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This work was enabled usine ihe input of a number of colleagues. especially M. Dreger. T. Ixallinger. M. Mareoni. A. Pamyatnykh. E. Paunzen and V. Ripepi.
This work was enabled using the input of a number of colleagues, especially M. Breger, T. Kallinger, M. Marconi, A. Pamyatnykh, E. Paunzen and V. Ripepi.
systematies which need to be controlled in non-interferometric CMB experiments.
systematics which need to be controlled in non-interferometric CMB experiments.
DASI is designed to measure the angular power spectrum of the CMB and produce high signal-to-noise Images on angular scales corresponding to the first three acoustic peaks predicted by adiabatic inflationary models for a flat Universe.
DASI is designed to measure the angular power spectrum of the CMB and produce high signal-to-noise images on angular scales corresponding to the first three acoustic peaks predicted by adiabatic inflationary models for a flat Universe.
In this. the second of three papers describing the results of the first season of DASI observations. we focus on the determination of the CMB power spectrum from the calibrated data.
In this, the second of three papers describing the results of the first season of DASI observations, we focus on the determination of the CMB power spectrum from the calibrated data.
This paper includes a discussion of the data analysis method. potential sources of astronomical and terrestrial contamination in the data. tests of data consistency. and the resulting angular power spectrum.
This paper includes a discussion of the data analysis method, potential sources of astronomical and terrestrial contamination in the data, tests of data consistency, and the resulting angular power spectrum.
Detailed descriptions of the instrument. observations. and data calibration are. given in LLeitch (2001. hereafter Paper 1).
Detailed descriptions of the instrument, observations, and data calibration are given in \markcite{leitch01}L Leitch (2001, hereafter Paper I).
The extraction of the cosmological parameters from the DASI angular power spectrum is presented in a third paper. PPryke (2001. hereafter Paper ILD.
The extraction of the cosmological parameters from the DASI angular power spectrum is presented in a third paper, \markcite{pryke01}P Pryke (2001, hereafter Paper III).
An extensive discussion of the instrument design and its performance are given in Paper LE: here we provide a brief overview stressing the aspects of the instrument that are particularly relevant to the measurement of the CMB angular power spectrum.
An extensive discussion of the instrument design and its performance are given in Paper I; here we provide a brief overview stressing the aspects of the instrument that are particularly relevant to the measurement of the CMB angular power spectrum.
Interferometers offer several unique features which make them well suited for sensitive measurements of CMB anisotropy: 1) the Fourier transform of the sky plane is measured directly: the effective sky brightness differencing is instantaneous. 2) 180° phase switching at the receivers with synchronized post-detection demodulation ts used to reduce instrumental offsets (DASI uses both fast (40 KHz) phase switching. with hardware demodulation and slow (1 Hz) phase switching with software demodulation to reduce offsets to well below the j/K level). and 3) the effective window function in /-space is well understood and uncertainties in the primary beam response do not lead to uncertainties in the resultant power spectrum that become large at small angular scales.
Interferometers offer several unique features which make them well suited for sensitive measurements of CMB anisotropy: 1) the Fourier transform of the sky plane is measured directly; the effective sky brightness differencing is instantaneous, 2) $180\degr$ phase switching at the receivers with synchronized post-detection demodulation is used to reduce instrumental offsets (DASI uses both fast (40 KHz) phase switching with hardware demodulation and slow $\sim 1$ Hz) phase switching with software demodulation to reduce offsets to well below the $\mu$ K level), and 3) the effective window function in $l$ -space is well understood and uncertainties in the primary beam response do not lead to uncertainties in the resultant power spectrum that become large at small angular scales.
An interferometer directly samples the Fourier transform of the sky brightness distribution.
An interferometer directly samples the Fourier transform of the sky brightness distribution.
The response pattern on the sky for a given pair of antennas is a sinusoidal fringe pattern attenuated by the primary beam of the individual antennas.
