of the atmospheric turbulence. This is the main reason why we
believe that achieving excellent correction over at least part of the
optical domain is crucial for the scientific excellence of this instrument.
In the previous section we reported our considerations on how to obtain the
desired adaptive correction. Here we will describe what we wish to obtain
in imaging.

2.1.1 Interferometric Imaging
Table 2 reports single pupil and interferometric angular resolution,
angular pixel sizes and F/n of the individual beams, based on assumed pixel
sizes. We expect we could correct a field of approximately 30 arcsec
diameter with high Strehl ratio in the R band. A look at Table 2 shows that
we need about 20,000 by 20,000 pixels to cover that field in R in
interferometric mode, if we sample at the minimum possible sampling to
avoid loss of fringe contrast, ~ 4 pixels per fringe.



Table 2: Angular resolution, angular pixel size, field and beam numerical
aperture for single pupil and for interferometric imaging. The telescope
effective diameter D is 8.25 m and the max baseline b is 22.65 m. All F/n
refer to single pupil beams.
|Band |Pixel |N of |(/D |(/2D |(/2D |(/2D |(/b |(/4b| ( |( /4b|
| |size |Pix | | |Field|F/n | | |/4b | |
| |(m | |mas |mas | | |mas | |Field|F/n |
| | | | | |arcse| | |mas | | |
| | | | | |c | | | |arcse| |
| | | | | | | | | |c | |
|R |13 |40002 |17.5|8.7 |35 |37.1 |6.4 |1.6 |6.4 |203.9|
|I |13 |40002 |22.5|11.2 |45 |28.9 |8.2 |2.0 |8 |158.6|
|J |18 |20002 |31.2|15.6 |31 |20.8 |11.4|2.8 |5.6 |114.2|
|H |18 |20002 |40.0|20.0 |40 |16.3 |14.6|3.6 |7.2 |89.2 |
|K |18 |20002 |54.9|27.5 |55 |11.8 |20.0|5.0 |10 |64.9 |



The number of pixels necessary to cover the well corrected interferometric
field in the other bands is of the same order, due to the increase of the
corrected field at longer wavelengths. The detectors are therefore likely
to be a significant part of the instrument cost, although it is certainly
possible to start with a single chip and a reduced interferometric field,
as reported in Table 2. A single InSb array could in principle cover the
entire range R-K, but this solution is not advisable for various reasons
(first of all sensitivity to thermal background). We consider as baseline
choice that of (at least) one 4000x4000 CCD for R and I and of a 2000x2000
HgCdTe detector for J, H and K. Detectors of this size can fully exploit
the field available at (/D resolution and provide a much larger
interferometric field tha available at any other large telescope
interferometer.
Excellent, but not impossible, performances in readout noise and dark
current are necessary in R and I, where, as shown in Table 3,
interferometric imaging could be easily limited by detector performances.
In J, H and K the background is higher and the detectors, using multiple
readout during integration, can more easily reach the required
performances. In the K band we have assumed a background ~0.8 mag brighter
than the sky background to account for the contribution of ambient
temperature optics.











Table 3: Noise equivalent magnitude per pixel for the case of
interferometric imaging. Assumptions are: 1000 s of integration time, 30%
transmission for both, background and signal photons.
|Ban|Sky |angul|BG |BG |R/O |BG+R/O|BG |R/O |Noise |
|d |bright|ar |flux | |Noise | |Noise |Noise |equivalent|
| |n. |pixel| |Noi| |noise |equival|equival|mag |
| | | | |se | | |ent |ent | |
| | |size | | | | |mag |mag | |
| |mag/sq|mas |ph |ph |ph |ph |mag/pix|mag/pix|mag/pix |
| |-arcse| | | | | | | | |
| |c | | | | | | | | |
|R |20.80 |1.6 |1.0E+|3.2|2 |3.80 |35.3 |35.8 |35.1 |
| | | |01 | | | | | | |
|I |19.30 |2.0 |6.9E+|8.3|2 |8.52 |33.8 |35.4 |33.8 |
| | | |01 | | | | | | |
|J |15.00 |2.8 |6.9E+|83.|10 |83.91 |31.0 |33.3 |31.0 |
| | | |03 |3 | | | | | |
|H |13.00 |3.6 |7.2E+|267|10 |268.04|29.0 |32.5 |29.0 |
| | | |04 |.8 | | | | | |
|K |12.00 |5.0 |3.4E+|583|10 |583.79|27.9 |32.3 |27.9 |
| | | |05 |.7 | | | | | |



