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izmiran.troitsk.ru/flars/LARSbtml) ancl others.
izmiran.troitsk.ru/lars/LARS.html) and others.
However. they have the same disadvantage that we noticed. above in respect to à too small effective area. of such antennas.
However, they have the same disadvantage that we noticed above in respect to a too small effective area of such antennas.
No wonder that the instruments do not register weal solar events in the decameter wavelength. range. but this does not mean their lack.
No wonder that the instruments do not register weak solar events in the decameter wavelength range, but this does not mean their lack.
Sometimes. the impossibility. of their detection because. of a small antenna aperture may. be fatal for the analysis of solar events.
Sometimes, the impossibility of their detection because of a small antenna aperture may be fatal for the analysis of solar events.
Πας. such kind. of solar events is desirable to study by more comprehensive antennas with ellective back-ends.
Thus, such kind of solar events is desirable to study by more comprehensive antennas with effective back-ends.
In. closing it. should. be added: that. for. the firm belief. we applied the aforesaid analvsis for the data. represented in the internet site from the measurements. obtained. by the instrument GBSRLBS.
In closing it should be added that for the firm belief we applied the aforesaid analysis for the data, represented in the internet site from the measurements, obtained by the instrument GBSRBS.
Though the instrument GBSRBS is located in the North America. and UTI-2 is in East Europe. some solar events were observed simultaneously by both raciotelecopes.
Though the instrument GBSRBS is located in the North America, and UTR-2 is in East Europe, some solar events were observed simultaneously by both radiotelecopes.
1t
It
In this paper we discussed the high energy 5-rav enission of the FR I large scale jets.
In this paper we discussed the high energy $\gamma$ -ray emission of the FR I large scale jets.
We used the X-ray observations compiled with data from the other spectral bands in order to reconstruct energy. distribution of ultrarelativistic electrons present in the considered objects.
We used the X-ray observations compiled with data from the other spectral bands in order to reconstruct energy distribution of ultrarelativistic electrons present in the considered objects.
Our phenomenological approach to recover electron spectrum [rom observations rather than from theories of particle acceleration is dictated by the fact. Chat such theories are still insulliciently developed to enable quantitative analvsis.
Our phenomenological approach to recover electron spectrum from observations rather than from theories of particle acceleration is dictated by the fact, that such theories are still insufficiently developed to enable quantitative analysis.
Next. we analyzed possible origin ol the seed photons contributing to the inverse-Compton emission of the obtained electron spectrum. including nuclear jet radiation as well as ambient. stellar ancl cireumstellar enussion of the host galaxies. ancl cliscussecl in detail spectral properties of (he resulüing 5- radiative output.
Next, we analyzed possible origin of the seed photons contributing to the inverse-Compton emission of the obtained electron spectrum, including nuclear jet radiation as well as ambient, stellar and circumstellar emission of the host galaxies, and discussed in detail spectral properties of the resulting $\gamma$ -ray radiative output.
The approach take into account effects connected with relativistic bulk motion of the emitting region. correcting earlier computations presented in the literature.
The approach take into account effects connected with relativistic bulk motion of the emitting region, correcting earlier computations presented in the literature.
We also clearly indicated limitation of the adopted approach due the Ixlein-Nishina regime effects.
We also clearly indicated limitation of the adopted approach due the Klein-Nishina regime effects.
Other restrictions of the presented model are connected with only roughly evaluated ealactie radiation fields and hardly known parameters of the hidden blazar radiation.
Other restrictions of the presented model are connected with only roughly evaluated galactic radiation fields and hardly known parameters of the hidden blazar radiation.
Also. the effects of 5-rav absorption on CIB radiation are ouly shortly mentioned in this paper.
Also, the effects of $\gamma$ -ray absorption on CIB radiation are only shortly mentioned in this paper.
Ilowever. even with all the aforementioned uncertainties. one can conclude that the detection of the hieh energv 5-rav enission bv [future ground-based and space telescopes from. al least. some of the FR I large scale jets is possible.
However, even with all the aforementioned uncertainties, one can conclude that the detection of the high energy $\gamma$ -ray emission by future ground-based and space telescopes from, at least, some of the FR I large scale jets is possible.
