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The uncertainties in the estimation of these amplitudes are 450 px at. 13 and 15 (11. and ~SOO ply at 17 Cllz. | The uncertainties in the estimation of these amplitudes are $\sim 450$ $\mu$ K at 13 and 15 GHz, and $\sim 800$ $\mu$ K at 17 GHz. |
Although there are evidence of the presence of the ὃς 286 source in the map at 17 Cllz. due to its small amplitude as compared with the noise in this map. we decided to quote a 95 upper limit. | Although there are evidence of the presence of the 3C 286 source in the map at 17 GHz, due to its small amplitude as compared with the noise in this map, we decided to quote a 95 upper limit. |
In the region of the Galactic plane several point-like and extended sources are detected at the three frequencies. | In the region of the Galactic plane several point-like and extended sources are detected at the three frequencies. |
Good agreement between the positions and Ηχος of these sources have been found when comparing our data with the low-frequeney maps at 408 MllIz (Llaslam et al. ( | Good agreement between the positions and fluxes of these sources have been found when comparing our data with the low-frequency maps at 408 MHz (Haslam et al. ( |
1982) ancl 1420 Alllz (Reich 1982: Reich Reich 1986). | 1982) and 1420 MHz (Reich 1982; Reich Reich 1986). |
This is illustrated in Fie. | This is illustrated in Fig. |
9. where we display these surveys and our data in the region of Cve N. (νο AX is the source on the left of the plot. while most of the other structure corresponds to the (νο X complex. | 9, where we display these surveys and our data in the region of Cyg X. Cyg A is the source on the left of the plot, while most of the other structure corresponds to the Cyg X complex. |
A detailed analysis of the structure detected in our maps and a comparison between these datasets and the Galactic plane survey at 8.35 anc 14.35 CllIz (Langston et al. | A detailed analysis of the structure detected in our maps and a comparison between these datasets and the Galactic plane survey at 8.35 and 14.35 GHz (Langston et al. |
2000) will be presented in a forthcoming paper. | 2000) will be presented in a forthcoming paper. |
1) We have presented a new eround-based CALB experiment working at 13. 15 and 17 Cllz and an angular resolution of 1 degree based on a circular scanning strategy. | 1) We have presented a new ground-based CMB experiment working at 13, 15 and 17 GHz and an angular resolution of 1 degree based on a circular scanning strategy. |
The performance of the svstem and the reliabilitv of the data obtained during the first. months of commissioning. and operation have been discussed. | The performance of the system and the reliability of the data obtained during the first months of commissioning and operation have been discussed. |
We have demonstrated the possibility. of removing the atmospheric ancl differential eround pickup cllectively by the use of simple techniques that take into account the variation on angular ancl time scales of such components. | We have demonstrated the possibility of removing the atmospheric and differential ground pickup effectively by the use of simple techniques that take into account the variation on angular and time scales of such components. |
Daily maps covering 6000 square degrees of the sky are routinely obtained with sensitivities of SOO (dx. per beam area (ENWILM ~ 17)) at each frequency. | Daily maps covering 6000 square degrees of the sky are routinely obtained with sensitivities of $\sim800$ $\mu$ K per beam area (FWHM $\sim $ ) at each frequency. |
2) Observations at two ciüllerent. elevations have been performed. providing a stacked map of 9000 square degrees αἲ 13. 15 and 17 Cllz with mean sensitivities of 140. 150 ancl 250 ply per beam area respectively. | 2) Observations at two different elevations have been performed providing a stacked map of 9000 square degrees at 13, 15 and 17 GHz with mean sensitivities of 140, 150 and 250 $\mu$ K per beam area respectively. |
Phese stacked maps show no evidence of svstematies or striping due to 17f noise or residual atmospheric Iuctuations. | These stacked maps show no evidence of systematics or striping due to $1/f$ noise or residual atmospheric fluctuations. |