The sensitivity achievable on unresolved or barely resolved sources is
affected by many parameters, first of all the quality of the adaptive
correction and of the co-phasing.
In case of perfect optics, correction and co-phasing, ~80% of the photons
from an unresolved source would fall within the first zero of the Airy
function of a single pupil, corresponding to about 550 pixels at the
sampling adopted in Table 3. About one half of this light is concentrated
on about 100 pixels. An average signal to noise ratio of 10 in the fringes
could therefore be obtained in only 1000 seconds on point sources about 6.5
mag brighter than the noise equivalent magnitudes (per pixel) reported in
Table 3. At this level of S/N the interference pattern would be clear
enough to allow PSF de-convolution algorithms to derive full two-
dimensional resolution and better S/N from multiple exposures.
In 1000 seconds the PSF rotation is typically of 4 to 10 degrees, and the
loss of angular resolution is well acceptable. About a night of effective
integration is therefore necessary (with our theoretical assumptions) to
obtain a full angular resolution measurement of a point source of magnitude
~30 in the R band. The signal to noise degradation due to non perfect
adaptive correction and co-phasing is essentially proportional to the
Strehl ratio at high SR values. A global SR of ( 0.4 is possible on the
entire wavelength range considered, as shown in the previous section, i.e.,
less than one magnitude loss in limiting magnitude compared with the above
theoretical values. Specialized data reduction techniques can further
increase the sensitivity in special cases.
It is clear that the about six magnitudes of difference in sensitivity and
the factor of about 2 in angular resolution between R and H justify the
effort of pushing the interferometer to the shortest possible wavelength.
Red-shifted ultraviolet and optical features of high z galaxies are
certainly a prime candidate for interferometric observations in R and I.
The sky coverage in interferometric imaging is not as good as in (/D
imaging, because the accuracy required in the low order correction is
correspondingly higher. Still, a sky coverage of a few to several percent
at high Galactic latitude is to be expected, more than sufficient to obtain
interferometric images of large samples of essentially any type of sources.
The sensitivity achievable on very well resolved sources, that can be
considered to be uniform on scales much larger than the fringe separation,
is only depending on source extension, because in the absence of features
on an angular scale of ~ ( /b one can bin many pixels. Table 3 shows that
even in the R band we can operate in almost background limited conditions,
therefore the binned data are only slightly noisier than data obtained with
pixels of larger angular extension. A uniform source at S/N = 1/10 (per
pixel) would be clearly detected at several ( if it fills about 2000 pixels
(this corresponds to about twice the Airy disk in diameter, ~85 mas in R.
The corresponding surface brightness would be, in R, about 23rd mag/sq.
arcsec). By processing in a different way the same data obtained for
highest angular resolution one can then recover information on larger
scales at lower surface brightness.
A recent deep R image obtained by VLT with excellent seeing (0.37 arcsec
FWHM) in 75 minutes of effective integration reaches a 3( limit at mR 28
for point objects. Essentially all the objects in that field could
therefore be resolved with 6 mas resolution in the same band by the LBT
interferometer if, in real life, we will not loose more than a magnitude
with respect to the above theoretical limits. The 60 times higher
resolution would provide general morphology and identification of other
features such as star formation sites, supernovae, large globular clusters.
In a similar way the majority of the objects appearing in the Hubble Deep
Field could be imaged at about 10 times higher resolution.

