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All these galaxies have been previously studied: by other authors. who have determined PC fractions using the more traditional procedure of adopting the spectrum. of an elliptical galaxy (or the bulge of a normal spiral) as a starlight template.
All these galaxies have been previously studied by other authors, who have determined FC fractions using the more traditional procedure of adopting the spectrum of an elliptical galaxy (or the bulge of a normal spiral) as a starlight template.
According to Tran (1995a). for instance. the FC accounts for 25 per cent of t1e light. at ASSOOA iin àrk1210. while in ΔΗΚΟ88 it contributes with 27 per cent.
According to Tran (1995a), for instance, the FC accounts for 25 per cent of the light at $\lambda$ in Mrk1210, while in Mrk348 it contributes with 27 per cent.
Quir results for Mrk1210 show that the FC contributes with only =10 per cent of the ight at ADSTOA.. while stars with 10 vr or younger. contribute with 11 per centnote that hese percentage contribuions are approximately constant over the 5500.5870 iinerval.
Our results for Mrk1210 show that the FC contributes with only $\approx 10$ per cent of the light at $\lambda$, while stars with 10 Myr or younger, contribute with 11 per cent—note that these percentage contributions are approximately constant over the 5500–5870 interval.
For MrR348 we find only z3 per cent contribution rom an PC and z2 per cent [rom voung stars.
For Mrk348 we find only $\approx 3$ per cent contribution from an FC and $\approx 2$ per cent from young stars.
For 3633. Ixoski (1978) found a 19 per cent FC contribution to the ight at A5500A.. while the svnthesis vields less than 1 per cent from either FC or voung stars.
For 3C33, Koski (1978) found a 19 per cent FC contribution to the light at $\lambda$, while the synthesis yields less than $1$ per cent from either FC or young stars.
For Mrk573. Mrk607 and GC1358. Ίναν (1994) estimated an FC fraction at AJ400:N ool 20 per cent. 10 per cen and 17 per cent. respectively.
For Mrk573, Mrk607 and NGC1358, Kay (1994) estimated an FC fraction at $\lambda$ of 20 per cent, 10 per cent and 17 per cent, respectively.
Our results show only a 1 per cent FC for Mrk573 and even ess for NGC1858 and Mrk607. while voung stars contribute Byh 2 per cent to the spectrum of Mrk573 at ALS and <1 per cent for NGC135s8 and. ALrk607.
Our results show only a 1 per cent FC for Mrk573 and even less for NGC1358 and Mrk607, while young stars contribute with $\approx$ 2 per cent to the spectrum of Mrk573 at $\lambda$ and $<1$ per cent for NGC1358 and Mrk607.
1n order to check if the smaller FC contribution found by us is real. or if it is just due to the spectral synthesis technique used. we repeated the spectral synthesis allowing only FC contributions. at A5SSTOA.. larger than 25 per cent [or Mrk1210 and. δικός. 20 per cent for 3633 and Mrk573. 15 per cent for NGCT1358 ane 10 per cent. for. Mrk607.
In order to check if the smaller FC contribution found by us is real, or if it is just due to the spectral synthesis technique used, we repeated the spectral synthesis allowing only FC contributions, at $\lambda$, larger than 25 per cent for Mrk1210 and Mrk348, 20 per cent for 3C33 and Mrk573, 15 per cent for NGC1358 and 10 per cent for Mrk607.
Suppose the magnetic field is distributed on the disk from r=ry tor =rs. where rac»Dorgy then At the inner boundary of the disk. where r=ry.< we have g=gms.
Suppose the magnetic field is distributed on the disk from $r=r_1$ to $r=r_2$ , where $r_2>r_1\ge r_{ms}$, then At the inner boundary of the disk, where $r=r_{ms}\le r_1$, we have $g =g_{ms}$.
For a (hin Keplerian disk. (he “no-torque inner boundary condition is an excellent approximation (NovikovandThorne1973:MuchotrzebPaczviski1982:AbramowiezIxato1989:2000;Armitage.Revnold.andChiang 2000).