Thus. the future observations will provide important constraints on the unknown jet parameters. like (he magnetic field intensity and the jet Doppler factor. confirming or excluding possibilities of D«D, and 9z1 discussed in the literature.
Thus, the future observations will provide important constraints on the unknown jet parameters, like the magnetic field intensity and the jet Doppler factor, confirming or excluding possibilities of $B \ll B_{eq}$ and $\delta \neq 1$ discussed in the literature.
Our analvsis provides constraints for y-ray emission of the nearby FR [sources Centaurus A and M 87.
Our analysis provides constraints for $\gamma$ -ray emission of the nearby FR I sources Centaurus A and M 87.
For Centaurus A we predict measurable bv future 5-rav missions = fluxes at 10 GeV and 0.1—1 TeV photon energies due to comptonisation of the blazar radiation aud the svuchrotron sel[-Compton process. respectvelv.
For Centaurus A we predict measurable – by future $\gamma$ -ray missions – fluxes at $10$ GeV and $0.1 - 1$ TeV photon energies due to comptonisation of the blazar radiation and the synchrotron self-Compton process, respectively.
In (the case of M 87 we show that recently detected VILE. emission can result from comptonisation of the stellar ancl cireumstellar inlrared photons of the host galaxy.
In the case of M 87 we show that recently detected VHE emission can result from comptonisation of the stellar and circumstellar infrared photons of the host galaxy.
Possibility that the jets and not the active centers im these two FR I radio galaxies can dominate production of high enerey 5-ravs was not discussed previously in the literature.
Possibility that the jets – and not the active centers – in these two FR I radio galaxies can dominate production of high energy $\gamma$ -rays was not discussed previously in the literature.
Some other authors. e.g. BaianclLee(2001).. Protheroeetal.(2003). or DoneaandProtheroe(2003).. suggested and studied production of very hieh energy. 5-ravs in the nuclear regions of the considered. sources by thescale.. blazar-like jets.
Some other authors, e.g. \citet{bai01}, \citet{pro03} or \citet{don03}, suggested and studied production of very high energy $\gamma$ -rays in the nuclear regions of the considered sources by the, blazar-like jets.
This constitutes the main difference with our work.
This constitutes the main difference with our work.
Unfortunatellv. the angular resolution of theLAC systems will not allow [or separation of the kiloparsec scale jet. 2-rav. radiation from the eventuall core component.
Unfortunatelly, the angular resolution of the systems will not allow for separation of the kiloparsec scale jet $\gamma$ -ray radiation from the eventuall core component.
Llowever. the core vs. the kpc-scale jet origin of 5-ravs can be
However, the core vs. the kpc-scale jet origin of $\gamma$ -rays can be
aliphatic chain CIT; (Fig 1e.d). even if uo livdroxyl is attached: 8.21 to 8.59 jan. These modes are characterized by the motion of a carbon atom of the backbone. from perpendicular to the backbone at the shorterer waveleugtlis. to parallel at the longer oues.
aliphatic chain $_{3}$ (Fig 1c,d), even if no hydroxyl is attached; 8.24 to 8.59 $\mu$ m. These modes are characterized by the motion of a carbon atom of the backbone, from perpendicular to the backbone at the shorterer wavelengths, to parallel at the longer ones.
The iuteusities are relatively strong.
The intensities are relatively strong.
8 to 10 yan Iu this baud. coroneue aud derivatives display weak lines associated with plane C-II bending vibrations.
8 to 10 $\mu$ In this band, coronene and derivatives display weak lines associated with in-plane C-H bending vibrations.
10 to 13 jan As expected. this baud mostly displavs stroug. out-of-plaue beudiugs of the C-II bouds of PAIIs: as the wavelength increases. the motion involves iore and more neighbouring bonds.
10 to 13 $\mu$ As expected, this band mostly displays strong, out-of-plane bendings of the C-H bonds of PAHs: as the wavelength increases, the motion involves more and more neighbouring bonds.
Iu isolated benzene aud pyrene. the vibration of all C-II bouds in phase gives rise to a strong ir line near L1 722. which is not observed in the sky.
