A new estimate of the Local Standard of Rest from data on young Galactic objects
aa r X i v : . [ a s t r o - ph . GA ] N ov Baltic Astronomy, vol. 99, 999–999, 2014
A NEW ESTIMATE OF THE LOCAL STANDARD OF REST FROMDATA ON YOUNG OBJECTS
V.V. Bobylev , and A.T. Bajkova Central (Pulkovo) Astronomical Observatory of RAS, 65/1 Pulkovskoye Ch.,St. Petersburg, Russia; [email protected] Sobolev Astronomical Institute, St. Petersburg State University, Bibliotechnayapl.2, St. Petersburg, Russia; [email protected]
Received: 2014 December 99; accepted: 2014 December 99
Abstract.
To estimate the peculiar velocity of the Sun with respect to theLocal Standard of Rest (LSR), we used young objects in the Solar neighborhoodwith distance measurement errors within 10%–15%. These objects were thenearest Hipparcos stars of spectral classes O–B2.5, masers with trigonometricparallaxes measured by means of VLBI, and two samples of the youngest andmiddle-aged Cepheids. The most significant component of motion of all thesestars is induced by the spiral density wave. As a result of using all these samplesand taking into account the differential Galactic rotation, as well as the influenceof the spiral density wave, we obtained the following components of the vectorof the peculiar velocity of the Sun with respect to the LSR: ( U ⊙ , V ⊙ , W ⊙ ) LSR =(6 . , . , . ± (0 . , . , .
3) km s − . We have found that the Solar velocitycomponents ( U ⊙ ) LSR and ( V ⊙ ) LSR are very sensitive to the Solar radial phase χ ⊙ in the spiral density wave. Key words:
Masers – Galaxy: kinematics and dynamics – galaxies: individ-ual: local standard of rest.1. INTRODUCTIONThe peculiar velocity of the Sun with respect to the LSR ( U ⊙ , V ⊙ , W ⊙ ) LSR playsan important role in analysis of the kinematics of stars in the Galaxy. To properlyanalyze Galactic orbits, this motion should be removed from the observed velocitiesof stars, since it characterizes only the Solar orbit — namely, its deviation fromthe purely circular orbit. In particular, to build a Galactic orbit of the Sun, it isdesirable to know the components ( U ⊙ , V ⊙ , W ⊙ ) LSR .There are several ways to determine the peculiar velocity of the Sun withrespect to the LSR. One of them is based on using the Str¨ o mberg relation. Themethod consists in finding such values ( U ⊙ , V ⊙ , W ⊙ ) LSR that correspond to zerostellar velocity dispersions (Dehnen & Binney, 1998; Bobylev & Bajkova, 2007;Co¸ckuno˘glu et al., 2011; Golubov et al., 2013). This method was addressed, forexample, in the work by Sch¨onrich et al. (2010), where the gradient of metallicityof stars in the Galactic disk was taken into account, and the velocity obtained is( U ⊙ , V ⊙ , W ⊙ ) LSR = (11 . , . , . ± (0 . , . , .
4) km s − .Another method implies a search for such ( U ⊙ , V ⊙ , W ⊙ ) LSR that lead to min- new estimate of the local standard of rest U ⊙ , V ⊙ , W ⊙ ) LSR = (7 . , . , . ± (1 . , . , .
1) km s − .To obtain this, 20000 local stars with known line-of-site velocities were used. Us-ing the improved database of proper motions and line-of-site velocities of Hippar-cos stars the same authors found the following new values: ( U ⊙ , V ⊙ , W ⊙ ) LSR =(14 . , . , . ± (1 . , . , .
1) km s − (Francis & Anderson, 2012); a great effortwas made to get rid of the influence of inhomogeneous distribution of velocities ofstars caused by kinematics of stellar groups and streams.Another method is based on transferring stellar velocities towards their ori-gin. Following this approach, Koval’ et al. (2009) derived the following values of( U ⊙ , V ⊙ , W ⊙ ) LSR = (5 . , . , . ± (0 . , . , .
