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According to the evolutionary models for low-mass objects from ?,, ? (2003) the absolute J, H, and K,-band magnitudes of BB are consistent with a brown dwarf with 28*!? Mj,,,, at an age of 1215 MMyr, for the 8 PPic moving group.
According to the evolutionary models for low-mass objects from \cite{chabrier2000}, \cite{baraffe2002} (2003) \nocite{baraffe2003} the absolute J, H, and $\rm K_{s}$ -band magnitudes of B are consistent with a brown dwarf with $28^{+12}_{~-4}$ $M_{Jup}$, at an age of $12^{+8}_{-4}$ Myr, for the $\beta$ Pic moving group.
Even in the case that the TTel system would be older than MMyr the mass of BB would still be within the substellar mass-regime, e.g. Mj, at MMyr.
Even in the case that the Tel system would be older than Myr the mass of B would still be within the substellar mass-regime, e.g. $M_{Jup}$ at Myr.
The model predicted color J—K; for such a brown dwarf companion is J—K,=0.78+0.03 mmag (CIT-2MASScolortransformationapplied, ?),, which agrees well with the color of BB, derived from our NACO photometry J—K,=0.50-:0.12 mmag.
The model predicted color $J-K_{s}$ for such a brown dwarf companion is $J-K_{s}=0.78\pm0.03$ mag \citep[CIT$-$2MASS color transformation applied,][]{carpenter2001}, which agrees well with the color of B, derived from our NACO photometry $J-K_{s}=0.80\pm0.12$ mag.
According to the used models the effective temperature of BB should range between 2500 and KK, which corresponds to a spectral type M6-8 (?)..
According to the used models the effective temperature of B should range between 2500 and K, which corresponds to a spectral type M6-8 \citep{golimowski2004}. .
With only about 15AAU of projected separation PZTTelBB is the second closest substellar companion imaged directly around a young star after 6 bb (?),, i.e. much closer than e.g. the substellar companions of LLup (?),, CCha (?),, or 77329 (??)..
With only about AU of projected separation B is the second closest substellar companion imaged directly around a young star after $\beta$ b \citep{lagrange2010}, i.e. much closer than e.g. the substellar companions of Lup \citep{neuhaeuser2005}, Cha \citep{schmidt2008}, or 7329 \citep{lowrance2000, guenther2001}.
With 6 PPic (?) and 77329 (?), PZTTel is the youngest star with both a substellar companion and a debris disk, which are all members of the 8 PPic moving group, which shows a high rate of stars with both debris disks and substellar companions.
With $\beta$ Pic \citep{smith1984} and 7329 \citep{smith2009}, Tel is the youngest star with both a substellar companion and a debris disk, which are all members of the $\beta$ Pic moving group, which shows a high rate of stars with both debris disks and substellar companions.
BB is the first substellar companion imaged directly,where orbital motion is detected (as deceleration).
B is the first substellar companion imaged directly,where orbital motion is detected (as deceleration).
fact was successfully performed in IKowal et al. (
fact was successfully performed in Kowal et al. (
2009). where the plana effects were simulated through the introduction of the anomalous resistivity.
2009), where the plasma effects were simulated through the introduction of the anomalous resistivity.
Uulike LOO09. where the acceleration of energetic particles in the reversus field of the heliosheathn was considered. we cousider the magnetic field reversals in the heliotail.
Unlike LO09, where the acceleration of energetic particles in the reversing field of the heliosheath was considered, we consider the magnetic field reversals in the heliotail.
While the field reversals iu the heliosheath arise from) Suu's rotation with the maguetic axis beiug nof parallel to the rotation axis. the reversals iu the heliotail arise from the 11 vear solar eveles;
While the field reversals in the heliosheath arise from Sun's rotation with the magnetic axis being not parallel to the rotation axis, the reversals in the heliotail arise from the 11 year solar cycles.
As a result. the scale of the reversals is expected to be much larger.
As a result, the scale of the reversals is expected to be much larger.
This provides a possibility of accelerating higher cuerey particles.