For a thin Keplerian disk, the “no-torque inner boundary condition” is an excellent approximation \citep{nov73,muc82,abr89,pac00,arm00}.
. Though recently this boundary condition has been challenged when there is à magnetic connection between the disk and the malerial in (he transition region (Ixrolik1999:ILawleyandIxrolik2000).. in this paper we do not consider such a magnetic connection thus we adopt the “no-torque inner boundary condition”.
Though recently this boundary condition has been challenged when there is a magnetic connection between the disk and the material in the transition region \citep{kro99,haw00}, in this paper we do not consider such a magnetic connection thus we adopt the “no-torque inner boundary condition”.
The appearance of a magnetic coupling between the black hole aud (he disk does nol introduce a torque at the inner boundary.
The appearance of a magnetic coupling between the black hole and the disk does not introduce a torque at the inner boundary.
This is because of (hat the integration in equation (15)) always vanishes at r=rj, for anv Lf.
This is because of that the integration in equation \ref{int}) ) always vanishes at $r=r_{ms}$ for any $H$.
In other words. the torque produced bv the magnetic coupling propagates outward only in the disk.
In other words, the torque produced by the magnetic coupling propagates outward only in the disk.
Therefore. in the following discussion we take gy.=0.
Therefore, in the following discussion we take $g_{ms} = 0$.
Then we have E ⊏≺⇂∏≀↧↴∐∪∐⊔≺≨↕⊔↖≺↽↔↴↕∖↽≼↲⋟∖⊽⊔∐↲↕⋅≀↧↴≺∐≀↕↴∐∪∐∐∏⇀↸∪↓≯⊔∐↲≺∐⋟∖⊽↳↽⋅≼↲≺⇂∏≀↧↴∐∪∐⊔⊺↕⋝↕⋟≸↽↔↴↕∖⇁≼↲⋟∖⊽⊔∐↲↕∐∩↲↕⋅∐≀↧↴↥ viscous torque of the disk.
Then we have and Equation \ref{flux1}) ) gives the radiation flux of the disk, equation \ref{torque1}) ) gives the internal viscous torque of the disk.
The first terms on the right hand sides of equation (16)) and equation (17)) are the familiar results for a standard accretion disk (PageaudThorne1974).
The first terms on the right hand sides of equation \ref{flux1}) ) and equation \ref{torque1}) ) are the familiar results for a standard accretion disk \citep{pag74}.
. The integration of equation (8)) gives the power (1.e.. the Iuminositv) of the disk. which is (heenerey radiated bv (he disk per unit lime as measured by an observer al infinity (Thorne1974)
The integration of equation \ref{con_ener}) ) gives the power (i.e., the luminosity) of the disk, which is theenergy radiated by the disk per unit time as measured by an observer at infinity \citep{tho74}
133 eulaxies with spectroscopic redshifts iu the GOODS CDFS field. using the photometric redshift nethod of itez 2000 which incorporates Bayesian maguitude/redslitt xiors.
433 galaxies with spectroscopic redshifts in the GOODS CDF–S field, using the photometric redshift method of tez 2000 which incorporates Bayesian magnitude/redshift priors.
We find an overall performance o(A)=0,11. with better performance for some subsamples.
We find an overall performance $\sigma (\Delta) = 0.11$, with better performance for some subsamples.
The raction of catastrophic redshift outliers is <1lO%.. and is substantially simaller for galaxies with high values of he Bavesiaui ODDS parameter.
The fraction of catastrophic redshift outliers is $<$, and is substantially smaller for galaxies with high values of the Bayesian ODDS parameter.
We see uo strong trend in the performance of the photometric redshifts versus uaenitudes. except for an increase in the outlier rate.
We see no strong trend in the performance of the photometric redshifts versus magnitudes, except for an increase in the outlier rate.
Eaploving a redshift/uaguitude prior in this process secus fo be crucial iji reducing the scatter between photometric aud specticscopic redshifts.
Employing a redshift/magnitude prior in this process seems to be crucial in reducing the scatter between photometric and spectroscopic redshifts.
We have applied the amethod to two subsamples of galaxies of particular interest: EROs and ACNs.