2) km s − . Experience in using the Str¨ o mberg relation showed that the youngest starssignificantly deviate from a linear dependence when analyzing ( V ⊙ ) LSR . Thisoccurs when the dispersions σ <
17 km s − (Dehnen & Binney, 1998). Therefore,the youngest stars are not normally used in this method. Cepheids and otheryoungest objects fall into this area, which allows us to include them in our analysis.This behavior of velocities of the youngest stars is primarily connected with theeffect of the Galactic spiral density wave (Lin & Shu, 1964). For instance, theanalysis of kinematics of 185 Galactic Cepheids by Bobylev & Bajkova (2012)demonstrated that perturbation velocities inferred by the spiral density wave canbe determined with high confidence.The purpose of this paper is to estimate the velocity ( U ⊙ , V ⊙ , W ⊙ ) LSR usingspatial velocities of the youngest objects in the Solar neighborhood that haveparallax errors not larger than 10%–15%. For these stars, we consider not onlythe impact of the differential rotation of the Galaxy, but also the influence of thespiral density wave.2. METHODAssuming that the angular rotation velocity of the Galaxy (Ω) depends only onthe distance R from the axis of rotation, Ω = Ω( R ), the apparent velocity V ( r )of a star at heliocentric radius r can be described in vectorial notation by thefollowing relation V ( r ) = − V ⊙ + V θ ( R ) − V θ ( R ) + V ′ , (1)where V ⊙ ( U ⊙ , V ⊙ , W ⊙ ) is the mean stellar sample velocity due to the peculiarSolar motion with respect to the LSR (hence its negative sign), the velocity U isdirected towards the Galactic center, V is in the direction of Galactic rotation, W is directed to the north Galactic pole; R is the Galactocentric distance of theSun; R is the distance of an object from the Galactic rotation axis; V θ ( R ) is thecircular velocity of the star with respect to the center of the Galaxy, V θ ( R ) isthe circular velocity of the Sun, while V ′ are residual stellar velocities.From the above relation (1), one can write down three equations in components( V r , V l , V b ), the so-called Bottlinger’s equations (Eq. 6.27 in Trumpler & Weaver,1953): V r = (Ω − Ω ) R sin l cos b,V l = (Ω − Ω ) R cos l − Ω r cos b,V b = − (Ω − Ω ) R sin l sin b. (2)02 V.V. Bobylev, A.T. Bajkova
These are exact formulas, and the signs of Ω follow Galactic rotation. Afterexpanding Ω into Taylor series against the small parameter R − R , then expandingthe difference R − R , where the distance R is R = r cos b − R r cos b cos l + R , and then substituting the result into Eq. (2), one gets the equations of the Oort–Lindblad model (Eq. 6.34 in Trumpler & Weaver, 1953).Our approach departs from the above in that the distances r are known quitewell. In this case, there is no need to expand R − R into series, since the distance R is calculated using the distances r . Furthermore, our approach implies an extraassumption that the observed stellar velocities include perturbations due to thespiral density wave V sp ( V R , ∆ V θ ), with a linear dependence on both V sp and V ⊙ .This allows us to write: − V ⊙ = − V ⊙ LSR + V sp . Perturbations from the spiraldensity wave have a direct influence on the peculiar Solar velocity V ⊙ LSR . Thenthe relation (1) takes the following form: V ( r ) = − V ⊙ LSR + V sp + V θ ( R ) − V θ ( R ) + V ′ , (3)which, considering the expansion of the angular velocity of Galactic rotation Ωinto series up to the second order of r/R reads V r = − U ⊙ cos b cos l − V ⊙ cos b sin l − W ⊙ sin b + R ( R − R ) sin l cos b Ω ′ + 0 . R ( R − R ) sin l cos b Ω ′′ +∆ V θ sin( l + θ ) cos b − V R cos( l + θ ) cos b, (4) V l = U ⊙ sin l − V ⊙ cos l + ( R − R )( R cos l − r cos b )Ω ′ +( R − R ) ( R cos