This provides a possibility of accelerating higher energy particles.
Tn this paper we combine data from different experiments to prove that there exists a statistically sienificant excess of cherectic particles in the direction of the solar system magnetotail.
In this paper we combine data from different experiments to prove that there exists a statistically significant excess of energetic particles in the direction of the solar system magnetotail.
We proposed an explanation of this excess as arising from the acceleration of cheregctic particles iu reconnection regions aloug the maegnetotail.
We proposed an explanation of this excess as arising from the acceleration of energetic particles in reconnection regions along the magnetotail.
These regions arise as oppositely directed magnetic field of the solar wind is pressed together in the magnetotail.
These regions arise as oppositely directed magnetic field of the solar wind is pressed together in the magnetotail.
The chanec of the magnetic field polarity arises in our scenario due to the well established solar cveles.
The change of the magnetic field polarity arises in our scenario due to the well established solar cycles.
We thank our colleagues in Ice Cube collaboration. in particular. Francis EHalzeu. for numerous fruitful discussions.
We thank our colleagues in Ice Cube collaboration, in particular, Francis Halzen, for numerous fruitful discussions.
AL acknowledges the support of the NSF eraut. AST Osdstls. NASA erant. N5L66201101 and of the NSF-spousored Ceuter for AMaenetic SelfOreanization.
AL acknowledges the support of the NSF grant AST 0808118, NASA grant X5166204101 and of the NSF-sponsored Center for Magnetic Self-Organization.
PD ackuowledges the support from the U.S. National Science Foundation-Office of Polar Programs.
PD acknowledges the support from the U.S. National Science Foundation-Office of Polar Programs.
While the LV99 model provides fast reconnection. 1.0. the reconnection that does not depeud ou resistivity. without appealing to auv collisionless plasiina effects. for some of the Μπακοα]ο eveuts. c.e. for the acceleration of low energy electrons the mucroplivsics of reconnection may be important.
While the LV99 model provides fast reconnection, i.e. the reconnection that does not depend on resistivity, without appealing to any collisionless plasma effects, for some of the small-scale events, e.g. for the acceleration of low energy electrons the microphysics of reconnection may be important.
Iu the recounectiou process described by LV99 model the reconnection speed is given by Eq. (5)).
In the reconnection process described by LV99 model the reconnection speed is given by Eq. \ref{eq:recon1}) ).
The same speed cau be obtained if oue considers the local reconnection of flux tubes.
The same speed can be obtained if one considers the local reconnection of flux tubes.
It was shown in LV99 that the probability of magnetic field lines which oues eutered the reconnection laver to reenter the reconnection laver again is low.
It was shown in LV99 that the probability of magnetic field lines which ones entered the reconnection layer to reenter the reconnection layer again is low.
As a result it was shown in LV99 that the elobal reconnection rate is where ρω Is the velocity of reconnection within the small-scale local Sweet-Parker recounection events. 6. reconmection events on scale Ay depicted on the lower pauel of Figure 5.
As a result it was shown in LV99 that the global reconnection rate is where $V_{rec, local}$ is the velocity of reconnection within the small-scale local Sweet-Parker reconnection events, i.e. reconnection events on scale $\lambda_{\|}$ depicted on the lower panel of Figure \ref{fig:recon1}.
It ix easy to show that assmuing that if and Ay are related. through the (5595 critical balance. namely λενAre. Eq. CAT))
It is easy to show that assuming that if $\lambda_{\|}$ and $\lambda_{\bot}$ are related through the GS95 critical balance, namely $\lambda_{\|}/V_A\approx \lambda_{\bot}/v$, Eq. \ref{global}) )
results iu too high reconnectionAj rates even for Swect-Parker recounectiou at scales A4.
results in too high reconnection rates even for Sweet-Parker reconnection at scales $\lambda_{\bot}$ .
The naively obtained rates much exceed those provided by Eq. (5)).
