- Email: [email protected]
The spectrum of Cosmic Microwave Background at low frequency
New Astronomy Reviews 43 (1999) 207–214 www.elsevier.nl / locate / newar
The spectrum of Cosmic Microwave Background at low frequency Massimo Gervasi Dipartimento di Fisica, Universita` degli studi di Milano, Via Celoria 16, 20133, Milan, Italy
Abstract The spectrum of Cosmic Microwave Background at frequency close or below 1 GHz is currently known with a large uncertainty. The difficulties arise both from the calibration technique and the presence of large foreground signals. In this spectral region large deviations from the Planck distribution are still permitted. In this paper the current status of our knowledge is described. Experiments involved and techniques employed are also discussed. 1999 Elsevier Science B.V. All rights reserved. Keywords: Cosmic Microwave Background: spectrum PACS: 98.70.v
1. Introduction The spectrum of Cosmic Microwave Background (CMB) has been accurately measured in the millimeter region by the COBE-FIRAS experiment (Mather et al., 1990). As expected from the standard Big Bang theory, FIRAS has shown that in the frequency range n 5 60–600 GHz the CMB spectrum is very close to a black-body (BB) distribution. The temperature best fit of the CMB is T CMB 5 2.72860.004 K (Fixsen et al., 1996). In this spectral region a possible distorsion of the BB distribution is very marginal. In fact rms deviations are , 5 3 10 25 with respect to the peak intensity. Outside the frequency range covered by FIRAS other measurements confirmed this indication, but the accuracy obtained is definitely worse. The error bar of experimental measurements of CMB absolute temperature is rapidly increasing going to low frequency, as it is evident in Fig. 1. Here the most representative measurements in the different spectral regions are shown. In Fig. 1 the COBE-FIRAS
measurement covers the 60–600 GHz region but its error bar is not distinguishable. Nevertheless distorsions of the CMB spectrum are expected as a consequence of energy dissipation of processes occurring in the early Universe. Examples of these processes are particle–antiparticle annihilation and nucleosynthesis at higher redshifts, structure formation and density wave or turbulence dissipation at a more recent epoch. Energy injection can both add new photons to the pre-existing radiation field and increase the energy content of the plasma particles. After every process of energy dissipation the plasma–radiation mixture is subject to a process of thermalization through the Compton scattering and the Bremsstrahlung mechanisms. For a complete review see Zeldovich & Sunyaev (1969), Sunyaev & Zeldovich (1970), Zeldovich et al. (1972), Illarionov & Sunyaev (1975a), Illarionov & Sunyaev (1975b), Danese & De Zotti (1982). Particles of the primordial plasma (mainly electrons) interact with the photons of the radiation field through Compton scattering. Therefore a fraction of
1387-6473 / 99 / $ – see front matter 1999 Elsevier Science B.V. All rights reserved. PII: S1387-6473( 99 )00013-5
208
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
Fig. 1. CMB absolute temperature measurements and expected spectral distorsions: Bose–Einstein (BE), Compton (C), Free–Free (FF) (see the text).
energy from electrons is transferred to photons. The final result is a shift of photons towards higher frequency, with a consequent depletion in the low frequency region. Bremsstrahlung processes involve interaction of electrons with protons. The final result is the production of photons at low energy (and frequency). The spectral distorsion produced is a function of the amount and of the epoch of the energy injection. A complete thermalization is expected for processes occurring at Z . 4 3 10 6 . For more recent processes we expect to find a signature on the CMB spectrum. For a complete review see Sunyaev & Zeldovich (1980), Danese & De Zotti (1980), Burigana et al. (1991a), Burigana et al. (1991b), Burigana et al. (1995). For processes occurring at epoch 4 3 10 6 . Z . 2 3 10 5 a Bose–Einstein (BE) spectrum is expected. This spectrum is characterized by a chemical potential ( m0 ) parameter. The BE spectrum shows a temperature dip with respect to the BB distribution. The amount of the temperature dip at the minimum is related to m0 , while the frequency position is related mainly to the baryon density: DT m 2 22 / 3 ]] . 6m0 (Vb h ) , T
lm . 6(Vb h 2 )22 / 3 (cm). We can estimate the region where we can find this minimum. With the current values accepted for the baryon density (Vb | 0.1–0.01) and the Hubble constant (h | 0.4–0.8) we obtain: lm | 30–250 cm. Therefore this effect can be important in the low frequency region. Taking into account the upper limit found by FIRAS ( m0 , 9 3 10 25 ) we can also estimate the temperature dip: DT m , 8–60 mK. This temperature depletion is far below the current accuracy, but a positive detection can give us a new independent estimation of the baryon density parameter. We cannot exclude the possibility of an even larger effect. In fact FIRAS measurement has been carried out at frequencies far from the region of the dip. For processes occurring at epoch 2 3 10 5 . Z . 1 3 10 3 a Comptonized spectrum is expected. This effect is important at high frequency, in the Wien region of the spectrum where FIRAS performed its observation. The upper limit found by FIRAS excludes large distorsions. Finally after the recombination (Z , 1 3 10 3 ) a reionization of the matter could generate a distorsion in the low frequency part of the
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
spectrum through the free–free interaction (Bartlett & Stebbins, 1991). The expected spectral distorsion is: DT YFF ]5] . T x2 Here YFF is the free–free parameter, while x 5 hn /kT is the adimensional frequency. This effect is rapidly increasing as the frequency decreases, therefore it is important at very long wavelenghts. For this reason there is no constraint from the COBE-FIRAS measurement. In Fig. 1 the spectral distorsions discussed in this section are also shown. Large distorted spectra are plotted in order to show the effect. It is evident from Fig. 1 that at low frequency measurements are consistent with large deviations from the BB distribution.