l − r cos b )0 . ′′ − r Ω cos b +∆ V θ cos( l + θ ) + V R sin( l + θ ) , (5) V b = U ⊙ cos l sin b + V ⊙ sin l sin b − W ⊙ cos b − R ( R − R ) sin l sin b Ω ′ − . R ( R − R ) sin l sin b Ω ′′ − ∆ V θ sin( l + θ ) sin b + V R cos( l + θ ) sin b, (6)where the following designations are used: V r is the line-of-sight velocity, V l =4 . rµ l cos b and V b = 4 . rµ b are the proper motion velocity components in the l and b directions, respectively, with the factor 4.74 being the quotient of the numberof kilometers in an astronomical unit and the number of seconds in a tropical year;the star’s proper motion components µ l cos b and µ b are in mas yr − , and the line-of-sight velocity V r is in km s − ; Ω is the angular velocity at the distance R from the rotation axis; parameters Ω ′ and Ω ′′ are the first and second derivativesof the angular velocity, respectively. To account for the influence of the spiraldensity wave, we used the simplest kinematic model based on the linear densitywave theory by Lin & Shu (1964), where the potential perturbation is in the formof a travelling wave. Then, V R = f R cos χ, ∆ V θ = f θ sin χ, (7)where f R and f θ are the amplitudes of the radial (directed toward the Galacticcenter in the arm) and azimuthal (directed along the Galactic rotation) velocityperturbations; i is the spiral pitch angle ( i < m is thenumber of arms (we take m = 2 in this paper); θ is the star’s position anglemeasured in the direction of Galactic rotation: tan θ = y/ ( R − x ), where x and new estimate of the local standard of rest y are the Galactic heliocentric rectangular coordinates of the object; radial phaseof the wave χ is χ = m [cot( i ) ln( R/R ) − θ ] + χ ⊙ , (8)where χ ⊙ is the radial phase of the Sun in the spiral density wave; we measure thisangle from the center of the Carina–Sagittarius spiral arm ( R ≈ λ , which is the distance along the Galactocentric radial direction betweenadjacent segments of the spiral arms in the Solar neighborhood (the wavelengthof the spiral density wave), is calculated from the relation 2 πR /λ = m cot( i ) . Wetake R = 8 . ± . V r − R ( R − R ) sin l cos b Ω ′ − . R ( R − R ) sin l cos b Ω ′′ − ∆ V θ sin( l + θ ) cos b − V R cos( l + θ ) cos b = − U ⊙ cos b cos l − V ⊙ cos b sin l − W ⊙ sin b, (9) V l − ( R − R )( R cos l − r cos b )Ω ′ − ( R − R ) ( R cos l − r cos b )0 . ′′ + r Ω cos b − ∆ V θ cos( l + θ ) − V R sin( l + θ ) = U ⊙ sin l − V ⊙ cos l, (10) V b + R ( R − R ) sin l sin b Ω ′ + 0 . R ( R − R ) sin l sin b Ω ′′ +∆ V θ sin( l + θ ) sin b + V R cos( l + θ ) sin b = U ⊙ cos l sin b + V ⊙ sin l sin b − W ⊙ cos b. (11)The system (9)–(11) can be solved by least-squares adjustment with respect tothree unknowns U ⊙ , V ⊙ , and W ⊙ . Another approach (which we follow) is tocalculate components of spatial velocities U, V, W of stars: U = V ′ r cos l cos b − V ′ l sin l − V ′ b cos l sin b,V = V ′ r sin l cos b + V ′ l cos l − V ′ b sin l sin b,W = V ′ r sin b + V ′ b cos b, (12)where V ′ r , V ′ l , V ′ b are left-hand parts of Eqs. (9)–(11) which are the observed stellarvelocities free from Galactic rotation and the spiral density wave. Then U = − U ⊙ ,V = − V ⊙ and W = − W ⊙ .3. DATA The sample of selected 200 massive ( < M ⊙ ) stars of spectral classes O–B2.5 isdescribed in detail in our previous paper (Bobylev & Bajkova, 2013a). It containsspectral binary O stars with reliable kinematic characteristics from the 3 kpc Solarneighborhood. In addition, the sample contains 124 Hipparcos (van Leeuwen,2007) stars of spectral types from B0 to B2.5 whose parallaxes were determinedto within 10% and better and for which there are line-of-sight velocities in thecatalog by Gontcharov (20006).04 V.V. Bobylev, A.T. Bajkova
Fig. 1.