The naively obtained rates much exceed those provided by Eq. \ref{eq:recon1}) ),
which proves that the Olunic resistivity effects are not the bottleneck of the LV99 model (see more details in LY99).
which proves that the Ohmic resistivity effects are not the bottleneck of the LV99 model (see more details in LV99).
We will use Eq. CÀ1)) (5))
We will use Eq. \ref{global}) ) \ref{eq:recon1}) )
to establish the scale of the local Sweet-Parker events A[ which results in the thickness of the Sweet-Parker lavers being Those Ανν correspond to Aper) in the lower paucl of Figure 5..
to establish the scale of the local Sweet-Parker events $\lambda_{\|}$: which results in the thickness of the Sweet-Parker layers being Those $\Delta_{turb}$ correspond to $\lambda_{perp}$ in the lower panel of Figure \ref{fig:recon1}.
The corresponding thickness is much smaller that the thickness of the lznninar Sweet-Parker laver. which is L51/2 which makes according to Eq. (3))
The corresponding thickness is much smaller that the thickness of the laminar Sweet-Parker layer, which is $L S^{-1/2}$, which makes according to Eq. \ref{constraint}) )
the collisionless effects important for the small scale heliotail recounection.
the collisionless effects important for the small scale heliotail reconnection.
This. however. is not going to change either the rate of maguctic reconnection or the acceleration of protons.
This, however, is not going to change either the rate of magnetic reconnection or the acceleration of protons.
The presence of collisionless effects can affect the acceleration of electrons. which is. however. not the process that we deal in this paper.
The presence of collisionless effects can affect the acceleration of electrons, which is, however, not the process that we deal in this paper.
We should add that the issue of fiuid being collisionless or collisional is not so simple in the presence of turbulence.
We should add that the issue of fluid being collisionless or collisional is not so simple in the presence of turbulence.
Collisionless fluids are subject to instabilities. which reduce the effective degree of their collisionality via inducing resonant scattering of particles.
Collisionless fluids are subject to instabilities, which reduce the effective degree of their collisionality via inducing resonant scattering of particles.
Such collisious mediated by magnetic field decrease the mean free path of particpa
Such collisions mediated by magnetic field decrease the mean free path of particles.
For instance. iu Lazarian Beresuval (2006) the problem of the collisious in a fluid of was treated self-cousisteutlv iu the presence of the gxroresonauce instability.
For instance, in Lazarian Beresnyak (2006) the problem of the collisions in a fluid of was treated self-consistently in the presence of the gyroresonance instability.
It was demonstrated there that turbulent compressious of the fluid on the mean free path induce instability at the evroradius which decrease the mean free path.
It was demonstrated there that turbulent compressions of the fluid on the mean free path induce instability at the gyroradius which decrease the mean free path.
For thermal plasma particles firchose and iuirrow instability should also be important (see Schelochihin et al.
For thermal plasma particles firehose and mirrow instability should also be important (see Schekochihin et al.
2010).
2010).
As a result. we can state that turbulence both decreases the scale over which clemeutary reconnection events take place. poteutially allowing the elementary recounectiou events to proceed im a collisionless fashion. but at the same the compressions arising frou the turbulent can decrease the mean free path of the particles.
As a result, we can state that turbulence both decreases the scale over which elementary reconnection events take place, potentially allowing the elementary reconnection events to proceed in a collisionless fashion, but at the same the compressions arising from the turbulent can decrease the mean free path of the particles.
The details of these interesting processes are nof iuportanut for the acceleration of cosmic ray protonsthat we deal in this paper.
The details of these interesting processes are not important for the acceleration of cosmic ray protonsthat we deal in this paper.
a consistent way with element diffusion as considered in this study is required for precise asteroseismology.
a consistent way with element diffusion as considered in this study is required for precise asteroseismology.
In this section we show that, in the frame of standard evolutionary computations for the formation of DA white dwarfs, the He content of these stars cannot be substantially smaller than that predicted by our calculations.
In this section we show that, in the frame of standard evolutionary computations for the formation of DA white dwarfs, the He content of these stars cannot be substantially smaller than that predicted by our calculations.