In isolated benzene and pyrene, the vibration of all C-H bonds in phase gives rise to a strong ir line near 14 $\mu$ m, which is not observed in the sky.
These two species were therefore excluded frou our selection.
These two species were therefore excluded from our selection.
The strongest lines in this band are due to solo C-IT bending iu coronene and its derivatives. and Le between 11.25 aud 11.3 gan. These structures also display lines of various intensities between 12.6 aud 12.9 jan. These two condensations of lines are reminiscent of the UID peaks at 11.3 aud 12.7 gan. while the lines due to the simaller aromatics. aud distributed over the rauge. remind us of the so-called 11-13-4522. plateau.
The strongest lines in this band are due to solo C-H bending in coronene and its derivatives, and lie between 11.25 and 11.3 $\mu$ m. These structures also display lines of various intensities between 12.6 and 12.9 $\mu$ m. These two condensations of lines are reminiscent of the UIB peaks at 11.3 and 12.7 $\mu$ m, while the lines due to the smaller aromatics, and distributed over the range, remind us of the so-called $\mu$ m plateau.
13 jan CII-chains have characteristic. weak modes between 12 and 13.3 jou. iu which the methylene eroups are rocking about an axis parallel to the chain's backbone.
13 $\mu$ $_{2}$ -chains have characteristic, weak modes between 12 and 13.3 $\mu$ m, in which the methylene groups are rocking about an axis parallel to the chain's backbone.
These must be taken iuto account iu a model where such chains connect the other elemeutaryv. structures.
These must be taken into account in a model where such chains connect the other elementary structures.
15 to 20 jan The larger compact PAIS (coronene and derivatives) have bulk modes ucar 17 xd LS gan. in which the motions of the carbon atous out of plane give rise to orderly ripples through the structure.
15 to 20 $\mu$ The larger compact PAHs (coronene and derivatives) have bulk modes near 17 and 18 $\mu$ m, in which the motions of the carbon atoms out of plane give rise to orderly ripples through the structure.
The strongest baud aceurs in coronene at 17.9 yan. when ouly the II atoms and the sub-peripheral € atoms (those that ave not linked to the atoms) are set in motion. the two groups beige out of phase.
The strongest band accurs in coronene at 17.9 $\mu$ m, when only the H atoms and the sub-peripheral C atoms (those that are not linked to the H atoms) are set in motion, the two groups being out of phase.
Trio structures have simular o.0.p modes near 17 aud 18.1 421: but they also contribute strong lines frou 15.8 to 16.6 jan anc near 19.5 pan. These three bands differ from one another by subtle changes in the relative pliases of the motions of the C atoms.
Trio structures have similar o.o.p modes near 17 and 18.1 $\mu$ m; but they also contribute strong lines from 15.8 to 16.6 $\mu$ m and near 19.5 $\mu$ m. These three bands differ from one another by subtle changes in the relative phases of the motions of the C atoms.
The motions aud ir iuteusities are cousiderablv chhanced when the central peutagon is capped with à N atom.
The motions and ir intensities are considerably enhanced when the central pentagon is capped with a N atom.
Structure | a. for instance bas a notable liue at 15.8 pau. involving the motions. out of phase. of the N atom aud its attached IT.
Structure 4 a, for instance has a notable line at 15.8 $\mu$ m, involving the motions, out of phase, of the N atom and its attached H.
model.
model.
Since we assume a fairly high magnetic field in this simulation and the Υ=5/3 case we use for comparison, the electron energy is limited by synchrotron losses.
Since we assume a fairly high magnetic field in this simulation and the $\gamma = 5/3$ case we use for comparison, the electron energy is limited by synchrotron losses.
Therefore in both models electrons have a similar value for Emax.
Therefore in both models electrons have a similar value for $E_{\rm max}$.
In this section we compare the results from the CSM and ISM models with B=3 µία. Bohm diffusion is assumed in both cases.
In this section we compare the results from the CSM and ISM models with $B=3\ \umu$ G. Bohm diffusion is assumed in both cases.