2.1.2 Phased but non Co-phased imaging
This is the basic mode with an adaptively corrected single pupil, and can
be obtained in the combined focal plane simply not applying the phase
correction, or physically displacing apart the two images on the same
detector or on different detectors. In all cases the sensitivity gain is
about the same that can be obtained by degrading of the same factor the
angular resolution in interferometric imaging by binning the pixels. The
difference is of course that one can cover a much larger field in this mode
if the detector is matched to a lower angular resolution.
Table 4 shows the noise equivalent magnitudes for both images that can be
formed by LBT when both pupils are phased but not co-phased. Summing the
photons by stacking the two images on top of each other on a single
detector or summing the digital signals from two separated images provides,
in photon noise limited observations, the same gain of 21/2 in signal to
noise ratio compared with single pupil imaging. About 80% of the photons of
a perfectly corrected point source image would fall within the first null
of a single pupil Airy pattern, ~ 18 pixels with the sampling adopted in
Table 4. About one half of the photons will fall on 4 pixels. A > 10(
detection could then be obtained in 1000 seconds on point sources about 4
magnitudes brighter than the Noise equivalent magnitudes reported in Table
4, with a sensitivity gain of about a factor of 10 compared to
interferometric imaging but, of course, with reduced angular resolution.
The (/D imaging mode with its "large" field and faint limiting magnitudes
should therefore allow LBT to obtain in a single night data over a field
wider (or deeper, if of the same size) than the Hubble deep field at
higher angular resolution in R or I. In a few nights one can obtain multi-
band photometry over the same field. Selected objects can then be observed
interferometrically to obtain even higher resolution.
Table 4: Noise equivalent magnitude per pixel for the case of single pupil
diffraction limited imaging ((/D). Assumptions are the same as for Table 2
and 3. The background is calculated for two telescopes added incoherently.
|Ban|Sky |angul|BG |BG |R/O|BG+R/O |BG |R/O |Noise |
|d | |ar |flux | |Noi|noise |Noise |Noise |equiva|
| | |pixel| |Noi|se | |equiva|equival|lent |
| | | | |se | | |lent |ent |mag |
| | |size | | | | |mag |mag | |
| |mag/ |mas |ph |ph |ph |ph |mag/pi|mag/pix|mag/pi|
| |sqarc| | | | | |x | |x |
| |sec | | | | | | | | |
|R |20.80|8.7 |7.8E+0|8.9|2 |12.7 |33.8 |35.8 |33.8 |
| | | |1 | | | | | | |
|I |19.30|11.2 |3.4E+0|18.|2 |26.3 |32.6 |35.4 |32.6 |
| | | |2 |5 | | | | | |
|J |15.00|15.6 |2.6E+0|160|10 |227.3 |29.9 |33.3 |29.9 |
| | | |4 |.6 | | | | | |
|H |13.00|20.0 |1.3E+0|358|10 |506.9 |28.3 |32.5 |28.3 |
| | | |5 |.4 | | | | | |
|K |12.00|27.5 |4.9E+0|702|10 |994.0 |27.3 |32.3 |27.3 |
| | | |5 |.8 | | | | | |



2.1.3 Non phased imaging.
This is only another name for Speckle Interferometry. A variety of
techniques allow to recover high angular resolution information and some of
them can reconstruct complex images. The common requirement is that of
freezing the effect of the atmospheric turbulence by using very short
integration times. It is clear from table 2 and 3 that in very short
integrations (<0.1 s) the detector noise determines the sensitivity, that
is therefore greatly reduced. On the other end these techniques can be used
in a total absence of adaptive correction or with partial adaptive
correction and no co-phasing. This makes of Speckle interferometry a very
valuable tool for early phases of the telescope life, when only moderate
adaptive correction will be available, and for later phases in areas of the
sky where natural guide stars are not available. In this case one can take
full advantage of the high order correction and this "single speckle" case
becomes another way of calling the "shift and add" method.
From a technical point of view the only specific requirement different from
those of the other imaging modes is that of a fast readout of the
detectors. This is easily obtainable with existing IR detectors, where only
a subset of the detector can be read at correspondingly higher rate, while
is not yet common on optical CCD detectors, but optical detectors with
addressable readout, similar to those uses in the NIR, are currently in
development.


2.2 Spectroscopic modes

Although images are of great importance in understanding astrophysical
phenomena, a quantitative study normally requires spectral information.
This is in principle obtainable at the same angular resolution as
interferometric images, but of course at a significant cost in sensitivity.
Nevertheless for sufficiently bright sources this provides a very rich and
detailed information. For faint resolved sources it is convenient to
increase the slit width to collect more photons, because it is clear from
Table 3 and 4 that even at a modest spectral resolution the sky background
is not going to be the dominant source of noise except in the K band. In
fact Tables 3 and 4 cannot be scaled easily with respect to spectral
resolution, because the background is largely concentrated in strong sky
emission lines, and it is therefore much lower over most of the spectrum
than average values scaled from the tables. Detector readout noise, dark
current and cosmic ray events will determine the optimum integration time
and sensitivity in most cases. We will not go through all possible
combinations of angular and spectral resolution at the various wavelengths,
but only discuss a few basic facts that are common to the various
combinations.