2. Low frequency absolute measurements Observations at low frequency are carried on using coherent receivers and large beamwidth antennas. A summary of the most recent measurements of CMB absolute temperature is shown in Table 1 and in Fig. 2. Points plotted in Fig. 2 are referred to in Table 1.
209
Table 1 Summary of low frequency CMB absolute temperature measurements
n (GHz)
T CMB (K)
Fig. 2
Reference
0.6 0.82 1.4 1.41 1.47 2.0 2.5 2.5 2.5 3.7 3.8
3.061.2 2.761.6 2.65 10.33 20.30 2.1160.38 2.2760.19 2.5560.14 2.6260.25 2.7960.15 2.5060.34 2.5960.13 2.6460.07
[1] [2] [3] [4] [8] [9] [5] [6] [7] [10] [11]
Sironi et al. (1990) Sironi et al. (1991) Staggs et al. (1996) Levin et al. (1988) Bensadoun et al. (1993) Bersanelli et al. (1994) Sironi et al. (1984) Sironi & Bonelli (1986) Sironi et al. (1991) De Amici et al. (1988) De Amici et al. (1990)
The accuracy of these absolute measurements is limited by the evaluation of systematics. In fact the receiver sensitivity is not a limiting parameter. At frequency around 1 GHz the observed signal is a combination of sky emission and environmental contribution. Sky emission itself contains a large foreground component, expecially from the galaxy. Therefore it is mandatory to take care of the observation strategy and to develop techniques for disentangling the different contributions.
Fig. 2. Low frequency measurements of the CMB absolute temperature. Reference of the points can be found in Table 1.
210
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
On the other hand due to the large dimension of the antenna used at these frequencies the absolute calibration of the observed signal is not trivial. As shown in Table 1 and in Fig. 2 the calibration accuracy becomes worse going to lower frequency. The calibration strategy is therefore fundamental together with a number of ancillary measurements. Despite the scientific motivation to detect spectral distorsions and the large uncertainties of the present measurements there are very few experiments operating at low frequency. The TRIS experiment (Bonelli et al., 1995; Zannoni et al., 1998) is the only one working below 1 GHz. TRIS is a system of three radiometers at 0.6, 0.82 and 2.5 GHz. The three radiometers adopt three scaled antennas (FWHM . 188) in order to look at the sky with the same beam. The experiment takes advantage of the frequency spectral analysis in order to disentangle the different components of the observed signal. In addition it is much easier to take into account the systematics introduced by instrumentation. TRIS is installed at Campo Imperatore – Gran Sasso (L’Aquila – Italy), at 2000 m a.s.l. Campo Imperatore is a relatively radio-quite site, supported by the logistics of LNGS-INFN.
3. Foreground signals The observed signal at frequency as low as 1 GHz can be written, in terms of antenna temperature, as follows: TA 5 T loc (n,u,f ) 1 T sky (n,a,d ). Several terms contribute to the local environment signal T loc : T loc 5 T atm (n,u,f ) 1 T gr (n,u,f ) 1 T RFI (n,u,f,t). Here atmospheric emission (T atm ), ground emission (T gr ), and radio frequency interferences (T RFI ) are considered. On the other hand the sky signal (T sky ) can be written as: T sky 5 T CMB (n ) 1 T gal (n,a,d ) 1 T egs (n ). The CMB signal (T CMB ) is added to the galactic (T gal ) emission and to the extra-galactic sources (T egs ).
3.1. Extra-galactic sources All the extra-galactic sources inside the antenna beam contribute to the signal. The extra-galactic signal can be written as the average of all the sources inside the beam. Inside a large beam a lot of sources contribute to the signal. Therefore T egs will become much more stable and will not change with the sky direction. Using a large beam it is not possible to directly measure the contribution of extra-galactic sources. We can adopt an estimation of T egs from observations with high resolution telescopes at n 5 178 MHz: T egs (n ) 5 (2363)(n (MHz) / 178)22.75 (K).