Radial ( V R , dark) and tangential (∆ V θ , light) perturbation velocities versusGalactocentric distances R . Location of the Sun is indicated by a dotted line. In this work we solve the problem of determining the peculiar velocity of theSun. This problem can be solved most reliable using the closest stars to the Sun.Therefore, from the database, including 200 stars, we have selected 161 stars fromthe Solar neighborhood of 0 . = 32 . ± . − kpc − , Ω ′ = − . ± .
19 km s − kpc − , Ω ′′ = 0 . ± .
42 km s − kpc − ,f R = − . ± . − , f θ = 7 . ± . − , χ ⊙ = − ◦ ± ◦ . For all samplesin the present work, we use the same value of the wavelength λ = 2 . ± . i = − . ± . ◦ for m = 2). We use coordinates and trigonometric parallaxes of masers measured by VLBIwith errors of less than 10% in average. These masers are connected with veryyoung objects (basically proto stars of high masses, but there are ones with lowmasses too; a number of massive super giants are known as well) located in activestar-forming regions.One of such observational campaigns is the Japanese project VERA (VLBIExploration of Radio Astrometry) for observations of water (H O) Galactic masersat 22 GHz (Hirota et al., 2007) and SiO masers (which occur very rarely amongyoung objects) at 43 GHz (Kim et al., 2008). Water and methanol (CH OH)maser parallaxes are observed in USA (VLBA) at 22 GHz and 12 GHz (Reid etal., 2009). Methanol masers are observed also in the framework of the EuropeanVLBI network (Rygl et al., 2010). Both these projects are joined together in theBeSSeL program (Brunthaler et al., 2011). VLBI observations of radio stars incontinuum at 8.4 GHz (Dzib et al., 2011) are carried out with the same goals. new estimate of the local standard of rest = 29 . ± . − kpc − , Ω ′ = − . ± .
20 km s − kpc − , Ω ′′ = 0 . ± .
166 km s − kpc − , f R = − . ± . − , f θ = 7 . ± . − . In this case, the values of the phase of the Sun in the spiral wave foundindependently from radial and tangential perturbations using Fourier analysis aredifferent: ( χ ⊙ ) R = − ◦ ± ◦ and ( χ ⊙ ) θ = − ◦ ± ◦ , respectively.Note that line-of-site velocities of masers given in the literature usually refer tothe standard apex of the Sun. So we fix such line-of-site velocities, making themheliocentric. We used the data on classical Cepheids with proper motions mainly from theHipparcos catalog and line-of-sight velocities from the various sources. The datafrom Mishurov et al. (1997) and Gontcharov (2006), as well as from the SIMBADdatabase, served as the main sources of line-of-sight velocities for the Cepheids. Forseveral long-period Cepheids, we used their proper motions from the TRC (Hog etal., 1998) and UCAC4 (Zacharias et al., 2013) catalogs. To calculate the Cepheiddistances, we use the calibration from Fouqu et al. (2007), h M V i = − . − .