2.2.1 Interferometric "long slit spectroscopy"
We discuss here a configuration in which a slit is placed in the
interferometric image plane, and re-imaged after dispersion on the
detector. The slit should be parallel to the fringe modulation direction,
so that a point source originates three parallel spectra, one for the
central PSF lobe and two satellite spectra on opposite sides for each of
the secondary lobes. In this way the angular resolution can be maintained
along the slit and essentially the same de-convolution algorithms used for
two dimensional imaging can be used for this mono-dimensional case. The
slit width is ( /D but, if we want to preserve a full angular resolution,
the pixel matching must remain that of interferometric imaging in the cross-
dispersion direction, ~( /4b. Pixel binning or anamorphic optics or a
combination of both can be considered to increase the sensitivity of this
peculiar spectroscopic mode that could be very interesting for specific
applications. It is worth noting that even without special tricks the
continuum of a point source about 12 magnitudes brighter than the column of
Table 3 reporting R/O noises per pixel could be detected in 1000 seconds at
a resolving power of ~ 1000 and S/n ~10, therefore in the range of
magnitude ~19 to 23 (from K to R).
The entire spectroscopic end of the instrument, from the slit to the
detector, has to be de-rotated, and again the maximum duration of each
integration is determined by PSF rotation in a similar way as in
interferometric imaging.

2.2.2 Phased spectroscopy with OH suppression.
This configuration differs from the previous one because the angular
resolution is now reduced to be of the order of ( /D or even significantly
less. There is therefore no need of co-phasing the pupils but only of
adaptive correction. The optical scheme and the de-rotation requirements
are similar to the previous case, but, of course, the re-imaging optics
needs to be considerably faster.
After post-detection "OH Suppression" (this here means more generally
removing strong atmospheric lines, independently of their origin) most of
the remaining portions of the spectrum are affected by very low background
and are detector noise limited even with a relatively large slit width,
because the remaining continuum background is dispersed. A spectral
resolving power of >2000 should be used to effectively remove the sky lines
in the Near IR, but post detection binning can then be used very
effectively to increase sensitivity, although with reduced resolving power.
The K band will only moderately profit of the OH suppression although a
much more considerable gain is to be expected in the K' band (2 to 2.3 (m).

Over most of the spectral range considered here detector noise limited
sensitivities should be obtained for slit widths of the order of 100 mas.
An example of this mode is reported in Table 5. Here it must be noted that
the effective resolving power is ~500, one half of the value reported in
the third column that refers to the spectral coverage of a single pixel.
The values of the background suppression factors, in the fifth column, are
only very crude estimates. The last column gives the magnitude of a source
filling a pixel whose continuum could be detected with S/N =1 in the usual
assumptions for integration time (1000 s) efficiency etc.
Assuming that the source fills the slit (2 pixels wide) and considering the
effective resolution (2 pix per resolution element at resolving power 500),
the detection of the continuum of a source with a surface brightness equal
to the one reported in column 5 of Table 5 would be obtained, at 3 (, in
only about 2000 s. No need to say that detecting emission lines is much
easier.
With longer integration times (several hours) one could therefore detect
the continuum of compact sources (~ 0.1 arcsec) at about mag 30 in R and I.
Even more spectacular are the performances in the NIR, due to the OH
suppression. This spectroscopic mode would then give access to spectroscopy
of essentially all the sources in the Hubble Deep Field, where, by the way,
there seems to be a cutoff right at angular sizes of about 0.1 arcsec.
Table 5: Example of a "wide slit" combined, but not co-phased,
spectroscopic configuration. The last column reports the magnitude,
integrated over the selected pixel, detected at S/N =1 in the continuum.
Both, background and signal are calculated with the full telescope aperture
(two mirrors). Other parameters are as in Table 3.
|Band |angula|(/(( |averag|BG |BG |R/O |BG+R/O |Noise |
| |r | |e BG |suppress|nois|noise|noise |equival|
| |pixel | | |ion |e | | |ent mag|
| |size | | |factor | | | | |
| |mas | |ph | |ph |ph |ph | |
|R |50.0 |1000.0|8.6 |2.0 |2.93|2 |3.54 |29.0 |
|I |50.0 |1000.0|10.1 |5.0 |4.48|2 |4.91 |28.4 |
|J |50.0 |1000.0|10.7 |50.0 |4.63|10 |11.02 |27.3 |
|H |50.0 |1000.0|13.7 |50.0 |5.24|10 |11.29 |26.7 |
|K |50.0 |1000.0|9.7 |2.0 |4.42|10 |10.93 |26.1 |