The response pattern on the sky for a given pair of antennas is a sinusoidal fringe pattern attenuated by the primary beam of the individual antennas.
For a pair of antennas with a physical separation (baseline) vector b. the center of the measured Fourier wavevector components. labeled u or (u.v). Is given by w=b,/À and v=h,/A. where b. and b, are the projections of the baseline normal to the line of sight. and A is the observing wavelength.
For a pair of antennas with a physical separation (baseline) vector $\mathbf{b}$, the center of the measured Fourier wavevector components, labeled $\mathbf{u}$ or $(u,v)$, is given by $u = b_x/\lambda$ and $v = b_y/\lambda$, where $b_x$ and $b_y$ are the projections of the baseline normal to the line of sight, and $\lambda$ is the observing wavelength.
The approximate conversion to multipole moment is given by /z25|u| ((White 19993), with a width A/ that is related to the diameter of the apertures in units of the observing wavelength (see refsec:formalism)).
The approximate conversion to multipole moment is given by $l \approx 2\pi\left|\mathbf{u}\right|$ \markcite{white99a}( (White 1999a), with a width $\Delta l$ that is related to the diameter of the apertures in units of the observing wavelength (see \\ref{sec:formalism}) ).
Both real and imaginary components of the Fourier plane can be measured with a complex correlator.
Both real and imaginary components of the Fourier plane can be measured with a complex correlator.
The real (even) component is measured by correlating the pair of signals without a relative phase delay: the imaginary (odd) component is measured by correlating the signals with a 90° phase shift introduced into one of the signal paths.
The real (even) component is measured by correlating the pair of signals without a relative phase delay; the imaginary (odd) component is measured by correlating the signals with a $90\degr$ phase shift introduced into one of the signal paths.
The averaged correlated output of the interferometer is called the (see eq. [1]])
The averaged correlated output of the interferometer is called the (see eq. \ref{eqn:vis}] ])
and is the fundamental data product.
and is the fundamental data product.
DASI is a l3-element interferometer operating at 26— GHz.
DASI is a 13-element interferometer operating at 26--36 GHz.
The 13 antenna elements are arranged 1n a threefold symmetric configuration on à common mount which can point in azimuth and elevation.
The 13 antenna elements are arranged in a threefold symmetric configuration on a common mount which can point in azimuth and elevation.
The mount is also able to rotate the array of horns about the line of sight to provide additional (η.v) coverage as well as the ability to perform consistency checks.
The mount is also able to rotate the array of horns about the line of sight to provide additional $(u,v)$ coverage as well as the ability to perform consistency checks.
The 13 elements provide 78 baselines with baseline lengths in the range 25-121 em.
The 13 elements provide 78 baselines with baseline lengths in the range 25–121 cm.
The configuration of the horns was chosen to provide dense coverage of the CMB angular power spectrum from 100«/900.
The configuration of the horns was chosen to provide dense coverage of the CMB angular power spectrum from $100 < l < 900$.
Each DASI antenna consists of à 20-cm aperture-diameter lensed corrugated horn which defines the ~374 FWHM field of view of the instrument.
Each DASI antenna consists of a 20-cm aperture-diameter lensed corrugated horn which defines the $\sim 3 \fdg 4$ FWHM field of view of the instrument.
The receivers use cooled low-noise high electron mobility transistor (HEMT) amplifiers ((Pospiezalski 1995). and have system noise temperatures referred to above the atmosphere ranging from 18 K to 35 K at the center of the band.
The receivers use cooled low-noise high electron mobility transistor (HEMT) amplifiers \markcite{pospieszalski95}( (Pospiezalski 1995), and have system noise temperatures referred to above the atmosphere ranging from 18 K to 35 K at the center of the band.
The receivers downeonvert the 26-36 GHz RF band to a 2-12 GHz IF band.
The receivers downconvert the 26–36 GHz RF band to a 2–12 GHz IF band.
