Äîêóìåíò âçÿò èç êýøà ïîèñêîâîé ìàøèíû. Àäðåñ îðèãèíàëüíîãî äîêóìåíòà : http://www.mrao.cam.ac.uk/ppeuc/astronomy/papers/jones.ps Äàòà èçìåíåíèÿ: Fri Jul 4 14:20:13 1997 Äàòà èíäåêñèðîâàíèÿ: Tue Oct 2 10:11:36 2012 Êîäèðîâêà: Ïîèñêîâûå ñëîâà: south pole

Primordial Experiments II
Michael E. Jones 1
Mullard Radio Astronomy Observatory 2 , Cavendish Laboratory,
Madingley Road, Cambridge CB3 0HE, U.K.
CMB experiments using amplifiers cover the full range of angular scales of interest, from
l = 2 to l = 6000, and provide constraints on the CMB power spectrum with similar sens­
itivity to bolometric measurements. Total­power measurements have been made possible
by careful experimental design, including sophisticated scanning strategies and optical
arrangements and use of exceptionally dry sites. Interferometric measurements offer sup­
pression of several systematic effects and have been made on smaller angular scales and
from less extreme sites; several new experiments are being constructed covering the ex­
pected angular scales of the Doppler peaks. Measurements on smaller scales at centimetre
wavelengths are prone to contamination from extragalactic sources, and higher­resolution
data are necessary to cope with this.
1 Introduction
The previous paper described the current state and future prospects of CMB experiments based on
bolometric receivers. Here I intend to do the same for what are often loosely termed heterodyne
receivers. The fundamental distinction is between systems that measure the total energy collected by
the telescope, and those that convert the electric field of the incoming radiation to voltage which can
be directly amplified. (The term `heterodyne' strictly refers to systems in which the signal is down­
converted to an intermediate frequency.) The practical consequences of this distinction are that the two
classes of experiment have different advantages and disadvantages, and suffer from different systematic
effects. Bolometers are intrinsically broad­band and insensitive to polarization; both frequency and
polarization response are defined by external filters. Amplifiers are intrinsically polarized, and one
struggles to make a single amplifier with more than 50% bandwidth.
Bolometers are `intrinsically' more sensitive, largely due to their higher bandwidths, but in practice
amplified systems obtain comparable sensitivity because much longer integration times are available.
This is because the relative transparency of the atmosphere at centimetre wavelengths allows ground­
based operation (albeit often at extreme sites). There is a further division of `electric field measuring'
instruments: the field may simply be squared, to measure the received power; or the fields from two
different anntennas may be multiplied, leading to the field of interferometry.
2 Receiver technology
The usual current technology for low­noise amplification is High Electron­Mobility Transistors (HEMTs)
(also known as HFETs, Heterostructure Field­Effect Transistors). These are widely used at frequen­
cies up to 50 GHz (e.g. Lai et al. (1994)) and have been demonstrated up to 140 GHz (Wang et al.
1995). A typical amplifier might have a noise temperature of 0.5 K/GHz, 20 dB of gain and 30­50%
bandwidth. The relatively high operating temperature (10­15 K) compared to bolometers (0.05­0.3 K)
means that the cryogenic technology required is much simpler.
An alternative to HEMTs is superconductor­insulator­superconductor (SIS) mixers. These are
often used for (non­CMB) heterodyne observations at frequencies 100--500 GHz, and provide gain and
down­conversion in a single package. It is now possible to have SIS receivers with T sys = 60K at 200
GHz with a bandwidth esentially limited only by the IF amplifier. SIS mixers have been used for CMB
astronomy (Meinhold & Lubin (1991)) but bolometers have generally been superior for total­power
1 mailto:mike@mrao.cam.ac.uk
2 http://www.mrao.cam.ac.uk/
1