The sky coverage of this "long and wide slit" mode is also likely to be
significantly larger than that of the various diffraction limited imaging
and spectroscopy modes. The requirements on the natural star used for the
correction of low orders are in fact drastically reduced. Instead of
correcting, for instance, tip-tilt to a small fraction of ( /D, in this
case we can afford correcting to a few ( /D. As a well corrected field of
about one arc-minute is to be expected in the NIR, a multi-slit mode can be
considered.
This deep spectroscopic mode can therefore be used to obtain a large number
of spectra (up to the order of 102 per night in multi-slit mode) over deep
fields obtained in (/D imaging mode, or to obtain spectroscopy of extremely
faint objects (mR~30) in long integrations.

2.2.3 Very high resolution stellar spectroscopy
A very large telescope with very good adaptive correction in principle is
an ideal instrument for extremely high resolution work. The main advantage
of the adaptive correction is that a fairly compact spectrograph can
provide an extremely high resolving power on a nearly diffraction limited
slit. Moreover in most cases the source itself can be used for low order
correction, so that the sky coverage is essentially complete. We didn't
even attempt a conceptual design of such an instrument, which differs
significantly from the other low dispersion spectrographs discussed. On the
other hand an optical layout compatible with the proposed configuration of
the corrected-combined focus doesn't seem to be impossible even if cross
dispersion is needed.


2.3. Other fancy stuff

We spend here only a few words on more specialized instrumental modes that
could greatly profit of the advanced correction and/or of the binocular
nature of the telescope and that seem to be compatible with the general
layout we assumed for the imaging and spectroscopic modes.

2.3.1 Nulling Interferometry
This mode is likely to be implemented in the LWI specifically to study the
zodiacal light emission in the thermal IR of nearby stars, an important
parameter for planning future space missions devoted to the search of earth-
like planets. In principle there is no problem in implementing nulling
interferometry at shorter wavelength in the SWI. The nulled field would be
reduced in proportion, while the light "leaks" would be larger for various
reasons, but still this way of implementing a coronographic mode is likely
to be much more effective than conventional coronography. The
implementation of this mode is conceptually simple, because in principle it
is sufficient to pick up the two beams after the combiner/corrector or even
after the field re-imaging optics and make them cross at the position of a
beam-splitter.

2.3.2 Polarimetry
Although a non negligible amount of linear polarization is certainly
introduced by the two reflections at 45 degrees on the tertiary and on the
combiner, there are various ways of obtaining accurate linear polarization
measurements at high angular resolution in all imaging modes. In non co-
phased imaging modes a fixed bi-refringent prism can be used, for instance,
in the converging beam and the telescope rotation could be used to explore
90 degrees. The signal of a field star (we always have in the isoplanatic
field at least a "bright" one for adaptive low order correction) could be
used to remove transmission variations.
In the interferometric (co-phased) imaging mode the double modulation
(fringe rotation and polarization angle) might be difficult to disentangle
and it could be better to rotate the polarizer in a number of positions for
each interferometric sub-integration. Several images at different polarizer
angles would then be reconstructed and compared.
Similar consideration can also be extended to spectro-polarimetry. Although
these considerations are extremely rough, we want to stress the importance
of considering a high angular resolution polarimeter option


3. Scientific Objectives



4. A conceptual layout

We have considered a variety of observing modes at the combined LBT focus
and an advanced adaptive optics system. The problem is now to find a way of
satisfying the contrasting requirements of the various components for all,
or at least for most, of the observing modes we described. The basic
requirements emerging from the above discussion of observing modes and of
adaptive requirements are the following:
1. We need to form an image of a high atmospheric layer on a deformable
mirror
2. We need to image the field on the detectors at very different scales
(from about F/6 to about F/200)
3. We need to have an intermediate image of the field for spectroscopy
(diffraction limited slits diffract!)
4. We need to find appropriate positions for two beam splitters (for
laser and natural star wave front sensing)
5. We need to de-rotate detectors and spectrographs
6. We have to build up the instrumentation we discussed gradually,
starting from the simplest configurations.