Each receiver IF ts further split and downconverted to ten | GHz wide bands centered at 1.5 GHz.
Each receiver IF is further split and downconverted to ten 1 GHz wide bands centered at 1.5 GHz.
An analog correlator ((Padin 2001b) processes the 1 GHz bands into 780 complex visibilities.
An analog correlator \markcite{padin00}( 2001b) processes the 1 GHz bands into 780 complex visibilities.
The stability of the instrument. its location at the South Pole. and the fact that its mount is fully steerable. have given us great flexibility in designing and adapting our observing strategy.
The stability of the instrument, its location at the South Pole, and the fact that its mount is fully steerable, have given us great flexibility in designing and adapting our observing strategy.
We are able to choose fields to avoid foreground contamination, balance sensitivity and sky coverage. and observe in patterns that reject ground and other spurious signals while producing datasets containing correlations which are computationally tractable.
We are able to choose fields to avoid foreground contamination, balance sensitivity and sky coverage, and observe in patterns that reject ground and other spurious signals while producing datasets containing correlations which are computationally tractable.
CMB fields were observed during the period spanning 05 May-07 November 2000.
CMB fields were observed during the period spanning 05 May–07 November 2000.
The data presented here comprise 97 days of observation. representing an observing efficiency of better than (of the days devoted exclusively to CMB. observations). with the remainder lost to hardware maintenance and repairs.
The data presented here comprise 97 days of observation, representing an observing efficiency of better than (of the days devoted exclusively to CMB observations), with the remainder lost to hardware maintenance and repairs.
Observations were never prevented due to weather. and only of data were lost due to weather based edits. confirming previous assessments of the exceptional quality of the site ((Lay 2000: Chamberlin. Lane. Stark 1997).
Observations were never prevented due to weather, and only of data were lost due to weather based edits, confirming previous assessments of the exceptional quality of the site \markcite{lay98,chamberlin97}( 2000; Chamberlin, Lane, Stark 1997).
The presence of near-field ground contamination at a level well above the CMB signal limits our ability to track single fields over a wide range in azimuth.
The presence of near-field ground contamination at a level well above the CMB signal limits our ability to track single fields over a wide range in azimuth.
Repeated tracks over the full azimuth range show a strong variation of the ground with direction. with amplitudes of tens of Jy on some of the shortest baselines. but there is little evidence for time variability on periods as long as five days.
Repeated tracks over the full azimuth range show a strong variation of the ground with direction, with amplitudes of tens of Jy on some of the shortest baselines, but there is little evidence for time variability on periods as long as five days.
One advantage of observing near the South Pole ts that sources track at a constant elevation. which enables us to observe several sources at constant elevation over a given range in azimuth.
One advantage of observing near the South Pole is that sources track at a constant elevation, which enables us to observe several sources at constant elevation over a given range in azimuth.
Observations were divided among 4 constant declination (elevation) rows of 8 fields. on a regular hexagonal grid spaced by Ih in right ascension. and 67 in declination.
Observations were divided among 4 constant declination (elevation) rows of 8 fields, on a regular hexagonal grid spaced by 1h in right ascension, and $6\degr$ in declination.
The grid center was selected to avoid the Galactic plane and to coincide with a global minimum in the IRAS 100. μπα map of the southern sky.
The grid center was selected to avoid the Galactic plane and to coincide with a global minimum in the IRAS 100 $/mu$ m map of the southern sky.
Each field in à row was observed over the same azimuth range. leading to a nearly identical ground contribution.
Each field in a row was observed over the same azimuth range, leading to a nearly identical ground contribution.
The elevation of the rows are 617.677.557.497, which we label the A. B. C and D rows for the order in which they were observed (see Paper I for full coordinates).
The elevation of the rows are $61\degr, 67\degr, 55\degr, 49\degr$, which we label the A, B, C and D rows for the order in which they were observed (see Paper I for full coordinates).
The field separation of Ih in RA represents a compromise between immunity to time variability of the ground signal and a desire to minimize inter-field correlations.