3.2. Galactic emission Galactic emission dominates the sky signal in the region below 1 GHz, as it is shown in Fig. 3. There are two mechanisms contributing to the galactic emission: the syncrothron radiation of electrons moving inside the magnetic field of the galaxy and the free–free emission. Both contributions can be described as a power low of frequency. The spectral index for free–free emission is gff . 2.1, for syncrothron it is gsy . 2.8. Below 20 GHz syncrothron radiation is dominant and the galactic temperature is: T gal (n,a,d ) 5 T gal (n0 ,a,d )(n /n0 )2g. The amplitude changes with the sky position, because the galactic emission is highly anisotropic. The spectral index also changes with position. In fact free–free contribution is concentrated on the galactic disk, while syncrothron radiation is present also towards the galactic halo. The accuracy of the existing maps of galactic emission (Haslam et al., 1982; Reich & Reich, 1988) at low frequency is not sufficient. Therefore it is necessary to disentangle galactic emission from the measured sky temperature. The anisotropic behaviour of galactic emission can be used to separate it from the rest of the sky signal. If we get observations at different frequencies and cover large regions of the sky, we can make use of the T–T plot technique.
3.3. Atmospheric emission Atmospheric emission can be evaluated observing
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
211
Fig. 3. Low frequency galactic and extragalactic (EGS) emission together with CMB. Both syncrothron (Sync) and free–free (ff) radiation are estimated. Maximum and minimum syncrothron signals are estimated for an antenna looking at d 5 428 and a FWHM 5 188.
the sky signal at different elevation. If z is the zenith angle, atmospheric emission can be modeled as: T atm (z) 5 T atm (0) sec(z). This technique (dip scan) is effective when the sky temperature (T sky ) is uniform and the major changes are due to the air mass emission. This is true at frequencies above few GHz. Unfortunately at lower frequency the galactic emission is much stronger and the sky temperature is largely anisotropic. On the other hand at these low frequencies the atmospheric emission is stable and predictable. In fact the emission of the oxigen band dominates over the water vapour. There are several models able to predict the atmospheric emission (Ajello et al., 1995) when the characteristics of the observation site are known (mainly geographic position and elevation a.s.l.). In any case a dry and elevated place is recommended as observation site. The antenna temperature of the atmosphere at the zenith from Campo Imperatore is shown in Fig. 4.
3.4. Radio frequency interferencies The spectral region near 1 GHz is polluted by
many artificial signals. A lot of devices emit radiation in this frequency band. Most of them are broadcasting signals (i.e. TV channels, cellular telephones). RFI signals are variable with time both in amplitude and in frequency. When the intensity is large their presence is easy to be detected, but faint signals are much more dangerous. All these RFIs have to be avoided in order to observe the cosmic radiation. Therefore it is particularly important to observe from a radio-quite place, shielded from RFIs. Furthermore a spectrometer receiver with narrow band channels and a stable central frequency is preferred. In this way the frequency band can be explored and a clean channel can be selected. This is possible with a downconversion section using a phase locked loop technique. In Fig. 5 the frequency interval between 810 MHz and 830 MHz measured by TRIS from Campo Imperatore is shown.
3.5. Ground and Sun emission Ground radiation entering through the beam sidelobes can also contribute to the measured antenna temperature. A precise knowledge of the beam profile far from the axis is necessary in order to
212
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
Fig. 4. Estimated atmospheric antenna temperature at the zenith from Campo Imperatore.
evaluate T gr . A beam profile with low sidelobes and backlobes is preferred together with an observation site with a flat horizon profile. The ground contribution can be heavily reduced by using metallic shields reflecting the sky radiation into the sidelobes. Shields around the antenna can also help to experimentally evaluate the ground signal. If measurements are carried on during day-time the contribution from the Sun has to be evaluated. Drift scans at the zenith from Campo Imperatore have shown a large signal of the Sun in summertime, while in winter-time the contribution is not easily detectable. In order to avoid to reject day-time observations it is necessary to cross-check drift scans measured at a distance of several months.
4. Calibration techniques
Fig. 5. An example of signal (and standard deviation) observed by TRIS in the several channels of the 820 MHz band. Frequency intervals dominated by RFI are well evidentiated.