Work done in the past on a similar interferometric configuration by P.
Bayard and D. Bonaccini [4] has shown that relatively simple transmission
optics (spherical doublets) can provide a very high Strehl ratio over a
wide field and a very extended wavelength range, provided re-focussing is
foreseen for the different bands. Although very little optical design has
been done yet to verify the actual feasibility of this instrument, we have
adopted a layout entirely based on transmission optics, sketched in Fig.2,
because we expect it to represent a viable solution. Of course different
alternatives will have to be explored, especially for the field re-imaging
optics
Figure 2: A sketch of the combined focus layout.
In the scheme represented in Fig. 2 the beams emerging from the beam
combiner act as a new "a-focal station" of the telescope, where the beams
are already adaptively corrected and homothetically oriented, but not yet
really combined. The list of modes of operation we have presented in
section 2 could be considered as a list of possible instruments for this
new focal station.
Each instrument can therefore be optimized indipendently after the CC. In
fact it is likely that even the collimator could be changed to optimize
optical quality and transmission, depending on wavelength and application.
The complex and expensive CC, the laser star sensors, the optical support
structure and the "instrument de-rotator" are the components that need to
be designed for common use of all the instruments. In Fig. 2. the de-
rotator is placed not far from the pupil image to provide the option of
intercepting the parallel beams in the "instrument" at the most convenient
position.
In most cases the low order sensing is done at the same plate-scale as the
science object, and therefore no much other optics is needed if the
instrument optics has a sufficiently wide field. In some cases, a typical
example is that of Phased Spectroscopy with OH suppression, the image scale
is too fast for low order sensing and a separate optics is needed for this
purpose. This can either be arranged in the Instrument or after the beam
combiners in the parallel beam by inserting there appropriate beam-
splitters (no co-phasing is required in this case).
Several important components are missing from Fig.3, in particular ADC's,
filters, shutters, baffles. They could be placed in various parts of the
"combining" optics or of the "instrument". The instrument normally includes
the optics, the detectors and the appropriate low order wave front sensors
with their positioning devices, as these depend strongly on the band and on
the application. A conceptual sketch of one of the instruments, the "(/D
spectrograph", is reported in Fig. 2 only to illustrate the above
discussion.
We only made a preliminary optical check on the collimator to understand
its imaging quality for the high atmospheric layer to be imaged on the CC.
At the first attempt and without any optimization the collimator gave very
good results, summarized in table 5. The design wavelength range was
initially 1 to 1.5 (m, but the same collimator gives similar performances
in J, H and K provided a slight refocus is admitted.
Table 7: Results of preliminary optical analysis of a doublet collimator
|Parameter |Value |
|Type of system |Spherical doublet|
|Material |IRg2/BaF2 |
|Focal length |1.2 m |
|Diameter (field 1'x1') |140 mm |
|(( |1 to1.5 (m |
|Pupil Image Diameter |80 mm |
|6 Km image Diameter (+1' |80 (+20) mm |
|field) | |
|Field at SR>0.95 |1.0 arc-min |
|80% encircled energy diameter |< 0.3 mm |
|at CC | |
|90% encircled energy diameter |< 30 (m |
|at pupil image | |

The image quality at the CC is adequate if we consider that the corrector
has elements of size ~ 3x5 mm. The collimator described in Table 7 is the
one reported in Fig. 2.
The two pupil images, 80 mm in diameter, are separated by about 60 mm to
preserve pupil geometry. If we allow for about 1 arcmin field, the field re-
imaging optics (presumably another doublet, or a triplet, that we will call
"camera" in the following) would be ~ 280 mm X 120 mm if placed in the
pupil image plane. The camera doublet represented in Fig. 2 is only for
illustration of the concept and has not been designed, as a number of
different cameras are likely to be needed for different instruments.


5. Steps to the goal

The implementation of the above capabilities must take into account a
number of constraints including the fact that in an initial period of about
two years only one of the primary mirrors will be available and the various
stages of the evolution of the adaptive system. Table 7 summarizes the
instrumental options available at different stages of adaptive optics. Each
cell reports the bands in which useful work can be done and the telescope
configuration that can be used.