3) The strongest radio sources at high Galactic latitudes have been detected at the levels expected and the structure seen in the Galactic plane is in £ood agreement with the low-frequency surveys at 408 and. 1420 MIEL. | 3) The strongest radio sources at high Galactic latitudes have been detected at the levels expected and the structure seen in the Galactic plane is in good agreement with the low-frequency surveys at 408 and 1420 MHz. |
4) Several improvements of the system have been discussed. | 4) Several improvements of the system have been discussed. |
They. include updating of the filters and optics. which will improve the sensitivity of the daily data by a | They include updating of the filters and optics, which will improve the sensitivity of the daily data by a |
the inuer regions of the salaxy: fixine Yy- to more reasonable values worscus this discrepancy. | the inner regions of the galaxy; fixing $\Upsilon_V$ to more reasonable values worsens this discrepancy. |
The NEW halo piriuueters (corresponding to Y, 0) are c—6.5x1.1 aud CoU66.0€ 13.0. | The NFW halo parameters (corresponding to $\Upsilon_V=0$ ) are $\pm1.4$ and $v_{200}=66.0\pm13.0$ . |
The value of 6.5 for the concentration paraueter is much lower than the rauge of NEW concentration parameters predicted by LCDM sinulatious (Bullock et al. | The value of $6.5$ for the concentration parameter is much lower than the range of NFW concentration parameters predicted by LCDM simulations (Bullock et al. |
2001). | 2001). |
It is worth mentioniue here that even for the Ta rotation curve alone. isotheriial halos provide a substantially better fit than an NEW halo (de Blok et al. | It is worth mentioning here that even for the $\alpha$ rotation curve alone, isothermal halos provide a substantially better fit than an NFW halo (de Blok et al. |
2001). | 2001). |
As discussed. i section 21 for KIN9S 251. we SC he I baud scale length derived by Earacheutsev ct al. ( | As discussed in section \ref{ssec:opt_obs} for KK98 251, we use the I band scale length derived by Karachentsev et al. ( |
2000). | 2000). |
For isothermal halo models. keeping Y; as a free maralueter in the fit eave uceative values for Y;. | For isothermal halo models, keeping $\Upsilon_I$ as a free parameter in the fit gave negative values for $\Upsilon_I$. |
Further. if one keeps Y, fixed. the \? continously decreases as YT, is decreased. | Further, if one keeps $\Upsilon_I$ fixed, the $\chi^2$ continuously decreases as $\Upsilon_I$ is decreased. |
Fig. 9|[ | Fig. \ref{fig:massmodel2}[ [ |
C] shows the best fit mass model orKI98 251 for a coustant density halo using Y;=0.5 Qvlich correspouds to the observed «V-I color. frou he low inetallicity Druzual & Charlot SPS model using a modified Salpeter IME). | C] shows the best fit mass model forKK98 251 for a constant density halo using $\Upsilon_I=0.5$ (which corresponds to the observed $<$ $>$ color, from the low metallicity Bruzual $\&$ Charlot SPS model using a modified Salpeter IMF). |
The derived halo parameters using various values of Y; are given in Table 2.. | The derived halo parameters using various values of $\Upsilon_I$ are given in Table \ref{tab:halo}. . |
As can be secu frou Table 2.. the halo parameters are relatively iuxeusitive to the assumed value of Yj. | As can be seen from Table \ref{tab:halo}, the halo parameters are relatively insensitive to the assumed value of $\Upsilon_I$. |
The total dynamical mass (at the last measured poiut of the rotation curve) is [τοῦ «105A... | The total dynamical mass (at the last measured point of the rotation curve) is $\rm{M_T}$ $\times 10^8 \rm{M_\odot}$. |
For IKK98 251. we found that even for Y;=0. with uo value of c and cog could a eood fit using an NEW halo be obtained. | For KK98 251, we found that even for $\Upsilon_I=0$, with no value of c and $v_{200}$ could a good fit using an NFW halo be obtained. |
The best fit NFW 1iodel is also shown in Fig. 9|[ | The best fit NFW model is also shown in Fig. \ref{fig:massmodel2}[ [ |
C]: as can be seen. he clata deviate substautially from the model. | C]; as can be seen, the data deviate substantially from the model. |