We have applied the method to two sub–samples of galaxies of particular interest: EROs and AGNs.
The results for EROs are inore accurate {o(A)=0.051) than for faint galaxies in general
The results for EROs are more accurate $\sigma(\Delta) = 0.051$ ) than for faint galaxies in general.
They are somewhat less accurate or the N-rav sources (ic. ACNS)- (0CA)= 0.101). but eood enough to be useful for many applications.
They are somewhat less accurate for the X-ray sources (i.e. AGNs)- $\sigma (\Delta)=0.104$ ), but good enough to be useful for many applications.
which is a simple mocification of the analogous equation (5).
which is a simple modification of the analogous equation (5).
On the righthaud side of equation (12). the factor f,ntUo7 is necessarily. greater than unity.
On the righthand side of equation (12), the factor $f_m^{-2}$ is necessarily greater than unity.
" Ou ∖∖∕↓∣∣⋅⋅ ⋅ {δι
On the other hand, is, in practice, so small that.
Consider. for example. the models described iu Sectioi 2.2. which lad a stellar mass distribtion approorate Lo‘the infant Pleiacles.
Consider, for example, the models described in Section 2.2, which had a stellar mass distribution appropriate for the infant Pleiades.
Here.1.
Here,.
.. lu our cluster. we find empirically that tle mininiuin Wass in auy newly formed jnaryv ds:: the correspouxine [ναι eis Sκ10.
In our cluster, we find empirically that the minimum mass in any newly formed binary is; the corresponding $f_m$ -value is $8\times 10^{-3}$.
Equation (12) then prediets that0.2.. iu good agreemeltw ih our numerical residts.
Equation (12) then predicts that, in good agreement with our numerical results.
ΤΙis derivation is. of course. highly siuplifiecd. aid the quautitative result above should not be even {«)0 Latch weigh.
This derivation is, of course, highly simplified, and the quantitative result above should not be given too much weight.
The relative velociy of au ei"OUIer liu the coο willty dically be larger than the rest of the ¢Uster. alle the core cleusity will ye hieher than the average.
The relative velocity of an encounter in the core will typically be larger than in the rest of the cluster, and the core density will be higher than the average.
A more complete erivatiou of /, wotk also cousider the pivsical basis fo‘the minimll nass it.
A more complete derivation of $t_b$ would also consider the physical basis for the minimum mass $m$.
Presumably. this init Is set by the "ulte at which οναι]‘al friction alle)ws stars of various mass to drift iuwaid.
Presumably, this limit is set by the rate at which dynamical friction allows stars of various mass to drift inward.
Ve wil --ot. emubellish the areument aloiϱ these lien. jut simply tote that equation (12) adds stification for our Nall polis: (1) The rate of bilary* formation is very seusitive to the stellar jass Cisibution. and (2) even iu hypotletical clusters composed eitirely of single stars. binaries orn relatively quickly.
We will not embellish the argument along these lines, but simply note that equation (12) adds justification for our main points: (1) The rate of binary formation is very sensitive to the stellar mass distribution, and (2) even in hypothetical clusters composed entirely of single stars, binaries form relatively quickly.
It is oIv by adopting the extreme assumption that these single stars have identical mass that binary [oration can be delayed to the point of core A hard binary that resides within a cluster. no matter how i formecl. adds euergy to the whole system.
It is only by adopting the extreme assumption that these single stars have identical mass that binary formation can be delayed to the point of core A hard binary that resides within a cluster, no matter how it formed, adds energy to the whole system.
The process. like the creation of new pairs. is a three-body interaction.
The process, like the creation of new pairs, is a three-body interaction.
As a result of the eucounter. the binary usually tightens aud releases energy.
As a result of the encounter, the binary usually tightens and releases energy.
This leating accounts for tlie expansion of both the mass and muuber shells in Figure 3. forf,.
This heating accounts for the expansion of both the mass and number shells in Figure 3, for.
. Expansion driven by binaries is global. aud ciffers qualiatively from the dual contraction aud expausiοι seen in the single-mass moclel (Figure 1).