Table 7: Bands in which the main observing modes can be used at different
stages of the development of the adaptive system. Of course there are
significant differences in terms of accuracy, sensitivity and sky coverage
for the same band and mode depending on adaptive configuration. "Ngs" and
"lgs" stand, respectively, for laser and natural guide star, "corr" stands
for corrector. The numbers in brackets refer to the telescope
configuration: 1 stands for single primary, 2 for the binocular
configuration. See text for further comments.
| |No AO |1 corr, |1 corr,1 |1 corr, 4|2 corr, 4|
| | |ngs |lgs |lgs |lgs |
|Speckle |K to R |K to R |K to R |K to R |K to R |
|imaging |(1,2) |(1,2) |(1,2) |(1,2) |(1,2) |
|(/D imaging | |K to R |K to J |K to R |K to R |
| | |(1,2) |(1,2) |(1,2) |(1,2) |
|(( /D | |K to R |K to J |K to R |K to |
|spectroscopy | |(1,2) |(1,2) |(1,2) |R(1,2) |
|( /b imaging | |K to R |K to J |K to R |K to R |
| | |(2) |(2) |(2) |(2) |
|VHR | |K to R |K to J |K to R |K to R |
|Spectroscopy | |(1,2) |(1.2) |(1,2) |(1,2) |

Apart from the obvious need of a binocular configuration for all
interferometric work, all the other instruments can work in the NIR (and in
some cases also in the "optical" range), although at reduced performances
and sky coverage, with a single mirror and in essentially all the adaptive
configurations. This means that we can adopt a step by step evolutionary
scheme that reduces both risk and cost. The possible steps are the
following:
1. Soon after first light we can build a combined station (with at least
one collimator, at least one of the non-adaptive combiner flats and
the camera), the instrument de-rotator and one or two instruments (NIR
( /D imager and spectrograph). This system can be used with the
baseline adaptive system and with bright natural stars and can be used
with more advanced AO without significant changes, except updates of
the wave front sensors and completion of the second arm (collimator
and combiner) when the second primary is installed.
2. After introduction and successful test of the multiple laser guide
stars with the NIR combined instruments, a new set of "( /D"
instruments working at shorter wavelength can be installed and used
because focus anisoplanatism can now be removed. As we are still using
non-coherent instruments, this step could be done either with a single
pupil or with a double pupil telescope.
3. After introducing multiple lgs we can introduce double conjugate
correction by replacing the beam combiner flats with deformable
mirrors. This greatly enhances the capabilities of the already built
instrumentation, and also represents the final development necessary
for optimum interferometry if the telescope is already binocular.
In each of the above steps one has the possibility of producing valuable
science and of testing subsystems that are needed for the further step. In
step 1 the laser system is tested before multiplying the number of lasers,
in step 2 the multiple laser system can be used for eliminating focus
anisoplanatism problems, but will provide real data to be used for
optimization of the double conjugate scheme. In a similar way, as soon as
the second telescope becomes operative one can start testing and debugging
the interferometric mode.
Clearly the above program overlaps in part with the program for other
instruments, in particular with the diffraction limited modes of operation
of the Near IR Spectrograph/camera. On the other end this can be used to
divide functions among instruments, making all of them simpler.
Although the precise time sequence is difficult to define at the present
pre-conceptual stage, the goal of having the entire multi-laser/double-
conjugate/multi-instrument system available and working soon after the
installation of the second mirror is not in principle impossible. This will
depend from a good coordination of the activities more than from technical
or financial problems.