It has been recently sueeested that dwarf ealaxics deviate systematically from the Tully-Fisher (TF) relation defined by bright ealaxies (ie. Stil (1999).. Swaters(1999). AMeGanehletal. (2000))). with small ealaxies being underhuuiuous compared to what would be expected had they followed the same TF relation as L.. galaxies. | It has been recently suggested that dwarf galaxies deviate systematically from the Tully-Fisher (TF) relation defined by bright galaxies (i.e. \cite{stil99}, , \cite{swaters99}, \cite{mcgaugh00}) ), with small galaxies being underluminous compared to what would be expected had they followed the same TF relation as $\sim L_*$ galaxies. |
Contrary to this sugecstion. ο&Tuffa(1999) found no evidence of auv break. in the near-IR TF relation for faint chwart galaxies. | Contrary to this suggestion, \cite{pierini99} found no evidence of any break in the near-IR TF relation for faint dwarf galaxies. |
The reason for this discrepancy is unclear. | The reason for this discrepancy is unclear. |
We note however that high resolution III iuages are probably crucial in studies of the ΤΕ reation for vorv faint galaxies. | We note however that high resolution HI images are probably crucial in studies of the TF relation for very faint galaxies. |
Tus is because () the inclinaJon nav often be difficult to obtain from iniages of fain nireenlar galaxies. aud (1) the HI velocity. wicth aay not be a good iudicator of the rotation velocity im faint dwarf galaxies where rancdoni motions are conrxuwable to thepeak rotatioial velocities (c.g. Camelopardalis D: Beemetal. (2003).. DDO?2?210: Beemu&Chneugalur (2001))). | This is because (i) the inclination may often be difficult to obtain from images of faint irregular galaxies, and (ii) the HI velocity width may not be a good indicator of the rotation velocity in faint dwarf galaxies, where random motions are comparable to thepeak rotational velocities (e.g. Camelopardalis B; \cite{begum03}, , DDO210; \cite{begum04}) ). |
Forsuch ealaxies. it is nuüportaut to accurately correct for he pressure support (vasvunuetric drift”correction). which | Forsuch galaxies, it is important to accurately correct for the pressure support (“asymmetric drift”correction), which |
At bolometric luminosities above LOL... infrared (IR) galaxies become the dominant population of extragalactic objects in the local universe (2x0.3). | At bolometric luminosities above $10^{11} L_ {\odot}$, infrared (IR) galaxies become the dominant population of extragalactic objects in the local universe $z \leq 0.3$ ). |
These galaxies are subdivided into three categories: luminous (LIRGs. Ly,>LOML. ). ultraluminous (ULIRGs. Lip> LOPL.). and byperluminous (IvLIRGs. Ly,>LOML.: Sanders&Mirabel1996)). | These galaxies are subdivided into three categories: luminous (LIRGs, $L_{\rm IR} > 10^{11} L_ {\odot}$ ), ultraluminous (ULIRGs, $L_{\rm IR} >10^{12} L_ {\odot}$ ), and hyperluminous (HyLIRGs, $L_{\rm IR} > 10^{13} L_ {\odot}$; \citealp{SM96}) ). |
Even though these IR galaxies are relatively rare. comprising less than of the total IR enerev density in the local Universe (Soiler&Neugebauer1991).. some studies suggest (hat the majority of galaxies with Lp>LOML. eo through a stage of intense IR. emission LOST). | Even though these IR galaxies are relatively rare, comprising less than of the total IR energy density in the local Universe \citep{SN91}, some studies suggest that the majority of galaxies with $L_{\rm B} > 10^{11}
L_ {\odot}$ go through a stage of intense IR emission \citep{SOI87}. |
Most IR galaxies with Ly,<10L. are single. gas-rich spirals. and their IR emission ean be accounted for by star formation. | Most IR galaxies with $L_{\rm IR} < 10^{11} L_{\odot}$ are single, gas-rich spirals, and their IR emission can be accounted for by star formation. |
In the Iuminosity range 10!<Lip1051... most of the galaxies are interacting/mereineg svstems wilh enormous «quantities of molecular eas (~LOMM. ). | In the luminosity range $10^{11} < L_{\rm IR} < 10^{12} L_ {\odot}$, most of the galaxies are interacting/merging systems with enormous quantities of molecular gas $\sim 10^{10}~M_ {\odot}$ ). |
At the lower end of this range. the bulk of the IR luminosity is due to warm dust. grains heated bv a nuclear starburst. while active galactic nuclei (AGN) become increasingly important at higher huninosities. | At the lower end of this range, the bulk of the IR luminosity is due to warm dust grains heated by a nuclear starburst, while active galactic nuclei (AGN) become increasingly important at higher luminosities. |
Galaxies with Ly,>LOL. are believed to be advanced mergers powered by a combination of starburst and AGN 1996 ). | Galaxies with $L_{\rm IR} >10^{12}