As we already saw in the analytical calculations in Sect. 3.2,,
As we already saw in the analytical calculations in Sect. \ref{sec:time-scalesanalytical},
the maximum particle energy depends strongly on the shock velocity V;.
the maximum particle energy depends strongly on the shock velocity $V_{\rm s}$.
The density of the medium in which the SNR expands affects the shock velocity and therefore leads to differences in the cosmic-ray acceleration rate between the CSM and the ISM models.
The density of the medium in which the SNR expands affects the shock velocity and therefore leads to differences in the cosmic-ray acceleration rate between the CSM and the ISM models.
A blast wave expanding into the CSM hits a relatively dense medium early on in its evolution.
A blast wave expanding into the CSM hits a relatively dense medium early on in its evolution.
As a result, the initial velocity is smaller but the deceleration proceeds more slowly, as the swept-up mass increases with radius as May R, as opposed to Ma,οςR? in the ISM case.
As a result, the initial velocity is smaller but the deceleration proceeds more slowly, as the swept-up mass increases with radius as $M_{\rm sw} \propto R$ , as opposed to $M_{\rm sw} \propto R^3$ in the ISM case.
Ultimately this results in a shock with a higher velocity at the end of the simulation (1500 yr).
Ultimately this results in a shock with a higher velocity at the end of the simulation $1500$ yr).
In the ISM model the initial shock velocity is higher, but the deceleration much more severe, and the maximum attainable particle energy is lower, at least for the model parameters used here.
In the ISM model the initial shock velocity is higher, but the deceleration much more severe, and the maximum attainable particle energy is lower, at least for the model parameters used here.
In Fig.
In Fig.
11 we show the evolution of the radius of the forward and the reverse shock in the top panel, and in the bottom panel the velocity of the blast wave for the SNR in CSM, ISM.
\ref{fig:rvshocktimemulti} we show the evolution of the radius of the forward and the reverse shock in the top panel, and in the bottom panel the velocity of the blast wave for the SNR in CSM, ISM.
The injection rate of particles is taken to be proportional to the mass that is swept up per unit time by the blast wave.
The injection rate of particles is taken to be proportional to the mass that is swept up per unit time by the blast wave.
As a result the age distribution of cosmic rays in the CSM and ISM models differs.
As a result the age distribution of cosmic rays in the CSM and ISM models differs.
This difference affects both the maximum energy and the shape of the overall spectrum.
This difference affects both the maximum energy and the shape of the overall spectrum.
This is shown in Fig. 12..
This is shown in Fig. \ref{fig:ismcsmFptime}.
The proton/electron spectrum in the CSM model, represented as p?F(p), is slightly concave.
The proton/electron spectrum in the CSM model, represented as $p^2 \: F(p)$, is slightly concave.
This arises because of the higher fraction of ‘old’ particles in the CSM model compared to the ISM model.
This arises because of the higher fraction of `old' particles in the CSM model compared to the ISM model.
On average, these older particles have a higher energy and (for Bohm diffusion) a larger diffusivity.
On average, these older particles have a higher energy and (for Bohm diffusion) a larger diffusivity.
Low-energy particles on the other hand are more rapidly swept away from the shock.
Low-energy particles on the other hand are more rapidly swept away from the shock.
This simulation shows the differences that arise in a time-dependent calculation with respect to the results at a steady (unchanging) shock.
This simulation shows the differences that arise in a time-dependent calculation with respect to the results at a steady (unchanging) shock.
In Fig.
In Fig.
13 we show the maximum energy of the particles as extracted from the simulations.
\ref{fig:pmaxismcsm} we show the maximum energy of the particles as extracted from the simulations.
Since the slope of the overall spectrum in this case is about q=2.15, we define Emax as the e-folding energy, where p?!?F(p) for the cumulative spectrum decreases to a value smaller than 1/e times the value at lower energies.
Since the slope of the overall spectrum in this case is about $q=2.15$, we define $E_{\rm max}$ as the e-folding energy, where $p^{2.15} \: F(p)$ for the cumulative spectrum decreases to a value smaller than $1/e$ times the value at lower energies.
The higher average velocity of the blast wave when it evolves into a CSM increases Emax by a factor 2—4 for the CSM models.