6. A preliminary study phase

Even a preliminary feasibility study of the many aspects of this proposal
represents a massive work that requires activity of different groups with
different expertise. We simply list here in random order a number of
activities that, at a preliminary study level, are sufficiently independent
from each other to make it possible to divide them among different groups.
1. Characterization of the turbulence parameters at Mt. Graham. The main
goal is to obtain some statistics around the year on the vertical
distribution of the turbulence.
2. Exploratory optical design of the entire system including a number of
instruments. The scope should be, in this phase, to check the
compatibility of the parameters of the combined focus (or of the a-
focal combination) with those of the various instruments.
3. Comparison of various techniques for wave front sensing for low
(includes piston) and for high orders
4. Simulations of multi-laser, double-conjugate systems.
5. Optical model of the entire interferometer and interferometric
tolerance analysis.
6. Exploratory mechanical (and thermal) design of various options (vacuum
at ambient or low temperature, Helium at controlled temperature) for
the combined focal station.
7. Preliminary study of the CC.
8. Preliminary study of algorithms for best extraction of information
from the various modes of operation.
9. Study of Sodium light notch filters and reflectors and, in general, of
stray light suppression.
The above list is certainly not complete and doesn't include activities
that are already going on such as the development of the lasers and that of
the adaptive secondary mirrors.
It is only after the completion of this study phase, possibly in about one
year, that a well defined plan can be formulated and precise
responsibilities assigned for the next phase of design and construction of
the various sub-systems and instruments.


7. How much will it cost

There is certainly no precise answer to this question at this time. Still
it is useful to attempt to define reasonable goals.
1. The Lasers: the current development at UoA is aimed to produce Sodium
lasers at 250 K$ each. Taking into account the large number we need
(say 10 in total, but 2 are already in the telescope baseline), that
the cost is likely to go down with time and that the projection optics
is already in the baseline AO system, we can assume that the cost of
the upgrade from 2 to 10 lasers will be < 2 M$.
2. The Correcting Combiner: It cannot be more expensive than the two
adaptive secondary mirrors, because the optics is much easier to
produce and the mechanics is much simpler. The electronics could be
essentially the same, while the actuators (most likely piezos) could
be somewhat more expensive. A cost of about 3/4 of that of the two
secondary mirrors, therefore ~1.5 M$, seems to be a reasonable
estimate.
3. Optics and Mechanics of the Combined Focus: the rest of the hardware
necessary for beam combination, including the 8 laser wave front
sensors and their detectors, the collimators, one of the cameras and
the instrument de-rotator is likely to cost about 1 M$.
4. Instruments: these can become very expensive if one wants to exploit
the enormous field we can provide in principle. For (relatively) small
field instruments based on a single large chip the average cost of the
instrument can be low, considering the relatively simple optics and
mechanics. Including the low order wave front sensors an average cost
of about 1/2 M$ per instrument (here this means per function, or, per
mode, as more than one of the modes of operation could be implemented
in a single instrument) doesn't seem to be out of reality.
A complete system to do (relatively) small field (/b imaging and ( /D
imaging and spectroscopy, according to our above rough estimate could be
done with about 6 M$. It isn't really cheap, but it can do great science no
one else can do.



8. What would the Italian Groups wish (and be able) to do?

We first of all reaffirm our desire to collaborate with our LBT partners in
all phases of the development of the high angular resolution potential of
LBT, starting from conceptual definition and division of work in the
preliminary phase. The Italian community has obviously interest in
essentially all the instruments that we have mentioned, as it is probably
the case for our non-Italian LBT partners. The following considerations are
therefore only a preliminary attempt to identify priorities in scientific
interest and areas of technical experience among the group of Institutes
presenting this proposal.

8.1 Scientific Priorities



8.2 Technical expertise (To be completed after preliminary distribution)

Sodium lasers
Correctors
WFS
Optics
Detectors
Mechanics
...
...

8.3 ???? da Torino?



9. References

[1] D. G. Sandler, S. Stahl, J. R. P. Angel, M. Lloyd-Hart and D. McCarthy,

``Adaptive optics for diffraction-limited infrared imaging with 8-m
telescopes,''
J. Opt. Soc. Am. A ,11, 925-945 (1994).
[2] M. Tallon and R. Foy "Adaptive telescope with laser probe: isoplanatism
and cone effect"
Astron. Astrophys. 235, 549-557 (1990)
[3] J. M. Beckers "Increasing the Size of the Isoplanatic Patch with
Multiconjugate Adaptive Optics"
in ESO Conf. on "Very Large Telescopes and their Instrumentation", M. H.
Ulrich ed., p 693 (1988)
[4] P. Bayard and D. Bonaccini " Optical design for interferometry with the
Large Binocular Telescope"
Proceedings of SPIE conference on Amplitude and Intensity Spatial
Interferometry II, 2200, (1994) p. 446
-----------------------


Figure 1: A cut through the calculated PSF in the I band . The assumed
total r0 is 20 cm at 0.5 (m.