L_ {\odot}$ are believed to be advanced mergers powered by a combination of starburst and AGN \citep{SM96}. |
. Previous observations of I. galaxies have revealed very. broad absorption lines in ULIBRGs. indicating rotation plus large amounts of turbulent gas (Mirabel1982). | Previous observations of IR galaxies have revealed very broad absorption lines in ULIRGs, indicating rotation plus large amounts of turbulent gas \citep{MIR82}. |
. High angular resolution Very Large Array (VLA) and Very Long Baseline Interferometry (VLBI) observations show that (hese galaxies have the absorbing Tsituated in the inner few hundred parsecs along the line of sight to the nuclear continuum sources (Baan 2003).. | High angular resolution Very Large Array (VLA) and Very Long Baseline Interferometry (VLBI) observations show that these galaxies have the absorbing situated in the inner few hundred parsecs along the line of sight to the nuclear continuum sources \citep{BGSM87,MOM03}. . |
several OII 18 cin absorpüon and meeamaser (hereafter OIIM). emissionsurvevs of ULIRGs have alsobeen published (Baan1989;Darling&Giovanelli2000.2001.2002). | Several OH 18 cm absorption and megamaser (hereafter OHM) emissionsurveys of ULIRGs have alsobeen published \citep{B89,DG00,DG01,DG02}. |
. Baan(1989) concluded that the OILM emission usually occus in galaxies with higher (FIR) huninosities and flatter 100.25 yam spectra rather than in those with OIL 13 em absorption leatures. | \citet{B89} concluded that the OHM emission usually occurs in galaxies with higher far-IR (FIR) luminosities and flatter 100–25 $\mu$ m spectra rather than in those with OH 18 cm absorption features. |
llere. we report results [rom an on-going spectroscopic survey with the Arecibo Raclic targeting the 21 em and the main and satellite OIL 18 cm lines of 85 IR ealaxies [rom the 2-Jv IRAS-NVSS sample (Yun.Reddy&Condon 2001). | Here, we report results from an on-going spectroscopic survey with the Arecibo Radio targeting the 21 cm and the main and satellite OH 18 cm lines of 85 IR galaxies from the 2-Jy IRAS-NVSS sample \citep{YRC01}. . |
In this paper. we adopt fy — T1 km s!Mpe |. Q4; = 0.27. and O4 = 0.73. | In this paper, we adopt $H_0$ = 71 km $\rm s^{-1}$$\rm Mpc^{-1}$ , $\Omega_{M}$ = 0.27, and $\Omega_{\Lambda}$ = 0.73. |
2008 EVS were made in January 2009 with the 2018«20148 U12 CCD. while the LOT observations were all made duriug the observation ruus iu the dark period of January 2010 except (1123651) 2003 QOLOL which was observed in April 2009. with the 1310«1300 PL-1300B CCD and a 0.5. focal reducer. | 2008 EV5 were made in January 2009 with the $2048\times2048$ U42 CCD, while the LOT observations were all made during the observation runs in the dark period of January 2010 except (143651) 2003 QO104, which was observed in April 2009, with the $1340\times1300$ PI-1300B CCD and a $0.5\times$ focal reducer. |
A broad-band Bessell BVRI filter svstem was used on both telescopes. with the wavelengths centered at 12. 510. 617 ancl 786 nii respectively. | A broad-band Bessell $BVRI$ filter system was used on both telescopes, with the wavelengths centered at 442, 540, 647 and 786 nm respectively. |
Landolt standard stars (Landolt1992) are referred in optical photometry because thev cau oeuarautee Lighest accuracy (better than 0.01-0.02 nag) in mist cases. | Landolt standard stars \citep{lan92} are preferred in optical photometry because they can guarantee highest accuracy (better than 0.01-0.02 mag) in most cases. |
However. as Laudolt standard stars are ouly available to a very limited region. iu uost occasions one needs to know the atinosphleric extinction cocfiicicnt by observing staudard stars in different aimmmass to work ou targets locate far ποιι Landolt fields. | However, as Landolt standard stars are only available to a very limited region, in most occasions one needs to know the atmospheric extinction coefficient by observing standard stars in different airmass to work on targets locate far from Landolt fields. |
As we were unable to observe sufficient. Laudolt standard stars to ect a secure extinction coefficient. through every obscrving velit. au alternative approach introduced by Warner(2007) is eniploved. | As we were unable to observe sufficient Landolt standard stars to get a secure extinction coefficient through every observing night, an alternative approach introduced by \citet{war07} is employed. |