To do this, we compute the evolution of a 2Mo-star from the ZAMS until the thermally-pulsing phase on the AGB.
To do this, we compute the evolution of a $2 M_{\odot}$ -star from the ZAMS until the thermally-pulsing phase on the AGB.
The only way we envisage in which the star may experience a substantial decrease in its content of He is by undergoing a large number of thermal pulses.
The only way we envisage in which the star may experience a substantial decrease in its content of He is by undergoing a large number of thermal pulses.
In order for the model star to experience the largest possible number of thermal pulses, and thus, the content of He decreases as much as possible, we switched off mass loss during this stage in our evolutionary code.
In order for the model star to experience the largest possible number of thermal pulses, and thus, the content of He decreases as much as possible, we switched off mass loss during this stage in our evolutionary code.
The results of this experiment are depicted in Fig. 2,,
The results of this experiment are depicted in Fig. \ref{max-he},
in which we show the He content in the region limited by the boundaries of the He-free core (HeFC) and the HFC in terms of time during the thermally pulsing phase (upper panel), and the surface luminosity and the H- and He-burning luminosities for each pulse in that phase (lower panel).
in which we show the He content in the region limited by the boundaries of the He-free core (HeFC) and the HFC in terms of time during the thermally pulsing phase (upper panel), and the surface luminosity and the H- and He-burning luminosities for each pulse in that phase (lower panel).
We stopped the experiment when the object experienced about 30 thermal pulses, which is enough for our purposes.
We stopped the experiment when the object experienced about 30 thermal pulses, which is enough for our purposes.
We found that the He content of the object decreased from Mnc/Mc=3.34x10? (before the first thermal pulse) to Mue/Mo=8.6x10? (before the thirtieth thermal pulse).
We found that the He content of the object decreased from $M_{\rm He}/M_{\odot}= 3.34 \times 10^{-2}$ (before the first thermal pulse) to $M_{\rm He}/M_{\odot}= 8.6 \times 10^{-3}$ (before the thirtieth thermal pulse).
Thus, the decrease (in solar masses) of the He content of the HFC is of a factor 3.89.
Thus, the decrease (in solar masses) of the He content of the HFC is of a factor $3.89$.
However, it should be kept in mind that this reduction is due mainly to the increase of the mass of the future white dwarf, that grows from Mrwp=0.523Mo to Mpwp=0.7114Mg between the thermal pulses 1 to 30.
However, it should be kept in mind that this reduction is due mainly to the increase of the mass of the future white dwarf, that grows from $M_{\rm fWD}= 0.523 M_{\odot}$ to $M_{\rm fWD}= 0.7114 M_{\odot}$ between the thermal pulses 1 to 30.
Our experiment shows that the He mass left in a DA white dwarf could be as much as afactor 3—4 lower than the values predicted by standard evolutionary computations, but not 2 or 3magnitude lower, which would be necessary for g-mode periods to be substantially affected.
Our experiment shows that the He mass left in a DA white dwarf could be as much as a $3-4$ lower than the values predicted by standard evolutionary computations, but not 2 or 3 lower, which would be necessary for $g$ -mode periods to be substantially affected.
We conclude that we can safety ignore the variation of Mue in our asteroseismological analysis of ZZ Ceti stars.
We conclude that we can safety ignore the variation of $M_{\rm He}$ in our asteroseismological analysis of ZZ Ceti stars.
The DA white dwarf models employed in this study are the result of full evolutionary calculations of progenitor stars for solar-like metallicity (Z— 0.01).
The DA white dwarf models employed in this study are the result of full evolutionary calculations of progenitor stars for solar-like metallicity $Z= 0.01$ ).
The complete evolution of eleven evolutionary sequences with initial stellar mass in the range 1—5Mo has been computed from the ZAMS through the thermally-pulsing and mass-loss phases on the AGB and finally to the domain of planetary nebulae.