678 log P, where the period P is in days. Given h M V i , taking the period-averagedapparent magnitudes h V i and extinction A V = 3 . E ( h B i − h V i ) mainly fromAcharova et al. (2012) and, for several stars, from Feast & Whitelock (1997), wedetermine the distance r from the relation r = 10 − . h M V i − h V i − A V ) (13)and then assume that the relative error of Cepheid distances determined by thismethod is 10%. We divided the entire sample into two parts, depending on thepulsation period, which well reflects the mean Cepheid age ( t ). We use the cal-ibration from Efremov (2003), log t = 8 . − .
65 log P, obtained by analyzingCepheids in the Large Magellanic Cloud. Parameters of the Galactic rotation andspiral density wave depend on the age of the Cepheids. Therefore, for each sampleof Cepheids of the given age, these effects should be addressed individually. We usethe values of the parameters found in the work by Bobylev & Bajkova (2012) forthree age groups. The youngest Cepheids with periods of P ≥ d are characterizedby the average age of 55 Myr, middle-aged Cepheids with periods of 5 d ≤ P < d have the average age of 95 Myr, while the oldest Cepheids with periods of P < d have that of 135 Myr. In the present work, a sample of old Cepheids is not usedbecause there are very few of them in the Solar neighborhood, and their kinematicparameters are not very reliable.According to Bobylev & Bajkova (2012), for the youngest Cepheids with peri-ods of P ≥ d Ω = 26 . ± . − kpc − , Ω ′ = − . ± .
13 km s − kpc − ,Ω ′′ = 0 . ± .
10 km s − kpc − , f R = − . ± . − , χ ⊙ = − ◦ ± ◦ ,and the value of velocity perturbations in the tangential direction f θ is assumedto be zero. For middle-aged Cepheids with (5 d ≤ P < d ) Ω = 30 . ± . − kpc − , Ω ′ = − . ± .
13 km s − kpc − , Ω ′′ = 0 . ± .
14 km s − kpc − ,06 V.V. Bobylev, A.T. Bajkova
Table 1.
Components of the peculiar velocity of the Sun with respect to the LSR,calculated considering the differential Galactic rotation only.
Stars U ⊙ V ⊙ W ⊙ N ⋆ distancekm s − km s − km s − kpcO–B2.5 10 . ± . . ± . . ± . < . . ± . . ± . . ± . < . P ≥ d . ± . . ± . . ± . < d ≤ P < d . ± . . ± . . ± . < Table 2.
Components of the vector of the peculiar velocity of the Sun with respectto the LSR, calculated considering both the differential Galactic rotation and the spiraldensity wave.
Stars U ⊙ V ⊙ W ⊙ N ⋆ distancekm s − km s − km s − kpcO–B2.5 4 . ± . . ± . . ± . < . . ± . . ± . . ± . < . P ≥ d . ± . . ± . . ± . < d ≤ P < d . ± . . ± . . ± . < . ± . . ± . . ± . f R = − . ± . − , f θ = 2 . ± . − . The values for the phase of the Sunin the spiral wave found separately from radial and tangential perturbations by pe-riodogram analysis based on Fourier transform slightly differ: ( χ ⊙ ) R = − ◦ ± ◦ and ( χ ⊙ ) θ = − ◦ ± ◦ .
4. RESULT AND DISCUSSION R Here we describe the results obtained at fixed value of R = 8 kpc, assuming theparameters of differential Galactic rotation and the spiral density wave calculatedearlier independently for each stellar sample.In Figure 1, there are radial ( V R ) and tangential (∆ V θ ) velocities of pertur-bations vs Galactocentric distance R , induced by the spiral density wave. Thesevelocities are calculated according to the formulas (7) and (8) assuming θ = 0 ◦ ,and the amplitudes of perturbations f R and f θ defined in the data description(Section 3). As it can be seen from this figure, at R = R , the perturbationsachieve about 5 km s − in the radial direction. In the tangential direction, thesame value is achieved for two samples: of youngest O–B2 stars and of masers.In the case of young Cepheids, perturbations in the tangential direction are notsignificant. In the case of middle-aged Cepheids, perturbations in the tangentialdirection at R = R are close to zero. Note that a very small Solar neighborhood( R → R ) is crucial to determine the velocity ( U ⊙ , V ⊙ , W ⊙ ) LSR .In Table 1, the components of the peculiar velocity of the Sun with respect tothe LSR ( U ⊙ , V ⊙ , W ⊙ ) LSR are given. They were obtained only taking into accountthe influence of the differential Galactic rotation. Components of this vector, givenin Table 2, were calculated considering both the effects of the differential Galacticrotation and of the spiral density wave. new estimate of the local standard of rest
Table 3.