The higher average velocity of the blast wave when it evolves into a CSM increases $E_{\rm max}$ by a factor $2-4$ for the CSM models.
Due to the low magnetic field strength, the synchrotron loss time for this model is significantly longer than the running time of the simulations (~4x10* yr versus 1500 yr).
Due to the low magnetic field strength, the synchrotron loss time for this model is significantly longer than the running time of the simulations $\sim 4\times 10^4$ yr versus $1500$ yr).
Therefore there is no significant difference between the proton and the electron
Therefore there is no significant difference between the proton and the electron
- Hz.
- $\mu$ Hz.
One would expect large values of the EACF in this range close to vj. Which is not the case.
One would expect large values of the EACF in this range close to $\numax$, which is not the case.
Since the variations of Av(v) are not greater in this region than in others. one can explain the low values of the EACF either by a perturbation caused by a mixed mode or by a low value of the mode lifetime.
Since the variations of $\deltanunu$ are not greater in this region than in others, one can explain the low values of the EACF either by a perturbation caused by a mixed mode or by a low value of the mode lifetime.
The presence of a mixed mode ts in fact quite unlikely: such a mode with a much longer lifetime than normal p mode should appear with a large amplitude. which is not the case.
The presence of a mixed mode is in fact quite unlikely: such a mode with a much longer lifetime than normal p mode should appear with a large amplitude, which is not the case.
Thus. one has to favor the hypothesis of the varying mode lifetime.
Thus, one has to favor the hypothesis of the varying mode lifetime.
The comparison of the EACF with the bolometric amplitude per radial mode (see Sect. ??))
The comparison of the EACF with the bolometric amplitude per radial mode (see Sect. \ref{secampl}) )
indicates à strong increase of the mode lifetime below ;;Hz and a strong decrease above yHz (Fig. 14)).
indicates a strong increase of the mode lifetime below $\mu$ Hz and a strong decrease above $\mu$ Hz (Fig. \ref{comp_amp_ampl}) ).
This behaviour is in agreement with the measurement made in a similar star with a very close large separation. such as HD 49385 (?).. in spite of a different effective temperature.
This behaviour is in agreement with the measurement made in a similar star with a very close large separation, such as HD 49385 \citep{2010arXiv1003.4368D}, in spite of a different effective temperature.
However. it is not yet possible to quantify the lifetime.
However, it is not yet possible to quantify the lifetime.
To compute the maximum bolometric amplitude per radial mode we need to smooth the p-mode hump and correct from the background in this frequency range.
To compute the maximum bolometric amplitude per radial mode we need to smooth the p-mode hump and correct from the background in this frequency range.
To do so. we have used different methods (seeforanextensiveexplana-tion 222).
To do so, we have used different methods \citep[see for an extensive explanation][]{2008ApJ...682.1370K,2010MNRAS.402.2049H,2010A&A...511A..46M}.
Then to correct from the instrumental response function we use the method developped by ?..
Then to correct from the instrumental response function we use the method developped by \citet{2009A&A...495..979M}.
Thus. we have obtained a bolometric amplitude per radial mode of Apott_o x pppm at ~ 1070 μΗΖ.
Thus, we have obtained a bolometric amplitude per radial mode of $_{\rm{bol}, l=0}$ $\pm$ ppm at $\sim$ 1070 $\mu$ Hz.
However. among the different teams and depending on the method and on the way we fit the background. the value for the maximum amplitude per radial mode varied within the range 2.4 to 2.9 ppm indicating that the error bar of 0.2 ppm could be underestimated.
However, among the different teams and depending on the method and on the way we fit the background, the value for the maximum amplitude per radial mode varied within the range 2.4 to 2.9 ppm indicating that the error bar of 0.2 ppm could be underestimated.
So we have estimated a new value for the error of 0.6 ppm. which ts derived from the scatter of the smoothed power spectrum about the background fit outside the oscillation range (?)..
So we have estimated a new value for the error of 0.6 ppm, which is derived from the scatter of the smoothed power spectrum about the background fit outside the oscillation range \citep{2009CoAst.160...74H}.
The top panel of Fig.
The top panel of Fig.