The field separation of 1h in RA represents a compromise between immunity to time variability of the ground signal and a desire to minimize inter-field correlations.
A given field row was observed daily over two azimuth
A given field row was observed daily over two azimuth
onto white cwarfs is reviewed in some detail ancl updated with old and new theoretical and empirical considerations.
onto white dwarfs is reviewed in some detail and updated with old and new theoretical and empirical considerations.
oth lxoester(1976)and Wesemael(1970). were among the first to. consider the issue of interstellar accretion onto helium-rich. white dwark. both concluding that the atmospheres of these stars (both type DB and DC) were with Boncli-Llovle accretion of interstellar hyvelrogen.
Both \citet{koe76} and \citet{wes79} were among the first to consider the issue of interstellar accretion onto helium-rich white dwarfs, both concluding that the atmospheres of these stars (both type DB and DC) were with Bondi-Hoyle accretion of interstellar hydrogen.
That is. in most cases a single interstellar cloud encounter lasting around. 107 vvr would sullice to transform a DD (or DC) star into à DA (Ixoester1976).
That is, in most cases a single interstellar cloud encounter lasting around $ ^5$ yr would suffice to transform a DB (or DC) star into a DA \citep{koe76}.
. Therefore. the existence of more than 1000 hyvdrogen-poor white cdwarls in the SDSS (Eisensteinctal.2006). argues strongly. against the Boncli-Llovle tvpe accretion of interstellar hydrogen.
Therefore, the existence of more than 1000 hydrogen-poor white dwarfs in the SDSS \citep{eis06} argues strongly against the Bondi-Hoyle type accretion of interstellar hydrogen.
Regardless of particulars relevant only to the metal-enriched. white dwarf varieties. the lack of hydrogen in helium atmosphere white chvarfs is of great. observational significance.
Regardless of particulars relevant only to the metal-enriched white dwarf varieties, the lack of hydrogen in helium atmosphere white dwarfs is of great observational significance.
lligh resolution and high signal-to-noise spectroscopy reveals roughly. equal numbers of DB stars where hydrogen is either uncletected or. detected but deficient by a few to several orders of magnitude relative to helium (Vossetal.2007).
High resolution and high signal-to-noise spectroscopy reveals roughly equal numbers of DB stars where hydrogen is either undetected or detected but deficient by a few to several orders of magnitude relative to helium \citep{vos07}.
. The highest mass of hydrogen seen in a DB (metal-free) atmosphere is around δι1Y? gg ina IxIx star. while the older and cooler counterparts studied here (the DZA stars) contain up to 41073 gg. This higher mass of atmospheric hvdrogen would be accreted in under 10" ver at the interstellar Bondi-Hloyle rate assumed bv Dupuisetal. (1993a).. and hence this cannot be the
The highest mass of hydrogen seen in a DB (metal-free) atmosphere is around $8\times10^{22}$ g in a K star, while the older and cooler counterparts studied here (the DZA stars) contain up to $4\times10^{24}$ g. This higher mass of atmospheric hydrogen would be accreted in under $^6$ yr at the interstellar Bondi-Hoyle rate assumed by \citet{dup93a}, , and hence this cannot be the
In the light of recent new results on the behaviour of the Eddington ratio ο às a function of the SMBH mass and of the redshift. we have revisited the predictions of the model presented in Papl. This model naturally predicts an anti-correlation between the fraction fis. of X-ray obscured. Compton-thin AGN. and Myj. provided there is interstellar molecular gas with appropriate values of surface density X approximately within 25 to 450 pe (depending on Λάρι) from the active nucleus.
In the light of recent new results on the behaviour of the Eddington ratio $\lambda$ as a function of the SMBH mass and of the redshift, we have revisited the predictions of the model presented in PapI. This model naturally predicts an anti-correlation between the fraction $f_{abs}$ of X-ray obscured, Compton-thin AGN, and $M_{BH}$, provided there is interstellar molecular gas with appropriate values of surface density $\Sigma$ approximately within 25 to 450 pc (depending on $M_{BH}$ ) from the active nucleus.