Expansion driven by binaries is global, and differs qualitatively from the dual contraction and expansion seen in the single-mass model (Figure 1).
This clifferei is also apparent when we view the evolution of the clusters ageMDregate ellerey.
This difference is also apparent when we view the evolution of the cluster's aggregate energy.
First. we need to iuguishénferiel audlop-level euergies.
First, we need to distinguish and energies.
In the first category is the gravitational biudiug euergy ο Ἡ binary. ancl the kinetic enerey of both σοιiponent stars with respect to their center of mass.
In the first category is the gravitational binding energy of each binary, and the kinetic energy of both component stars with respect to their center of mass.
I itle top-level category are the center-of-1nass kinetic energies of all bound stellar systems. whetlier siele or multiple. and the gravitational potenial energy of tliis array.
In the top-level category are the center-of-mass kinetic energies of all bound stellar systems, whether single or multiple, and the gravitational potential energy of this array.
Thus. the kinetic energy fV, c«)isklered previously was actually a top-level quantity.
Thus, the kinetic energy $K_r$ considered previously was actually a top-level quantity.
The clusters total energy Ey is the sum of the two coutributious: Here. we are ignoring the relatively sinall amount of energy. carried olf by escaping stars.
The cluster's total energy $E_0$ is the sum of the two contributions: Here, we are ignoring the relatively small amount of energy carried off by escaping stars.
Iu
In
We therefore choose the same approach to interpret the low X-ray flux states presented in Fig. 8,
We therefore choose the same approach to interpret the low X-ray flux states presented in Fig. \ref{SED0},
including the SED from 15 January 2006, the center of the core period of our multiwavelength campaign described here.
including the SED from 15 January 2006, the center of the core period of our multiwavelength campaign described here.
For this purpose, we use the equilibrium version of the leptonic SSC + EC model of Bóttcher&Chiang(2002).
For this purpose, we use the equilibrium version of the leptonic SSC + EC model of \cite{bc02}.
. In this model, a relativistic electron population is assumed to be in an equilibrium state between on-going injection, escape, and radiative cooling in a spherical emission region of radius R, moving relativistically with bulk Lorentz factor I along the jet.
In this model, a relativistic electron population is assumed to be in an equilibrium state between on-going injection, escape, and radiative cooling in a spherical emission region of radius $R$, moving relativistically with bulk Lorentz factor $\Gamma$ along the jet.
Electrons are injected with a power-law distribution specified by an injection spectral index q and low- and high-energy cutoffs y; and y», respectively.
Electrons are injected with a power-law distribution specified by an injection spectral index $q$ and low- and high-energy cutoffs $\gamma_1$ and $\gamma_2$, respectively.
The radiative cooling is evaluated self-consistently with the radiative output via synchrotron, SSC, and external-Compton emission.
The radiative cooling is evaluated self-consistently with the radiative output via synchrotron, SSC, and external-Compton emission.
Particles are assumed to escape on an energy-independent time scale fese=ηΚ/ο, where 7>1 is a free parameter.
Particles are assumed to escape on an energy-independent time scale $t_{\rm esc} \equiv \eta R/c$, where $\eta > 1$ is a free parameter.
Particle escape will become relevant for particles with energies below a critical break energy y; defined through the condition that {εις=foot, the radiative cooling time scale.
Particle escape will become relevant for particles with energies below a critical break energy $\gamma_b$ defined through the condition that $t_{\rm esc} = t_{\rm cool}$, the radiative cooling time scale.
Escaping particles are leaving the emission region in random directions and will therefore primarily end up in a slower sheath that may be present around the fast spine producing the high-energy emission.
Escaping particles are leaving the emission region in random directions and will therefore primarily end up in a slower sheath that may be present around the fast spine producing the high-energy emission.
The magnetic field B is a free parameter of the model, as are the parameters characterizing the external radiation field.
The magnetic field $B$ is a free parameter of the model, as are the parameters characterizing the external radiation field.