A proper calibration technique is fundamental to obtain accurate absolute measurements. In fact the sky temperature is computed by comparing the observed sky signal with a reference source. The best strategy is to observe a calibration source with the
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
antenna under the same conditions of the sky measurements. This requirement cannot be satisfied for antennas with a beamwidth as large as 108 or even more, because there are not diffuse celestial sources of calibration. Therefore the reference source has to be artificial. It is necessary to have an accurate evaluation of both the zero level and the gain of the radiometer. This is possible with a calibration source at a well-known temperature close to the sky temperature. For this reason a stable reference source at cryogenic temperature is needed. A source of calibration filling the antenna beam can be used when the antenna dimensions are reasonably small. A liquid helium reference source has been developed by the LBL-Berkeley group (Bensadoun et al., 1992) for absolute temperature calibration of CMB at several frequencies. The antennas have been rotated upside-down and aimed at the aperture of the reference source. Despite the large mouth of the calibration source radiometers at low frequency (n 5 1.4 GHz) have been calibrated without the horn extension. When the antenna mouth is large (i.e. at low frequency) the horn cannot be directly calibrated. The reference signal has to be injected into the radiometer through a switch. Several calibration loads have been developped to inject the signal both through the waveguide (Roll et al., 1967; Penzias, 1968; Gervasi et al., 1995) and the cable connecting the receiver (Limon et al., 1989). The switch used to alternatively connect the receiver with the horn and the reference load is a potential source of systematic effects. The switch contributes to the signal through its insertion loss. Particular attention has to be payed to the asymmetries and to the repetibility. A room temperature switch requires a symmetry and stability extremely high to obtain and difficult to measure. Therefore a switch cooled down at cryogenic temperature is preferred. On the other hand active devices (as a solid state GaAs switch) have to be avoided, because of the elevated excess noise introduced and the poor repetibility in different thermal cycles. An interesting solution has been adopted by the Princeton group (Staggs et al., 1996). They used a cryogenic hybrid circuit inside a correlation receiver. For TRIS a cryogenic fully passive switch has been developped
213
(Zannoni et al., 1998), which takes advantage of the properties of a l / 4 transmission line. Finally it is important to evaluate the contribution of all the parts of the antenna that cannot be directly calibrated. Some components (like the horn and the waveguide) can be easily modelled. Another possibility is to perform ancillary measurements on these elements. A way to reduce the contribution of the lossy components is to cool them down at liquid helium temperature. This technique is effective particularly for the components which have unpredictable losses (i.e. joints and waveguide to cable transition).
Acknowledgements The author acknowledges Mario Zannoni for useful suggestions and fundamental support during the preparation of this contribution.
References Ajello, C., et al., 1995, ApJS, 96, 643. Bartlett, J.G. & Stebbins, A., 1991, ApJ, 371, 8. Bensadoun, M., et al., 1992, Rev. Sci. Instr., 63, 4377. Bensadoun, M., et al., 1993, ApJ, 409, 1. Bersanelli, M., et al., 1994, ApJ, 424, 517. Bonelli, G., et al., 1995, ApL&C, 32, 15. Burigana, C., et al., 1991a, A&A, 246, 49. Burigana, C., et al., 1991b, ApJ, 379, 1. Burigana, C., et al., 1995, A&A, 303, 323. Danese, L. & De Zotti, G., 1980, A&A, 84, 364. Danese, L. & De Zotti, G., 1982, A&A, 107, 39. De Amici, G., et al., 1988, ApJ, 329, 556. De Amici, G., et al., 1990, ApJ, 359, 219. Fixsen, D.J., et al., 1996, ApJ, 473, 576. Gervasi, M., et al., 1995, Rev. Sci. Instr., 66, 4798. Haslam, C.T.G., et al., 1982, A&AS, 47, 1. Illarionov, A.F. & Sunyaev, R.A., 1975a, SvA, 18, 413. Illarionov, A.F. & Sunyaev, R.A., 1975b, SvA, 18, 691. Levin, S.M., et al., 1988, ApJ, 334, 14. Limon, M., et al., 1989, J. Phys. E, 22, 963. Mather, J.C., et al., 1990, ApJ, 354, L37. Penzias, A.A., 1968, IEEE Trans. MTT, 16, 608. Reich, P. & Reich, W., 1988, A&AS, 74, 7. Roll, P.G., et al., 1967, Annals Phys., 44, 289. Sironi, G. & Bonelli, G., 1986, ApJ, 311, 418. Sironi, G., et al., 1984, PhRvD, 29, 2686. Sironi, G., et al., 1990, ApJ, 357, 301.
214
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
Sironi, G., et al., 1991, ApJ, 378, 550. Staggs, S.T., et al., 1996, ApJ, 458, 407. Sunyaev, R.A. & Zeldovich, Y.B., 1970, Ap&SS, 7, 20. Sunyaev, R.A. & Zeldovich, Y.B., 1980, Ann. Rev. Astron. Astrophys., 18, 537.
Zannoni, G., et al., 1998, Proc. Intern. Conf. on 3K Cosmology, AIP Conf. Ser., in press. Zeldovich, Y.B. & Sunyaev, R.A., 1969, Ap&SS, 4, 301. Zeldovich, Y.B., et al., 1972, Soviet Phys. JETP, 35, 643.
