Download Title: Optical and NEar IR Interferometric or

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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
1
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
2
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 multilaser 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 m V  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
3
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.
4
Figure 1: A cut through the calculated PSF in the I band . The assumed total r is 20
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
5
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:
a. a spherical corrector illuminated by the telescope F/15 beam (two or three added reflections)
b. a double pass off-axis parabolic collimator illuminating a flat corrector (four added reflections).
c.
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.
6
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
7
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
8
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.
9
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
size
N of Pix
/D
/2D
m
/2D
/2D
Field
F/n
mas
mas
arcsec
/b
/4b
 /4b
 /4b
Field
F/n
mas
mas
arcsec
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.
10
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.
Band
Sky
angular
BG
BG
brightn.
pixel
flux
Noise
R/O Noise BG+R/O BG Noise R/O Noise
noise
size
mag/sqarcsec
mas
ph
ph
ph
ph
equivalent equivalent
mag
mag
mag/pix
mag/pix
Noise
equivalent
mag
mag/pix
R
20.80
1.6 1.0E+01
3.2
2
3.80
35.3
35.8
35.1
I
19.30
2.0 6.9E+01
8.3
2
8.52
33.8
35.4
33.8
J
15.00
2.8 6.9E+03
83.3
10
83.91
31.0
33.3
31.0
H
13.00
3.6 7.2E+04 267.8
10
268.04
29.0
32.5
29.0
K
12.00
5.0 3.4E+05 583.7
10
583.79
27.9
32.3
27.9
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.
11
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
12
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.
Band
Sky
angular
BG
BG
pixel
flux
Noise
ph
ph
R/O
Noise
BG+R/O
noise
size
mag/
mas
ph
ph
BG Noise R/O Noise
Noise
equivale equivalent equivale
nt
nt mag
mag
mag
mag/pix
mag/pix
mag/pix
sqarcsec
R
20.80
8.7
7.8E+01
8.9
2
12.7
33.8
35.8
33.8
I
19.30
11.2
3.4E+02
18.5
2
26.3
32.6
35.4
32.6
J
15.00
15.6
2.6E+04 160.6
10
227.3
29.9
33.3
29.9
H
13.00
20.0
1.3E+05 358.4
10
506.9
28.3
32.5
28.3
K
12.00
27.5
4.9E+05 702.8
10
994.0
27.3
32.3
27.3
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,
13
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 reimaged 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
14
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
angular
/
average
BG
pixel
BG
BG
noise
R/O
noise
suppressio
n factor
BG+R/O
Noise
noise
equivalent
mag
size
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.
15
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 (m R~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.
16
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:
a. We need to form an image of a high atmospheric layer on a deformable mirror
b. We need to image the field on the detectors at very different scales (from about F/6 to about
F/200)
c.
We need to have an intermediate image of the field for spectroscopy (diffraction limited slits
diffract!)
d. We need to find appropriate positions for two beam splitters (for laser and natural star wave front
sensing)
e. We need to de-rotate detectors and spectrographs
f.
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
17
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
Type of system
Material
Focal length
Diameter (field 1'x1')

Pupil Image Diameter
6 Km image Diameter (+1' field)
Field at SR>0.95
Value
Spherical doublet
IRg2/BaF2
1.2 m
140 mm
1 to1.5 m
80 mm
80 (+20) mm
1.0 arc-min
80% encircled energy diameter at CC
< 0.3 mm
90% encircled energy diameter at pupil
image
< 30 m
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.
18
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, ngs
1 corr,1 lgs
1 corr, 4 lgs
2 corr, 4 lgs
K to R (1,2)
K to R (1,2)
K to R (1,2)
K to R (1,2)
K to R (1,2)
/D imaging
K to R (1,2)
K to J (1,2)
K to R (1,2)
K to R (1,2)
 /D
spectroscopy
K to R (1,2)
K to J (1,2)
K to R (1,2)
K to R(1,2)
K to R (2)
K to J (2)
K to R (2)
K to R (2)
K to R (1,2)
K to J (1.2)
K to R (1,2)
K to R (1,2)
Speckle imaging
 /b imaging
VHR
Spectroscopy
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
19
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.
20
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
21
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
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