Warner applied third-order polvuonüals ou his optical observatious of 128 carefully-chosen Laudolt staudiud stars to find the conversion terms between the 2\TASS JI syste (Skrutskiectal.2006) and the Laudolt system. the errors of Warners method are O.031 for BV aud V. fand 0.021 for VR as dudicated in his paper. | Warner applied third-order polynomials on his optical observations of 128 carefully-chosen Landolt standard stars to find the conversion terms between the 2MASS $JHK$ system \citep{skr06} and the Landolt system, the errors of Warner's method are 0.034 for $B-V$ and $V-I$ and 0.021 for $V-R$ as indicated in his paper. |
The feld-ofwiew for LOT and the 0.H-1u. telescope are 22&22) and TsAT! respectively. which are laree enough to include sufficient 2\LASS catalog stars for setting up a good in-field transformation. | The field-of-view for LOT and the 0.41-m telescope are $22'\times22'$ and $47'\times47'$ respectively, which are large enough to include sufficient 2MASS catalog stars for setting up a good in-field transformation. |
To assess the accuracy of Warners method. we observed a few stars in AIGF (NGC 2682) and derived their colors following the procedure described by Warner. then compared them with another high precision VRZ photometry (Tavlorctal.2008). | To assess the accuracy of Warner's method, we observed a few stars in M67 (NGC 2682) and derived their colors following the procedure described by Warner, then compared them with another high precision $VRI$ photometry \citep{tay08}. |
. The resulting accuracy is better than 0.02 mag in V.AR aud VfF (Table 2)). | The resulting accuracy is better than 0.02 mag in $V-R$ and $V-I$ (Table \ref{tbl-2}) ). |
The targets are all observed with an airiiass of Xx2 with the predicted visual magnitude brighter than 19.0. | The targets are all observed with an airmass of $X\leq2$ with the predicted visual magnitude brighter than 19.0. |
Observational details as well as basic information of cach target are shown iu Tables 3. aud Las for NEAs/high ecceutricity (inclination) objects and paired-asteroid candidates. | Observational details as well as basic information of each target are shown in Tables \ref{tbl-3} and \ref{tbl-4} as for NEAs/high eccentricity (inclination) objects and paired-asteroid candidates. |
The exposure sequence for all targets is D-R-V- I to nünnmuze the error produced by siguificaut brightness variation. | The exposure sequence for all targets is $B$ $R$ $V$ $I$ to minimize the error produced by significant brightness variation. |
nage frames are then bias-subtracted aud fat-fielded. aud the frineine effects in J-banucdl images are also removed. | Image frames are then bias-subtracted and flat-fielded, and the fringing effects in $I$ -band images are also removed. |
The raw observations are then inspected manually o exclude the bad cases. such as the target asteroid crossing over or passing very close to ckeround stars. or low signal-to-ratio (SNR) caused by unstable weather couditious. | The raw observations are then inspected manually to exclude the bad cases, such as the target asteroid crossing over or passing very close to background stars, or low signal-to-ratio (SNR) caused by unstable weather conditions. |
Photometric ueasuremieuts are then performed with Warner's software MPO Canopus. | Photometric measurements are then performed with Warner's software $MPO$ $Canopus$. |
At least ten backeround stars with known 2MASS magnitudes are used o derive the transtormation coefficient between iustruimental maguitude and standard magnitude or each field. | At least ten background stars with known 2MASS magnitudes are used to derive the transformation coefficient between instrumental magnitude and standard magnitude for each field. |
In very few cases the Waited imiuber of background stars cannot guarantee a eood trausfori to be derived. so the coefficients derived. frou observations obtained iu the same üeht with a similar airmass (AN<~0.05) are used instead. | In very few cases, the limited number of background stars cannot guarantee a good transform to be derived, so the coefficients derived from observations obtained in the same night with a similar airmass $\Delta X<\sim0.05$ ) are used instead. |