We used the newly found (Mp.2) dependence (z from the local universe to 0.7) and a distribution of X. irrespective of the morphological type of the host galaxies and invariant with z. to convert that prediction into one versus Ly.
We used the newly found $\lambda(M_{BH},z)$ dependence $z$ from the local universe to 0.7) and a distribution of $\Sigma$, irrespective of the morphological type of the host galaxies and invariant with $z$, to convert that prediction into one $\it versus$ $L_X$.
The results appear to reproduce. at least qualitatively. the anti-correlation between fp, and Ly claimed by several authors. but not the still controversial claim of an increase in fup, With z.
The results appear to reproduce, at least qualitatively, the anti-correlation between $f_{abs}$ and $L_X$ claimed by several authors, but not the still controversial claim of an increase in $f_{abs}$ with $z$.
The latter increase. particularly in the form claimed by LF2005 on which we did concentrate. might possibly be reproduced by assuming that the typical values of X do increase with the redshift and. furthermore. that in the early type galaxies. at least in the more massive ones. the presence of this quantity 15 correlated in time with the activity of their nucleus.
The latter increase, particularly in the form claimed by LF2005 on which we did concentrate, might possibly be reproduced by assuming that the typical values of $\Sigma$ do increase with the redshift and, furthermore, that in the early type galaxies, at least in the more massive ones, the presence of this quantity is correlated in time with the activity of their nucleus.
At the highest luminosities. however. according to our model. fp, can hardly exceed For the sake of completeness. it should be noted that the results just summarised are sensitive to the presence of a cut-off in the ACMpij.z) distribution. and they correspond to the case where the cut-off is placed at the physical limit ΞΙ. Steps forward. that are relevant not only to the model discussed in this paper can be made at least in two ways.
At the highest luminosities, however, according to our model, $f_{abs}$ can hardly exceed For the sake of completeness, it should be noted that the results just summarised are sensitive to the presence of a cut-off in the $\lambda(M_{BH},z)$ distribution, and they correspond to the case where the cut-off is placed at the physical limit $\lambda$ Steps forward, that are relevant not only to the model discussed in this paper can be made at least in two ways.
The first and most obvious one requires systematic. investigation of the amount of molecular gas in an appropriately selected sample of early type galaxies. both active and inactive. out to redhifts of 0.5 at least.
The first and most obvious one requires systematic investigation of the amount of molecular gas in an appropriately selected sample of early type galaxies, both active and inactive, out to redhifts of 0.5 at least.
The line strength and profile may not tell us enough about the spatial distribution. but much about the amount and the kinematics of the gas.
The line strength and profile may not tell us enough about the spatial distribution, but much about the amount and the kinematics of the gas.
The former property needs to wait for observations with a large interferometric array of telescope. such as ALMA.
The former property needs to wait for observations with a large interferometric array of telescope, such as ALMA.
The second step is somewhat indirect and therefore requires interpretation.
The second step is somewhat indirect and therefore requires interpretation.
The method ts described and applied to a handful of broad-line objects in Matolino et al. (
The method is described and applied to a handful of broad-line objects in Maiolino et al. (
2007).
2007).
It allows an estimate ofC to be inferred for each object on the basis of measurements of infrared spectral features due to dust reradiation of the ultraviolet from the active nucleus.
It allows an estimate of $C$ to be inferred for each object on the basis of measurements of infrared spectral features due to dust reradiation of the ultraviolet from the active nucleus.
It is worth noting that Maiolino et al. (
It is worth noting that Maiolino et al. (
2007) find an anti-correlation between C
2007) find an anti-correlation between $C$
these polarisation parameters was of z120°.
these polarisation parameters was of $\approx 120^\circ$.