The spectrum of Cosmic Microwave Background at low frequency Massimo Gervasi Dipartimento di Fisica, Universita` degli studi di Milano, Via Celoria 16, 20133, Milan, Italy
Abstract The spectrum of Cosmic Microwave Background at frequency close or below 1 GHz is currently known with a large uncertainty. The difficulties arise both from the calibration technique and the presence of large foreground signals. In this spectral region large deviations from the Planck distribution are still permitted. In this paper the current status of our knowledge is described. Experiments involved and techniques employed are also discussed. 1999 Elsevier Science B.V. All rights reserved. Keywords: Cosmic Microwave Background: spectrum PACS: 98.70.v
1. Introduction The spectrum of Cosmic Microwave Background (CMB) has been accurately measured in the millimeter region by the COBE-FIRAS experiment (Mather et al., 1990). As expected from the standard Big Bang theory, FIRAS has shown that in the frequency range n 5 60–600 GHz the CMB spectrum is very close to a black-body (BB) distribution. The temperature best fit of the CMB is T CMB 5 2.72860.004 K (Fixsen et al., 1996). In this spectral region a possible distorsion of the BB distribution is very marginal. In fact rms deviations are , 5 3 10 25 with respect to the peak intensity. Outside the frequency range covered by FIRAS other measurements confirmed this indication, but the accuracy obtained is definitely worse. The error bar of experimental measurements of CMB absolute temperature is rapidly increasing going to low frequency, as it is evident in Fig. 1. Here the most representative measurements in the different spectral regions are shown. In Fig. 1 the COBE-FIRAS
measurement covers the 60–600 GHz region but its error bar is not distinguishable. Nevertheless distorsions of the CMB spectrum are expected as a consequence of energy dissipation of processes occurring in the early Universe. Examples of these processes are particle–antiparticle annihilation and nucleosynthesis at higher redshifts, structure formation and density wave or turbulence dissipation at a more recent epoch. Energy injection can both add new photons to the pre-existing radiation field and increase the energy content of the plasma particles. After every process of energy dissipation the plasma–radiation mixture is subject to a process of thermalization through the Compton scattering and the Bremsstrahlung mechanisms. For a complete review see Zeldovich & Sunyaev (1969), Sunyaev & Zeldovich (1970), Zeldovich et al. (1972), Illarionov & Sunyaev (1975a), Illarionov & Sunyaev (1975b), Danese & De Zotti (1982). Particles of the primordial plasma (mainly electrons) interact with the photons of the radiation field through Compton scattering. Therefore a fraction of
1387-6473 / 99 / $ – see front matter 1999 Elsevier Science B.V. All rights reserved. PII: S1387-6473( 99 )00013-5
208
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
Fig. 1. CMB absolute temperature measurements and expected spectral distorsions: Bose–Einstein (BE), Compton (C), Free–Free (FF) (see the text).
energy from electrons is transferred to photons. The final result is a shift of photons towards higher frequency, with a consequent depletion in the low frequency region. Bremsstrahlung processes involve interaction of electrons with protons. The final result is the production of photons at low energy (and frequency). The spectral distorsion produced is a function of the amount and of the epoch of the energy injection. A complete thermalization is expected for processes occurring at Z . 4 3 10 6 . For more recent processes we expect to find a signature on the CMB spectrum. For a complete review see Sunyaev & Zeldovich (1980), Danese & De Zotti (1980), Burigana et al. (1991a), Burigana et al. (1991b), Burigana et al. (1995). For processes occurring at epoch 4 3 10 6 . Z . 2 3 10 5 a Bose–Einstein (BE) spectrum is expected. This spectrum is characterized by a chemical potential ( m0 ) parameter. The BE spectrum shows a temperature dip with respect to the BB distribution. The amount of the temperature dip at the minimum is related to m0 , while the frequency position is related mainly to the baryon density: DT m 2 22 / 3 ]] . 6m0 (Vb h ) , T
lm . 6(Vb h 2 )22 / 3 (cm). We can estimate the region where we can find this minimum. With the current values accepted for the baryon density (Vb | 0.1–0.01) and the Hubble constant (h | 0.4–0.8) we obtain: lm | 30–250 cm. Therefore this effect can be important in the low frequency region. Taking into account the upper limit found by FIRAS ( m0 , 9 3 10 25 ) we can also estimate the temperature dip: DT m , 8–60 mK. This temperature depletion is far below the current accuracy, but a positive detection can give us a new independent estimation of the baryon density parameter. We cannot exclude the possibility of an even larger effect. In fact FIRAS measurement has been carried out at frequencies far from the region of the dip. For processes occurring at epoch 2 3 10 5 . Z . 1 3 10 3 a Comptonized spectrum is expected. This effect is important at high frequency, in the Wien region of the spectrum where FIRAS performed its observation. The upper limit found by FIRAS excludes large distorsions. Finally after the recombination (Z , 1 3 10 3 ) a reionization of the matter could generate a distorsion in the low frequency part of the
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
spectrum through the free–free interaction (Bartlett & Stebbins, 1991). The expected spectral distorsion is: DT YFF ]5] . T x2 Here YFF is the free–free parameter, while x 5 hn /kT is the adimensional frequency. This effect is rapidly increasing as the frequency decreases, therefore it is important at very long wavelenghts. For this reason there is no constraint from the COBE-FIRAS measurement. In Fig. 1 the spectral distorsions discussed in this section are also shown. Large distorted spectra are plotted in order to show the effect. It is evident from Fig. 1 that at low frequency measurements are consistent with large deviations from the BB distribution.