The complete evolution of eleven evolutionary sequences with initial stellar mass in the range $1-5 M_{\odot}$ has been computed from the ZAMS through the thermally-pulsing and mass-loss phases on the AGB and finally to the domain of planetary nebulae.
The values of the stellar mass of our set of models is shown in the upper row of Table 1..
The values of the stellar mass of our set of models is shown in the upper row of Table \ref{table1}.
The range of stellar mass covered by our computations comfortably accounts for the stellar mass of most of the observed pulsating DA white dwarfs.
The range of stellar mass covered by our computations comfortably accounts for the stellar mass of most of the observed pulsating DA white dwarfs.
Our asteroseismological approach basically consists in the employment of detailed white dwarf models characterized by very accurate physical ingredients.
Our asteroseismological approach basically consists in the employment of detailed white dwarf models characterized by very accurate physical ingredients.
These models are obtained by computing the complete evolution of the progenitor stars.
These models are obtained by computing the complete evolution of the progenitor stars.
We have applied successfully this approach to the hot DOVs or GW Vir stars (see Córrsico et al.
We have applied successfully this approach to the hot DOVs or GW Vir stars (see Córrsico et al.
2007a, 2007b, 2008, 2009).
2007a, 2007b, 2008, 2009).
Since the final chemical
Since the final chemical
several authors have examined the sharp eutolls observed in the II] disks of Sgalaxies in (he context of using these signatures to infer the local ionizing background. (Maloney 1993.093. Corbelli Salpeter 1993. DDove ShShull 1994).
Several authors have examined the sharp cutoffs observed in the HI disks of galaxies in the context of using these signatures to infer the local ionizing background (Maloney 1993, Corbelli Salpeter 1993, Dove Shull 1994).
994). The truncationsicalions are modeled as arisingrising primarily from photoionization of the disk gas by (the local extragalactic background radiation field.
The truncations are modeled as arising primarily from photoionization of the disk gas by the local extragalactic background radiation field.
Using5 21 em observations (Corbelli. Scheider. Salpeter 1989. van Gorkom 1993) to constrain these models. limits on the local ionizing background are placed at ? !. where and where J,= lor an isotropic radiation field.
Using 21 cm observations (Corbelli, Scheider, Salpeter 1989, van Gorkom 1993) to constrain these models, limits on the local ionizing background are placed at $10^{4} < \Phi_{{\rm ion}} < 5 \times 10^{4}$ $^{-2}$ $^{-1}$ , where and where $J_{\nu}=I_{\nu}$ for an isotropic radiation field.
Acdcditionallv. narrow-band ancl Fabry-Perot observations of Ho emission from intergalactic clouds (Stocke et 11991. Blancd-Hawthorn et 11994. Vogel et 11995. Donahue. Aldering. Stocke 1995) place limits of Pj,€10! ? 1. or η)<7.6x1075 eres stem ? ! sr? for ag=L8. while results from measurements of Galactic high velocity clouds (Νπίντον Revnolds 1989. Songaila. Bryant. Cowie 1939. Tulte. Revnolds. llaffner 1998) imply d,οσοςlObem 7s +. though the ionization of high velocity clouds may be contaminated by a Galactic stellar contribution.
Additionally, narrow-band and Fabry-Perot observations of $\alpha$ emission from intergalactic clouds (Stocke et 1991, Bland-Hawthorn et 1994, Vogel et 1995, Donahue, Aldering, Stocke 1995) place limits of $\Phi_{{\rm ion}} \lesssim 10^{4}$ $^{-2}$ $^{-1}$, or $J(\nu_{0}) < 7.6 \times 10^{-23}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$ for $\alpha_{s}=1.8$, while results from measurements of Galactic high velocity clouds (Kutyrev Reynolds 1989, Songaila, Bryant, Cowie 1989, Tufte, Reynolds, Haffner 1998) imply $\Phi_{{\rm ion}} \lesssim 6 \times 10^{4}$ $^{-2}$ $^{-1}$, though the ionization of high velocity clouds may be contaminated by a Galactic stellar contribution.