Analysis then proceeded by applying this sct of values as a starting point for the fit to each individual night.
Analysis then proceeded by applying this set of values as a starting point for the fit to each individual night.
The parameters of the constant component. were allowed to vary within the same error range of those of the variable component as quoted in Table 1. since they indicate the [limiting accuracy of the model fitting.
The parameters of the constant component were allowed to vary within the same error range of those of the variable component as quoted in Table 1, since they indicate the limiting accuracy of the model fitting.
We looked for the values of the variable component on each night which minimise the residuals while keeping our pre-determined bounds for the constant component. i.e. Locus2 19-21 my. posu.8 1-5 ο and Gua& 110-1407.
We looked for the values of the variable component on each night which minimise the residuals while keeping our pre-determined bounds for the constant component, i.e. $I_{\rm{cons}} \approx$ 19-21 mJy, $p_{\rm{cons}} \approx$ 1-5 $\%$ and $\theta_{\rm{var}} \approx$ $^\circ$.
An accept:e solution was found for every night. and the residuals were kept below for each dataset.
An acceptable solution was found for every night, and the residuals were kept below for each dataset.
A σους indication of the appropriateness of this moclel in describing the entire dataset. is that a reasonable fit was obtained [for each night without the need [for the parameters of the constant component to depart. from the boundaries mentioned above.
A good indication of the appropriateness of this model in describing the entire dataset is that a reasonable fit was obtained for each night without the need for the parameters of the constant component to depart from the boundaries mentioned above.
Such boundaries are regarded as indicating the range of accuracy within which the components parameters can be regarded. as “constant” since they rellect the intrinsic uncertainty of the fitting as defined by the errors in the parameters of Table 1.
Such boundaries are regarded as indicating the range of accuracy within which the component's parameters can be regarded as “constant”, since they reflect the intrinsic uncertainty of the fitting as defined by the errors in the parameters of Table 1.
Final confidence intervals for the polarisation parameters of the constant component were estimated [rom the range of night-to-night variations in its best-fit parameters. and are given by poons—42A and Goon.=1302:107.
Final confidence intervals for the polarisation parameters of the constant component were estimated from the range of night-to-night variations in its best-fit parameters, and are given by $p_{\rm{cons}} = 4 \pm 1\%$ and $\theta_{\rm{cons}} = 130^\circ \pm 10^\circ$.
They are therefore compatible with a set of constant. parameters throughout the campaing within the observational errors.
They are therefore compatible with a set of constant parameters throughout the campaing within the observational errors.
This best-fit model is shown in Figure 3.
This best-fit model is shown in Figure 3.
For all nights we had fo.1. indicating that the background. component dominates the photometric Dux emission.
For all nights we had $I_{\rm{v/c}} < 1$, indicating that the background component dominates the photometric flux emission.
Phe values of duu derived for cach individual night are presented in Table 1. corresponding to Liu.
The values of $I_{\rm{var}}$ derived for each individual night are presented in Table 1, corresponding to $I_{\rm{cons}}$.
The derived: parameters for the constant component are found to match the regular values of the polarisation compiled by Tommasictal.(2001). for PISS 2155-304. suggesting its association with a persistent optical jet component.
The derived parameters for the constant component are found to match the regular values of the polarisation compiled by \cite{tommasi} for PKS 2155-304, suggesting its association with a persistent optical jet component.
The degree of polarisation posu. is also similar to the minimum values measured. for this source at. 43 Cillz and in historical optical data. ancl the corresponding position angle is well aligned. with the racio-core EVRA as determined. by Pineretal.(2008).
The degree of polarisation $p_{\rm{cons}}$ is also similar to the minimum values measured for this source at 43 GHz and in historical optical data, and the corresponding position angle is well aligned with the radio-core EVPA as determined by \cite{b21}.
. This. coincidence also attests to the presence of a field component in the jet which is common both to the radio ancl optical wavelengths and persistent in time. and. whose direction is transverse to the How. às expected from a shock-compressed tangled field.