Although cach target was planned ο he observed 3-5 times. various reasons (such as arect/star encounter. unstable weather. aud/or iustruinent problems) may prevent us to do so. and for some targets only one observation of cach filter was obtained. | Although each target was planned to be observed 3-5 times, various reasons (such as target/star encounter, unstable weather, and/or instrument problems) may prevent us to do so, and for some targets only one observation of each filter was obtained. |
These results shall be used with care. | These results shall be used with care. |
Observations of 53 asteroids were obtained and reduced following the procedure described iu Section 2. including 35NEAs*.. 6 high eccentricityinclination asteroids, and 12 main-belt pairecd-asteroid caucdidates. | Observations of 53 asteroids were obtained and reduced following the procedure described in Section 2, including 35, 6 high eccentricity/inclination asteroids, and 12 main-belt paired-asteroid candidates. |
The objects are then classified using the Daudxetal. (2003)7s derivation of the Tholen taxouomy (Tholen198L).. intent for optical broa-baudometry?. | The objects are then classified using the \citet{dan03}' 's derivation of the Tholen taxonomy \citep{tho84}, intent for optical broad-band. |
. We note that the spectra Yppearances of C-. Be. F-. Ge aud X-class are fairly close. uakiug it difficult to classify them uniquely. by xoacd-baud photometry. so all objects with colors simular to these classes are considered as N-class. | We note that the spectral appearances of C-, B-, F-, G- and X-class are fairly close, making it difficult to classify them uniquely by broad-band photometry, so all objects with colors similar to these classes are considered as X-class. |
The only exception is that the object shown to je particularly blue (5Vox 0.75) aud can | The only exception is that the object shown to be particularly blue $B-V\ll0.75$ ) and can |
magnetically dominated. plasma may. be considered. in ALD approximation. | magnetically dominated plasma may be considered in MHD approximation. |
In the MIID regime. the wave may decay as a result. of nonlinear steepening and. subsequent formation of multiple shocks. | In the MHD regime, the wave may decay as a result of nonlinear steepening and subsequent formation of multiple shocks. |
The wave steepening occurs because the wave velocity depends on the plasma density and therefore the compressive part of the wave moves faster than the expansive one. | The wave steepening occurs because the wave velocity depends on the plasma density and therefore the compressive part of the wave moves faster than the expansive one. |
In the strongly magnetized plasma. even a large amplitude wave is only weakly nonlinear because the wave velocity varies with the density only bv a [actor of about ασ. | In the strongly magnetized plasma, even a large amplitude wave is only weakly nonlinear because the wave velocity varies with the density only by a factor of about $1/\sigma$. |
The reason is that. as it was shown above. the conductivity current. which only introduces nonlinearity into the MIID. equations. is small in the magnetically dominated. plasma. | The reason is that, as it was shown above, the conductivity current, which only introduces nonlinearity into the MHD equations, is small in the magnetically dominated plasma. |
The cilference in the velocities between the compressive and expansive parts of the wave may be estimated. from ]5q.(2) as The shock forms when the leading edge of the wavelront becomes vertical (e... Landau Lifshitz 1959). | The difference in the velocities between the compressive and expansive parts of the wave may be estimated from Eq.(2) as The shock forms when the leading edge of the wavefront becomes vertical (e.g., Landau Lifshitz 1959). |
Dhis occurs within a characteristic nonlinear time it takes from the Compressive part to shift. relative the expansive one by lfw. | This occurs within a characteristic nonlinear time it takes from the compressive part to shift relative the expansive one by $\sim 1/\omega$. |
Therefore the shock forms after the wave travels a distance ‘This estimate is confirmed by a rigorous derivation in Appendix 1. | Therefore the shock forms after the wave travels a distance This estimate is confirmed by a rigorous derivation in Appendix 1. |