Table 1: Solutions to systematic problems in some beam­switching experiments.
Systematic problem
Experiment Optics Groundspill Atmosphere Receiver stability
COBE
Twin horns at
60 ffi , spun and
precessed
Shield around
horns, point away
from Earth
Above it! Dicke switch
between horns
South Pole
Gregorian
offset­fed
paraboloid,
chopping
secondary
Ground shield,
over­sized
primary
Dry site, single
difference chop Fast chop
Saskatoon
Offset­fed
paraboloid
feeding chopping
flat
Ground shield,
over­sized flat
Dry site, variable
weighting of data
can produce
high­order
differences
Fast chop
Tenerife
Twin corrugated
horns feeding
chopping flat
Horns fixed,
over­sized flat
Dry site,
triple­beam
(double
difference) chop
Fast chop
measurements due to their large bandwidths. SIS mixers may become important when interferometric
CMB measurements are pushed to higher frequencies.
3 Beam­switching experiments
In beam­switching experiments the objective is to scan a beam across the sky, synchronously detecting
the component of the total power that is in phase with the scan (and hence in the sky), whilst
rejecting signals due to receiver power fluctuations, the atmosphere, and groundspill. These spurious
signal sources are extremely difficult to suppress to the high accuracy required, and are part of the
reason why so many years passed between the discovery of the CMB itself and the measurement of
its anisotropy. They are also the reason why, after early attempts to use `general­purpose' telescopes
(such as the Owens Valley 40­m) for CMB work, almost all successful CMB observations have been
made with `custom­built' instruments with stringent suppression of systematics as the over­riding
design criterion.
Table 1 shows how a variety of switched­beam experiments have overcome these problems. Ground­
spill is a particularly insidious systematic as it necessarily occurs at the switching frequency, and
requires careful screening and under­illumination of mirrors. A related problem, the varying of the
illumination of the primary by a switching secondary, can be circumvented, as in the Saskatoon
experiment (Wollack et al. (1997)), by placing a switching flat mirror in front of all the focussing
optics. All ground­based experiments chop in azimuth to minimize differential ground­spill and change
of airmass.
4 Interferometers
An alternative to simply squaring the detected voltage to measure the total received power is to
multiply the signals from two antennas pointed at the same part of the sky and measure the correlated
power. Since this measures the correlation between two different points on the incoming wavefront,
the correlated signal contains information about the spatial structure of the source. By sampling the
wavefront at different spacings the source structure can be reconstructed (see e.g. Saunders (1997) for
an introduction to the theory of interferometric imaging).
The fact that only radiation entering both antennas produces an output signal removes many
systematic effects. Receiver stability is no longer an issue as gain fluctuations are uncorrelated between
receivers. Most of the emission from the ground and the atmosphere that enters the antennas is
2

multiply and
integrate
multiply and
integrate
0
90
90
0
Real Imaginary
front­end amplifier
LO
path compensation
quadrature splitting
correlation and phase­
switch demodulation
scan control
synchronous
demodulation
front­end amplifier
band­pass filter
power detection
detected signal
complex correlated
signal
generator
phase switch
Figure 1: Comparative block diagrams of a single interferometer baseline and a beam­switching system.
uncorrelated, and the component that is correlated is reduced by two further effects. If the path
difference from the emission source to the two inputs to the correlator is more than the coherence
length (= c=\Deltaš) then the signal is strongly attenuated. Also, the rotation of the earth causes the phase
of the correlated signal from the sky to vary at a characteristic rate, which is used to synchronously
detect the sky signal; all other signals are also strongly attenuated. This means that interferometers
can operate in atmospheric conditions which would be impossible for total­power systems (the CAT
(O'Sullivan et al. (1995)) is at sea level and yet can observe for 60% of the time). An interferometer
array also makes economical use of its front­end amplifiers, the main noise source in the system; an
array with N antennas measures N (N \Gamma 1)=2 complex quantites.
The price for this immunity to many systematics is greatly increased complexity (see Fig 1). An
interferometer requires several antennas which can track the sky, each with its own receiver; a path
compensation system to keep the paths to each correlator matched to much less than the coherence
length; and the correlators themselves, which multiply and integrate the inputs, two per antenna
pair (since the correlated signal is complex i.e. has two orthogonal components). It is usual to also
phase­modulate the signal as soon as possible and synchronously demodulate it in the correlator, to
eliminate any spurious signals in the RF and IF systems. The receiving system must also maintain
phase stability over the signal path for at least the time interval between calibrations.
This complexity has meant that relatively few CMB measurements have been made so far with
interferometers. The VLA, a `general­purpose' interferometer, has been used to search for arcminute­
scale CMB fluctuations (with some recent success, see Richards et al. (1997), Partridge (1997)) but
is not well optimised for CMB work, having a relatively poorly filled aperture. This concept of filling
factor of the synthesised aperture, roughly equal to the total physical area of the antennas divided by
the area swept out by the longest baseline, is crucial to the design of CMB interferometers since it
enters directly into the temperature sensitivity of the array.
The only fully custom­built CMB interferometer that has so far produced results is the Cos­
mic Anisotropy Telescope (CAT) (Robson et al. (1993), Scott et al. (1996)). The Ryle Telescope,
which is used for Sunyaev­Zel'dovich astronomy as well as primordial CMB work (see section 6) and
source­subtraction for the CAT, is a `semi­custom' experiment, in which an existing general­purpose
instrument, the 5­km telescope, was converted for CMB work with new receiver and correlator sys­
tems. However, there are at least four dedicated CMB interferometers curently being commissioned
or constructed: the Tenerife 33 GHz interferometer (NRAL/IAC) (Watson (1997)), the Very Small
Array (MRAO/NRAL/IAC), the Cosmic Background Imager (CalTech) and the Very Compact Array
(Chicago). All these work in the 26­36 GHz band and between them will cover angular scales of 6
3