This coincidence also attests to the presence of a field component in the jet which is common both to the radio and optical wavelengths and persistent in time, and whose direction is transverse to the flow, as expected from a shock-compressed tangled field.
From the second night of the campaign onwards. the position angle of the variable component rotated continuously from 70 (i.e. approximately 90° mis-aligned with the jet-projected PA.) to 2120°. in close alignment with the direction. of the persistent jet. component.
From the second night of the campaign onwards, the position angle of the variable component rotated continuously from $70^\circ$ (i.e. approximately $90^\circ$ mis-aligned with the jet-projected P.A.) to $\approx 120^\circ$, in close alignment with the direction of the persistent jet component.
The rotation of C4, could be interpreted as the gradual alignment of the field of a new “blob” of material. as it encounters a shock in the core that re-organises its field.
The rotation of $\theta_{\rm{var}}$ could be interpreted as the gradual alignment of the field of a new “blob” of material, as it encounters a shock in the core that re-organises its field.
The masiniun value observed for the source's polarisation degree coincides with the epochs of greatest alignment between the two fields. and the start of the rotation in 6, marks the onset of the increase on the baseline photometric Hux seen towards the end of the campaign.
The maximum value observed for the source's polarisation degree coincides with the epochs of greatest alignment between the two fields, and the start of the rotation in $\theta_{\rm{var}}$ marks the onset of the increase on the baseline photometric flux seen towards the end of the campaign.
Such a scenario. where both optical position angles €, and Goon. tend to align with the direction of the radio EVPA when the observed polarisation is high. was considered. before by Valtaojaetal.(1991b). for the quasar 3€ 273 during a radio-to-optical Hare.
Such a scenario, where both optical position angles $\theta_{\rm{var}}$ and $\theta_{\rm{cons}}$ tend to align with the direction of the radio EVPA when the observed polarisation is high, was considered before by \cite{valb} for the quasar 3C 273 during a radio-to-optical flare.
In such à scenario a correlation is expected. between the optical and polarised Iuxes which is marginally observed in our dataset. and more observations at more active source states are necessary to better establish the validitv of the correlation for this object.
In such a scenario a correlation is expected between the optical and polarised fluxes which is marginally observed in our dataset, and more observations at more active source states are necessary to better establish the validity of the correlation for this object.
As noted. before. the observed flux variability happens on two cilferent timescales. its amplitude being dominated. by intranight variability. superimposed on a background. level that steadily increases towards the end of the campaign. and which we have associated to the evolution of the variable (or shocked) component in the model of the previous section.
As noted before, the observed flux variability happens on two different timescales, its amplitude being dominated by intranight variability, superimposed on a background level that steadily increases towards the end of the campaign, and which we have associated to the evolution of the variable (or shocked) component in the model of the previous section.
‘To try to identify the physical origin of these variations and in particular the nature of the (ux microvariabilitv. we observe that the intranight flux changes were accompanied bv changes in the spectral index.
To try to identify the physical origin of these variations and in particular the nature of the flux microvariability, we observe that the intranight flux changes were accompanied by changes in the spectral index.
The source presented colour variations both in intranight. timescales and in the nightly averages.
The source presented colour variations both in intranight timescales and in the nightly averages.
Phe intranight (VoZ) colours varied in the range ο. 0.27. with ercatest amplitude in the third night of the campaign. when variability was the greatest.
The intranight $(V-I)$ colours varied in the range 0.12 – 0.27, with greatest amplitude in the third night of the campaign, when variability was the greatest.
Colour variations can be linked to radiative cooling of electrons in à magnetisecd plasma. implving svnchrotron Lifetimes of the order of the intraday. timescales of a few hours.
Colour variations can be linked to radiative cooling of electrons in a magnetised plasma, implying synchrotron lifetimes of the order of the intraday timescales of a few hours.