Components of the vector of the peculiar velocity of the Sun with respectto the LSR, calculated considering both the differential Galactic rotation and the spiraldensity wave for the different values of χ ⊙ . χ ⊙ − ◦ − ◦ − ◦ U ⊙ V ⊙ U ⊙ V ⊙ U ⊙ V ⊙ km s − km s − km s − km s − km s − km s − O–B2.5 6 . ± . . ± . . ± . . ± . . ± . . ± . As it is seen from Tables 1 and 2, considering the effect of the spiral densitywave for O–B2.5 stars and for masers leads to a considerable variation of thecomponents ∆ U ⊙ and ∆ V ⊙ by ≈ − . In addition, this gives smaller errors ofthe velocity ( U ⊙ , V ⊙ , W ⊙ ) LSR , which is especially noticeable for masers.The velocity ( V ⊙ ) LSR (Table 1) found from the data on masers differs from( V ⊙ ) LSR = 12 . − (Sch¨onrich et al., 2010) by ≈ − , which is in accor-dance with the result of analysis of masers in the Local arm (Xu et al., 2013).The following average values of the parameters ( U ⊙ , V ⊙ , W ⊙ ) LSR found in thepresent work are, essentially, more accurate than the estimate ( U ⊙ , V ⊙ , W ⊙ ) LSR =(5 . , . , . ± (2 . , . , .
2) km s − obtained from 28 masers by Bobylev & Ba-jkova (2010) considering the influence of the spiral density wave. The averagevalue of ( V ⊙ ) LSR (Table 2) is in a good agreement with the result by Sch¨onrich etal. (2010). There is a discrepancy in the ( U ⊙ ) LSR component with Sch¨onrich etal. (2010), and especially with Francis & Anderson (2012).Note that the revised Str¨ o mberg relation applied to the experimental RAVEdata gives an absolutely different velocity ( V ⊙ ) LSR ≈ − (Golubov et al.,2013). Using another approach to analysis of RAVE data Pasetto et al. (2012)obtained the following velocities: ( U ⊙ , V ⊙ ) LSR = (9 . , . ± (0 . , .
29) km s − .Recently kinematic analysis of RAVE and the GCS (Nordstr¨om et al. 2004)surveys was made by Sharma et al. (2014). To constrain kinematic parameters,were used analytic kinematic models based on the Gaussian and Shu distributionfunctions. Sharma et al. (2014) obtained the following velocities: ( U ⊙ , V ⊙ , W ⊙ ) LSR = (10 . +0 . − . , . +0 . − . , . +0 . − . ) km s − .Thus different methods give different results, and a final agreement on thevalues of the velocity ( U ⊙ , V ⊙ , W ⊙ ) LSR is not achieved till now. We consider ourestimates most reliable as they are based on the youngest stars characterized bya small velocity dispersion and by small Galactic orbit eccentricities as well.
Here we describe the results obtained for three particular values of R = 7 . , . , and 8 . f R and f θ , as well as the values of the Solar phase χ ⊙ in the spiral wave, arechosen as above in Section 4.1.For this purpose, we took a sample of masers (55 masers, σ π /π < , r < . R ,08 V.V. Bobylev, A.T. Bajkova found the following parameters of the Galactic rotation curve:Ω = 30 . ± . − kpc − , Ω ′ = − . ± .