For any given set of parameters, the code also evaluates the ratios of co-moving energy densities in the magnetic field, u, and the equilibrium electron distribution, u;, the corresponding kinetic (L,) and Poynting flux (Lg) powers, and their ratio, ep=u,/u,Lg/Le.
For any given set of parameters, the code also evaluates the ratios of co-moving energy densities in the magnetic field, $u'_B$ and the equilibrium electron distribution, $u'_e$, the corresponding kinetic $L_e$ ) and Poynting flux $L_B$ ) powers, and their ratio, $\epsilon_B \equiv u'_B / u'_e = L_B / L_e$.
For a more detailed description of this equilibrium model see Acciarietal.(2009).
For a more detailed description of this equilibrium model see \cite{acciari09}.
. The model fits only the near-infrared to gamma-ray emission, which is believed to be produced at the sub-pc scales of the jet, where the source is still optically thick at radio wavelengths.
The model fits only the near-infrared to gamma-ray emission, which is believed to be produced at the sub-pc scales of the jet, where the source is still optically thick at radio wavelengths.
The optically thick radio emission is believed to be produced on > pc scales.
The optically thick radio emission is believed to be produced on $>$ pc scales.
Above ~100 GGHz however, the emission comes from sub-parsec scales of a region which is located at or very near the base of the VLBI jet (size of VLBI core at 100 GHz x1015 ccm).
Above $\sim 100$ GHz however, the emission comes from sub-parsec scales of a region which is located at or very near the base of the VLBI jet (size of VLBI core at 100 GHz $\leq 10^{18}$ cm).
We find that the optical through X-ray SED of 15 January 2006 can be well represented with the following model parameters: L,=1.2x10? ergs s“!, y=15x107, y.=10,¢ =3.7,B =08 GR =3.5x10'° cm, =15, and an external radiation field energy density vex,=2.2x107 ergs cm.
We find that the optical through X-ray SED of 15 January 2006 can be well represented with the following model parameters: $L_e = 1.2 \times 10^{45}$ ergs $^{-1}$, $\gamma_1 = 1.5 \times 10^3$, $\gamma_2 = 10^5$, $q = 3.7$, $B = 0.8$ G, $R = 3.5 \times 10^{16}$ cm, $\Gamma = 15$, and an external radiation field energy density $u_{\rm ext} = 2.2 \times 10^{-4}$ ergs $^{-3}$.
In order to reduce the number of free parameters, we choose the observing angle in a way that the Doppler factor D=(T[1—BrcosOops])!-[.
In order to reduce the number of free parameters, we choose the observing angle in a way that the Doppler factor $D \equiv \left( \Gamma [ 1 - \beta_{\Gamma} \cos\theta_{\rm obs} ] \right)^{-1} = \Gamma$.
The magnetic field specified above corresponds to a value slightly below equipartition with eg—0.56.
The magnetic field specified above corresponds to a value slightly below equipartition with $\epsilon_B = 0.56$.
The fit is illustrated by the red solid curve in Fig. 9..
The fit is illustrated by the red solid curve in Fig. \ref{SEDfits}.
While the X-ray spectrum is always produced through SSC emission, the y-ray emission is strongly dominated by the EC emission.
While the X-ray spectrum is always produced through SSC emission, the $\gamma$ -ray emission is strongly dominated by the EC emission.
The model does, indeed, slightly over-predict the soft y-ray (ISGRI) flux.
The model does, indeed, slightly over-predict the soft $\gamma$ -ray (ISGRI) flux.
However, the ISGRI error bars are large, and our fit can still be considered marginally consistent with the ISGRI spectrum.
However, the ISGRI error bars are large, and our fit can still be considered marginally consistent with the ISGRI spectrum.
Remarkably, the other two low-X-ray states can be modeled with essentially identical parameters, changing only the low-energy cutoff of the electron distribution.
Remarkably, the other two low-X-ray states can be modeled with essentially identical parameters, changing only the low-energy cutoff of the electron distribution.
The fits to the P2 and June 2003 spectra shown in Fig.
The fits to the P2 and June 2003 spectra shown in Fig.
9 have been generated choosing γι=750 for P2 and γι=550 for June 2003.