The svnchrotron lifetime in the observer's Frame. written in terms of the observed photon frequency in units of CGllz. &e;gz. is eiven by (Pacholezvk1970):: For Fac equal to the timescales of intranight variations in the Ro band. and using typical Doppler factors for PINS 2155-304 of about 9~30 (e.g. as for the compact components in Watarzviskietal. 2008)) we obtain a magnetic field 60.5 G for the variable component.
The synchrotron lifetime in the observer's frame, written in terms of the observed photon frequency in units of GHz, $\nu_{GHz}$ , is given by \citep{pac}: For $t_{\rm{sync}}$ equal to the timescales of intranight variations in the R band, and using typical Doppler factors for PKS 2155-304 of about $\delta \sim 30$ (e.g. as for the compact components in \citealt{kat}) ) we obtain a magnetic field $B \lesssim 0.5$ G for the variable component.
The fact that we see such changes in colour simultaneously with the intranight variations. suggests they can be taken as a direct. signature of particle acceleration and cooling at the source. with Gio9faac.
The fact that we see such changes in colour simultaneously with the intranight variations, suggests they can be taken as a direct signature of particle acceleration and cooling at the source, with $t_{\rm{acc}} < t_{\rm{sync}}$.
An upper limit to the size of the acceleration. region can then be set. to ro<Olsacc/(l|2)zmL10i"em~5«10Spe.
An upper limit to the size of the acceleration region can then be set to $r_{\rm{s}} < \delta t_{\rm{sync}}c/(1+z) \approx 10^{16} \rm{cm} \sim 5 \times 10^{-3} \rm{pc}$.
This limit is in accordance with predictions for the thickness of shocks given by Marscher&Gear(1985). and points to an origin for the flux microvariability as the result of particle acceleration taking place at a shock front. with high magnetic field due to plasma compression.
This limit is in accordance with predictions for the thickness of shocks given by \cite{shock} and points to an origin for the flux microvariability as the result of particle acceleration taking place at a shock front, with high magnetic field due to plasma compression.
Magnetic fields of this order have also been considered by Marscher&Gear(1985). as typical estimates for the field intensity in blazar cores. and are of the same order of magnitude of those recently found to explain he low-activity state of FPermi-detected blazars Collab. 2010)..
Magnetic fields of this order have also been considered by \cite{shock} as typical estimates for the field intensity in blazar cores, and are of the same order of magnitude of those recently found to explain the low-activity state of Fermi-detected blazars \citep{Fermib}. .
In the SED model of Ixatarzvüskietal.(2008) such values for the B-field ancl Doppler factor are also associated with the variable shocked components. as opposed to the extended jot which had lower values for bothαςΠΙΟΤΟΥ».
In the SED model of \cite{kat} such values for the -field and Doppler factor are also associated with the variable shocked components, as opposed to the extended jet which had lower values for bothparameters.
Our values forD. are in fact an order of magnitude ügher than those derived by Aharonianetal.(2000). from
Our values for$\mathbf{B}$ are in fact an order of magnitude higher than those derived by \cite{b1c} from
specifically for the cold component in the inter-arm regions in Fig.
specifically for the cold component in the inter-arm regions in Fig.
7.
7.
Clearly. the inter-arm gas in the multi-phase simulation reaches densities over a magnitude higher than for the single phase simulation.
Clearly the inter-arm gas in the multi-phase simulation reaches densities over a magnitude higher than for the single phase simulation.
From local simulations which include cooling. gas with densities 710 ? is found to be cold. i.e. 100-200 Ix or less (Fig.
From local simulations which include cooling, gas with densities $>10$ $^{-3}$ is found to be cold, i.e. 100-200 K or less (Fig.
1. 22).
1, \citealt{Glover2006a}) ).
We therefore expect a significant proportion of the cold gas entering the spiral arms to remain cold in a multi-phase medium.
We therefore expect a significant proportion of the cold gas entering the spiral arms to remain cold in a multi-phase medium.