15 shows the comparison of the amplitude found for HD 170987 with the published values of four other CoRoT solar-like targets (??) as a function of ων.
\ref{abol} shows the comparison of the amplitude found for HD 170987 with the published values of four other CoRoT solar-like targets \citep{2008Sci...322..558M, 2009A&A...506...33M} as a function of $\nu_{\rm max}$.
The bottom panel shows the amplitudes compared to theoretical values for each star calculated using the GCMYTAS scaling relation from Eq.
The bottom panel shows the amplitudes compared to theoretical values for each star calculated using the $(L/M)^{s} T_{\rm eff}^{-0.5}$ scaling relation from Eq.
3 of ?. with 2200.7 (2)..
3 of \citet{kjeldsen95} with 0.7 \citep{2007A&A...463..297S}.
It can be seen that while in all cases the amplitudes are systematically about Using (Av)IlοuAlo+ O.8uHz. fin 300μῃῆζ. fua.=1200 Hz (see Sect.
It can be seen that while in all cases the amplitudes are systematically about Using $\langle \Delta \nu \rangle = 55.2\,\pm\,0.8\,\mu$ Hz, $f_{\rm min} = 400\,\mu$ Hz, $f_{\rm max} = 1200\,\mu$ Hz (see Sect.
5.1 and 5.2). and the spectroscopic information loge=4.20+0.14 dex. |[M/H]=-0.20+0.15 dex and τη=0540+80 K (see Sect. ??)).
5.1 and 5.2), and the spectroscopic information $\log g = 4.20\,\pm\,0.14$ dex, $[M/H] = -0.20\,\pm\,0.15$ dex and $T_{\rm eff} = 6540\,\pm\,80$ K (see Sect. \ref{secspec}) ),
we compare these values with stellar models to determine the radius and the mass of HD 170987.
we compare these values with stellar models to determine the radius and the mass of HD 170987.
The stellar models that we use are the Aarhus Stellar Evolution Code (ASTEC) coupled with an adiabatic pulsation code (ADIPLS) (Christensen-Dalsgaard 2008a.b).
The stellar models that we use are the Aarhus Stellar Evolution Code (ASTEC) coupled with an adiabatic pulsation code (ADIPLS) (Christensen-Dalsgaard 2008a,b).
These codes need as input the stellar parameters of mass. age. chemical composition. and mixing-length parameter. and return stellar observables B;. such as radius. effective temperature. and the frequencies of the oscillation modes.
These codes need as input the stellar parameters of mass, age, chemical composition, and mixing-length parameter, and return stellar observables $B_i$ , such as radius, effective temperature, and the frequencies of the oscillation modes.
The parameters that best describe the observables are obtained by minimizing a y functio=) where v; and e; are the /=1.2.....M observations and errors.
The parameters that best describe the observables are obtained by minimizing a $\chi^2$ function; where $y_i$ and $\epsilon_i$ are the $i=1,2,...,M$ observations and errors.
Here. we have Mz-44.
Here, we have $M$ 4.
The Levenberg-Marquardt algorithm is used for the optimization. and this incorporates derivative information to guess the next set of parameters that will reduce the value of y.
The Levenberg-Marquardt algorithm is used for the optimization, and this incorporates derivative information to guess the next set of parameters that will reduce the value of $\chi^2$.
Naturally. an initial guess of the parameters is needed and these are obtained from a small grid of stellar evolution tracks.
Naturally, an initial guess of the parameters is needed and these are obtained from a small grid of stellar evolution tracks.
Because there are few observations and just as many parameters. there are inherent correlations between mass. age and chemical composition.
Because there are few observations and just as many parameters, there are inherent correlations between mass, age and chemical composition.
To help avoid local minima problems. we minimize the y function beginning at several initial guesses of the parameters (mainly varying in mass and age). and these initial guesses are estimated from the grids.
To help avoid local minima problems, we minimize the $\chi^2$ function beginning at several initial guesses of the parameters (mainly varying in mass and age), and these initial guesses are estimated from the grids.
We therefore obtain several sets of parameters with a corresponding y value that match the observations as best as possible.
We therefore obtain several sets of parameters with a corresponding $\chi^2$ value that match the observations as best as possible.