Tundinson et ((1999) have reanalyzed the 3C273/NGC3067 Ποια using the Πα imaging data from Stocke οἱ ((1991) as well as new GIIRS spectra of 3C273. in order {ο model the ionization balance in the absorbing gas in the halo of NGC3067.
Tumlinson et (1999) have reanalyzed the 3C273/NGC3067 field using the $\alpha$ imaging data from Stocke et (1991) as well as new GHRS spectra of 3C273, in order to model the ionization balance in the absorbing gas in the halo of NGC3067.
From this analysis. they derive the limits. 2600—4j,—«10! ? f. or LO7<Jim)3.8x1077 eres ? tsr tat 2=0.0047.
From this analysis, they derive the limits, $2600 < \Phi_{{\rm ion}} < 10^{4}$ $^{-2}$ $^{-1}$, or $10^{-23} < J(\nu_{0}) < 3.8 \times 10^{-23}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$ at $z = 0.0047$.
Wevimnann et (2001) have recently reported an upper limit of j,,<LOIxI0 7s to or JQ)<3.84x107 eress ? ! ! from Fabry-Perot observations of the intergalactie HI cloud. 12254-01. for a [ace-on disk geometry.
Weymann et (2001) have recently reported an upper limit of $\Phi_{{\rm ion}} < 1.01 \times 10^{4}$ $^{-2}$ $^{-1}$, or $J(\nu_{0}) < 3.84 \times 10^{-23}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$ from Fabry-Perot observations of the intergalactic HI cloud, 1225+01, for a face-on disk geometry.
If an inclined disk geometry is assumed. this lower limit becomes (νι)<9.6x107! eres stem ? F !.
If an inclined disk geometry is assumed, this lower limit becomes $J(\nu_{0}) < 9.6 \times 10^{-24}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$.
These results are sumniarized in Figure 16..
These results are summarized in Figure \ref{fig:allzcomp}.
It is encouraging that the proximitv effect. value is consistent with (he limits on the backgroundset by these more direct estimates which are possible locally.
It is encouraging that the proximity effect value is consistent with the limits on the backgroundset by these more direct estimates which are possible locally.
Tlaardt Macau (1996) calculated the spectrum of the UV background as a function of frequency and redshBift using a model based on the integrated emission [rom QSOs alone.
Haardt Madau (1996) calculated the spectrum of the UV background as a function of frequency and redshift using a model based on the integrated emission from QSOs alone.
The
The
be evidence for a possible connection between the helical movement of Cl and the total flux-density periodicity.
be evidence for a possible connection between the helical movement of $C1$ and the total flux-density periodicity.
The jet might precess and therefore due to variability of the Doppler boosting we will see periodic flares in the light curves as well as changes of position of a bend in the jet.
The jet might precess and therefore due to variability of the Doppler boosting we will see periodic flares in the light curves as well as changes of position of a bend in the jet.
We plot the position angles of all features da—d7. including the quasi-stationary component Cl versus the core separation in Fig. 15..
We plot the position angles of all features $da$ $d7$, including the quasi-stationary component C1 versus the core separation in Fig. \ref{0605_pa}.
The trajectories of the components d1. d2. do. d7. and Cl follow significantly curved trajectories. whereas the trajectories of the components 40 and d3 follow a linear path.
The trajectories of the components $d1$, $d2$, $d6$, $d7$, and $C1$ follow significantly curved trajectories, whereas the trajectories of the components $d0$ and $d3$ follow a linear path.
This can be an evidence for a possible co-existence of two types of jet component trajectories. similar to the source 3C 345 (Klare et al.
This can be an evidence for a possible co-existence of two types of jet component trajectories, similar to the source 3C 345 (Klare et al.
2005).
2005).
The spread of position angles of all components is changing along the jet.
The spread of position angles of all components is changing along the jet.
In the inner part. within the first | mas. the position angles are in the range of 123-150 degrees. in the middle part of the jet between | — 5 mas. the position angles range from 110° to 1257. whereas in the outer part from 5 mas to 10 mas. from 95° to 1157.