\ref{SEDfits} have been generated choosing $\gamma_1 = 750$ for P2 and $\gamma_1 = 550$ for June 2003.
As already alluded to before, the higher value of the low-energy cutoff may be responsible for the break in the X-ray spectrum seen in the Chandra data of our January 2006 campaign.
As already alluded to before, the higher value of the low-energy cutoff may be responsible for the break in the X-ray spectrum seen in the Chandra data of our January 2006 campaign.
The change in y affects the peak frequencies of all radiation components às yyος y γεςοςVi» and vssc« yt as long as Compton scattering occurs in the Thomson regime near the Compton peaks, which is the case for the parameters adopted here.
The change in $\gamma_1$ affects the peak frequencies of all radiation components as $\nu_{\rm sy} \propto \gamma_1^2$ , $\nu_{\rm EC} \propto \gamma_1^2$, and $\nu_{\rm SSC} \propto \gamma_1^4$ , as long as Compton scattering occurs in the Thomson regime near the Compton peaks, which is the case for the parameters adopted here.
In addition, increasing y, also has the effect of increasing the synchrotron radiation energy density, hence leading to an increasing SSC contribution in the high-energy SED.
In addition, increasing $\gamma_1$ also has the effect of increasing the synchrotron radiation energy density, hence leading to an increasing SSC contribution in the high-energy SED.
In Fig.
In Fig.
9 we included the individual radiation components - synchrotron, SSC, EC (BLR), EC (disk), and the disk emission as dashed, dot-dashed, double-dot-dashed, double-dash-dotted and dotted curves for the 15 January 2006 SED fit.
\ref{SEDfits} we included the individual radiation components - synchrotron, SSC, EC (BLR), EC (disk), and the disk emission as dashed, dot-dashed, double-dot-dashed, double-dash-dotted and dotted curves for the 15 January 2006 SED fit.
This illustrates the dominant SSC contribution to the soft y-ray regime in this fit.
This illustrates the dominant SSC contribution to the soft $\gamma$ -ray regime in this fit.
In the states fitted with lower values of y;, the SSC component still dominates in the X-ray band, but plays virtually no role in the production of the y-ray emission.
In the states fitted with lower values of $\gamma_1$, the SSC component still dominates in the X-ray band, but plays virtually no role in the production of the $\gamma$ -ray emission.
In the model calculations for 3C 279, all our cases are in the fast-cooling regime.
In the model calculations for 3C 279, all our cases are in the fast-cooling regime.
In the fast-cooling regime, particles will cool below γι to a cutoff energy determined by a balance between radiative cooling and escape.
In the fast-cooling regime, particles will cool below $\gamma_1$ to a cutoff energy determined by a balance between radiative cooling and escape.
In an internal shock scenario, the value of y, may be related to the relative Lorentz factor of two colliding shells, and the efficiency of transferring swept-up proton energy to accelerated electrons.
In an internal shock scenario, the value of $\gamma_1$ may be related to the relative Lorentz factor of two colliding shells, and the efficiency of transferring swept-up proton energy to accelerated electrons.
Because we do not fully understand thedetails of this transfer process,
Because we do not fully understand thedetails of this transfer process,
13/06/2007, sep=4.008+0.008 aarcsec, PA=166.3+0.2°, AK,=10.7+0.1 mmag).
13/06/2007, $sep = 4.008\pm0.008$ arcsec, $PA=166.3\pm0.2^{\circ}$, $\Delta K_{s}=10.7\pm0.1$ mag).
In the case that this object is a non-moving background object, we expect to find it at sep=3.824+0.009 aarcsec, and PA=165.80+0.21° in September 2009.
In the case that this object is a non-moving background object, we expect to find it at $sep=3.824\pm0.009$ arcsec, and $PA=165.80\pm0.21^{\circ}$ in September 2009.
Our NACO astrometry of the candidate agrees well with the predicted values, hence this source is a slowly moving object, clearly not related with the TTel system.
Our NACO astrometry of the candidate agrees well with the predicted values, hence this source is a slowly moving object, clearly not related with the Tel system.