Figure 2: Radio to sub­mm spectrum of RXJ1459.9+3337. The line is a spectrum of ff = 2, i.e. that of
the CMB. Other similar sources could peak at even higher frequencies. The surface density of sources
such as this is not well known.
arcmin to 4 degrees.
5 Foregrounds
Foreground emission can be divided into that which is spatially similar to the CMB (but may have a
different frequency spectrum), e.g. diffuse Galactic emission, and emission which is spatially distinct
from the CMB, e.g. extragalactic point sources. In the frequency range in which HEMT experiments
operate Galactic emission is dominated by diffuse synchrotron and free­free radiation, both of which
have frequency spectra quite distinct from the CMB (flux spectral indices ff = \Gamma0:5 to \Gamma1:0 and
ff = \Gamma0:1 respectively, where S / š ff , compared to ff = 2 for the CMB.
Galactic emission can thus be separated from CMB structure if multi­frequency information is avail­
able, but only at the cost of loss of signal­to­noise ratio. The variance on the foreground­subtracted
data is increased both by an factor which depends on the separation of the frequency channels used,
and by an additional term which depends on how well the spectrum of the foreground is known
(Dodelson (1995)):
oe 2
subtracted =F 2 oe 2
0 + oe 2
model
The foreground degredation factor F usually dominates over the contribution from the model uncer­
tainty oe 2
model . As an example, in an experiment with two frequency channels at 28 and 35 GHz and a
foreground with flux spectral index ff = \Gamma0:5, F = 3:8. That is, the noise level is increased by a factor
of nearly two by the act of subtraction the foreground component. In the design of any experiment,
serious consideration must be given to F ; it may be that, for given technical and cost constraints, it
4

Table 2: Flux sensitivities and rough densities of confusing sources for some CMB experiments.
Experiment Beamsize Flux sensitivity Density of sources
COBE 10 ffi 80 Jy few/sky
Tenerife 5 ffi 0.5 Jy few/total field
South Pole 1 ffi 0.3 Jy 1/10 beams
Saskatoon 1 ffi 0.3 Jy 1/10 beams
CAT 0:5 ffi 10 mJy 1/beam
Ryle 0:03 ffi 0.1 mJy 1/few beams
VSA 0:2 ffi 5 mJy 1/few beams
is better to have fewer, more sensitive channels to minimise oe 0 rather than add channels that do not
significantly reduce F .
Extragalactic point sources present a rather different problem. Their flux spectral indices can
vary in the range \Gamma2 ! ff ! +2:5 due to a combination of synchrotron emission and self­absorbtion.
Most sources have negative spectral index, but not all; some 15% of sources have ff 5GHz
1:4GHz ? 0 and
the fraction with `inverted' spectra increases at higher frequency (O'Sullivan (1995)). Also, individual
source spectra are not very predictable; in a sample of 31 sources with measured fluxes at 15 GHz
(from the CAT1 field) the true 15 GHz fluxes differed from those extrapolated from 1.4 to 5 GHz by
factors between 0:5 and 1:6.
Point sources can of course be removed in the same way as diffuse foregrounds via their spectral
differences, but at the cost in signal­to­noise described above. Much better is to measure their fluxes at
higher resolution and subtract the exact contribution of the sources from the CMB data. Experiments
that operate in the regime of sensitivity and resolution where sources are a significant problem must
have access to higher­resolution source data.
Some sources can have spectra which mimic the CMB. Fig.2 shows the spectrum of the source
RXJ1459.9+3337, which has a flux spectral index of ff = +1:8 between 1.4 and 22 GHz. The only way
sources such as this can be distinguished from CMB fluctuations is by having higher resolution data
at the same frequency as the CMB observations. Sources can also be variable; the brightest source
in the CAT2 field (Baker (1997)) varied by a factor of two during the time of the observations. The
only answer for this is simultaneous monitoring, requiring a telescope with higher resoltion and flux
sensisivity than the CMB telescope.
Point sources will become a more serious problem as experiments move to smaller angular scales
and higher sensitivity. Table 2 shows the flux sensitivity per pixel of various CMB experiments,
along with a very rough estimate of the density of sources above that flux level. These estimates
are necessarily rough as there is little information on source counts above 8 GHz, and there are no
high­resolution all­sky surveys between 5 and 3000 GHz.
6 Results
Fig.3 shows some CMB power­spectrum results (adapted from Rocha (1997)). The apparent smooth
delineation of the the Doppler peak is perhaps a little misleading as by chance there is little overlap
in l­coverage of the experiments plotted; when bolometric measurements are included the picture
becomes a little more obscure! It appears clear, however, that the high power measured by Saskatoon
combined with the lower power measured by CAT points to the existence of a real peak in the power
spectrum.
This plot includes one new point, from the Ryle Telescope (RT), at l = 6300 \Sigma 2000. This is from
observations of the Lynx 2 field, which has been the subject of several long integrations on different
radio telescopes, including the VLA (Windhorst et al. (1993)). This field was observed at 15.2 GHz
with the RT compact array on 54 occasions between 1994 February and 1997 April, resulting in a
map with a noise level of 26¯Jybeam \Gamma1 and a resolution of 33 \Theta 55 arcsec. After removal of two
point sources the visibility data were analysed for excess power using the same procedures as used for
CAT data (Scott et al. (1996)). This results in a preliminary limit for flat­band power in the interval
4300 ! l ! 8300 of \DeltaT ! 18¯K (67% confidence), further confirming the result that CMB fluctuations
5