2. Low frequency absolute measurements Observations at low frequency are carried on using coherent receivers and large beamwidth antennas. A summary of the most recent measurements of CMB absolute temperature is shown in Table 1 and in Fig. 2. Points plotted in Fig. 2 are referred to in Table 1.
209
Table 1 Summary of low frequency CMB absolute temperature measurements
n (GHz)
T CMB (K)
Fig. 2
Reference
0.6 0.82 1.4 1.41 1.47 2.0 2.5 2.5 2.5 3.7 3.8
3.061.2 2.761.6 2.65 10.33 20.30 2.1160.38 2.2760.19 2.5560.14 2.6260.25 2.7960.15 2.5060.34 2.5960.13 2.6460.07
[1] [2] [3] [4] [8] [9] [5] [6] [7] [10] [11]
Sironi et al. (1990) Sironi et al. (1991) Staggs et al. (1996) Levin et al. (1988) Bensadoun et al. (1993) Bersanelli et al. (1994) Sironi et al. (1984) Sironi & Bonelli (1986) Sironi et al. (1991) De Amici et al. (1988) De Amici et al. (1990)
The accuracy of these absolute measurements is limited by the evaluation of systematics. In fact the receiver sensitivity is not a limiting parameter. At frequency around 1 GHz the observed signal is a combination of sky emission and environmental contribution. Sky emission itself contains a large foreground component, expecially from the galaxy. Therefore it is mandatory to take care of the observation strategy and to develop techniques for disentangling the different contributions.
Fig. 2. Low frequency measurements of the CMB absolute temperature. Reference of the points can be found in Table 1.
210
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
On the other hand due to the large dimension of the antenna used at these frequencies the absolute calibration of the observed signal is not trivial. As shown in Table 1 and in Fig. 2 the calibration accuracy becomes worse going to lower frequency. The calibration strategy is therefore fundamental together with a number of ancillary measurements. Despite the scientific motivation to detect spectral distorsions and the large uncertainties of the present measurements there are very few experiments operating at low frequency. The TRIS experiment (Bonelli et al., 1995; Zannoni et al., 1998) is the only one working below 1 GHz. TRIS is a system of three radiometers at 0.6, 0.82 and 2.5 GHz. The three radiometers adopt three scaled antennas (FWHM . 188) in order to look at the sky with the same beam. The experiment takes advantage of the frequency spectral analysis in order to disentangle the different components of the observed signal. In addition it is much easier to take into account the systematics introduced by instrumentation. TRIS is installed at Campo Imperatore – Gran Sasso (L’Aquila – Italy), at 2000 m a.s.l. Campo Imperatore is a relatively radio-quite site, supported by the logistics of LNGS-INFN.
3. Foreground signals The observed signal at frequency as low as 1 GHz can be written, in terms of antenna temperature, as follows: TA 5 T loc (n,u,f ) 1 T sky (n,a,d ). Several terms contribute to the local environment signal T loc : T loc 5 T atm (n,u,f ) 1 T gr (n,u,f ) 1 T RFI (n,u,f,t). Here atmospheric emission (T atm ), ground emission (T gr ), and radio frequency interferences (T RFI ) are considered. On the other hand the sky signal (T sky ) can be written as: T sky 5 T CMB (n ) 1 T gal (n,a,d ) 1 T egs (n ). The CMB signal (T CMB ) is added to the galactic (T gal ) emission and to the extra-galactic sources (T egs ).
3.1. Extra-galactic sources All the extra-galactic sources inside the antenna beam contribute to the signal. The extra-galactic signal can be written as the average of all the sources inside the beam. Inside a large beam a lot of sources contribute to the signal. Therefore T egs will become much more stable and will not change with the sky direction. Using a large beam it is not possible to directly measure the contribution of extra-galactic sources. We can adopt an estimation of T egs from observations with high resolution telescopes at n 5 178 MHz: T egs (n ) 5 (2363)(n (MHz) / 178)22.75 (K).