After the shock formation. the wave continue to distort at the characteristic scale yy such that eventually all the wave energy. dissipates in theshock. | After the shock formation, the wave continue to distort at the characteristic scale $x_{\rm nl}$ such that eventually all the wave energy dissipates in theshock. |
However both 0 and w may change because the plasma heats up and accelerates (or decelerates) considerably in the course of the wave dissipation. | However both $\sigma$ and $\omega$ may change because the plasma heats up and accelerates (or decelerates) considerably in the course of the wave dissipation. |
T'herefore numerically the dissipation scale may differ from the shock formation scale even though the same expression (7) is valid for both scales. | Therefore numerically the dissipation scale may differ from the shock formation scale even though the same expression (7) is valid for both scales. |
Far bevond the light evlinder. the wind may be considered as purely radial. whereas the magnetic field. as purely toroidal. | Far beyond the light cylinder, the wind may be considered as purely radial, whereas the magnetic field as purely toroidal. |
Phe wind is assumed to be super-EFMS. 5>να. | The wind is assumed to be super-FMS, $\gamma>\sqrt{\sigma}$. |
"Transforming the electromagnetic fields from the wind frame into the pulsar frame. one can see that the condition. (5) reduces to the condition 0«2, (the factor of two appears »eause in. the proper plasma frame⋅ £;−∕=0 whereas in. je pulsar frame Le,=e;D,c By). | Transforming the electromagnetic fields from the wind frame into the pulsar frame, one can see that the condition (5) reduces to the condition $\delta B<B_*$ (the factor of two appears because in the proper plasma frame $E'_*=0$ whereas in the pulsar frame $E_*=v_*B_*\approx B_*$ ). |
So FAIS waves may o: generated. by the rotating oblique magnetosphere at veh latitudes where the magnetic Ποιά does not change 1e sign. | So FMS waves may be generated by the rotating oblique magnetosphere at high latitudes where the magnetic field does not change the sign. |
In the equatorial belt (its width depends on the nele between the magnetic and the rotation axes of the »ulsar) an entropy wave is generated. in the form of the striped wind. | In the equatorial belt (its width depends on the angle between the magnetic and the rotation axes of the pulsar) an entropy wave is generated in the form of the striped wind. |
Generally a superposition of an entropy. and FAIS waves is generated here (one can talk about. the superposition because even a large. amplitude FAIS wave is weakly nonlinear in the magnetically dominated. plasma) however an important point is that an entropy. wave arises here inevitably because there is no other MILD wave that may transfer alternating magnetic field in a high-o plasma. | Generally a superposition of an entropy and FMS waves is generated here (one can talk about the superposition because even a large amplitude FMS wave is weakly nonlinear in the magnetically dominated plasma) however an important point is that an entropy wave arises here inevitably because there is no other MHD wave that may transfer alternating magnetic field in a $\sigma$ plasma. |
Now let us consider validity. of ΑΔΗΠΟ approximation (126.€6)) for the FAIS wave. | Now let us consider validity of MHD approximation (Eq.(6)) for the FMS wave. |
The magnetic field in the pulsar wind is predominantlv toroidal and may be presented: as where By=¢f/Q25-10"P em is the light evlinder radius. £2 the pulsar period. | The magnetic field in the pulsar wind is predominantly toroidal and may be presented as where $R_L=c/\Omega=5\cdot 10^9 P$ cm is the light cylinder radius, $P$ the pulsar period. |
The magnetic field at the light cvlinder may be estimated as ∖∖⋎⇂↥≺⋅↓⋅≺⋅∕∣∶↓∪⋟∕∣⊽⋝⊔≺⋅⋅≼∙⊔↓⋟↓⊳∖∣⇂⊔⋅⊔↓⋜↧⋏∙≟⊔≺⋅∣⊔∙⊔↓∪⊔↓⋖⋅⊔↿∪⇂ "m EN . ↿↓↥∢⊾⊳∖⋯↓⋅⊳↾ | The magnetic field at the light cylinder may be estimated as where $\mu=10^{30}\mu_{30}$ $\cdot$ $^3$ is the magnetic moment of the star. |
∐↥∢⊾∖∖⋎⋜↧∖⇁∢⊾∐⋅∢⊾⊏↥⋯⊾⊔≼∼∙∖⇁⊲↓⊔↿↓↥⋖⋅↓≻⊔↓⊳∖⋜⊔⋅⇂⋅↓⋅⋜⋯↓⋖⋅↕⊳∖⇂↓↕∢⊾↓≻⊔↓⊳∖⋜⊔⋅ ↓⋅⋖⋟↿⋜↧⇂↕∢≱↓↥∐⋅∢⊾⊏↥⋯⋅⊔≼∼∙∖⇁∶↕↓↕⇂↓↕⋖⊾∖∖⋰↓⊔∠⊔⋅↓⋅⋜⋯↓⋖⊾∣⇂↥⋖⊾⇂⋅↓⋅⋖⋅⊏↥⋯⋅↓∐⇍∙∖⇁⊲ | The wave frequency in the pulsar frame is the pulsar rotation frequency; in the wind frame the frequency is $\omega=\Omega/(2\gamma)$ . |
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