In the inner part, within the first 1 mas, the position angles are in the range of 123–150 degrees, in the middle part of the jet between 1 – 5 mas, the position angles range from $110^\circ$ to $125^\circ$ , whereas in the outer part from 5 mas to 10 mas, from $95^\circ$ to $115^\circ$.
The probable periodical variability of the radio total densities of B0605—085 and the helical motion of the quasi-stationary Jet component C] can be explained by Jet precession.
The probable periodical variability of the radio total flux-densities of $-$ 085 and the helical motion of the quasi-stationary jet component $C1$ can be explained by jet precession.
We used a precession model. described in Abraham Carrara (1998) and Caproni Abraham (2004) for fitting the helical path of the jet component C1.
We used a precession model, described in Abraham Carrara (1998) and Caproni Abraham (2004) for fitting the helical path of the jet component $C1$.
In this model. the rectangular coordinates of a precession cone in the source frame are changing with time ¢’ due to precession: where Q is the semi-aperture angle of the precession cone. dy is the angle between the precession cone axis and the line of sight and it is actually the average viewing angle of the source. and aq is the projected angle of the cone axis onto the plane of the sky (see Fig. 16)).
In this model, the rectangular coordinates of a precession cone in the source frame are changing with time $t'$ due to precession: where $\Omega$ is the semi-aperture angle of the precession cone, $\phi_{0}$ is the angle between the precession cone axis and the line of sight and it is actually the average viewing angle of the source, and $\eta_{0}$ is the projected angle of the cone axis onto the plane of the sky (see Fig. \ref{0605_model_plot}) ).
The angular velocity of the precession is w=1/P. where P is the period.
The angular velocity of the precession is $\omega = 1 / P$, where $P$ is the period.
We take a P value of 7.9-year (timescale found in the total flux-density radio light curves and in the helical movement of CI).
We take a $P$ value of 7.9-year (timescale found in the total flux-density radio light curves and in the helical movement of $C1$ ).
The time in the source frame () and the frame of the observer (7) are related by the Doppler factor ὁ as However. the only time-dependent. term in. the equations (6)) and (7)) is wr. which does not depend on this time corrections since w-|/r.
The time in the source frame $t'$ ) and the frame of the observer $t$ ) are related by the Doppler factor $\delta$ as However, the only time-dependent term in the equations \ref{eq:x_t}) ) and \ref{eq:y_t}) ) is $\omega t$, which does not depend on this time corrections since $\omega \sim 1/t$.
Therefore. we can fit the trajectory in the rectangular coordinate X(1) of the stationary jet component C1. using the formula 6..
Therefore, we can fit the trajectory in the rectangular coordinate $X(t)$ of the quasi-stationary jet component $C1$, using the formula \ref{eq:x_t}.
We assume that the motion of C1 reflects the movement of the jet.
We assume that the motion of $C1$ reflects the movement of the jet.
The non-linear least squares method was used for the precession model fitting.
The non-linear least squares method was used for the precession model fitting.
The angular velocity of the precession ω and the redshift of the source z are known and we fix them during the fit. whereas the aperture angle of the precessing cone Q. dy. and jo are used as free parameters.
The angular velocity of the precession $\omega$ and the redshift of the source $z$ are known and we fix them during the fit, whereas the aperture angle of the precessing cone $\Omega$, $\phi_{0}$, and $\eta_{0}$ are used as free parameters.
The results of the fit are shown in Fig.
The results of the fit are shown in Fig.
17. (Left) and the parameters of the precession model for the parsee jet of B0605—085 are listed in Table 6..
\ref{0605_precession_model} (Left) and the parameters of the precession model for the parsec jet of $-$ 085 are listed in Table \ref{0605_precession}.
The precession model fits well the trajectory of the quasi-stationary jet component C1.
The precession model fits well the trajectory of the quasi-stationary jet component $C1$.
The same parameters were used to describe the motion of C1 in declination.
The same parameters were used to describe the motion of $C1$ in declination.