1 10 100 1000
0
Figure 3: CMB power­spectrum results from several experiments. From left to right: COBE; Tenerife;
South Pole; five Saskatoon points; two CAT points (the new CAT2 data (Baker (1997)) have been
added to the previous CAT1 results); Australia Telescope; Ryle Telescope.
on arcminute scales have an order of magnitude less power than on degree scales.
7 Conclusions
CMB experiments using amplifiers in the 15­90 GHz frequency range have provided significat results
on all angular scales using both beam­switching and interferometric techniques. Interferometers, which
so far have given results on smaller angular scales (l = 300 to l = 6000), are poised to move into the
crucial regime of l – 80 where the detail of the Doppler peaks is expected to resolve many astrophysical
problems. However, foregrounds remain a problem. The relatively small frequency range available to
HEMTs exacts a heavy price in signal­to­noise of subtraction of diffuse foregrounds, and point sources
demand simultaneous, high­resolution observations to complement the CMB data.
References
Baker, J.C., 1997, in `Particle Physics and the Early Universe', eds Bately, R., Jones, M.E. and Green,
D.A., http://www.mrao.cam.ac.uk/ppeuc/proceedings/
Dodelson, S., 1995, astro­ph/9512021
Lai, R., Bautista, J.J., Fujiwara, B., Tan, K.L., Ng, G.I., Dia, R.M., Streit, D., Liu, P.H., Freudenthal,
A., Laskar, J., Pospieszalski, M.W., 1994, IEEE Micro. Guided Wave Lett., 4, 329
Meinhold, P., Lubin, P., 1991, ApJ, 370, L1
6

O'Sullivan, C.M.M., 1995, PhD thesis, University of Cambridge.
O'Sullivan, C., Yassin, G., Scott, P.F., Saunders, R., Robson, M., Pooley, G., Lasenby, A.N., Kender­
dine, S., Jones, M.E., Hobson, M.P., Duffett­Smith, P.J., 1995, MNRAS, 274, 861
Partridge, R.B., 1997, in `Particle Physics and the Early Universe', eds Bately, R., Jones, M.E. and
Green, D.A., http://www.mrao.cam.ac.uk/ppeuc/proceedings/
Richards,E.A., Fomalont, E.B., Kellermann, K.I., Partridge, R.B., Windhorst R.A., 1997, 113, 1475
Robson, M., Yassin, G., Woan, G., Wilson, D.M.A., Scott, P.F., Lasenby, A.N., Kenderdine, S., Duffet­
Smith, P.J., 1993, A&A, 286, 1028
Rocha, G., 1997, PhD thesis, University of Cambridge
Saunders, R., 1997, in `Particle Physics and the Early Universe', eds Bately, R., Jones, M.E. and
Green, D.A., http://www.mrao.cam.ac.uk/ppeuc/proceedings/
Scott, P.F., Saunders, R., Pooley, G., O'Sullivan, C., Lasenby, A.N., Jones, M.E., Hobson, M.P.,
Duffet­Smith, P.J., Baker, J.C., 1996, ApJ, 461, L1
Wang, H., Lai, R., Lo, D.C.W., Streit, D.C., Liu, P.H, Dia, R.M., Pospieszalski, M.W., Berenz, J.,
1995, IEEE Microw. Guided Wave Lett., 5 190
Watson, R., 1997, in `Particle Physics and the Early Universe', eds Bately, R., Jones, M.E. and Green,
D.A., http://www.mrao.cam.ac.uk/ppeuc/proceedings/
Windhorst, R.A., Fomalont, E.B., Partridge, R.B., Lowenthal, J.D., 1993, ApJ, 405, 498
Wollack E.J., Devlin M.J., Jarosik N., Netterfield C.B., Page L., Wilkinson D., 1997, ApJ, 476, 440
7