21 km s − kpc − , Ω ′′ = 1 . ± .
180 km s − kpc − , R = 7 . = 29 . ± . − kpc − , Ω ′ = − . ± .
20 km s − kpc − , Ω ′′ = 0 . ± .
166 km s − kpc − , R = 8 . = 29 . ± . − kpc − , Ω ′ = − . ± .
18 km s − kpc − , Ω ′′ = 0 . ± .
154 km s − kpc − . R = 8 . Here we describe our results for several model values of the Solar phase χ ⊙ inthe spiral density wave for a sample of O–B2.5 stars (161 stars, r < . U ⊙ and V ⊙ arevery sensitive to the above parameter ( W ⊙ velocities are not shown in the Tableas they are practically not affected by the density wave).It is easy to understand these results by analyzing the corresponding panel ofFig. 1 and Table 1. For instance, for χ ⊙ = − ◦ , the radial perturbation curve( V R ) is near its maximum, so the influence to the U ⊙ component is most prominent( U ⊙ = 0 . ± . − ). On the contrary, the tangential perturbation curve (∆ V θ )is about zero, so there is no effect on the V ⊙ component ( V ⊙ = 12 . ± . − ).For χ ⊙ = − ◦ , the radial perturbation curve ( V R ) is near zero, so there is no effecton the U ⊙ component, U ⊙ = 11 . ± . − . The tangential perturbation curve(∆ V θ ) is near to minimum, so there is no significant effect on the V ⊙ component, V ⊙ = 7 . ± . − .We must note that, in our previous paper (Bobylev & Bajkova, 2013a), theuncertainty of R was not taken into account when determining the Solar phase inthe spiral density wave χ ⊙ = − ± ◦ . We have redone Monte Carlo simulationand obtained the following results:1. If we consider only the error σ R = 0 . σ χ ⊙ = 0 . ◦ .The explanation for this is that when you change R , the length of a wavestretches like a rubber band, but the phase of the Sun in the spiral wavepractically does not change.2. If we consider the errors of all observed parameters of stars – parallaxes,proper motions, line-of-site velocities – along with the uncertainty σ R , thenthe Solar phase in the spiral density wave becomes χ ⊙ = − ± ◦ .Based on the data from Table 3, we may conclude that, in the range of phasevalues from − ◦ to − ◦ (which is even above the 1 σ level), the Solar velocities new estimate of the local standard of rest U ⊙ = 6 − − and V ⊙ =8 − − range.5. CONCLUSIONSFor evaluation of the peculiar velocity of the Sun with respect to the Local Stan-dard of Rest, we used young objects from the Solar neighborhood with distanceerrors of not larger than 10%–15%. These are the nearest Hipparcos stars ofspectral classes O–B2.5, masers with trigonometric parallaxes measured by meansof VLBI, and two samples of the youngest and middle-aged Cepheids. The wholesample consists of 297 stars. A significant fraction of motion of these stars is causedby the Galactic spiral density wave, because the amplitudes of perturbations inradial ( f R ) and tangential ( f θ ) directions reach ≈
10 km s − . For each sample of stars, the impact of differential Galactic rotation and ofthe Galactic spiral density wave was taken into account. It was shown that, forthe youngest objects – namely, stars of spectral classes O–B2.5 and masers –considering the effect of the spiral density wave leads to a change in the values ofthe components of the peculiar velocity of the Sun with respect to the LSR ∆ U ⊙ and ∆ V ⊙ by ≈ − . Cepheids are less sensitive to the influence of the spiraldensity wave.Average values of the peculiar velocity of the Sun with respect to the LSR arecalculated according to the results of analysis of four samples of stars; they havethe following values: ( U ⊙ , V ⊙ , W ⊙ ) LSR = (6 . , . , . ± (0 . , . , .
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