3.2. Galactic emission Galactic emission dominates the sky signal in the region below 1 GHz, as it is shown in Fig. 3. There are two mechanisms contributing to the galactic emission: the syncrothron radiation of electrons moving inside the magnetic field of the galaxy and the free–free emission. Both contributions can be described as a power low of frequency. The spectral index for free–free emission is gff . 2.1, for syncrothron it is gsy . 2.8. Below 20 GHz syncrothron radiation is dominant and the galactic temperature is: T gal (n,a,d ) 5 T gal (n0 ,a,d )(n /n0 )2g. The amplitude changes with the sky position, because the galactic emission is highly anisotropic. The spectral index also changes with position. In fact free–free contribution is concentrated on the galactic disk, while syncrothron radiation is present also towards the galactic halo. The accuracy of the existing maps of galactic emission (Haslam et al., 1982; Reich & Reich, 1988) at low frequency is not sufficient. Therefore it is necessary to disentangle galactic emission from the measured sky temperature. The anisotropic behaviour of galactic emission can be used to separate it from the rest of the sky signal. If we get observations at different frequencies and cover large regions of the sky, we can make use of the T–T plot technique.
3.3. Atmospheric emission Atmospheric emission can be evaluated observing
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
211
Fig. 3. Low frequency galactic and extragalactic (EGS) emission together with CMB. Both syncrothron (Sync) and free–free (ff) radiation are estimated. Maximum and minimum syncrothron signals are estimated for an antenna looking at d 5 428 and a FWHM 5 188.
the sky signal at different elevation. If z is the zenith angle, atmospheric emission can be modeled as: T atm (z) 5 T atm (0) sec(z). This technique (dip scan) is effective when the sky temperature (T sky ) is uniform and the major changes are due to the air mass emission. This is true at frequencies above few GHz. Unfortunately at lower frequency the galactic emission is much stronger and the sky temperature is largely anisotropic. On the other hand at these low frequencies the atmospheric emission is stable and predictable. In fact the emission of the oxigen band dominates over the water vapour. There are several models able to predict the atmospheric emission (Ajello et al., 1995) when the characteristics of the observation site are known (mainly geographic position and elevation a.s.l.). In any case a dry and elevated place is recommended as observation site. The antenna temperature of the atmosphere at the zenith from Campo Imperatore is shown in Fig. 4.
3.4. Radio frequency interferencies The spectral region near 1 GHz is polluted by
many artificial signals. A lot of devices emit radiation in this frequency band. Most of them are broadcasting signals (i.e. TV channels, cellular telephones). RFI signals are variable with time both in amplitude and in frequency. When the intensity is large their presence is easy to be detected, but faint signals are much more dangerous. All these RFIs have to be avoided in order to observe the cosmic radiation. Therefore it is particularly important to observe from a radio-quite place, shielded from RFIs. Furthermore a spectrometer receiver with narrow band channels and a stable central frequency is preferred. In this way the frequency band can be explored and a clean channel can be selected. This is possible with a downconversion section using a phase locked loop technique. In Fig. 5 the frequency interval between 810 MHz and 830 MHz measured by TRIS from Campo Imperatore is shown.
3.5. Ground and Sun emission Ground radiation entering through the beam sidelobes can also contribute to the measured antenna temperature. A precise knowledge of the beam profile far from the axis is necessary in order to
212
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
Fig. 4. Estimated atmospheric antenna temperature at the zenith from Campo Imperatore.
evaluate T gr . A beam profile with low sidelobes and backlobes is preferred together with an observation site with a flat horizon profile. The ground contribution can be heavily reduced by using metallic shields reflecting the sky radiation into the sidelobes. Shields around the antenna can also help to experimentally evaluate the ground signal. If measurements are carried on during day-time the contribution from the Sun has to be evaluated. Drift scans at the zenith from Campo Imperatore have shown a large signal of the Sun in summertime, while in winter-time the contribution is not easily detectable. In order to avoid to reject day-time observations it is necessary to cross-check drift scans measured at a distance of several months.
4. Calibration techniques
Fig. 5. An example of signal (and standard deviation) observed by TRIS in the several channels of the 820 MHz band. Frequency intervals dominated by RFI are well evidentiated.