Hence, on basis of the achieved detection limit, we can conclude that there is no further companion of AA within a separation of aarcsec (~ AAU) around the star.
Hence, on basis of the achieved detection limit, we can conclude that there is no further companion of A within a separation of arcsec $\sim$ AU) around the star.
However, we cannot exclude additional very close companions.
However, we cannot exclude additional very close companions.
? list multi-epoch radial-velocity data of PZTTel with a scatter of 4.4kkm/s, which could possibly be induced by such a close companion or by the spots on the stellar surface (?)..
\citet{malaroda2000} list multi-epoch radial-velocity data of Tel with a scatter of km/s, which could possibly be induced by such a close companion or by the spots on the stellar surface \citep{barnes2000}.
Within all observing epochs, from June 2007 to May 2010, the separation and position angle of the BB change both linearly over time with slopes of sép=34.7+ 1.0mmas/yr and PA=-041+ 0.08°/yr.
Within all observing epochs, from June 2007 to May 2010, the separation and position angle of the B change both linearly over time with slopes of $\dot{sep} = 34.7\pm1.0$ mas/yr and $\dot{PA} = -0.41\pm0.08$ $^{\circ}$ /yr.
As PZTTelBB is located at an angular separation of ~ aarcsec in average, which corresponds to a projected separation of about AAU at the distance of PZTTel, the expected escape velocity is mmas/yr, i.e. 1.3 times larger than the measured motion of the companion relative to its primary.
As B is located at an angular separation of $\sim\,$ arcsec in average, which corresponds to a projected separation of about AU at the distance of Tel, the expected escape velocity is $\sim$ mas/yr, i.e. 1.3 times larger than the measured motion of the companion relative to its primary.
Hence, the high relative motion of BB (detected on a significance level > 370) is consistent with orbital motion of the companion around AA. Between June 2007 and September 2009 the separation and the position angle of BB change with slopes of PA=-0.5+ 0.3°/yr and mmas/yr (see also section 2), while we measure slopes of sép=30.9+2.3 mmas/yr and PA=—0.2+0.4°/yr between September 2009 and May 2010, i.e. smaller absolute values of both velocities in the second time interval.
Hence, the high relative motion of B (detected on a significance level $> 37\,\sigma$ ) is consistent with orbital motion of the companion around A. Between June 2007 and September 2009 the separation and the position angle of B change with slopes of $\dot{PA} = -0.5\pm0.3$ $^{\circ}$ /yr and mas/yr (see also section 2), while we measure slopes of $\dot{sep} = 30.9\pm2.3$ mas/yr and $\dot{PA} = -0.2\pm0.4$ $^{\circ}$ /yr between September 2009 and May 2010, i.e. smaller absolute values of both velocities in the second time interval.
The relative motion of the companion slightly slows down (detected on a ~2c significance level), while its separation to its primary increases, consistent with orbital motion.
The relative motion of the companion slightly slows down (detected on a $\sim2\,\sigma$ significance level), while its separation to its primary increases, consistent with orbital motion.
Due to the slow variation in PA compared with that of sep the system is likely to be seen nearly edge-on.
Due to the slow variation in $PA$ compared with that of $sep$ the system is likely to be seen nearly edge-on.
With the obtained NACO detection limit we conclude that BB can be imaged with NACO (even without PSF subtraction) down to a separation of aarcsec around its much brighter primary.
With the obtained NACO detection limit we conclude that B can be imaged with NACO (even without PSF subtraction) down to a separation of arcsec around its much brighter primary.
By tracing back the trajectory of the companion (derived from its relative motion) we expect that PZTTelBB should have been detectable with NACO only since September 2004.
By tracing back the trajectory of the companion (derived from its relative motion) we expect that B should have been detectable with NACO only since September 2004.
Indeed, AA was observed with NACO in July 2003, as reported by ? who also took deep K;- images of the star, but could not detect its companion.
Indeed, A was observed with NACO in July 2003, as reported by \cite{masciadri2005} who also took deep $\rm K_{s}$ -band images of the star, but could not detect its companion.