A proper calibration technique is fundamental to obtain accurate absolute measurements. In fact the sky temperature is computed by comparing the observed sky signal with a reference source. The best strategy is to observe a calibration source with the
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
antenna under the same conditions of the sky measurements. This requirement cannot be satisfied for antennas with a beamwidth as large as 108 or even more, because there are not diffuse celestial sources of calibration. Therefore the reference source has to be artificial. It is necessary to have an accurate evaluation of both the zero level and the gain of the radiometer. This is possible with a calibration source at a well-known temperature close to the sky temperature. For this reason a stable reference source at cryogenic temperature is needed. A source of calibration filling the antenna beam can be used when the antenna dimensions are reasonably small. A liquid helium reference source has been developed by the LBL-Berkeley group (Bensadoun et al., 1992) for absolute temperature calibration of CMB at several frequencies. The antennas have been rotated upside-down and aimed at the aperture of the reference source. Despite the large mouth of the calibration source radiometers at low frequency (n 5 1.4 GHz) have been calibrated without the horn extension. When the antenna mouth is large (i.e. at low frequency) the horn cannot be directly calibrated. The reference signal has to be injected into the radiometer through a switch. Several calibration loads have been developped to inject the signal both through the waveguide (Roll et al., 1967; Penzias, 1968; Gervasi et al., 1995) and the cable connecting the receiver (Limon et al., 1989). The switch used to alternatively connect the receiver with the horn and the reference load is a potential source of systematic effects. The switch contributes to the signal through its insertion loss. Particular attention has to be payed to the asymmetries and to the repetibility. A room temperature switch requires a symmetry and stability extremely high to obtain and difficult to measure. Therefore a switch cooled down at cryogenic temperature is preferred. On the other hand active devices (as a solid state GaAs switch) have to be avoided, because of the elevated excess noise introduced and the poor repetibility in different thermal cycles. An interesting solution has been adopted by the Princeton group (Staggs et al., 1996). They used a cryogenic hybrid circuit inside a correlation receiver. For TRIS a cryogenic fully passive switch has been developped
213
(Zannoni et al., 1998), which takes advantage of the properties of a l / 4 transmission line. Finally it is important to evaluate the contribution of all the parts of the antenna that cannot be directly calibrated. Some components (like the horn and the waveguide) can be easily modelled. Another possibility is to perform ancillary measurements on these elements. A way to reduce the contribution of the lossy components is to cool them down at liquid helium temperature. This technique is effective particularly for the components which have unpredictable losses (i.e. joints and waveguide to cable transition).
Acknowledgements The author acknowledges Mario Zannoni for useful suggestions and fundamental support during the preparation of this contribution.
References Ajello, C., et al., 1995, ApJS, 96, 643. Bartlett, J.G. & Stebbins, A., 1991, ApJ, 371, 8. Bensadoun, M., et al., 1992, Rev. Sci. Instr., 63, 4377. Bensadoun, M., et al., 1993, ApJ, 409, 1. Bersanelli, M., et al., 1994, ApJ, 424, 517. Bonelli, G., et al., 1995, ApL&C, 32, 15. Burigana, C., et al., 1991a, A&A, 246, 49. Burigana, C., et al., 1991b, ApJ, 379, 1. Burigana, C., et al., 1995, A&A, 303, 323. Danese, L. & De Zotti, G., 1980, A&A, 84, 364. Danese, L. & De Zotti, G., 1982, A&A, 107, 39. De Amici, G., et al., 1988, ApJ, 329, 556. De Amici, G., et al., 1990, ApJ, 359, 219. Fixsen, D.J., et al., 1996, ApJ, 473, 576. Gervasi, M., et al., 1995, Rev. Sci. Instr., 66, 4798. Haslam, C.T.G., et al., 1982, A&AS, 47, 1. Illarionov, A.F. & Sunyaev, R.A., 1975a, SvA, 18, 413. Illarionov, A.F. & Sunyaev, R.A., 1975b, SvA, 18, 691. Levin, S.M., et al., 1988, ApJ, 334, 14. Limon, M., et al., 1989, J. Phys. E, 22, 963. Mather, J.C., et al., 1990, ApJ, 354, L37. Penzias, A.A., 1968, IEEE Trans. MTT, 16, 608. Reich, P. & Reich, W., 1988, A&AS, 74, 7. Roll, P.G., et al., 1967, Annals Phys., 44, 289. Sironi, G. & Bonelli, G., 1986, ApJ, 311, 418. Sironi, G., et al., 1984, PhRvD, 29, 2686. Sironi, G., et al., 1990, ApJ, 357, 301.
214
M. Gervasi / New Astronomy Reviews 43 (1999) 207 – 214
Sironi, G., et al., 1991, ApJ, 378, 550. Staggs, S.T., et al., 1996, ApJ, 458, 407. Sunyaev, R.A. & Zeldovich, Y.B., 1970, Ap&SS, 7, 20. Sunyaev, R.A. & Zeldovich, Y.B., 1980, Ann. Rev. Astron. Astrophys., 18, 537.
Zannoni, G., et al., 1998, Proc. Intern. Conf. on 3K Cosmology, AIP Conf. Ser., in press. Zeldovich, Y.B. & Sunyaev, R.A., 1969, Ap&SS, 4, 301. Zeldovich, Y.B., et al., 1972, Soviet Phys. JETP, 35, 643.